{
    "0807/0807.0809_arXiv.txt": {
        "abstract": "Studies of nucleosynthesis in neutrino-driven winds from nascent neutron stars show that the elements from Sr through Ag with mass numbers $A\\sim 88$--110 are produced by charged-particle reactions (CPR) during the $\\alpha$-process in the winds. Accordingly, we have attributed all these elements in stars of low metallicities (${\\rm [Fe/H]}\\lesssim -1.5$) to low-mass and normal supernovae (SNe) from progenitors of $\\sim 8$--$11\\,M_\\odot$ and $\\sim 12$--$25\\,M_\\odot$, respectively, which leave behind neutron stars. Using this rule and attributing all Fe production to normal SNe, we previously developed a phenomenological two-component model, which predicts that ${\\rm [Sr/Fe]}\\geq -0.32$ for all metal-poor stars. The high-resolution data now available on Sr abundances in Galactic halo stars show that there is a great shortfall of Sr relative to Fe in many stars with ${\\rm [Fe/H]}\\lesssim -3$. This is in direct conflict with the above prediction. The same conflict also exists for two other CPR elements Y and Zr. The very low abundances of Sr, Y, and Zr observed in stars with ${\\rm [Fe/H]}\\lesssim -3$ thus require a stellar source that cannot be low-mass or normal SNe. We show that this observation requires a stellar source leaving behind black holes and that hypernovae (HNe) from progenitors of $\\sim 25$--$50\\,M_\\odot$ are the most plausible candidates. Pair-instability SNe from very massive stars of $\\sim 140$--$260\\,M_\\odot$ that leave behind no remnants are not suitable as they are extremely deficient in producing the elements of odd atomic numbers such as Na, Al, K, Sc, V, Mn, and Co relative to the neighboring elements of even atomic numbers, but this extreme odd-even effect is not observed in the elemental abundance patterns of metal-poor stars.  If we expand our previous phenomenological two-component model to include three components (low-mass and normal SNe and HNe) and use for example, the observed abundances of Ba, Sr, and Fe to separate the contributions from these components, we find that essentially all of the data are very well described by the new model. This model provides strong constraints on the evolution of [Sr/Fe] with [Ba/Fe] in terms of the allowed domain for these abundance ratios. This model also gives an equally good description of the data when any CPR element besides Sr (e.g., Y or Zr) or any heavy $r$-process element besides Ba (e.g., La) is used. As the stars deficient in Sr, Y, and Zr are dominated by contributions from HNe, they define the self-consistent yield pattern of that hypothecated source. This inferred HN yield pattern for the low-$A$ elements from Na through Zn ($A\\sim 23$--70) including Fe is almost indistinguishable from what we had previously attributed to normal SNe. As HNe are plausible candidates for the first generation of stars and are also known to be ongoing in the present epoch, it is necessary to re-evaluate the extent to which normal SNe are substantial contributors to the Fe inventory of the Galaxy. We conclude that HNe are important contributors to the abundances of the low-$A$ elements over the history of the universe. We estimate that they contributed $\\sim 24\\%$ of the bulk solar Fe inventory while normal SNe contributed only $\\sim 9\\%$ (not the usually assumed $\\sim 33\\%$). This implies a greatly reduced role of normal SNe in the chemical evolution of the low-$A$ elements. ",
        "introduction": "In this paper we consider that the elements from Sr through Ag in metal-poor stars represent the products of nucleosynthesis in neutrino-driven winds from forming neutron stars. This approach allows us to obtain information on the stellar sources that contributed to the chemical enrichment of the interstellar medium (ISM) in the Galaxy and the intergalactic medium (IGM) at early and recent  times. We previously proposed a phenomenological two-component model (\\citealt{qw07}; hereafter QW07) to account for the abundances of heavy elements in metal-poor stars. That model focused on the elements commonly considered to be produced by the generic ``$r$-process''. It specifically attributed all the elements from Sr through Ag in metal-poor stars to the charged-particle reactions (CPR) in the neutrino-driven winds from nascent neutron stars and used this as a diagnostic of the sources for these CPR elements. In contrast, the true $r$-process elements (e.g., Ba and higher atomic numbers) are produced by extensive rapid neutron capture. It was assumed in the two-component model that Fe was only produced by normal supernovae (SNe) from progenitors of $\\sim 12$--$25\\,M_\\odot$, which leave behind neutron stars, and that the heavy $r$-process elements ($r$-elements) with mass numbers $A>130$ were formed in low-mass SNe from progenitors of $\\sim 8$--$11\\,M_\\odot$, which also leave behind neutron stars but produce no Fe. Thus the CPR elements would be produced by both low-mass and normal SNe and the corresponding yields were estimated (QW07). It follows that if these SNe were the only sources, then the presence of Fe should always be associated with that of the CPR elements. A more extensive study of the available observational data shows that some low-metallicity stars have Fe but essentially no Sr. In particular, \\citet{fulbright04} found a star with ${\\rm [Fe/H]}=\\log{\\rm (Fe/H)}-\\log{\\rm (Fe/H)}_\\odot=-2.88$ and $\\log\\epsilon({\\rm Sr})=\\log({\\rm Sr/H})+12<-2.6$ in the dwarf galaxy Draco. From the two-component model we would have estimated $\\log\\epsilon({\\rm Sr})=-0.28$ for this star, which is far above the observational upper limit. These results clearly indicate that if the CPR elements are always produced during the formation of neutron stars, then there must be an additional stellar source contributing Fe that does not leave behind neutron stars, or else the above model for the production of the CPR elements is in error.  Utilizing a more extensive data base than QW07 and especially treating the data on stars very deficient in Sr, Y, and Zr relative to Fe at ${\\rm [Fe/H]}<-3$, the present paper will show that a third source in addition to the two sources (low-mass and normal SNe) in the model of QW07 is required to account for the elemental abundances in metal-poor stars. It will be argued that the third source producing Fe but no CPR elements is most likely associated with hypernovae (HNe) from progenitors of $\\sim 25$--$50\\,M_\\odot$, which leave behind black holes instead of neutron stars.  It is then shown that essentially all of the stellar data on elemental abundances at ${\\rm [Fe/H]}\\lesssim -1.5$ can be decomposed in terms of three distinct types of sources. This decomposition also identifies a yield pattern for the elements from Na through Zn including Fe that is attributable to HNe. An important conclusion is that this HN yield pattern is almost indistinguishable from what is attributed to normal SNe. Further, the discovery of extremely energetic HNe associated with gamma-ray bursters (e.g., \\citealt{galama,iwa98}) in the present universe requires that contributions from this source must be considered both in early epochs and on to the present. This leads to a reassessment of the contributions from different sources to the Galactic Fe inventory, which shows that ongoing HNe must play an important role and that the usual attribution of $\\sim 1/3$ of the solar Fe inventory to normal SNe is not valid.  We aim to present a phenomenological three-component (low-mass and normal SNe and HNe) model for the chemical evolution of the early Galaxy that may provide a quantitative, self-consistent explanation for many of the results from stellar observations. We focus on three groups of elements: the low-$A$ elements from Na through Zn ($A \\sim 23$--70), the CPR elements from Sr through Ag ($A \\sim 88$--110), and the heavy $r$-elements ($A > 130$, Ba and higher atomic numbers). In \\S\\ref{sec-2cm} we give a brief outline of the two-component model of QW07 with low-mass and normal SNe represented by the $H$ and $L$ sources, respectively. In \\S\\ref{sec-data} we present the data on abundances of Sr and Ba as well as Y and La for a large sample of metal-poor stars, and show that the two-component model fails at ${\\rm [Fe/H]}\\lesssim -3$ and that an additional source producing Fe but no Sr or heavier elements is required to account for the data at such low metallicities. This source is identified with HNe. It is then shown that the extended three-component model with HNe, $H$, and $L$ sources gives a good representation of nearly all the data on the CPR elements Sr, Y, and Zr, but leads to the conclusion that the HN yield pattern is indistinguishable from that of the $L$ source for all the low-$A$ elements. Considering that HNe not only represent the first massive stars (Population III stars) but also must continue into the present epoch, we reinterpret the yields attributed to the hypothetical $L$ source as the combined contributions from normal SNe, which we designate as the $L^*$ source, and HNe. In \\S\\ref{sec-3cm} we show that the three-component model with HNe, $H$, and $L^*$ sources gives a very good representation of essentially all the data on the CPR elements Sr, Y, and Zr and further discuss the characteristics of these sources and their roles in the chemical evolution of the universe. We give our conclusions in \\S\\ref{sec-con}. ",
        "conclusions": "\\label{sec-con} The two-component model of QW07 with the $H$ and $L$ sources provided a good description of the elemental abundances in metal-poor stars of the Galactic halo for $-2.7<{\\rm [Fe/H]}\\lesssim -1.5$. A key ingredient of that model is the attribution of the elements from Sr through Ag in metal-poor stars to the charged-particle reactions in neutrino-driven winds from nascent neutron stars but not to the $r$-process. However, that model cannot explain the great shortfall in the abundances of Sr, Y, and Zr relative to Fe for stars with ${\\rm [Fe/H]}\\lesssim -3$. The observations on these three CPR elements require that there be an early source producing Fe but no Sr or heavier elements. It is shown that if such a third source is assumed, then the data can be well explained by an extended three-component model. From considerations of the abundance patterns of the low-$A$ elements (from Na through Zn), it is concluded that this third source is most likely associated with HNe from massive stars of $\\sim 25$--$50\\,M_\\odot$ that do not leave behind neutron stars. We here consider the third source to be HNe.  It is shown that the available data on the evolution of [Sr/Fe] with [Ba/Fe], that of [Y/Fe] with [La/Fe], and that of [Zr/Fe] with [Ba/Fe] are well described by the extended model with HNe, $H$ and $L$ sources, which also provides clear constraints on the abundance ratios that should be seen. It is further shown that the abundance patterns of the low-$A$ elements for HNe and the $L$ sources are not distinguishable. Considering that HNe are observed to be ongoing events in the present universe, we are forced to conclude that the $L$ source, which was assumed to have provided $\\sim 1/3$ of the solar Fe inventory (the rest attributed to SNe Ia), is in fact a combination of normal SNe (from progenitors of $\\sim 12$--$25\\,M_\\odot$), which we define as the $L^*$ source, and HNe. The net Fe contributions from HNe are found to be $\\sim 3$ times larger than those from normal SNe.  Using the three-component model with HNe, $H$, and $L^*$ sources, we obtain a very good quantitative description of essentially all the available data. In particular, this model provides strong constraints on the evolution of [Sr/Fe] with [Ba/Fe] in terms of the allowed domain for these abundance ratios. It gives an equally good description of the data when any CPR element besides Sr (e.g., Y or Zr) or any heavy $r$-element besides Ba (e.g., La) is used. The model is also compatible with the non-$s$-process contributions to the solar abundances of all the CPR elements. The anomalous abundance patterns of the low-$A$ elements observed in a small number of stars appear to fit the description of faint SNe (e.g., \\citealt{iwamoto}), which are a rarer type of events from the same progenitor mass range as HNe but with even weaker explosion energies and smaller Fe yields than normal SNe (e.g.,  \\citealt{nomoto06}). The anomalous abundance patterns observed reflect the fact that faint SNe produce very little of the Fe group elements but an abundant amount of the elements from hydrostatic burning in their outer shells. This gives rise to the extremely high abundances of Na, Mg, and Al relative to Fe observed in the anomalous stars. The quasi-uniform abundance patterns of the elements from Si through Zn in all cases (including the stars with anomalous abundances of Na, Mg, and Al) appear to reflect some robustness in the outcome of explosive burning that may arise from the limited range of conditions required for such nucleosynthesis.  In this paper we used the elemental yield patterns for three prototypical model sources to calculate the abundances of an extensive set of elements (relative to hydrogen) for metal-poor stars with ${\\rm [Fe/H]}\\lesssim -1.5$. As the yield patterns adopted for the assumed prototypical sources are taken from the data on two template stars, they must represent the results of stellar nucleosynthesis. The full version of the three-component model appears very successful in calculating the abundances of the elements ranging from Na through Pt in stars with ${\\rm [Fe/H]}\\lesssim -1.5$. In contrast to this phenomenological approach, there are extensive studies of Galactic chemical evolution (GCE) that use the various theoretical results on the absolute yields of metals for different stellar types. These theoretical yields are not calculated from first principles, but are dependent on the parametrization used in the various stellar models. In those GCE studies, the elemental abundances for an individual star are not predicted. Instead, general trends for the elemental abundances are calculated assuming different sources, the rates at which they contribute, and a model of mixing in the ISM for different regions of the Galaxy. These results give a good broad description for typical elemental abundances in the general stellar population at higher metallicities of ${\\rm [Fe/H]}>-1.5$. This is a regime in which the observational data are quite convergent with only limited variability. However, as anticipated by \\citet{gilroy} and supported by the considerable scatter in the abundances of heavy elements observed in stars with ${\\rm [Fe/H]}\\lesssim -1.5$, the chemical composition of the ISM in the early Galaxy was extremely inhomogeneous. For ${\\rm [Fe/H]}<-2$ there are gross discrepancies between the observations and the smoothed model of GCE. In no case does that model give the elemental abundances for an individual star. It is our view that the simple phenomenological model used here permits a clearer distinction between the different stellar sources contributing to the ISM and the IGM at early times. This model also gives specific testable predictions, which can be used to further check its validity.  In conclusion, we consider that the general three-component model with HNe, $H$, and $L^*$ sources provides a quantitative and self-consistent description of nearly all the available data on elemental abundances in stars with ${\\rm [Fe/H]}\\lesssim -1.5$. Further, HNe may be not only explosions from the first massive stars (i.e., the Population III stars much sought after by many) that provided a very early and variable inventory to the IGM through ejection of enriched gas from small halos, but also are important ongoing contributors to the chemical evolution of the universe."
    },
    "0807/0807.2384_arXiv.txt": {
        "abstract": "LS 5039, a possible black hole x-ray binary, was recently observed with Giant Meterwave Radio Telescope. The observed spectrum presented here shows that the spectrum is inverted at the low frequency. When combined with the archival data with orbital phase similar to the present observations, it shows a clear indication of a spectral turnover. The combined data are fitted with a broken power-law and the break frequency signifies a possible spectral turnover of the spectrum around 964 MHz. Truly simultaneous observations in radio wavelength covering a wide range of frequencies are required to fix the spectrum and the spectral turn over which will play a crucial role in developing a deeper understanding of the radio emitting jet in LS 5039. ",
        "introduction": "LS5039 is a X-ray binary system consisting of a compact object and a companion star of ON6.5V(f) type with mass 22.9M$_{\\odot}$. The mass of the compact object is not accurately known. But recent study shows that the mass of the compact object is 3.7M$_{\\odot}$ (\\cite{cas05}). Although there are lot of debates regading the true nature of LS 5039 (\\cite{mira06}), the dynamical mass estimate indicates that it is a possible black hole x-ray binary candidate. The orbital period of the binary system is 3.9 days (\\cite{cas05}). The binary orbit is highly eccentric with an eccentricity of 0.35 and has an inclination angle 24$^0$.9$\\pm$2$^0$.8 (\\cite{cas05}).  LS5039 was proposed to be a very high energy gamma-ray emitter by \\cite{parde00} and was recently seen at TeV $\\gamma$-ray energies by gamma-ray telescope High Energy Stereoscopic System (\\cite{aha06}). The gamma-ray light curve shows an orbital modulation of $\\sim3.9$ days. The gamma-ray flux as well as the spectral hardness vary with the orbital phase. The flux variation with orbital phase can be accounted for by considering the absorption of gamma-rays due to $\\gamma\\gamma$ pair production. This also implies that $\\gamma$-photons are emitted from a region within 1 AU of the compact object (\\cite{aha06}). The observed spectrum is very steep at the superior conjunction and becomes harder as the compact object moves away from it and the spectral index goes below 2.0 as the compact object reaches the inferior conjuntion. But this variation of spectral hardness can not be adequately explained by the gamma-ray absorption via pair production alone.  This source was observed in X-ray by $ROSAT$ (\\cite{mot97}), $RXTE$ (\\cite{ribo99}), $ASCA$ (\\cite{mart05}), $BeppoSAX$ (\\cite{reig03}), $Chandra$ (\\cite{bosc05}) and $XMM-Newton$ (\\cite{mart05}) missions during 1996 to 2003. These observations were effectively carried out during different orbital phases of the binary system. All these observations revealed different flux levels and different spectral indices in the X-ray energy band, but no x-ray spectral state transition were reported in these observations. More recently \\cite{bosc05} reported results of $RXTE$ observations when the source was observed consecutively for four days in July 2003. The X-ray flux is found to vary with the orbital phases and it maximizes near periastron passage. This phase dependence of X-ray flux is due to the motion of the compact object accreting matter from the wind of the massive companion while moving in a highly eccentric orbit (\\cite{bosc05}). But the observed anti-correlation between the photon index and the X-ray flux does not fit to the scenario where X-rays are considered to be produced in and around an accretion disc. \\cite{bosc05} argued that the x-ray emission might be due to inverse Compton/synchrotron process in a relativistic jet which might possibly explain the observed anti-correlation between the photon index and the observed flux. But the photon indices reported in these observations are very similar to that generally observed in case of hard state spectrum of black hole x-ray binaries.   LS5039 was first observed by \\cite{marti98} in radio with Very Large Array (VLA) and the observation resulted a power-law spectrum with a negative power-law index. This indicates an optically thin synchrotron emission by non-thermal electrons. This observation also revealed a moderate variability in the radio flux. Later, \\cite{ribo99} carried out an observation campaign in radio with VLA at 2.0, 3.6, 6.0 and 20.0 cm wavelengths and with Green Bank Interferometer (GBI) at 3.6 and 13.3 cm wavelengths. The observed spectrum was a power-law with a power-law index of -0.46$\\pm$0.01 supporting the previous observation of \\cite{marti98}. \\cite{pare02} first resolved the source by observing with European VLBI Network (EVN) and the Multi-Element Radio-Linked Interferometer Network (MERLIN) and it was found that LS 5039 consists of an asymmetric two-sided jet with an maximum extension of one side $\\sim$ 1000 AU. Recently \\cite{ribo08} studied the radio morphology of LS5039 using Very Long Baseline Array (VLBA) and it was found that the radio morphology changed into asymmetric from a very symmetric configuration within five days, but the radio emission from the core did not vary appreciably. {\\it But what is most interesting found in these radio observations is the optically thin non-thermal radio spectrum of the source which is in direct contradiction to the generally observed trend of flat or inverted radio spectrum for persistent black hole x-ray binaries in the hard state~(\\cite{fen01,gal03}). Therefore it is important to study LS5039 in the low-frequency radio band.}   We observed LS5039 at wavelengths 21 cm, 50 cm and 128 cm (1280 MHz, 614 MHz and 234 MHz respectively) using Giant Meterwave Radio Telescope~(GMRT). In this letter we discuss the results of the observation. In Section 2 we discuss the observation and data reduction, results are discussed in Section 3 and we conclude the paper in Section 4. ",
        "conclusions": "Here, we report the low frequency radio spectrum of a X-ray binary LS5039 as observed with GMRT. The spectrum is inverted and shows an indication of spectral turnover. With the assumption that the variability of radio emission from the source is not appreciable, present data are plotted with the archival data. The fitting of the data with a broken power-law reveals a possible indication of a spectral turnover at 964 MHz and an inverted spectrum with spectral index 0.749. This estimation of turnover frequency is indeed obtained using the archival VLA data (\\cite{marti98}) with the assumption that the source is not variable. But this estimation may depend on the variability of the source. To have an idea about the sensitivity of the spectral break frequency on flux variability, we introduced a moderate ($\\sim 10\\%$) flux variation to VLA data and fitted the spectrum. The spectral fitting shows that the break frequency $\\nu_b$ does not change appreciably and it is well within the uncertainities of the statistical fit.  As LS 5039 is a possible black hole x-ray binary persistently in low/hard x-ray spectral state, the results of the present observation are consistent with the trend of inverted radio spectrum observed for persistent black hole x-ray binary and this is the first source of this kind for which the indication of a spectral turnover is obtained. The spectral turnover can be used to determine the magnetic field in the jet provided the size of the source is known at that frequency. But it is to be noted that the magnetic field depends on the fourth power of the angular size. Therefore, a reliable measurement of magnetic field is possible if and only if the angular size of the source at the turnover frequency is determined with sufficient accuracy. This result will further help to constrain the theoretical models to explain the broadband spectrum from the source.   To have a deeper understanding on the radio spectrum and its dependence on the orbital phases it is important to have simultaneous observations over the entire radio band. Nevertheless, we are conducting a long term campaign to study the spectral behaviour at frequencies 234, 614 and 1280 MHz over different orbital phases of LS5039. This will allow us to constrain the low frequency end of the spectrum of LS5039."
    },
    "0807/0807.2451_arXiv.txt": {
        "abstract": "Emission from Active Galactic Nuclei is known to vary strongly over time over a wide energy band, but the origin of the variability and especially of the inter-band correlations is still not well established. Here we present the results of our X-ray and optical monitoring campaign of the quasar \\mr , covering a period of 2.5 years. The X-ray 2--10 keV flux is remarkably well correlated with the optical B, V and R bands, their fluctuations are almost simultaneous with a delay consistent with 0 days and not larger than 4 days in either direction. The amplitude of variations shows an intriguing behaviour: rapid, large amplitude fluctuations over tens of days in the X-rays have only small counterparts in the optical bands, while the long-term trends over hundreds of days are \\emph{stronger} in the B band than in X-rays. We show that simple reprocessing models, where all the optical variability arises from the variable X-ray heating, cannot simultaneously explain the discrepant variability amplitudes on different time-scales and the short delays between X-ray and optical bands. We interpret the variability and correlations, in the optically-thick accretion disc plus corona scenario, as the result of intrinsic accretion rate variations modulating both X-ray and optical emission, together with reprocessing of X-rays by the accretion disc. ",
        "introduction": "Optical continuum variability is a universal property of radio-quiet AGN which has been studied for several decades, but its origin is still unclear.  The optical continuum emission almost certainly originates from the accretion disc \\citep{koratkarblaes99}, so it is natural to assume that disc variability (e.g. through accretion rate fluctuations) drives the optical variability.  Models where the variability is {\\it intrinsic} to the disc, however, suffer from several problems.  Firstly, optical variability is seen on time-scales as short as a day, whereas accretion variability time-scales in standard accretion discs \\citep{shakura} should be long in the optically-emitting regions, comparable to the viscous time-scale which can be of the order of a year or more \\citep{trevesetal88}.  Secondly, variations are well-correlated in different optical bands, with only short (days) delays between bands, in the sense that longer-wavelength variations lag shorter-wavelength variations \\citep{wandersetal97,collieretal01,cackettetal07}.  Since local disc temperature $T$ should decrease with radius $R$ as $T\\propto R^{-3/4}$, shorter wavelengths originate from smaller radii. Therefore, in the case of inward-propagating fluctuations the long-wavelength variations lead, rather than lag, the short-wavelength variations. Also, the resulting delays would be large, comparable to the radial drift time-scale, contrary to the observed day-scale delays.  To account for the observed optical properties, \\citet{kroliketal91} suggested that the variability is driven by changes of the X-ray continuum which illuminates and thereby heats the disc, causing optical continuum variations.  Since the X-ray emission is likely to be centrally concentrated, blue bands should respond to the X-ray variations first, followed by red, with short delays corresponding to the light-travel time delay between the blue and red-emitting parts of the disc.  The predicted optical time-delay $\\tau$ scales with wavelength $\\lambda$ as $\\tau \\propto \\lambda^{4/3}$ \\citep{collieretal99}, and the observed dependencies of optical continuum lags on wavelength are consistent with this prediction \\citep{cackettetal07}.  A more stringent test of reprocessing models is to directly compare X-ray and optical variations, using multi-wavelength monitoring campaigns.  To date, a handful of these campaigns have been carried out, largely facilitated by the {\\it Rossi X-ray Timing Explorer} ({\\it RXTE}) in conjunction with various ground based observatories. The results have been mixed, with some AGN showing good evidence for correlated variability, as would be expected from reprocessing models (e.g.  NGC~4051, \\citealt{petersonetal00,shemmeretal03}; NGC~5548, \\citealt{uttleyetal03}; Mrk~509, \\citealt{marshalletal08}), while others show more complicated but possibly correlated behaviour (e.g. NGC~7469, \\citealt{nandraetal98,nandraetal00}) and one AGN shows no apparent correlation at all (NGC~3516, \\citealt{maozetal02}).  The best case for correlated variability to date is NGC~5548, but \\citet{uttleyetal03} note that it is difficult to reconcile the large fractional amplitude of optical variability on time-scales of months-years with reprocessing models, since the X-ray variability shows smaller-amplitude variations than the optical on these long time-scales.  One would expect smaller relative variations in optical since the intrinsic emission due to viscous disc heating would dilute any variable reprocessed component.  Furthermore, Gaskell (2007) notes that a simple energetics argument rules out reprocessing as the dominant source of optical variability in many AGN (including NGC~5548) which have `big blue bumps' dominating the total luminosity, significantly exceeding the X-ray luminosity available for reprocessing.  Detailed reprocessing calculations, produced by \\citet{reprocessing} for the case of NGC~3516 show that simple reprocessing of X-rays on their own could not produce the observed optical variability. Also, in some AGN with simultaneous X-ray and optical monitoring, there is evidence that the optical may {\\it lead} the X-ray variations, which might be expected if at least some of the optical variability is produced by intrinsic accretion fluctuations propagating through the disc \\citep{shemmeretal03,marshalletal08}.  Clearly, no single model provides a satisfactory explanation for all the data.  It is possible, however, that some combination of reprocessing and intrinsic accretion variations may explain the diverse range of optical/X-ray behaviour which is already observed in only a small sample of AGN.  The location of the optical emitting region probably plays a key role in determining the balance of intrinsic versus reprocessed variability.  It is governed by the disc temperature, which scales inversely with radius $R$ (in units of the Gravitational radius $R_g=GM/c^2$) but also scales with black hole mass $M_{\\rm BH}$ and accretion rate (as a fraction of the Eddington rate) $\\dot{m}$ as $T\\propto[(\\dot{m}/{M}_{\\rm BH})R^{-3}]^{1/4}$ \\citep{trevesetal88}.  Thus over the 3 decade range in black hole mass expected for AGN, we expect the radius corresponding to a given temperature to change by a factor of 10 (at the same fractional accretion rate), perhaps even more for different accretion rates. Since disc temperature governs the radius where peak optical emission is produced, \\citet{uttleyetal03} suggested that a diverse range of optical/X-ray behaviour might result from the range of masses and accretion rates expected in AGN \\citep[see also][]{li}.  For example, the AGN with the most massive black holes will have cool discs and very centrally concentrated optical-emitting regions, in terms of gravitational radii. Therefore, the disc variability time-scales will be short, comparable with the X-ray variability time-scales, so intrinsic disc variability as well as reprocessing may contribute significantly to the optical variations.  In contrast, AGN with lower-mass black holes will have hotter discs so optical emission will originate from larger radii, where disc variability time-scales are very long compared to the corresponding time-scales in the innermost regions, so that only reprocessing may contribute significantly to the rapid variability we observe.  To test composite models of optical variability, it is necessary to carry out combined optical/X-ray monitoring of AGN with widely varying masses and accretion rates.  To date, these campaigns have focused on Seyfert galaxies, which typically cover bolometric luminosities ranging from $10^{43}$ to $10^{45}$ erg~s$^{-1}$.  Here we report the first optical/X-ray correlation for a more luminous radio-quiet AGN, the quasar \\mr\\, which we have monitored simultaneously in X-ray and optical bands for the past 2.5 years.  We describe the data in Section~2 and show the cross correlations in Section~3. In Section~4 we explore reprocessing and propagating fluctuations scenarios to explain the X-ray/optical variability in this source and discuss the implications of our results in Section~5. ",
        "conclusions": "We have carried out a simultaneous monitoring campaign in X-ray and optical B, V and R bands on the quasar \\mr\\ over two and a half years. All bands are significantly variable and their fluctuations are well correlated. The cross correlation functions show peaks around a lag of zero days, significant over the 99\\% confidence level when compared to uncorrelated simulated data. All the delays between the bands are consistent with 0 days. We used daily sampled light curves in X-ray, B and V to constrain the lags between these bands, obtaining values of $0.6 \\pm 3.1$ days between X-ray and B, $1.3 \\pm 3.3$ day for X-ray vs V and $0.5 \\pm 2.4$ days between B and V lags, where positive values indicate higher energy band leading. The lag between X-ray and R band was determined using the long but more sparsely sampled light curve, obtaining a value of -4.5$\\pm$16.8 days.  The long term trends, over hundreds of days, are well matched by all the bands observed and are stronger in the B band than in X-rays, while the other optical bands show smaller-amplitude trends. The optical light curves are not corrected by host galaxy contribution, however, so the decreased amplitude of variability in lower energy bands might be partly due to a constant galactic contribution. On short time-scales, of tens of days, there are large X-ray flares which appear strongly attenuated in the optical bands. The intensively sampled light curves show that the rapid X-ray flares do have optical counterparts, but that these are very weak.  Pure reprocessing of X-rays cannot produce both the short and long time-scale optical variability: if the long term optical fluctuations were produced by reprocessing, we would expect similarly large rapid optical fluctuations, unless these could be smoothed by light travel effects. Any geometry of the reprocessor, however, cannot smooth out the fluctuations on time-scales of 20--50 days without introducing time-lags of a similar length, which are not observed. A similar conclusion was reached by \\citet{reprocessing} for the Seyfert galaxy NGC~3516, where the X-ray-UV variability could not be explained by simple reprocessing. We show that reprocessing of the observed X-rays cannot reproduce the optical variability in \\mr, either producing rapid fluctuations that were too large or long term trends that were too small. The short term fluctuations however could be reproduced reasonably well with the model if we limited the light curve lengths to $\\sim 100$ days, which reduces the contribution of long term trends to each light curve segment. Therefore, a moderate amount of reprocessing acting on an already long-term variable optical emission can reproduce the observed behaviour.  A simple way to decouple the amplitude of short and long term fluctuations is to have two distinct processes producing the optical variability, e.g.\\ accretion rate fluctuations plus reprocessing. The covering fraction of the disc, together with the ratio of intrinsic disc emission to X-ray heating determine the fraction of optical flux arising from reprocessed X-rays. Therefore, the size of the rapid optical fluctuations depends on these parameters and can be largely independent of the amplitude of the long term trends. To test this possibility, we generated simulated light curves and compared their statistical properties with the observed data. In the model, fluctuations propagate inward through the accretion flow, modulating the optical emission as they travel through the thermally emitting region and finally modulate the X-ray emission, assumed to be located in the centre. This process leads to the long term correlated variability. As the fluctuations are produced at every radius in the flow, on a radially dependent time-scale, the centrally emitted X-rays contain short time-scale fluctuations which are originally absent in the optical light curves.  The rapid fluctuations, only present in the X-ray emitting region are then introduced into the optical light curves through thermal reprocessing.  Lastly, although propagation times across the accretion disc can be long and produce a large lag between optical and X-ray fluctuations, reprocessing reduces the lag and can even revert it to a value of the order of the light crossing time to the reprocessor. For a black hole mass of $10^9 M_\\odot$ the light crossing time of one gravitational radius is 5000 s, so even for a large source height and inner truncation radius of 50 $R_g$ the crossing time would be 4 days, less for a less massive black hole, which is within the observed limits on the lag. We show using the simulated light curves, that small-amplitude rapid fluctuations in the optical band can shift the CCF peak towards a lag of zero days, even when the long term trends of the optical lead the X-rays by $\\sim$50 days.  As inward propagation of the fluctuations and reprocessing have opposite effects on the time lag, the relative amount of optical/UV variability produced by each process determines the sign of the lag. This might be the case in other monitored AGN where the optical/UV can either lead or lag the X-rays.  An interesting feature of our fits to the optical variability is that we expect the disc to subtend a small solid angle with respect to the X-ray source, in order to reduce the effects of X-ray reprocessing and produce the correct amplitude of optical rapid fluctuations.  This small solid angle is consistent with the small amount of reflection inferred from the broadband X-ray spectrum of MR~2251-178 observed by {\\sl BeppoSAX} \\citep{Orr}, which shows only a weak iron line with an equivalent width of $EW\\sim 70$~eV, and modest reflection continuum (reflected fraction $<0.4$).  It is therefore possible that MR~2251-178 represents a source in a state where the thin disc is slightly truncated and/or the X-ray emitting region is at a small height over the disc.  Optical/X-ray studies of other AGN should provide further useful constraints on their disc/corona geometries."
    },
    "0807/0807.3457_arXiv.txt": {
        "abstract": "We use time-evolutions of the linear perturbation equations to study the oscillations of rapidly rotating neutrons stars. Our models account for the buoyancy due to composition gradients and we study, for the first time, the nature of the resultant g-modes in a fast spinning star.  We provide detailed comparisons of non-stratified and stratified models. This leads to an improved understanding of the relationship between the inertial modes of a non-stratified star and the g-modes of a stratified system. In particular, we demonstrate that each g-mode becomes rotation-dominated, i.e. approaches a particular inertial mode, as the rotation rate of the star is increased. We also discuss issues relating to the gravitational-wave driven instability of the various classes of oscillation modes. ",
        "introduction": "The oscillations of rotating neutron stars are of great interest in astrophysics. Recent efforts to understand the various pulsation modes of compact stars have to a large extent been motivated by gravitational-wave astronomy~\\citep{1996PhRvL..77.4134A, 1998MNRAS.299.1059A}.  Different modes of oscillation depend on different pieces of internal physics and one may hope to use observations to probe, for example, the composition of the high-density region~\\citep{2008GReGr..40..945F,2007GReGr..39.1323B}. Work in this area, which has a long history \\citep[for a review see e.g.,][]{lrr-1999-2}, gained further momentum recently with the relatively successful matching between observed quasiperiodic oscillations in the tails of magnetar flares and calculated torsional oscillations of the neutron star crust~\\citep{1998ApJ...498L..45D, 2005ApJ...634L.153P, 2007AdSpR..40.1446W,2007MNRAS.374..256S}.  The excitement following these, the first ever, observations of likely neutron star vibrations is significant. It is clear that our models need to be improved significantly if we expect to carry out a quantitative astero-seismology programme for neutron stars, but the prospects for improvements look good. After all, we already have a reasonable understanding of the dynamics of the crust region \\citep[see][]{2006CQGra..23.5367C,2007MNRAS.374..256S} as well as the superfluid components in the core of the star~\\citep{lrr-2007-1}.  The present work is motivated by the need to consider more realistic models of rapidly rotating stars.  Our emphasis will be on the role of the internal stratification associated with composition variations. This leads to the presence of the so-called g-modes~\\citep{1992ApJ...395..240R}, and we want to investigate how these modes are affected by fast rotation (the analogous problem for thermal ocean g-modes of rapidly rotating neutron stars has already been considered, see \\citet*{1996ApJ...460..827B}).  In order to determine the rotational effects on the oscillation spectrum we study linear perturbations of axisymmetric stellar configurations, where the Coriolis force and the centrifugal flattening of the star are included in the background.  In this way, we go beyond the so-called slow rotation approximation, where the rotation itself is treated perturbatively and the mode-frequencies of non-rotating models are determined by a perturbation expansion with respect to the angular velocity of the star $\\Omega$. We are particularly interested in the relationship between the g-modes of a stratified stellar model and the inertial modes~\\citep{1999ApJ...521..764L}, for which the Coriolis force is the main restoring agent.  The inertial modes are important since they may be driven unstable by the emission of gravitational waves. The strongest instability is associated with the so-called r-modes~\\citep{2001IJMPD..10..381A,Andersson:2002ch}.  A key issue concerns the amplitude at which an unstable mode saturates. Detailed work has shown that an unstable r-mode saturates due to nonlinear couplings to other inertial modes~\\citep{2003ApJ...591.1129A}. However, the mode saturation has only been considered for slowly rotating, non-stratified models and it is important to ask to what extent these results will change if the models are made more realistic. In stratified neutron stars, the g-modes tend to become rotationally dominated beyond a certain rotation rate~\\citep[see e.g.,][]{2000ApJS..129..353Y, 2000A&A...354...86D} and may then be called inertia-gravity modes~\\citep{1989nos..book.....U}.  Effects due to the stratification are thus weakened, which indicates that the r-mode saturation estimates may still be reliable, at least for the fastest spinning systems.   Improving and extending a numerical code that has been used to study non-stratified stars~\\citep*{2002MNRAS.334..933J}, we approach the problem via time-evolutions of the equations that govern the linear perturbations of a, potentially rapidly, rotating Newtonian star. This strategy has the advantage that one does not have to deal with the many rotationally coupled multipoles, that tend to make a slow-rotation calculation less tractable.  On the other hand, a numerical evolution is not expected to have the precision of a frequency domain calculation. Neither should one expect it to yield the complete spectrum of modes. After all, the simulation results depend on the chosen initial data. In absence of a clear idea of the nature of the various modes one does not know how to excite specific oscillations. One way around this problem is to follow the ``recycling'' strategy developed by~\\citet*{Stergioulas:2003ep} and \\citet*{Dimmelmeier:2005zk}. Another, more pragmatic, approach is to simply study the modes that are excited by ``generic'' initial data. This is the attitude that we adopt here. ",
        "conclusions": "} \\label{sec:concl}  Information concerning the oscillation spectra of realistic neutron star models helps develop our understanding of the physics required to describe these objects. Any spectral property can in principle be attributed to particular physical quantities or configurations of the star. Astero-seismology studies, using either electromagnetic or gravitational signals (or indeed both) may thus help constrain the state of matter at supernuclear densities.  Of course, in order to facilitate this, we need to improve our theoretical models and clarify the origin of various stellar pulsation features.  In this paper, we have studied the pulsations of stratified and rapidly rotating neutron stars with the aim of understanding the dependence of the composition g-modes on the rotation rate of the star.  The stellar pulsations were studied using the linearised Euler and conservation equations on an axisymmetric background. In order to simplify the problem, we used the Cowling approximation where the gravitational potential perturbations are neglected. This approximation generates only a small error in the g-mode frequencies. The system of perturbation equations was evolved in time by a 2D numerical code based on a standard finite differencing scheme. This code was tested against various results available in literature and demonstrated good accuracy.  Since both Coriolis and buoyancy forces act on a perturbed fluid element of a rotating and stratified star, the low-frequency modes have a mixed inertia-gravity character. They typically behave as g-modes in the slowly rotating limit, while they assume the properties of the inertial modes when the star rotates rapidly.  By comparing the oscillation frequencies and the associated eigenfunctions of barotropic inertial modes and the inertia-gravity modes of non-barotropic stars, we have demonstrated how the two sets of modes are associated at fast rotation rates.  Our results show that it would be difficult to deduce the presence of composition gradients from the inertia-gravity mode spectrum of fast rotating stars.  However, as the neutron star ages and spins down, the dependence of the inertia-gravity modes on the star's rotation and their deviation from the inertial barotropic modes can, at least in principle, be used to estimate the Brunt-V\\\"{a}is\\\"{a}l\\\"{a} frequency and the degree of stratification.  As an extension of this work, we are currently developing a numerical code for studying the dynamics of rapidly rotating superfluid neutron stars. This is an important development since all mature neutron stars are expected to have superfluid components in the core. It is also well known that the associated multi-fluid dynamics leaves an imprint on the stellar oscillation spectrum, see \\citet{2007arXiv0709.0660L} for a recent discussion.  We hope to be able to report on the initial results of this investigation soon."
    },
    "0807/0807.4870_arXiv.txt": {
        "abstract": "To first approximation, a binary system conserves its angular momentum while it evolves to its state of minimum kinetic energy: circular orbit, all spins aligned, and components rotating in synchronism with the orbital motion. The pace at which this final state is achieved depends on the physical processes that are responsible for the dissipation of the tidal kinetic energy. For stars (or planets) with an outer convection zone, the dominant mechanism identified so far is the viscous dissipation acting on the equilibrium tide. For stars with an outer radiation zone, it is the radiative damping operating on the dynamical tide.  After a brief presentation of the tides, I shall review these physical processes; I shall discuss the uncertainties of their present treatment, describe the latest developments, and compare the theoretical predictions with the observed properties concerning the orbital circularization of close binaries. ",
        "introduction": "A fundamental property of closed mechanical systems is that they conserve their total momentum. This is true in particular for binary stars, star-planet(s) systems, whether they possess or not a circumstellar disc, if one can ignore the angular momentum that is carried away by winds and by gravitational waves. Through tidal interaction, kinetic energy and angular momentum are exchanged between the rotation of the components  their orbital motion and the disc. In the absence of such a disc, which is the case that we shall consider here, they evolve due to viscous and radiative dissipation to the state of minimum kinetic energy, in which the orbit is circular, the rotation of both stars is synchronized with the orbital motion, and their spin axis are perpendicular to the orbital plane. How rapidly the system tends to that state is determined by the strength of the tidal interaction, and thus by the separation of the two components: the closer the system, the faster its dynamical evolution. But it also depends on the efficiency of the physical processes that are responsible for the dissipation of the kinetic energy.  Provided these dissipation processes are well enough understood, the observed properties of a binary system can deliver important information on its evolutionary state, on its past history, and even on the conditions of its formation. The first step is thus to identify these physical processes, and it is surprising that this has not been seriously undertaken until the mid-sixties, while tidal theory as such had already reached a high degree of sophistication, starting with the pioneering work of Darwin (1879).  In his classical treatise, Kopal (1959) states from the onset that he is interested only in ``dynamical phenomena which are likely to manifest observable consequences in time intervals of the order of 10 or 100 years, and if so, tidal friction can be safely ignored\".  But stars live much longer than that, and this is why we shall consider here changes in the properties of binary systems that occur over their evolutionary time scale, and in particular the circularization of their orbits, which is both easy to observe and easy to interpret.  We shall deal mainly with binary stars, although much of what follows may be applied also to star-planet systems. ",
        "conclusions": "Let me summarize. The two tidal dissipation processes that have received most attention so far are the turbulent friction acting on the equilibrium tide, which was first described in the 60's (Zahn 1966b), and the radiative damping of the dynamical tide,  identified in the 70's (Zahn 1975). These processes operate  respectively in convection and radiation zones, and they have been quite successful in explaining the observed orbital circularization of binary stars. This is particularly true for the early-type MS binaries, for which we have now at our disposal very large samples gathered during the OGGLE and MACHO campaigns: their transition period is precisely defined and it agrees extremely well with that predicted by the theory of the dynamical tide, which is thus validated. However many of these binaries are circularized well above this transition period, as if they had experienced another, more efficient tidal dissipation mechanism. A very likely explanation for this behavior is that these binaries have undergone several episodes of resonance locking, as was described by Witte and Savonije (1999a, 1999b).  On the other hand, the equilibrium tide damped by turbulent dissipation accounts very well for the properties of binaries containing a red giant, as was demonstrated by Verbunt \\& Phinney (1995). It also explains the transition period of about 8 days observed in late-type binaries that are younger than about 1 Gyr: the explanation is that these have been circularized during the PMS phase, when they were much more voluminous and fully convective. The only serious discrepancy today seems to be the behavior of late-type main-sequence binaries older than 1~Gyr, whose transition period increases with age and is higher than that predicted when applying straightforward the theory of the equilibrium tide. Here again one may invoke the dynamical tide with resonance locking in the radiative core of these stars, as was shown by Witte and Savonije (2002).  Their mechanism appears thus highly promising, and it ought to be further explored. For instance, one should take into account that the tidal torque is applied primarily  to specific regions: the outer convection zone in late-type MS stars and the outermost part of the radiation zone in early-type stars.  These regions are synchronized more quickly then the rest of the star, and therefore differential rotation develops in their radiation zone. This increases the radiative damping, since the local tidal frequency tends then to zero as the tidal wave approaches the synchronized region, as we explained in \\S\\ref{grav-exc}.  For late-type binaries, a highly interesting alternative is offered by the damping of inertial waves in their convective envelope, which is being explored by Ogilvie and Lin (2007). This process is likely to play an important role also in giant planets (Ogilvie \\& Lin 2004). The difficulty in studying these waves is that they require highly resolved 2D numerical calculations, since the so-called traditional approximation is no longer applicable to render the problem separable.  Work is in progress on several other points, and I shall quote only a few.  Kumar and Goodman (1996) have studied the enhanced damping of the oscillations triggered in tidal-capture binaries, due to non-linear coupling between the eigenmodes, which is extremely strong in such highly eccentric orbits. Rieutord (2004, and in this volume) is examining the possibility that the so-called elliptic instability may occur in binary stars; this instability is observed in the laboratory when fluid is forced to rotate between boundaries that have a slight ellipticity, and it leads to turbulence. Even the equilibrium tide is being revisited, taking into account the differential rotation of the convection zone (Mathis \\& Zahn, in preparation).  Finally, it remains to explain why I made no attempt here to reconcile the theoretical predictions for the synchronization of the binary components with their observed surface rotation. The reason is that in most cases the tidal torque is applied mainly to the outermost part of the star, which is synchronized much more rapidly than the interior; therefore the interpretation of the surface rotation requires to model the transport of angular momentum within the star, and in particular where it proceeds slowest, i.e. in the radiation zones. This is a difficult task that only now begins  to be undertaken seriously: for recent accounts on this problem, see the reviews by Talon (2007) and Zahn (2007). I am confident that we will see much progress in solving this problem when the next school will be held on that theme, hopefully in a not too distant future!  \\bigskip"
    },
    "0807/0807.1452_arXiv.txt": {
        "abstract": "{The precision of radial velocity (RV) measurements to detect indirectly planetary companions of nearby stars has improved to enable the discovery of extrasolar planets in the Neptune and Super-Earth mass range. Detections of extremely low mass planets, even as small as 1 Earth mass or below, in short-period orbits now appears conceivable in ongoing RV planet searches. Discoveries of these Earth-like planets by means of ground-based RV programs will help to determine the parameter $\\eta_{\\oplus}$, the frequency of potentially habitable planets around other stars.} {In search of low-mass planetary companions we monitored Proxima Centauri (M5V) as part of our M dwarf program. In the absence of a significant detection, we use these data to demonstrate the general capability of the RV method in finding terrestrial planets. For late M dwarfs the classic liquid surface water habitable zone (HZ) is located close to the star, in which circumstances the RV method is most effective. We want to demonstrate that late M dwarfs are ideal targets for the search of terrestrial planets with the RV technique.} {Using the iodine cell technique we obtained differential RV measurements of Proxima Cen over a time span of 7 years with the UVES spectrograph at the ESO VLT. We determine upper limits to the masses of companions in circular orbits by means of numerical simulations.} {The RV data of Proxima Cen have a total rms scatter of $3.1~{\\rm m\\,s}^{-1}$ and a period search does not reveal any significant signals. In contrast to our earlier results for Barnard's star, the RV results for the active M dwarf Proxima Cen are only weakly correlated with H$_{\\alpha}$ line index measurements. As a result of our companion limit calculations, we find that we successfully recover all test signals with RV amplitudes corresponding to planets with $m \\sin i \\ge 2 - 3~M_\\oplus$ residing inside the HZ of Proxima Cen with a statistical significance of $>99\\%$. Over the same period range, we can recover 50\\% of the test planets with  masses of $m \\sin i \\ge 1.5 - 2.5~M_\\oplus$. Based on our simulations, we exclude the presence of any planet in a circular orbit with $m \\sin i \\ge 1~M_{\\rm Neptune}$ at separations of $a \\le 1$~AU.} {} ",
        "introduction": "Over the next decades we will be able to derive a first estimate of the frequency of stars with a potentially habitable Earth-like planet. This frequency is usually denoted by the parameter $\\eta_{\\oplus}$. CoRoT and Kepler are two space mission that have the capability to detect ``Super-Earths'' and even Earth analogs in short-period orbits (CoRoT) and at 1~AU (Kepler) using the transit method. The astrometry mission SIM Planetquest will achieve the sensitivity to detect Earth-like planets around a sample of nearby stars. Also, ground-based Doppler measurements have attained precision levels that make discoveries of planets with a few Earth masses possible (e.g. Lovis et al.~2006). Already the very first extrasolar planets found around the pulsar PSR 1257+12 (Wolszczan \\& Frail~1992), have such small masses, that they qualify as terrestrial planets. However, their formation is likely to have followed a different path than what is currently envisioned for the formation of terrestrial planets around main sequence stars. Besides its obvious astrobiological implications, and its crucial role for the design and overall costs of any future TPF/Darwin-type mission, the value of $\\eta_{\\oplus}$ will also impact our understanding of the formation of terrestrial planets in general (similar to the effect the extrasolar giant planets have on giant planet formation models.)  M dwarfs comprise the majority of stars in the solar neighborhood (e.g. Reid et al.~2004) and once their intrinsic faintness is overcome, they represent attractive targets for high precision Doppler surveys. Due to their lower masses the reflex motion of a planet of a given mass is higher than for a solar mass star. Thus, it is no surprise that so far the lowest mass extrasolar planets detected by the Doppler method are all orbiting M dwarfs (Rivera et al.~2005; Udry et al.~2007). The micro-lensing event reported by Beaulieu et al.~(2006) is also attributed to lensing by a very low-mass planetary companion to, most likely, an M dwarf host.  In this paper we present 7 years of high precision RV data for our closest neighbor in space: the M5V star Proxima Centauri. We demonstrate that with the data in hand we could have already detected planets with minimum masses as small as $1 - 2~{\\rm M}_{\\oplus}$. Constraints for giant planetary companions to Proxima Cen using HST Fine Guidance Sensor astrometry were reported by Benedict et al.~(1999) and the companion limits based on less precise RV data from the ESO Coud\\'e Echelle Spectrometer planet search were presented by K\\\"urster et al.~(1999). The combination of both studies already excluded all companions with (minimum) masses higher than $0.8~{\\rm M}_{\\rm Jupiter}$ for the period range $1$ to $600$ days. ",
        "conclusions": "\\begin{enumerate} \\item We present 7 years of high precision RV data for our closest neighbor in space, the M5V star Proxima Cen, obtained  with UVES + I$_2$ cell at the ESO VLT/UT2. We detect no significant periodicities (except close to the 1-yr peak in the window function) that can be attributed to orbiting companions. \\item Using the same set of spectra we measure an H$_{\\alpha}$ line index to estimate the magnetic activity level of Proxima Cen. These line indices show a large amount of scatter due to flaring activity and are only weakly correlated with the RV results. \\item Based on numerical simulations we demonstrate that we could have already detected all planets with $m \\sin i = 2 - 3~{\\rm M}_\\oplus$ on circular orbits inside the classic habitable zone of Proxima Cen (assuming a stellar mass of $0.12~{\\rm M}_{\\oplus}$).  \\end{enumerate}"
    },
    "0807/0807.3985_arXiv.txt": {
        "abstract": "We present evidence for anomalous microwave emission in the RCW175 \\hii~region. Motivated by 33~GHz $13\\arcmin$ resolution data from the Very Small Array (VSA), we observed RCW175 at 31~GHz with the Cosmic Background Imager (CBI) at a resolution of $4\\arcmin$. The region consists of two distinct components, G29.0-0.6 and G29.1-0.7, which are detected at high signal-to-noise ratio. The integrated flux density is $5.97\\pm0.30$~Jy at 31~GHz, in good agreement with the VSA. The 31~GHz flux density is $3.28\\pm0.38$~Jy ($8.6\\sigma$) above the expected value from optically thin free-free emission based on lower frequency radio data and thermal dust constrained by IRAS and WMAP data. Conventional emission mechanisms such as optically thick emission from ultracompact \\hii~regions cannot easily account for this excess. We interpret the excess as evidence for electric dipole emission from small spinning dust grains, which does provide an adequate fit to the data. ",
        "introduction": "In recent years there has been mounting observational evidence for a new diffuse component emitting at frequencies $\\approx 10-60$~GHz. The anomalous microwave emission was first detected at 14 and 32~GHz by \\cite{Leitch97}. Since then, a similar picture has emerged both at high latitudes \\citep{Banday03,deOliveira-Costa04,Fernandez-Cerezo06,Hildebrandt07,Bonaldi07} and from individual Galactic sources \\citep{Casassus04,Casassus06,Casassus07,Watson05,Scaife07,Dickinson07}, although negative detections have also been reported \\citep{Dickinson06,Scaife08}.  The spectral index between 20 and 40~GHz is $\\alpha \\approx -1.1$ \\citep{Davies06} with some evidence of flattening at $\\sim10-15$~GHz \\citep{Leitch97,deOliveira-Costa04,Hildebrandt07}. The emission appears to be very closely correlated with far-IR data suggesting a dust origin. Various emission mechanisms have been suggested, including hot ($T\\sim10^{6}$~K) free-free \\citep{Leitch97}, flat spectrum synchrotron \\citep{Bennett03b}, spinning dust \\citep{Draine98a,Draine98b} and magnetic dust \\citep{Draine99}. The overall picture is still very unclear and new data covering the range $10-60$~GHz are urgently needed.  RCW175 \\citep{Rodgers60} is a diffuse \\hii~region, which consists of a ``medium brightness'' optical filament (G29.1-0.7, S65) $\\sim 7\\arcmin \\times 5\\arcmin$ in extent, and a nearby compact source (G29.0-0.6), which is heavily obscured by dust. Although the filament is clearly seen in high resolution data, the compact counterpart is considerably brighter. The ionization is thought to be provided by a single B-1 II type star, which forms part of a 5-star cluster \\citep{Forbes89,Sharpless59} at a distance of 3.6~kpc.  Observations made with the Very Small Array (VSA) at 33~GHz \\citep{Watson03,Dickinson04} as part of a Galactic plane survey (Todorovi\\'c et al., in prep.) indicate that RCW175 is anomalously bright by a factor of $\\approx 2$, when compared with lower frequency data. In this {\\it Letter}, we present accurate Cosmic Background Imager (CBI) $31$~GHz observations of RCW175 and make a comparison with ancillary radio/FIR data. We find that the emission at $31$~GHz is significantly above what is expected from a simple model of free-free and vibrational dust emissions. ",
        "conclusions": "\\label{sec:conclusions}  Using data from the VSA and CBI, we have observed excess emission at $\\sim 30$~GHz from the \\hii~region RCW175. The flux density spectrum indicates that about half of the flux in this region is from optically thin free-free emission leaving about half unaccounted for. An upper limit at 94~GHz from WMAP data constrains the contribution of thermal dust emission. We have discarded optically thick free-free emission from UC~\\hii~regions and GPS sources as possible candidates for this excess; an upper limit at 94~GHz and high resolution radio data rule these out as the primary contributor at $\\sim 30$~GHz. We interpret the excess as electric dipole radiation from small rapidly spinning dust grains as predicted by \\cite{Draine98b}. These models provide a reasonable fit to the data that is consistent both in terms of spectral shape and emissivity. High resolution, multi-frequency data in the range $10-100$~GHz are needed to confirm this result and to investigate the nature of anomalous emission."
    },
    "0807/0807.3382_arXiv.txt": {
        "abstract": "In this letter, we determine the $\\kappa$-distribution function for a gas in the presence of an external field of force described by a potential U(${\\bf r}$). In the case of a dilute gas, we show that the $\\kappa$-power law distribution including the potential energy factor term can rigorously be deduced in the framework of kinetic theory with basis on the Vlasov equation. Such a result is significant as a preliminary to the discussion on the role of long range interactions in the Kaniadakis thermostatistics and the underlying kinetic theory. ",
        "introduction": "In this section we discuss briefly the standard case, i.e., the kinetic description of a classical gas under stationary conditions and immersed in a conservative force field, ${\\bf F}=-\\nabla U(r)$. Typical examples are a gas in the earth's gravitational field and ions in an external magnetic field \\cite{Huang,KT}. This kind of problem is important on their own because it permits to understand how the molecular motion is modified by force-fields different from those exerted by the containing vessel or even by the other particles of the gas. As widely known, its distribution function differs from the Maxwellian velocity statistics trough an extra exponential factor involving the potential energy whose general form reads \\begin{equation}\\label{e1} f({\\bf r},v)=n_0\\left({m\\over 2\\pi k_B T} \\right)^{3/2}\\exp \\left(-{{1\\over 2}m v^2 +U({\\bf r})\\over k_B T}\\right), \\end{equation} where $m$ is the mass of the particles, $T$ is the temperature and $n_0$ is the particle number density in the absence of the external force field. In addition, since this distribution function is normalized, it is easy to see that the number density is given by \\begin{equation}\\label{e2} n({\\bf r})=n_0 \\exp \\left[-{{U({\\bf r}) \\over k_B T}}\\right], \\end{equation} where the factor, $\\exp[-{U({\\bf r})/k_B T}]$, which is responsible for the inhomogeneity of $f({\\bf r},v)$, is usually called the Boltzmann factor. Expression ($\\ref{e1}$) follows naturally from an integration of the collisionless Boltzmann's equation \\begin{equation}\\label{e3} \\frac{\\partial{\\it f}}{\\partial{\\it t}} + {\\bf v}\\cdot\\frac{\\partial{\\it f}}{\\partial{\\bf r}} + \\frac{\\bf F}{m}\\cdot\\frac{\\partial{\\it f}}{\\partial{\\bf v}} = 0, \\end{equation} when stationary conditions are adopted along with the assumption that the total distribution can be factored \\begin{equation}\\label{e4} f({\\bf r},v)=f_0(v)\\chi({\\bf r}), \\end{equation} where $f_0(v)$ represents the Maxwell equilibrium distribution function, and $\\chi({\\bf r})$ is a scalar function of ${\\bf r}$. As one may show, after a simple normalization, the resulting expression for $\\chi({\\bf r})$ is exactly the Boltzmann factor for the potential energy of the external field, namely:  \\begin{equation}\\label{e2a} \\chi({\\bf r})=\\exp \\left[-{{U({\\bf r}) \\over k_B T}}\\right], \\end{equation} and combining this result with equation (\\ref{e4}) the Boltzmann stationary distribution (\\ref{e1}) is readily obtained. ",
        "conclusions": "In the last few years, several applications of the $\\kappa$-power law velocity distribution have been done in many disparate branches of physics [8-21]. However, such investigations are usually related with the $\\kappa$-velocity distribution function as given by equation (6).  On the other hand, many physical systems involve naturally the presence of a conservative force field as happens for example with ions in a magnetic field. Probably, the most popular problem of a gas in a force-field is the planetary atmosphere. In the standard simplified treatment, the temperature is uniform and the tree-dimensional motion occurs under the action of a constant gravitational field along the $z$-direction. To all this sort of problems, the extended $\\kappa$-distribution deduced here with basis on the Vlasov equation, namely  \\begin{eqnarray*}\\label{eq18} f_{\\kappa}({\\bf r},v) = A_\\kappa\\left[\\sqrt{1 +\\kappa^2\\left(-\\frac{m{\\bf v}^2}{2k_BT} - \\frac{U({\\bf r})}{k_BT}\\right)^2} + \\kappa \\left(-\\frac{m{\\bf v}^2}{2k_BT} - \\frac{U({\\bf r})}{k_BT}\\right) \\right]^{1/\\kappa}, \\end{eqnarray*} can be applied, and, might prove to be of extreme wide usefulness. Note also that a giroscopic term may also be added to the above power law distribution. In a rotating frame  with constant angular velocity, the whole effect is just to add a Coriolis term $-1/2m\\omega^{2}R^{2}$ to the potential energy, where $R$ is the perpendicular distance from the axis of rotation. In the Newtonian framework, such a term simulates a change in the potential energy due to gravity. Finally, it is worth mentioning that the present consistency between Vlasov equation and Kaniadakis thermostatistics also is valid in the context of Tsallis nonextensive framework \\cite{1}.    \\vspace{0.5cm}  \\noindent {\\bf Acknowledgments:} The authors are partially supported by the Conselho Nacional de Desenvolvimento Cient\\'{\\i}fico e Tecnol\\'ogico (CNPq - Brazil). JASL is also grateful to FAPESP No. 04/13668-0."
    },
    "0807/0807.1722_arXiv.txt": {
        "abstract": "% We study  the stability and the modes of non -- isothermal coronal loop models with  different intensity values of the equilibrium twisted magnetic field.We use an energy principle obtained via non -- equilibrium thermodynamic arguments. The principle is expressed in terms of Hermitian operators and allows to consider together the coupled system of equations: the balance of energy equation and the equation of motion, to obtain modes and eigenmodes in a spectrum ranging from  short to long--wavelength disturbances without having to  use weak varying  approximations of the equilibrium parameters. Long--wavelength perturbations introduce additional difficulties because the inhomogeneous nature of the medium determines  disturbances leading to continuous intervals of eigenfrequencies which cannot be considered as purely sinusoidal.We analyze the modification of periods, modes structure and stability when the helicity, the magnetic field strength and the radius of the fluxtube are varied.  The efficiency of the damping due to the resonant absorption mechanism is analyzed in a context of modes that can either impulsively release or storage magnetic energy.We find that the onset of the instability is associated to a critical value of the helicity and that the  magnetic energy content has a  determinant role on the instability of the system with respect to the stabilizing effect of the resonant  absorption mechanism. ",
        "introduction": "\\subsection{Variational principle}  A crucial requirement for any theoretical model of coronal structures is to give account of   the stability and evolution of   far--from--equilibrium states which are responsible of the characteristic rich topology and dynamics of the solar corona. This implies  to consider the coupling of thermal and mechanical  equations. Different stability analysis of solar structures can be found in the literature, generally restricted to special types of perturbations and specific equilibrium models. These includes, models that consider adiabatic configuration such as the ones analyzed via the classical criterion of \\citet{ber} or those that presuppose  static equilibrium and analyze thermal stability. In the application of Bernstein's criterion, the adiabatic assumption implies that the energy balance equation is not required and thus dissipation is impossible. Also the assumption of static models is a strong, and often unjustified, restriction for open systems.  In this paper we apply an energy principle to analyze the stability of solar coronal loops when helical modes are present. The principle was obtained in  previous papers (Paper I:  \\citep{cos1}; \\citep{cos2}; see also  \\citep{us0}) using a general procedure of irreversible thermodynamics -based on firmly established thermodynamic laws- that can be understood as an extension of  Bernstein's MHD principle to situations far from thermodynamic equilibrium.  In Paper I and in \\citep{cos2}  we showed how to obtain the variational principle for solar coronal structures from the equations that describe the dynamics of the system. The method consists of obtaining a Lyapunov function, also known as generalized potential, that represents the mathematical expression of the stability conditions.  The principle is subject to physically reasonable requirements of hermiticity and antihermiticity over the matrices. For a more detailed presentation see Paper I and the references therein.      \\subsection{Solar coronal loops}    MHD loop oscillations in the corona are known to be strongly damped, mostly having  decaying times of few periods $N_{p}\\approx 2 - 7 \\  periods.$ While thermal conduction, with the contribution of radiative cooling mechanisms, could be the main cause of the damping of pure MHD slow magnetoacoustic mode oscillations they are unimportant for the  MHD fast modes. Resonant absorption and phase mixing seem more promising in giving account of the rapid decay (\\citep{gh94}; hereafter HG, \\citep{go02})  of the ideal fast oscillations   of these strongly inhomogeneous and structured plasma systems. Inhomogeneous equilibrium distributions of plasma density and temperature varying continuously across the magnetic field led to plasma waves with continuous intervals of eigenfrequencies. The occurrence of the Alfv\\'en ideal MHD continuum in a thin edge layer is derived from the highly anisotropic character of  the fast   magnetoacoustic waves giving rise to a peak of the amplitudes where the perturbation develops large gradients and the absorption has maxims. However, there is another type of continuum commonly known as slow magnetosonic continuum associated to the inhomogeneity of the equilibrium parameters along the axis of the loop (see Paper I). This inhomogeneities are associated, for example, to changes in the % density concentration at the loop basis. If the magnetic field is twisted the inhomogeneities led to the coupling of Alfv\\'en and slow magnetosonic continuum modes (\\citep{bel}).  The resonant absorption mechanism of wave heating consists on the non--dissipative transference  of wave energy from the collective line-tied wave with fast discrete eigenvalues (kinetic energy of the fast radial component) to a local resonant  mode in the Alfv\\'en continuum, (kinetic energy of the azimuthal component), which is then dissipated in an enhanced manner. Then, the  continuum oscillations are converted into heat by dissipative processes; as the medium has large gradients in the Alfv\\'en speed, the oscillations of neighboring field lines become out of phase and  shear Alfv\\'en waves lead to enhanced viscous and ohmic dissipation (see \\citep{pr83} for the linear regime and \\citep{nak97} for the nonlinear one). The mode conversion from the collective to the local mode occurs in a  time that is  non--dissipative and generally much shorter than the second time scale which is related to the dissipative damping of the small--scale perturbations of the local mode in the resonance layer (\\citep{rob00}; \\citep{van04}).   The whole temporal pattern description  of modes that exhibit a combination of global (discrete line--tied fast eigenmode) and localize (Alfv\\'en continuum mode) behaviour is known as  quasi--mode. Moreover, the mixed nature of the modes is not only due to the temporal behaviour but also to the boundary value problem giving rise to a spatial behaviour which is also of a mixed nature, i.e. coronal loops with line--tying constraints cannot support pure waves: Alfv\\'en, slow or fast magnetoacoustic modes. HG studied the mixed spectral description of coronal loops (i.e. the resulting superposition of basic waves which adjust the line--tied condition) without assuming a straight magnetic field and forcing the loop to follow the photospheric velocity perturbations. They found that pure  Alfv\\'en and pure slow modes are  obtained as singular limiting cases of cluster spectra of Alfv\\'en--fast or slow--fast modes, where the fast components are localized in a photospheric boundary associated to the line--tied condition: the coronal part of the loop acting as a resonant cavity  of large   Alfv\\'en components  and  fast components, with a small but rapidly varying amplitude, located in the photospheric boundary layer. They found that heating of coronal loops by resonant absorption is due to the line--tied Alfv\\'en continuum which no  longer depends on the poloidal magnetic field and that the corresponding eigenmodes have a global ballooning feature which is characterized by  an accumulation point given by the Alfv\\'en frequency. In \\citep{gh93} (hereafter GH), a variational principle, based in Bernstein's principle, was obtained to  derive the Alfv\\'en and slow continuum frequencies in a line--tied inhomogeneous cylinder. Stability considerations led them to conclude the global  stability  of coronal loops.  In this paper, following results of Paper I  we apply our energy principle to consider the stability and mode structure of loop inhomogeneous coronal models with non--vanishing helicity.  Our principle has the advantages that it does not require a WKB approximation and that, as was mentioned, it  allows the consideration of the coupling of the thermal and mechanical equations that are necessary to analyze far from equilibrium states. ",
        "conclusions": "Convective motion of the photosphere  is believed to provide the energy that is storage in twisted magnetic coronal fields allowing the presence of long--lived coronal structures until it is released by instabilities  (\\citep{raa};  \\citep{vrs}). On the other hand, continuous spectra are generally associated to stability. An accepted conjecture  establishes that unstable modes have  a discrete spectrum (see \\citep{fre} or  \\citep{pri}). There are two types of  possible continuous spectra in this problem. The inhomogeneous character of the equilibrium parameters along the loop axis can lead to a continuum that couples to the Alfv\\'en continuum \\citep{bel}; e.g, when the disturbances considered are comparable to the inhomogeneous characteristic wavelength stable eigenvalues    can give rise to a  continuous spectrum  ($L/2$, the equilibrium structure in the $z$ component). This is the case studied in Paper I. On the other hand, GH established, for non vanishing helicity systems, that there is a continuous spectrum associated with the line-tied Alfv\\'en resonance leading to the  damping and heating by the resonant absorption mechanism and thus,  directly  relate to the stability of loops.  They  also pointed out how to  obtain  the resonant singular limit $\\omega$, from the class of physically permissible solutions, \\begin{equation} \\omega(r)=\\frac{nB_{z}(r)}{\\int_{-L}^{L}\\sqrt{\\rho(z)}dz}. \\label{11} \\end{equation} This resonance results because of the absence of an explicit dependence on the azimuthal magnetic field component ($B_{\\varphi}$).  Thus, in order  to understand in which conditions which  mechanism can dominate and give account of the different scenarios i.e., the driving of the instability or the damping of mode oscillations, it is critical to gain knowledge about the dynamics and energetic contribution of twisted  structures. Yet, the implications of the twisting in  theoretical and observational descriptions are poorly known; e.g., there is no clarity about  the modification of the dispersion relation  and observational data are   indirectly inferred.   In this paper we focused our attention to describe the changes in periods, stability and mode structure  of coronal loops when the helicity, the magnetic field intensity and the radius are varied. For loops with vanishing  helicity it is well established that  the Alfv\\'en line--tied resonance continuum is responsible of the damping of kink ($m=1$) quasi--modes via the transfer of energy from the radial component into the azimuthal one, i.e., from discrete global modes into the  local continuum modes where phase mixing can take place. Still, the twisting of the magnetic field leads to the coupling of MHD cylindrical modes making difficult to provide a classification in terms of the behavior of pure--like modes.   In order to calculate modes and frequencies we followed the schematic procedure described in Paper I and in \\citep{gal1}. We used a symbolic manipulation program to integrate the equations. $\\delta^{2}W_{p}$ and the perturbations were expanded in a six dimensional--Fourier basis on the independent coordinate $z$ that adjusts to border conditions, i.e., the four perturbated components  (eq.~\\ref{9}) were expanded in a six mode basis to obtain $24$ eigenvalues and eigenvectors for each of the  helicity and the magnetic field values. Only  the first eighteen eigenvalues were considered (the others are more than two order of magnitude smaller and accumulate at zero; the eigenvectors are also vanishingly small). Thus, a quadratic form for $\\delta^{2}W_{p}$ was obtained and was minimized with the Ritz variational procedure. A matrix discrete eigenvalue problem subject to a normalization constraint was obtained. From the resulting modes and   the generalized potential energy (eq.~\\ref{2}): $\\delta^{2}W_{p}\\geq0$ the stability of each mode was determined.   The coronal loop parameters used were: $L=10^{10}cm$ (or $L=100Mm$), $T_{b}=10^{4}K$ $T_{t}=10^{6}K$ $n_{e}=10^{8}cm^{-3}$ electron number density $p_{t} =2k_{B}T_{t} \\;$; $\\rho_{t}=m_{p}p_{t}/k_{B}T_{t}$. Frequencies and modes were calculated for two different values of the magnetic field: $B_{0}=10G$ and $B_{0}=100G$, and for three different values of the  helicity $b=(3.1 \\ 10^{-8}  ; \\ 3.1 \\ 10^{-7}; \\ 1.9  \\ 10^{-6})$  which correspond to the adimentional values: $b_{a}=(2.8; \\ 28; \\  170)$  with $N_{t}=(0.45;  \\ 4.3;  \\ 13.7)$, $N_{t}:$ the number of turns over the cylinder length. These helicity values defined as  weak, moderated and strong helicity respectively correspond to the classification given in  \\citep{asc4} (Chapter 5). The adimentional radius was initially chosen as $R=0.01$. In what follows we summarize the conclusions obtained from the data analysis which are displayed in  three tables.        Table 1 shows the periods (in minutes) for weak, moderate and strong helicity for two values of the magnetic field intensity ($B_{0}=10G$ and  $B_{0}=100G$ (left and right panel respectively).  S and U letters indicate the stable--unstable character of the modes. From the table we see that:  \\noindent I) Weak helicity modes are stable. This  is in accordance with the analytic results by  \\citet{rud} who studied nonaxisymmetric oscillations of a thin twisted magnetic tube with fixed ends in a zero-beta plasma.  \\noindent II) Higher modes have an accumulation point at zero, indicating  the presence of a continuum  spectra of stable modes (as in Paper I).  Note that,  calculus performed via discrete basis, as in our case, give spectra that are necessarily discrete. Thus, an accumulation of discrete eigenvalues suggests a stable continuum  spectrum.   \\noindent III) The  $B_{0}=10G$ case has larger periods than the $B_{0}=100G$ one. For moderate and strong helicity the eigenvalues of the first panel  follow a scaling law with that of the second one i.e., they scale with the magnetic field intensity exactly as  the Alfv\\'en speed does $P_{10G}\\simeq 10 P_{100G}$.  \\noindent IV) As HG and GH, we note a clustering of the spectra associated to the change from real to imaginary eigenvalues (and viceversa). There is a pronounced change (in the spacing of the periods or/and in the stability) from the sixth mode  to the seventh mode. This is noted by a double line in Table 1 and related to the importance of the parallel  component with respect to the perpendicular component (see Table 2). Up to period number ten real -- imaginary eigenvalues of the first panel ($B_{0}=10G$) correspond to real--imaginary  ones of the second panel ($B_{0}=100G$). Also, excepting large order periods $n>10$,  when the helicity is increased from weak to moderate  the imaginary stable eigenvalues turn to imaginary unstable ones. For $B_{0}=10G$  and weak and moderate helicity cases there are five different groups of periods ($P_{1}-P_{6};P_{7}-P_{10};P_{11}-P_{12};P_{13}-P_{14};P_{15}- P_{18}$) (see also Table 2). The clustering is more difficult to establish i.e., the differences are less pronounced, with increasing magnetic field intensity and larger order periods.  \\noindent In order to compare our results with those given by these authors we calculated  the expression eq.~\\ref{11} for our modes. We found that, all periods  excepting  $P_{1}-P_{6}$   weak helicity modes satisfy the relation and thus, they belong to the Alfv\\'en continuum spectrum justifying the scaling law described in III. As HG, we conclude that the change in the real--complex character of  the $P_{6}-P_{7}$ eigenvalues is associated to the existence of an accumulation point of the resonant Alfv\\'en continuum, however we find that this change is not necessarily related to a change in the stability as they claimed.  Note that all modes  with weak helicity are stable (even the imaginary ones); in all the other cases the imaginary character of the eigenvalues is associated to instability. Yet, the continuum stable eigenvalue conjecture is here still valid  \\citep{fre},  \\citep{pri}; it applies to a spectrum with an accumulation point in zero; we found stable modes for all the helicity values and for the two magnetic field values with $P_{n>14}$. Note that the analysis of stable modes  is still of interest because depending on the relative characteristic times of stable and unstable modes  the stable ones could   be active and accessible to observations.  \\noindent The presence of at least one  unstable mode means that the equilibrium state is unstable. Thus, taking into account the whole range of stable modes, we confirm previous results leading to conclude that field configurations with some degree of twisting  give a stabilizing effect allowing the storage of magnetic energy  \\citep{raa2}, i.e.,   when the helicity is augmented the stable weak case  turns to an unstable one suggesting a critical value.  \\bigskip  In Paper I,  we obtained  only one unstable mode classified as  slow magnetoacoustic mode due to the almost longitudinal character (parallel to the magnetic field) of the wavevector  perturbation and to the fact that the period did not changed with the intensity of the magnetic field,  resembling  acoustic waves with sound speeds, $v_{s}$,  independent of the magnetic field. The characteristic unstable time obtained in Paper I was $\\tau_{u}=36 \\ min$, corresponding to a typical slow magnetoacoustic fundamental  period with a characteristic wavelength  of the order of the loop length $L/2$.  Also, we obtained a continuous set of stable modes classified as fast magnetoacoustic modes due to their large value  component orthogonal to the magnetic field and to the fact that the eigenvalues scale with the intensity of the magnetic field  as   $P_{11G}\\simeq 10 P_{100G}$; thus resembling the dependence of the Alfv\\'en waves $v_{A} \\sim B_{0}$.   Table 2 (First Panel) displays the resulting features associated to the relative intensity of the parallel and perpendicular to the field components ($(\\xi_{\\parallel},\\xi_{\\perp},\\xi_{r}) $ is an orthogonal basis) and their classification  as  slow--like (S) or  fast--like (F). The relative phase between the  components is also indicated in the table by P (in phase) and IP (inverted phase). Table 2 (Second Panel) also shows the intensity relationship between the cylindrical components. In order to classify the modes and to compare with  the slow and fast magnetoacoustic modes obtained in Paper I, we calculated  the cylindrical mode components and also the tangential and normal to the field components ($\\xi_{\\parallel}= (Rb\\xi_{\\phi}+\\xi_{z})/\\Delta; \\xi_{\\perp}= (\\xi_{\\phi}-Rb\\xi_{z})/\\Delta$). Our interest in the $\\xi_{\\parallel}$, $\\xi_{z}$, $\\xi_{r}$, $\\xi_{\\perp}$ and $\\xi_{\\phi}$ components  resides in that: First,  when  the helicity is weak, the $\\xi_{\\parallel}$ component is expected to play the slow-mode role of $\\xi_{z}$ in Paper I. Second, the $\\xi_{r}$ component is related to the fast modes  and determines the resonant absorption mechanism when uniform cylindrical flux tubes are considered by the transferring of energy to the $\\xi_{\\phi}$  component. When helicity and inhomogeneous distribution of equilibrium parameters  are present it is worth investigating the transferring of energy from the $\\xi_{r}$ component to the others. In this case the resonant damping of global oscillations will occur by conversion of kinetic energy of the radial component into kinetic energy of the $\\xi_{\\parallel}$ and   $\\xi_{\\perp}$ components; both components forming the plane orthogonal to $\\xi_{r}$, and equal to the plane formed by $\\xi_{\\phi}$ and $\\xi_{z}$.   From the  analysis of the amplitude  of the components of the $P_{1}-P_{6}$ modes with respect to the $P_{7}-P_{18}$ ones in the weak helicity case i.e., real and imaginary eigenvector respectively, we could classify the first ones as slow-like modes because: I) their tangential components $\\xi_{\\parallel}$ are at least an order of magnitude larger than the normal ones $\\xi_{\\perp}$; II) as the helicity is weak $\\xi_{\\parallel}\\approx \\xi_{z}$ and $\\xi_{r}\\rightarrow 0; \\xi_{\\phi}\\rightarrow 0$, the wavevector is almost tangential to the magnetic field; III) they have a larger characteristic time and a shorter characteristic speed than the imaginary eigenvectors. On the contrary, imaginary eigenvalues  are associated to  large values of the $\\xi_{r}$ component and $\\xi_{\\perp}$  component (due to large values of   $\\xi_{\\phi}$ (see Table 2 Second Panel)) , and small values of the $\\xi_{\\parallel}$ and $\\xi_{z}$ components. As in Paper I,  when the eigenvalues change form real to imaginary the period strongly diminishes and  a change in the type of mode  from the slow  to fast magnetoacoustic type occurs. In opposition to Paper I where the acoustic mode has the same eigenvalue for both magnetic field intensities, here the modes are affected by the strengthening of the magnetic field leading to an-order-of-magnitude shorter  period than in the non--helicity case. The $\\xi_{\\parallel}$ and $\\xi_{\\perp}$ components are in an inverted phase for real eigenvector modes and in phase for imaginary eigenvector modes.     For moderate helicity the overall description is similar but  all the cases having  non vanishing $\\xi_{\\phi}$ component and  all the periods in the resonant line-tied continuum. As was mentioned, real--imaginary eigenvalues correspond to stable--unstable behavior.     In the strong helicity case, as the weak and moderate ones, we note for $P_{1}-P_{6}$ larger, but comparable,  values of  the $\\xi_{\\parallel}$ component with respect to the $\\xi_{\\perp}$ component. In this case  the two components of the mode are in phase. This relationship between the $\\xi_{\\parallel}$ and $\\xi_{\\perp}$ components of Table 2 (FP), and their associated phases is found again in the modes with $P_{15} - P_{18}$. In spite that  these features are associated to the slow magnetoacoustic characterization, Table 2 (SP) shows that as $\\xi_{z}$ is vanishingly small, the strong helicity case cannot be classify as a slow mode.  \\bigskip  When helicity is present the mixed character of the modes manifests  itself making difficult to identify the components that are involved in   the damping mechanism. However, taking into account  the resonant frequency of     eq.~\\ref{11}, we noted that (HG)  all  the modes, except those  with  $P_{1}-P_{6}$ periods of the weak helicity case, have  resonant frequencies suggesting that resonant absorption in helical modes is associated to modes with significant values of $\\xi_{\\perp}$ component. If this argument is correct we can affirm that  the damping mechanism of body helical modes is associated to the transfer of kinetic energy of the radial component into kinetic energy of the $\\xi_{\\perp}$ component which is not only related to the $\\xi_{\\phi}$ cylindrical contribution but also to the $\\xi_{z}$ one by the expression $ \\xi_{\\perp}=(\\xi_{\\phi}-Rb\\xi_{z}) \\Delta$.  \\bigskip  We also analyzed the change of the period as a function of the radius for different values of the helicity. We found that, for weak helicity,  the increasing of the radius leads to a decrease of the periods. This is in accordance with observations, e.g., observed sausage modes are associated with thicker and denser loop structures and lower periods; while in  other case (unstable cases) the increasing of the radius leads to an increase of the period.  \\bigskip  Table 3 -First and Second Panel- shows the variation of the radius $R$ with the twist $bR$ for weak and moderate helicity respectively. \\citep{rud} has conjectured that the line-tying condition at the tube ends should stabilize the tube and has suggested a critical value ($\\sim L b<q$; with $q$ a positive constant and $L$ the loop length) for the onset of instability. Also, \\citep{lin}  found that  when the  helicity grows beyond a critical value, the kink isolated twisted magnetic flux tubes below the photosphere become unstable. In fact,  Table 3 can be seen as  the variation of $R$ with  the  twist value: $bR$, for     constant values of the helicity $b$ in the two cases:  weak and moderate  respectively. Stability is guaranteed when the loop radius is varied  between $R=0.01$ and $R=0.1$ and the helicity is weak $b=0.05$ (for almost the same value of the length of the loop, $L$). However, when the helicity is incremented to $b=0.5$ even for the radius of $R=0.01$ the loop structure is unstable, thus, instability can be associated with the presence of helicity values larger than  a critical one.  \\bigskip  Figure~\\ref{fig:tres} shows the general potential energy for  $P_{6}$ and $P_{7}$ in the weak and moderate cases. Note the change of this function when the system turns from stable to unstable, as helicity is augmented i.e., from  $\\delta^{2} W_{p}>0$ to $\\delta^{2} W_{p}<0$. Figure~\\ref{fig:tres}a and  Figure~\\ref{fig:tres}c display    the total energy composed by the compressional, radiative, thermal and magnetic energy contributions of   $P_{6}$ mode in the weak and moderate case respectively. The same features but for the $P_{7}$ mode  are shown in  Figure~\\ref{fig:tres}e and  Figure~\\ref{fig:tres}g.  Figure~\\ref{fig:tres}b and  Figure~\\ref{fig:tres}d show the magnetic energy content alone for $P_{6}$ mode in the weak and moderate case respectively.  Figure~\\ref{fig:tres}f and  Figure~\\ref{fig:tres}h show the magnetic energy content for $P_{7}$ mode and for the weak and moderate case respectively.   It can be seen, in this and in all the other cases, that the magnetic energy content has a  determinant role on the stability--instability of the system, i.e., the stability changes when the magnetic generalized potential energy changes sign. Thus, a result of this analysis is that the stability of twisted coronal loops is fundamentally determined by the storing of magnetic energy,  being the other contributions less significant. Meanwhile, when the helicity is weak or vanishingly small and the magnetic contribution has a stabilizing effect  the other non-dominant contributions, as the non-adiabatic ones,  can play an important  role. This makes possible, for example, the damping of fast excitations due to resonant absorption. Yet, even when one of these contributions is unstable, stable modes could be active for a while if their characteristic periods  are shorter than the characteristic time  of the instability. This is the case of Paper I, where we obtained a slow   mode with an unstable characteristic times of $\\tau\\sim 36 \\ min$ coexisting with  stable fast modes with periods about $P\\sim 1 \\ min$; moreover, we showed that the instability can be nonlinearly saturated giving rise to a limit-cycle solutions, i.e., an oscillation between parallel plasma kinetic energy and plasma internal energy where the magnetic energy plays no relevant  role. Thus, the contribution to the stable--unstable character of the modes is mostly due to the magnetic energy content and not to other energetic contributions. Note that as the balance energy equation  takes into account non--adiabatic contributions, i.e., radiation, heat flow and heat function (with $L=0$ at the equilibrium), the resulting perturbations are not constrained to  the force--free condition. So, one result of the analysis is that the pertubation energy contribution is mainly due to magnetic forces. Thus, for these type of twisted magnetic field models, non--adiabatic perturbations (e.g. thermal perturbations) and resonant absorption seem unimportant to guarantee stability; a loop system with weak storage of magnetic energy (low values of the helicity)  could be  released if the helicity is suddenly increased, e.g., by footpoint motions. Meanwhile all the \"zoo\" of the coronal seismology can be active and accessible to observations.    \\begin{figure*} \\includegraphics[width=4.3cm]{P6ewp.eps} \\includegraphics[width=4.3cm]{P6emwp.eps} \\includegraphics[width=4.3cm]{P6emp.eps} \\includegraphics[width=4.3cm]{P6emmp.eps} \\includegraphics[width=4.3cm]{P7ewp.eps} \\includegraphics[width=4.3cm]{P7emwp.eps} \\includegraphics[width=4.3cm]{P7emp.eps} \\includegraphics[width=4.3cm]{P7emmp.eps} \\caption{Energy content of the sixth and seventh mode for $B_{0}=10G$. a) Total potential energy and b) magnetic potential energy respectively for the sixth mode $P_{6}=1.23 \\ min$ and for weak helicity. c) Total potential energy and d) magnetic potential energy respectively for the sixth mode $P_{6}=1.23 \\ min$ and for moderate helicity. e) Total potential energy and f) magnetic potential energy respectively for the seventh mode $P_{7}=0.07 \\ min$ and for weak helicity. g) Total potential energy and h) magnetic potential energy respectively for the sixth mode $P_{7}=0.07 \\ min$ and for moderate helicity.} \\label{fig:tres} \\end{figure*}   \\begin{table*} \\begin{tabular}{cccccccc} \\hline $P_{i}$&$weak $&$moderate$&$strong$&$ \\ \\ \\ $& $weak$&$moderate$&$strong$ \\\\  \\hline $P_{1}$&$1.921 \\ S $&$0.209 \\ S$&$0.525 \\ i \\ U$&$ \\ \\ \\ $&$0.159 \\ S$&$0.021 \\ S$&$0.052 \\ i \\ U$\\\\  \\hline $P_{2}$&$1.869 \\ S $&$0.204 \\ S$&$0.450 \\ S$&$ \\ \\ \\  $& $0.158 \\ S$&$0.020 \\ S$&$0.044 \\ S$\\\\  \\hline $P_{3}$&$1.535 \\ S $&$0.169 \\ S$&$0.430 \\ S$&$ \\ \\ \\  $& $0.154 \\ S$&$0.017 \\ S$& $0.042 \\ S$ \\\\  \\hline $P_{4}$&$1.533 \\ S $&$0.168 \\ S$&$0.424 \\ i \\ U$&$ \\ \\ \\  $&$0.153 \\ S$&$0.0167 \\ S$&$0.042 \\ i \\ U$\\\\  \\hline $P_{5}$&$1.306 \\ S  $&$0.143 \\ S$&$0.206 \\ i \\ U$& $ \\ \\ \\  $&$0.151 \\ S$&$0.014 \\ S$&$0.020 \\ i \\ U$\\\\  \\hline $P_{6}$&$1.228 \\ S  $&$0.135 \\ S$&$0.177  \\ i \\ U$&$ \\ \\ \\  $& $0.15 \\ S $&$0.013 \\ S$&$0.017 \\ i \\ U$\\\\  \\hline \\hline $P_{7}$&$0.068  \\ i \\ S$&$0.070 \\ i \\ U$&$0.125 \\ S$&$ \\ \\ \\  $& $0.0047 \\ i \\ S$&$0.007 \\ i \\ U$&$0.0125 \\ S$\\\\  \\hline $P_{8}$&$0.064  \\ i \\ S$&$0.066 \\ i  \\ U$&$0.122 \\ S$&$ \\ \\ \\  $& $0.0046 \\ i \\ S$&$0.006 \\ i  \\ U$&$0.012 \\ S$\\\\  \\hline $P_{9}$&$0.042  \\ i \\ S $&$0.044 \\ i  \\ U$&$0.101 \\ S$&$ \\ \\ \\  $& $0.0044 \\ i  \\ S$&$0.0043 \\ i \\ U$&$0.0101 \\ S$\\\\  \\hline $P_{10}$&$0.041  \\ i  \\ S$&$0.043  \\ i \\ U $&$0.100 \\ S$&$ \\ \\ \\  $&  $0.0043  \\ i  \\ S$&$0.0042 \\ i  \\ U$&$0.01 \\ S$\\\\  \\hline $P_{11}$&$0.033 \\ S  $&$0.036 \\ S$&$0.989 \\ S$&$ \\ \\ \\ $&  $0.0042 \\ i  \\ S$& $0.0036 \\ S$&$0.099 \\ S$\\\\  \\hline $P_{12}$&$0.032 \\ S  $&$0.035 \\ S$ &$0.096 \\ S$&$ \\ \\ \\  $&  $0.0041 \\ i  \\ S$&$0.0035 \\ S$ &$0.0096 \\ S$\\\\  \\hline $P_{13}$&$0.030  \\ i \\ S $&$0.031 \\ i \\ U$ &$0.085 \\ S$&$ \\ \\ \\  $&$0.003 \\ S$&$0.0031 \\ i \\ U$ &$0.0086 \\ S $\\\\  \\hline $P_{14}$ & $0.027  \\ i \\ S $&$0.029  \\ i  \\ U$&$0.081 \\ S $&$ \\ \\ \\  $&$0.0026 \\ S$&$0.003 \\ i  \\ U$&$0.0081 \\ S $\\\\  \\hline $P_{15}$&$0.025 \\ S $&$0.027 \\ S $&$0.077 \\ S  $&$ \\ \\ \\  $&$0.0025 \\ S $&$0.0027  \\ S  $&$0.0077 \\ S $\\\\  \\hline $P_{16}$&$0.024 \\ S $&$0.026 \\ S $&$0.076 \\ S  $&$ \\ \\ \\  $&$0.002 \\ S $&$0.003  \\ S  $&$0.0076 \\ S $\\\\  \\hline $P_{17}$&$0.02 \\ S $&$0.02 \\ S $&$0.063 \\ S  $&$ \\ \\ \\  $&$0.0024 \\ S $&$0.0021  \\ S  $&$0.0063 \\ S $\\\\  \\hline $P_{18}$&$0.018 \\ S $&$0.02 \\ S $&$0.059 \\ S  $&$ \\ \\ \\    $&$0.0024 \\ S $&$0.0025  \\ S  $&$0.006 \\ S $\\\\  \\hline \\end{tabular} \\caption{\\label{tab:table1} Eighteen first periods associated to stable (S) and unstable (U) eigenvalues (minutes) for A) Left panel: $B_{0}=10G$ with A1) left column: weak helicity, A2) middle column: moderate helicity, A3) right column: strong helicity and B) Right panel: $B_{0}=100G$ with B1, B2, B3 the same as in A. Larger order modes were discarded.} \\end{table*}   \\begin{table*} \\begin{tabular}{ccccccc} \\hline $P_{i}$&$weak $&$moderate$&$strong  \\ \\ \\ \\ $&$weak $&$moderate$&$strong$\\\\  \\hline $P_{1}$&$\\xi_{\\parallel} \\gg \\xi_{\\perp} \\mapsto 0  \\ S; \\ IP $&$\\xi_{\\parallel}  > \\xi_{\\perp}  \\ S; \\ IP  $&$\\xi_{\\parallel}  \\geq \\xi_{\\perp} \\ P \\ \\ \\ \\ $&$ \\xi_{z}\\gg  \\xi_{\\phi}   \\sim   \\xi_{r}\\mapsto 0  $&$ \\xi_{z}\\gg  \\xi_{\\phi}   \\sim   \\xi_{r} $&$  \\xi_{r} \\leq \\xi_{\\phi}; \\xi_{z}  \\mapsto 0  $\\\\  \\hline $P_{2}$&$\\xi_{\\parallel} \\gg \\xi_{\\perp} \\mapsto 0  \\ S; \\ IP $&$\\xi_{\\parallel}  > \\xi_{\\perp}  \\ S; \\ IP  $&$\\xi_{\\parallel}  \\geq \\xi_{\\perp}  \\ P \\ \\ \\ \\ $&$ \\xi_{z}\\gg  \\xi_{\\phi}   \\sim   \\xi_{r}\\mapsto 0  $&$ \\xi_{z}\\gg \\xi_{\\phi}   \\sim   \\xi_{r}  $&$   \\xi_{r} \\leq \\xi_{\\phi}; \\xi_{z}  \\mapsto 0  $\\\\  \\hline $P_{3}$&$\\xi_{\\parallel} \\gg \\xi_{\\perp} \\mapsto 0  \\ S; \\ IP $&$\\xi_{\\parallel}  > \\xi_{\\perp}  \\ S; \\ IP  $&$\\xi_{\\parallel}  \\geq \\xi_{\\perp}  \\ P \\ \\ \\ \\ $&$ \\xi_{z}\\gg  \\xi_{\\phi}   \\sim   \\xi_{r}\\mapsto 0  $&$ \\xi_{z}\\gg  \\xi_{\\phi}   \\sim   \\xi_{r} $&$  \\xi_{r} \\leq \\xi_{\\phi}; \\xi_{z}  \\mapsto 0 \\ $\\\\  \\hline $P_{4}$&$\\xi_{\\parallel} \\gg \\xi_{\\perp}  \\mapsto 0 \\ S; \\ IP $&$\\xi_{\\parallel}  > \\xi_{\\perp}  \\ S; \\ IP  $&$\\xi_{\\parallel}  \\geq \\xi_{\\perp}  \\ P \\ \\ \\ \\ $&$ \\xi_{z}\\gg  \\xi_{\\phi}   \\sim   \\xi_{r}\\mapsto 0  $&$ \\xi_{z}\\gg  \\xi_{\\phi}   \\sim   \\xi_{r}  $&$  \\xi_{r} \\leq \\xi_{\\phi}; \\xi_{z}  \\mapsto 0 $\\\\  \\hline $P_{5}$&$\\xi_{\\parallel} \\gg \\xi_{\\perp} \\mapsto 0 \\ S;  \\ IP $&$\\xi_{\\parallel}  > \\xi_{\\perp}  \\ S; \\ IP  $&$\\xi_{\\parallel}  \\geq \\xi_{\\perp}  \\ P  \\ \\ \\ \\ $&$ \\xi_{z}\\gg  \\xi_{\\phi}   \\sim   \\xi_{r}\\mapsto 0  $&$ \\xi_{z}\\gg  \\xi_{\\phi}   \\sim   \\xi_{r}  $&$ \\xi_{r} \\leq \\xi_{\\phi}; \\xi_{z}  \\mapsto 0  $\\\\  \\hline $P_{6}$&$\\xi_{\\parallel} \\gg \\xi_{\\perp}  \\mapsto 0 \\ S; \\ IP $&$\\xi_{\\parallel}  > \\xi_{\\perp}   \\ S; \\ IP  $&$\\xi_{\\parallel}  \\geq \\xi_{\\perp}  \\ P \\ \\ \\ \\ $&$ \\xi_{z}\\gg  \\xi_{\\phi}   \\sim   \\xi_{r}\\mapsto 0  $&$ \\xi_{z}\\gg  \\xi_{\\phi}   \\sim   \\xi_{r} $&$ \\xi_{r} \\leq \\xi_{\\phi}; \\xi_{z}  \\mapsto 0  $\\\\  \\hline \\hline $P_{7}$&$\\xi_{\\perp} \\gg \\xi_{\\parallel} \\mapsto 0  \\ F; \\ P $&$\\xi_{\\perp} > \\xi_{\\parallel}  \\ F;  \\ P$ &$\\xi_{\\perp}  \\geq \\xi_{\\parallel}  \\ IP \\ \\  \\ \\ $&$\\xi_{r}   \\sim   \\xi_{\\phi}\\gg   \\xi_{z}\\mapsto 0$&$\\xi_{r}   \\sim  \\xi_{\\phi}\\gg \\xi_{z} \\mapsto 0  $&$\\xi_{z} > \\xi_{r}  > \\xi_{\\phi}  $\\\\  \\hline $P_{8}$&$\\xi_{\\perp} \\gg \\xi_{\\parallel}\\mapsto 0   \\ F; \\ P $&$\\xi_{\\perp} > \\xi_{\\parallel}   \\ F; \\ P$&$\\xi_{\\perp}  \\geq \\xi_{\\parallel}  \\ IP \\ \\ \\ \\ $&$\\xi_{r}   \\sim   \\xi_{\\phi}\\gg   \\xi_{z}\\mapsto 0 $&$\\xi_{r}   \\sim   \\xi_{\\phi}\\gg  \\xi_{z}\\mapsto 0  $&$\\xi_{z} > \\xi_{r}> \\xi_{\\phi}  $ \\\\  \\hline $P_{9}$&$\\xi_{\\perp} \\gg \\xi_{\\parallel}\\mapsto 0  \\ F;  \\ P $&$\\xi_{\\perp} > \\xi_{\\parallel}   \\ F; \\ P$&$\\xi_{\\perp}  \\geq \\xi_{\\parallel}  \\ IP \\ \\ \\ \\ $&$\\xi_{r}  \\sim   \\xi_{\\phi}\\gg  \\xi_{z}\\mapsto 0  $&$\\xi_{r}  \\sim  \\xi_{\\phi}\\gg  \\xi_{z}\\mapsto 0  $&$\\xi_{z} > \\xi_{r} > \\xi_{\\phi}  $\\\\  \\hline $P_{10}$&$\\xi_{\\perp} \\gg \\xi_{\\parallel}\\mapsto 0 \\ F;   \\ P $&$\\xi_{\\perp} > \\xi_{\\parallel}  \\ F;  \\ P$&$\\xi_{\\perp}  \\geq \\xi_{\\parallel}  \\ IP \\ \\ \\ \\ $& $\\xi_{r}   \\sim   \\xi_{\\phi}\\gg  \\xi_{z}\\mapsto 0  $&$\\xi_{r}   \\sim   \\xi_{\\phi}\\gg  \\xi_{z}\\mapsto 0  $&$\\xi_{z}> \\xi_{r} > \\xi_{\\phi}  $\\\\  \\hline $P_{11}$&$\\xi_{\\perp} \\gg \\xi_{\\parallel}\\mapsto 0  \\ F;  \\ P $&$\\xi_{\\perp} > \\xi_{\\parallel}  \\ F;  \\ P$&$\\xi_{\\perp}  \\geq \\xi_{\\parallel}  \\ IP \\ \\ \\ \\ \\ $&$\\xi_{r}   \\sim   \\xi_{\\phi}\\gg   \\xi_{z}\\mapsto 0  $&$\\xi_{r}   \\sim   \\xi_{\\phi} \\gg  \\xi_{z}\\mapsto 0  $&$\\xi_{z} > \\xi_{r} > \\xi_{\\phi}  $ \\\\  \\hline $P_{12}$&$\\xi_{\\perp} \\gg \\xi_{\\parallel}\\mapsto 0  \\ F;  \\ P $&$\\xi_{\\perp} > \\xi_{\\parallel}   \\ F; \\ P$&$\\xi_{\\parallel}  \\geq \\xi_{\\perp}  \\ IP \\ \\ \\ \\ $&$\\xi_{r}   \\sim  \\xi_{\\phi}\\gg   \\xi_{z}\\mapsto 0  $&$\\xi_{r} \\sim  \\xi_{\\phi}\\gg \\xi_{z}\\mapsto 0  $ &$  \\xi_{r}>  \\xi_{\\phi} > \\xi_{z}   $ \\\\  \\hline $P_{13}$&$\\xi_{\\perp} \\gg \\xi_{\\parallel} \\mapsto 0  \\ F; \\ P $&$\\xi_{\\perp} > \\xi_{\\parallel}   \\ F; \\ P$&$\\xi_{\\perp}  \\geq \\xi_{\\parallel}  \\ IP \\ \\ \\ \\ $&$\\xi_{r}\\sim    \\xi_{\\phi}\\gg  \\xi_{z}\\mapsto 0  $&$\\xi_{r}   \\sim  \\xi_{\\phi}\\gg  \\xi_{z}\\mapsto 0 $&$\\xi_{z} > \\xi_{r} > \\xi_{\\phi}  $\\\\  \\hline $P_{14}$&$\\xi_{\\perp} \\gg \\xi_{\\parallel}\\mapsto 0   \\ F; \\ P $&$\\xi_{\\perp} > \\xi_{\\parallel}  \\ F;  \\ P$ &$\\xi_{\\perp}  \\geq \\xi_{\\parallel}  \\ IP \\ \\ \\ \\ $&$\\xi_{r}  \\sim   \\xi_{\\phi}\\gg   \\xi_{z}\\mapsto 0 $&$\\xi_{r}  \\sim  \\xi_{\\phi}\\gg  \\xi_{z}\\mapsto 0 $&$\\xi_{z} > \\xi_{r} > \\xi_{\\phi}  $\\\\  \\hline $P_{15}$&$\\xi_{\\perp} \\gg \\xi_{\\parallel} \\mapsto 0 \\ F;  \\ P $&$\\xi_{\\perp} > \\xi_{\\parallel}  \\ F;  \\ P$&$\\xi_{\\parallel}  \\geq \\xi_{\\perp}  \\ P \\ \\ \\ \\ $&$\\xi_{r}  \\sim   \\xi_{\\phi}\\gg \\xi_{z}\\mapsto 0 $&$\\xi_{r}  \\sim   \\xi_{\\phi}\\gg \\xi_{z}\\mapsto 0$&$  \\xi_{r}>  \\xi_{\\phi} > \\xi_{z}   $ \\\\  \\hline $P_{16}$&$\\xi_{\\perp} \\gg \\xi_{\\parallel} \\mapsto 0 \\ F;  \\ P $&$\\xi_{\\perp} > \\xi_{\\parallel}  \\ F;  \\ P$&$\\xi_{\\parallel}  \\geq \\xi_{\\perp}  \\ P \\ \\ \\ \\ $&$\\xi_{r}  \\sim   \\xi_{\\phi}\\gg \\xi_{z}\\mapsto 0  $&$\\xi_{r} \\sim    \\xi_{\\phi}\\gg  \\xi_{z}\\mapsto 0  $&$  \\xi_{r}>  \\xi_{\\phi} > \\xi_{z}   $ \\\\  \\hline $P_{17}$&$\\xi_{\\perp} \\gg \\xi_{\\parallel} \\mapsto 0  \\ F; \\ P $&$\\xi_{\\perp} > \\xi_{\\parallel}  \\ F;  \\ P$&$\\xi_{\\parallel}  \\geq \\xi_{\\perp}  \\ P \\ \\ \\ \\ $&$\\xi_{r} \\sim  \\xi_{\\phi}\\gg    \\xi_{z}\\mapsto 0$&$\\xi_{r}   \\sim  \\xi_{\\phi}\\gg  \\xi_{z}\\mapsto 0 $&$  \\xi_{r}>  \\xi_{\\phi} > \\xi_{z}   $ \\\\  \\hline $P_{18}$&$\\xi_{\\perp} \\gg \\xi_{\\parallel} \\mapsto 0  \\ F; \\ P $&$\\xi_{\\perp} > \\xi_{\\parallel}   \\ F; \\ P$ &$\\xi_{\\parallel}  \\geq \\xi_{\\perp}  \\ P \\ \\ \\ \\ $&$\\xi_{r}  \\sim   \\xi_{\\phi}\\gg  \\xi_{z}\\mapsto 0  $&$\\xi_{r}  \\sim  \\xi_{\\phi}\\gg  \\xi_{z}\\mapsto 0 $&$  \\xi_{r}>  \\xi_{\\phi} > \\xi_{z}   $\\\\  \\hline \\end{tabular} \\caption{\\label{tab:table2} First Panel: Intensity relationship between the tangential and normal to the field components of the eighteen   first periods for $B_{0}=10G$ and for  weak (first column), moderate (second column) and strong helicity (third column) cases.  The (P) indicates in phase and (IP) indicates inverted phase. Second Panel: Intensity relationship between the cylindrical  components of the eighteen   first periods for $B_{0}=10G$ and for  weak (first column), moderate (second column) and strong helicity (third column) cases.} \\end{table*}     \\begin{table*} \\begin{tabular}{cccccccc} \\hline $R$&$L $&$R/2L$&$Twist=bR \\ \\ \\ \\ \\ \\ \\ \\ \\ $&$R$&$L $&$R/2L$&$Twist=bR$ \\\\  \\hline $0.01$& $ 9.05 \\ 10^{7} $&$0.005 $&$  0.028\\ \\ \\ \\ \\ \\ \\ \\ \\ $&$0.01$& $ 8.07 \\ 10^{7} $&$0.005 $&$  0.28$\\\\  \\hline $0.02$&$ 9.04 \\ 10^{7}  $&$ 0.01  $&$ 0.057  \\ \\ \\ \\ \\ \\ \\ \\ \\ $&$0.015$&$ 8.32 \\ 10^{7}  $&$ 0.008  $&$ 0.43  $\\\\  \\hline $0.03$&$ 9.02 \\ 10^{7} $&$ 0.015 $&$ 0.085 \\ \\ \\ \\ \\ \\ \\ \\ \\ $&$0.02$&$ 7.86 \\ 10^{7} $&$ 0.011 $&$ 0.57$\\\\  \\hline $0.04  $&$ 8.99 \\ 10^{7}  $&$ 0.02 $&$ 0.11\\ \\ \\ \\ \\ \\ \\ \\ \\ $&$0.025  $&$ 7.38 \\ 10^{7}  $&$ 0.015 $&$ 0.71$\\\\  \\hline $0.05  $&$ 8.96 \\ 10^{7}  $&$ 0.025 $&$ 0.14\\ \\ \\ \\ \\ \\ \\ \\ \\ $&$0.03  $&$ 6.88 \\ 10^{7}  $&$ 0.02 $&$ 0.85$\\\\  \\hline $0.06  $&$ 8.9 \\ 10^{7}  $&$ 0.03 $&$ 0.17\\ \\ \\ \\ \\ \\ \\ \\ \\ $&$0.04  $&$ 5.97 \\ 10^{7}  $&$ 0.03 $&$ 1.13$\\\\  \\hline $0.1  $&$ 8.7 \\ 10^{7}  $&$ 0.05 $&$ 0.28\\ \\ \\ \\ \\ \\ \\ \\ \\ $&$0.05  $&$ 5.2 \\ 10^{7}  $&$ 0.04 $&$ 1.42$\\\\  \\hline \\hline \\end{tabular} \\caption{\\label{tab:table5} First Panel - Stable case: Variation of the Radius with the Twist for weak helicity $b=0.05$ and  $B_{0}=10G$. Second Panel - Unstable case: Variation of the Radius with the Twist for moderate helicity $b=0.5$ and  $B_{0}=10G$.} \\end{table*}"
    },
    "0807/0807.0935_arXiv.txt": {
        "abstract": "We present a study of flux-calibrated low-resolution optical spectroscopy of ten stars belonging to eight systems in the $\\sim 5$ Myr-old $\\epsilon$ Chamaeleontis ($\\epsilon$ Cha) pre-main-sequence (PMS) star cluster. Using synthetic broadband colours, narrow-band continuum, atomic and molecular lines derived from the spectra, we compare the $\\epsilon$ Cha stars to a slightly older PMS cluster, the $\\approx 8$ Myr-old $\\eta$ Cha cluster, and to main-sequence dwarfs. Using synthetic {\\it VRI\\,} colours and other indices, we find that the relationship between broadband colours and spectroscopic temperature indicators for $\\epsilon$ Cha cluster members is indistinguishable from that of Gyr-old dwarfs. This result is identical to that found earlier in $\\eta$ Cha. Gravity-sensitive line indices place the cluster between the giant and dwarf sequences, and there is clear evidence that $\\epsilon$ Cha stars have lower surface gravity than $\\eta$ Cha stars. This result is consistent with $\\epsilon$ Cha being the slightly younger PMS association, a few Myr younger according to the Hertzsprung-Russell (HR) diagram placement of these two clusters and comparison with PMS evolutionary grids. Late M-type $\\epsilon$ Cha cluster members show a $B$-band flux excess of $\\approx 0.2$ mag compared to observations of standard dwarfs, which might be related to enhanced magnetic activity. A similar level of excess $B$-band emission appears to be a ubiquitous feature of low mass members of young stellar populations with ages less than a few hundred Myr, a very similar timescale to the PMS phase of elevated relative X-ray luminosity. ",
        "introduction": "Analysis of flux-calibrated low-resolution optical spectra is a useful method for investigating the physical properties of stars in comparison with spectra of standard stars with well-defined properties. For nearby stars and stellar associations, such observations are readily obtained even with telescopes of modest aperture (e.g., Bessell 1991; Lyo, Lawson \\& Bessell 2004). Using synthetic broadband colours, narrow-band continuum, atomic and molecular line indices derived from calibrated spectra, we can study temperature, surface gravity and metallicity effects, e.g. Lyo et al. (2004) showed that the relationship between colours and spectroscopic temperature indicators for the $\\approx 8$ Myr-old $\\eta$ Cha cluster was indistinguishable from that for Gyr-old disk dwarfs, although the spectra showed some evidence for higher metallicity and clear evidence for lower gravity. The latter is a consequence of the elevated location of these stars, several mag above the zero main sequence in the HR diagram. The $\\eta$ Cha stars also displayed, like other young stellar populations, a $B$-band flux excess attaining $\\approx$ 0.2 mag for late-M cluster members.  By studying PMS clusters of different isochronal ages, we can use spectrophotometric techniques to address various stellar evolutionary issues, e.g. is there evidence in the spectroscopy for differences in the temperature-spectral type sequence as a function of age? Do gravity-dependent features scale with age as stellar low mass stars descend their Hayashi tracks? Is the observed $B$-band excess in late-M PMS stars a general property of young stellar populations, and does it vary as a function of spectral type and age?  In this paper, our target group is the young $\\epsilon$ Cha cluster associated with the early-type system $\\epsilon$ Cha AB and the codistant, comoving star HD104237 (DX Cha). HD104237 is the nearest-known Herbig Ae star and forms at least a quintet with low mass companions based upon {\\it Chandra X-ray Observatory} observations, optical/infrared imaging and spectroscopic study, some of which are likely multiple themselves (Feigelson et al. 2003; Grady et al. 2004). HD104237A itself is a spectroscopic binary with a K3 companion in an eccentric 20-d orbit (B\\\"ohm et al. 2004). The group is nearby ($d \\approx 114$ pc), compact in extent ($\\approx 1$ pc) and sparsely populated, containing eight stellar systems in its central region with stars ranging in spectral type from B9 to M5. HR diagram placement of the members ages the group at $3-5$ Myr (Feigelson et al. 2003) when compared to Siess, Dufour \\& Forestini (2000) evolutionary tracks, and $\\sim 6$ Myr (Luhman 2004) when using a mix of Palla \\& Stahler (1999), Baraffe et al. (1998) and Chabrier et al. (2000) tracks. Both the HR diagram comparisons of Feigelson et al. (2003) and Luhman (2004) assign younger ages to the early-type members of $2-3$ Myr. This discrepancy in age with spectral type might be a consequence of nearby low-mass companions elevating the luminosities of the early-type stars, or might serve to highlight differences between model isochrones and observational isochrones, or indicate a genuine age difference between the high-mass and low-mass stars in this cluster. Kinematic study indicates the group is another out-lying population of the Oph-Sco-Cen OB association (Feigelson et al. 2003), following groups and associations such as $\\eta$ Cha, the TW Hya association, and the $\\beta$ Pic moving group.  The $\\epsilon$ Cha cluster presents itself then, as a slightly younger analogue of the $\\eta$ Cha star cluster (Mamajek, Lawson \\& Feigelson 1999, 2000), an $\\approx 8$ Myr-old compact, sparse PMS group that was studied using spectrophotometric techniques by Lyo et al. (2004). The ages of PMS clusters derived by HR diagram comparisons using competing evolutionary grids are unreliable; differences can exceed several Myr at $\\sim 10$ Myr (see fig. 8 of Lyo et al. 2004). However, comparison between PMS groups using techniques such as spectrophotometry ought to reveal differences that can be interpreted as age-related trends. With an age difference of only a few Myr, comparison between the $\\epsilon$ Cha and $\\eta$ Cha groups is a demonstration of the sensitivity of such methods to rank in age various PMS populations. ",
        "conclusions": "Flux-calibrated low-resolution optical spectroscopy is an indispensable tool for investigating the physical properties of stars by comparison with other stellar groups and standard star calibrators. In this paper, we characterized the stellar population of the $\\sim 5$ Myr-old $\\epsilon$ Cha cluster via a number of temperature- and gravity-sensitive spectroscopic indicators, and compared its properties to the slightly older PMS group associated with $\\eta$ Cha, and to those of Gyr-old disk dwarfs.  Using synthetic broadband colours, narrow-band continuum, atomic and molecular line indices derived from the spectra, we find that the relationship between the broadband colours and spectroscopic temperature indicators for the $\\epsilon$ Cha cluster stars is indistinguishable from that of the Gyr-old dwarfs. We had previously reached the same conclusion for the slightly-older $\\eta$ Cha cluster (Lyo et al. 2004). However, there is a clear evidence that $\\epsilon$ Cha cluster stars have lower surface gravity than $\\eta$ Cha cluster stars from measurement of the gravity-sensitive $\\lambda\\lambda$8183, 8195 Na I doublet and FeH molecular spectral lines. This result is consistent with the $\\epsilon$ Cha cluster being slightly younger than $\\eta$ Cha, a few Myr younger according to the HR diagram placement of these two clusters and comparison with PMS evolutionary model grids.  We also found a $B$-band excess of $\\sim 0.2$ mag in the late M-type cluster members, similar to that found in $\\eta$ Cha and in other PMS populations with ages less than $\\sim 200$ Myr. This result suggests that the blue excess is an ubiquitous property of low-mass young stellar objects, with the most-likely origin being enhanced magnetic activity. The presence of significant excess blue emission appears to closely parallel the phase of high relative X-ray luminosity seen in low-mass $1-100$ Myr-old PMS populations."
    },
    "0807/0807.2271_arXiv.txt": {
        "abstract": "We report on the discovery of a bright \\lya\\ blob associated with the $z=3$ quasar SDSS~J124020.91$+$145535.6 which is also coincident with strong damped \\lya\\ absorption from a foreground galaxy (a so-called proximate damped \\lya\\ system; PDLA).  The one dimensional spectrum acquired by the Sloan Digital Sky Survey (SDSS) shows a broad \\lya\\ emission line with a FWHM $\\simeq 500\\mkms$ and a luminosity of $L_{\\rm Ly\\alpha} = 3.9 \\sci{43} {\\rm erg \\, s^{-1}}$ superposed on the trough of the PDLA.  Mechanisms for powering this large \\lya\\ luminosity are discussed. We argue against emission from \\ion{H}{2} regions in the PDLA galaxy since this requires an excessive star-formation rate $\\sim 500\\,\\msol \\rm yr^{-1}$ and would correspond to the largest \\lya\\ luminosity ever measured from a damped \\lya\\ system or starburst galaxy.  We use a Monte Carlo radiative transfer simulation to investigate the possibility that the line emission is fluorescent recombination radiation from the PDLA galaxy powered by the ionizing flux of the quasar, but find that the predicted \\lya\\ flux is several orders of magnitude lower than observed. We conclude that the \\lya\\ emission is not associated with the PDLA galaxy at all, but instead is intrinsic to the quasar's host and similar to the extended \\lya\\ ``fuzz'' which is detected around many AGN. PDLAs are natural coronagraphs that block their background quasar at \\lya\\ , and we discuss how systems similar to SDSS~J124020.91$+$145535.6 might be used to image the neutral hydrogen in the PDLA galaxy \\emph{in silhouette} against the screen of extended \\lya\\ emission from the background quasar. ",
        "introduction": "The ionizing radiation emitted by a luminous quasar can, like a flashlight, illuminate hydrogen clouds in its vicinity, teaching us about the size, kinematic structure, and density of the gas which surrounds it \\citep{rees88,hr01}. This is because recombinations from photoionized hydrogen ultimately produce Ly$\\alpha$ photons -- a fraction of the energy in the quasar's UV continuum is `focused' into fluorescent line radiation and re-emitted, allowing us to study the physical conditions in the emitting gas.  Extended \\lya\\ nebulae with luminosities up to $L_{\\rm Ly\\alpha}=5\\sci{42}~{\\rm erg~s^{-1}}$ and angular sizes as large $\\sim 10\\arcsec$ have been observed around many quasars \\citep[e.g.][]{djorgovski85,heckman91a,cjw+06}, but the physical mechanism powering this emission is still unclear.  In addition to fluorescent emission powered by photoionization, other mechanisms include emission from material shocked by radio-jets or starburst outflows or \\lya\\ cooling radiation from gravitational collapse.  The same set of physical mechanisms have been proposed to explain the analogous \\lya\\ nebulae observed around a subset of luminous radio galaxies \\citep{msd+90,villar07}, as well as the recently discovered \\lya\\ `blobs' \\citep{fmw99,steidel00,matsuda04}. The primary difference between the extended ``fuzz'' around quasars as compared to the nebulae around radio galaxies and \\lya\\ blobs is that a luminous quasar and hence the source of ionizing photons is directly detected in the former but not in the latter. Obscuration and orientation effects, as are often invoked in unified models of AGN, could be responsible for this difference. This explanation is plausible considering that the large energies in the jets $\\gtrsim 10^{60}$\\,erg \\citep[e.g.][]{Miley80} of luminous radio galaxies strongly suggest nuclear activity and evidence for an obscured AGN has been uncovered in several of the \\lya\\ blobs \\citep{csw+04,dbs+05,gsc+07}.  In the course of a survey for damped \\lya\\ absorption proximate to high $z$ quasars \\citep[][see also Ellison et al. 2002 and Russell et al. 2006]{phh08}, we discovered an extremely luminous \\lya\\ blob coincident with both the redshift of a $z=3$ proximate damped \\lya\\ system (PDLA)\\footnote{ A PDLA is defined as an absorber with $\\mnhi \\ge 2\\sci{20} \\cm{-2}$ located within $\\delta v = 3000 \\mkms$ of its background quasar} and its background quasar.  Although intervening DLAs ($\\delta v > 3000 \\mkms$) rarely show \\lya\\ emission \\citep{scb+89,mff04,kwy+06}, PDLAs appear to preferentially exhibit \\lya\\ emission superimposed on their \\lya\\ absorption trough \\citep{mw93,mwf98,eyh+02} which is $\\sim 5$ times brighter than the few intervening detections. Although based on poor statistics and heterogeneous samples, this putative discrepancy suggests that the \\lya\\ emission associated with PDLAs might be powered by its background quasar. In what follows, we analyze the basic characteristics of this absorber and discuss several physical scenarios for the origin of this \\lya\\ blob.  Finally, we emphasize several novel applications that exploit this unique configuration. Throughout this paper we use the best fit WMAP3 cosmological model of \\citet{wmap03}, with $\\Omega_m = 0.240$, $\\Omega_\\Lambda =0.76$, $h=0.73$.  % ",
        "conclusions": "(1) the gas is located within approximately 1\\,Mpc of the quasar; (2) the majority of PDLA systems lying within 1\\,Mpc of the quasar should also show significant \\ion{N}{5} absorption.  We explore the fluorescence scenario further by performing the following numerical simulation.  The PDLA system is modeled as a singular isothermal sphere within a halo of $10^{11}\\msol$, for which 5\\% of the mass is assumed to be in gas. The system is placed between the observer and the quasar with a distance 300\\,kpc to the quasar such that the PDLA's backside is illuminated by the quasar. We use the code developed in \\citet[][see also Cantalupo et al. 2005]{Juna08} to perform the self-shielding calculation, solving for the ionization structure of the gas, the recombination rate, and the resulting \\lya\\ luminosity. The radiative transfer of the fluorescent \\lya\\ photons is solved with a Monte Carlo method \\citep{zheng02,Juna08}.  In turn, we produce a map of the expected \\lya\\ surface-brightness as shown in Figure~\\ref{fig:theory}.  We have extracted several one-dimensional spectra from this surface-brightness image to illustrate the flux level and the expected emission line profiles.  Although the one-dimensional spectrum through the $3''$ SDSS fiber aperture has roughly the observed line width, it has an integrated flux that is several orders of magnitude lower than observed.  The low flux results from the fact that we are observing the illuminated galaxy in a `new moon' configuration; the optically thick regions which give rise to the \\lya\\ emission also serve to shield the radiation from our vantage point.  The observed flux is instead dominated by the outer regions of fiber which cover optically thin regions of gas in the halo surrounding the galaxy (aperture A2 in Figure~\\ref{fig:theory}). Based on this calculation, we conclude that the observed \\lya\\ emission in \\name\\ is unlikely to be fluorescent recombination radiation from the PDLA galaxy.   {\\bf Extended \\lya\\ ``Fuzz'' from the Quasar Host Halo:} We consider the most plausible hypothesis to be that the \\lya\\ emission is not associated with the PDLA galaxy at all, but instead is intrinsic to the quasar's host galaxy. Indeed, \\lya\\ emission has been detected as an extended `fuzz' around many quasars \\citep[e.g.][]{djorgovski85,heckman91a,cjw+06}, and with luminosities comparable to or greater than the $L_{\\rm Ly\\alpha} = 3.9\\sci{43} {\\rm erg \\, s^{-1}}$ measured for \\name. The velocity widths of this extended emission can be as large as $1000-1500$~\\kms consistent with our observed kinematic profile. The physical mechanism responsible for this emission could be fluorescent recombination radiation from gas in the quasar host halo illuminated by quasar's large ionizing flux \\citep{rees88,heckman91a,hr01,weidinger04}.  It is important to explain the distinction between the fluorescent \\lya\\ fuzz as it occurs in the quasar host halo to the fluorescence from the DLA galaxy which was considered in the previous section. For the former case the emission is presumed to be coming from small dense self-shielding clouds which permeate the quasar halo but have a small $\\sim 0.5\\%$ covering factor \\citep{msd+90,heckman91a} deduced from the fraction of the ionizing continuum which is being absorbed \\citep{heckman91a}. In the latter case, where a galaxy is illuminated, the emission comes from a large, kpc-scale self-shielding cloud and a large fraction $\\eta \\simeq 60\\%$ of the impingent ionizing continuum is converted into \\lya\\ recombination photons.  But the mechanism powering the extended emission line `fuzz' could be among other mechanisms put forth to explain the extended \\lya\\ halos around quasars \\lya\\ blobs, and high-redshift radio galaxies, such as gas shock-heated by a large scale outflow or cooling radiation from gravitational infall \\citep[e.g.][ and references therein]{matsuda04,ctf+06}.  In the halo fuzz scenario, the DLA-galaxy is in the foreground of the extended \\lya\\ emitting region of the quasar host, and should hence absorb the fraction of the \\lya\\ emission which it covers.  This configuration also offers a reasonable explanation for why PDLAs might preferentially exhibit \\lya\\ emission relative to intervening DLAs: the presence of a PDLA occults the bright quasar making it much easier to detect faint extended emission from the quasar halo."
    },
    "0807/0807.1272_arXiv.txt": {
        "abstract": " ",
        "introduction": "\\label{s:intro} Planetary systems are not just made up of \\textit{planets}, but are also composed of numerous small bodies ranging from asteroids and comets as large as 1000 km down to sub-$\\mu$m-sized dust grains. In the solar system the asteroids and comets are confined to relatively narrow rings known as the asteroid belt and the Kuiper belt (see chapters by Nakamura and Jewitt). These belts are the source of the majority of the smaller objects seen in the solar system, since such objects are inevitably created in collisions between objects within the belts (see chapter by Michel). Sublimation of comets as they are heated on approach to the Sun is another source of dust in the solar system.  It is known that extrasolar systems also host belts of planetesimals (a generic name for comets and asteroids) that are similar to our own asteroid belt and Kuiper belt. These were first discovered using far-IR observations of nearby stars, which showed excess emission above that expected to come from the stellar photosphere \\cite{auma84}. This emission comes from dust that is heated by the star and which re-radiates that energy in the thermal infrared, at temperatures between 40-200 K, depending on the distance of the dust from the star. The lifetime of the dust is inferred to be short compared with the age of the star, and so it is concluded that the dust cannot be a remnant of the proto-planetary disk that formed with the star (see chapter by Takeuchi), rather it must originate in planetesimal belts much in the same way that dust is created in the solar system \\cite{bp93}.  Over 300 main sequence stars are now known with this type of excess emission \\cite{mb98,su06,bryd06}, and such objects are either known as Vega-like (after the first star discovered to have this excess), or as debris disks. Statistical studies have shown that $\\sim 15$\\% of normal main sequence stars have debris disks, although it should be stressed that the disks which can be detected with current technology have greater quantities of dust than is currently present in the solar system by a factor of at least 10 \\cite{gwhd04}. Nevertheless this indicates that debris disks are common, more common in fact that extrasolar planets which are found around $\\sim 6$\\% of stars \\cite{fv05}. Studying these disks provides a unique insight into the structure of the planetary systems of other stars. Indeed, the nearest and brightest debris disks can be imaged, and such studies have provided the first images of nearby planetary systems. These images reveal the distribution of dust in the systems, which can in turn be used to infer the distribution of parent planetesimals, and also the architecture of the underlying planetary system. However, to do so requires an understanding of both the mechanism by which dust is produced in planetesimal belts and its consequent dynamical evolution, as well as of the dynamical interaction between planets and planetesimals and between planets and dust.  This chapter reviews our knowledge of debris disks from observations (\\S \\ref{s:obs}) and describes a simple model for planetesimal belt evolution which explains what we see (\\S \\ref{s:mod}), as well as how the detailed interaction between planets and planetesimals imposes structure on that planetesimal belt (\\S \\ref{s:plpl}), and how those perturbations translate into structures seen in the dust distribution (\\S \\ref{s:pldust}). Conclusions, including what has been learned about the planetary systems of nearby stars from studying these disks, are given in \\S \\ref{s:conc}. ",
        "conclusions": "\\label{s:conc} This chapter has considered the types of structures seen in the dusty debris disks of nearby stars (\\S \\ref{s:obs}) and how those structures can be used to determine the layout of their planetary systems, in terms of the distributions of both planetesimals and planets. The text has dwelled on the successes of the models at explaining the observed structures, because this illustrates the elements that are essential to any debris disk model if the observations are to be successfully explained (\\S \\ref{s:mod}), and because we are confident that we understand how a planet would perturb a planetesimal belt in an idealised system comprised of just one planet (\\S \\ref{s:plpl}) and to some extent how to extrapolate that to consider how the planet would affect the observed dust disk (\\S \\ref{s:pldust}). To summarise what we have learned: \\textbf{(i)} the axisymmetric structure of debris disks can mostly be explained by a model in which dust is created in collisions in a narrow planetesimal belt and is subsequently acted on by radiation forces; \\textbf{(ii)} the asymmetric structure of debris disks can mostly be explained by secular and resonant gravitational perturbations from unseen planets acting on the planetesimal belt and dust derived from it.  Knowing the radial location of the planetesimal belts is important because this demonstrates where in a protoplanetary disk grain growth must have continued to km-sized planetesimals \\cite{wd02}, and by analogy with the solar system there is reason to believe that the location of the planetesimal belts tells us indirectly the whereabouts of unseen planets, although it is worth bearing in mind that there may be alternative explanations for gaps in the planetesimal distribution related to the physics of the protoplanetary disk. Nevertheless, it appears that where we have the capability to look for detailed disk structure, there is good correspondence between the asymmetric structures observed with those expected if there are planets in these systems. The modelling is also sufficiently advanced that the disk structure can be used to infer information on the properties of the perturbing planets (such as the planet's mass, orbit and even evolutionary history). The planet properties which have been inferred in this way are particularly exciting when compared with those of exoplanets discovered using the radial velocity and transit techniques. Figure \\ref{fig:exopl} shows how the debris disk planets are similar to Uranus and Neptune in the solar system, occupying a unique region of parameter space. This is possible because the large size of debris disks means that the planets perturbing them are most often at large orbital radii, and it is easy for planets as small as Neptune to impose structure on a debris disk. There is also the tantalising possibility that in the future debris disk structures can be used to identify planets analogous to the Earth and Venus in extrasolar systems.  \\begin{figure} \\centering \\begin{tabular}{c} \\includegraphics[height=8.3cm]{wyattfig12.ps} \\end{tabular} \\caption{Distribution of planet masses and semimajor axes. Solar system planets are plotted as open circles, and those known from radial velocity and transit studies with a plus (taken from the list on http://exoplanets.eu dated 24 May 2007). The shaded region shows the current limits of radial velocity surveys for sun-like stars. Debris disk planets inferred from disk structure (all awaiting confirmation) are shown with filled circles. References for the plotted planet parameters are: HR4796 \\cite{wdtf99}, $\\epsilon$ Eridani \\cite{ogmt00}, Vega \\cite{wyat03}, HD141569 \\cite{wyat05b}, $\\eta$ Corvi \\cite{wgdc05}, Fomalhaut \\cite{quil06}, $\\beta$ Pictoris \\cite{fkl07}, although it should be noted that these parameters, particularly planet mass, are often poorly constrained.} \\label{fig:exopl} \\end{figure}  However, while it is incontrovertible that if there are planets present they would impose structure on a disk, the question of whether we have already seen these structures in extrasolar systems is still a matter for debate. In many cases the presence of an unseen planet is the only explanation for the observed structures, but that does not mean that it has to be the right explanation. The problem is that it is hard to confirm that the planets are there, since they lie beyond the reach of radial velocity studies (see Fig.~\\ref{fig:exopl}). Direct imaging could detect planets at this distance if they were a few times Jupiter mass \\cite{mhw03}, but not if they are Neptune mass. Thus the onus is on the models to make other testable predictions, and some of these have already been made (such as the orbital motion of the clumpy structures, and the disk structures expected to be seen at different wavelengths) and will be tested in the coming years. If these planets are confirmed, their addition onto plots like that shown in Figure \\ref{fig:exopl} will be invaluable for constraining planet formation models \\cite{il04}.  It is also important to remember that this theory cannot yet predict the quantities of small grains we would expect to see in any given disk. There are too many uncertainties regarding the dust production mechanisms, and it is possible that these processes may differ among stars with, e.g., different dust compositions. Applying dynamical models of the kind presented in \\S \\ref{s:mod} to a greater number of resolved disk observations will help to understand these differences. However, there is still the possibility that the problem is more fundamental in a way which is best illustrated by the archetypal debris disk Vega. The observed mass loss rate from $\\beta$ meteoroids in this system is $2M_\\oplus$/Myr, which indicates that this must be a transient, rather than a steady state, component \\cite{su05}. It is thus possible that the small grain population in debris disks is inherently stochastic, perhaps influenced by input from recent massive collisions \\cite{tele05}. Fortunately it appears that the large grain component of the majority of debris disks is evolving in steady state \\cite{wssr07} and so can be understood within the framework described in this chapter, and the same is likely also true for the small grain component (it is just the relative quantities of the different components that is less certain).  However, the possibility must be considered that in some systems the observed dust is transient in such a way that its origin will require a significant overhaul to the models presented here. For example, there are a few cases of sun-like stars surrounded by hot dust (e.g., \\S \\ref{ss:axi}) which cannot be maintained by steady state production in a massive asteroid given the age of the stars \\cite{wsgb07}. It is not clear what the origin of the transient event producing the dust is. However, it is known that the quantity of planetesimals in the inner solar system has had a stochastic component, notably involving a large influx $\\sim 700$ Myr after the solar system formed in an event known as the late heavy bombardment, the origin of which is thought to have been a dynamical instability in the architecture of the giant planets \\cite{gltm05}. So perhaps these systems are telling us about the more complex dynamics of their planetary systems. Given the complexity of planetary systems it seems inevitable that the models presented in this chapter are just the start of a very exciting exploration of the dynamics of extrasolar planetary systems."
    },
    "0807/0807.1758_arXiv.txt": {
        "abstract": "We investigate the formation of carbon-enhanced metal-poor (CEMP) stars via the scenario of mass transfer from a carbon-rich asymptotic giant branch (AGB) primary to a low-mass companion in a binary system. We explore the extent to which material accreted from a companion star becomes mixed with that of the recipient, focusing on the effects of thermohaline mixing and gravitational settling. We have created a new set of asymptotic giant branch models in order to determine what the composition of material being accreted in these systems will be. We then model a range of  CEMP systems by evolving a grid of models of low-mass stars, varying the amount of material accreted by the star (to mimic systems with different separations) and also the composition of the accreted material (to mimic accretion from primaries of different mass). We find that with thermohaline mixing alone, the accreted material can become mixed with between 16 and 88\\% of the pristine stellar material of the accretor, depending on the mass accreted and the composition of the material. If we include the effects of gravitational settling, we find that thermohaline mixing can be inhibited and, in the case that only a small quantity of material is accreted, can be suppressed almost completely. ",
        "introduction": "Carbon-enhanced, metal-poor (CEMP) stars are defined as stars with [C/Fe]\\footnote{[A/B] = $\\log (N_\\mathrm{A}/N_\\mathrm{B}) - \\log (N_\\mathrm{A}/N_\\mathrm{B})_\\odot$}$>$+1.0 \\citep{2005ARA&A..43..531B}, with [Fe/H]$<-2$ in most cases. These objects appear with increasing frequency at low metallicity \\citep{2006ApJ...652L..37L}. The study of CEMP stars is being used to probe conditions in the early universe. For example, CEMP stars have been used to infer the initial mass function in the early Galaxy \\citep[e.g.][]{2005ApJ...625..833L}. Chemical abundance studies have revealed that the majority of the CEMPs are rich in $s$-process elements like barium \\citep{2003IAUJD..15E..19A}, forming the so-called CEMP-s group. Recent survey work has detected a binary companion in around 68\\% of these CEMP-s stars and this is consistent with them all being in binary systems \\citep{2005ApJ...625..825L}.  Binary systems provide a natural explanation for these objects, which are of too low a luminosity to have been able to produce their own carbon. The primary of the system was an asymptotic giant branch (AGB) star which became carbon-rich through the action of third dredge-up\\footnote{Third dredge-up occurs when the convective envelope of an AGB star deepens after a thermal pulse and material that has experienced nuclear burning is brought to the surface.} and transferred material on to the low-mass secondary (most likely via a stellar wind). The primary became a white dwarf and has long since faded from view, with the carbon-rich secondary now being the only visible component of the system. It has commonly been assumed that the accreted material remains on the surface of the secondary until the star ascends the giant branch, at which point the deepening of the convective envelope (referred to as first dredge-up because material that has experienced CN-cycling in the stellar interior is brought to the surface) mixes the material with the interior of the star. However, the transferred material should become mixed with the interior of the accreting star via the process of thermohaline mixing \\citep[see e.g.][]{2007A&A...464L..57S, 2004MNRAS.355.1182C,2003NewA....8...23B}. This occurs when the mean molecular weight of the stellar gas increases toward the surface. A gas element displaced downwards and compressed will be hotter than its surroundings. It will therefore lose heat, become denser and continue to sink. This leads to mixing on thermal timescales until the molecular weight difference is eliminated \\citep{1980A&A....91..175K}. \\citet{2007A&A...464L..57S} showed that the inclusion of thermohaline mixing could result in the accreted material being mixed throughout 90\\% of the star.  Recent work has questioned the efficiency of thermohaline mixing in carbon-enhanced metal-poor stars. Using a sample of barium-rich CEMP stars, \\citet{2008ApJ...678.1351A} showed that the distribution of [C/H] values in turn-off stars (i.e. those stars that have reached the end of their main-sequence lives, are still of low-luminosity and have yet to become giants) was different from that in giants suggesting that significant mixing only happened at first dredge-up. A similar point was made by \\citet{2007arXiv0709.4240D} using the data of \\citet{2006ApJ...652L..37L}. These authors showed that the turn-off stars and giants were consistent with coming from the same distribution if first dredge-up resulted in the [C/H] value (and also the [N/H] value) being reduced by around 0.4 dex. They find that this result is consistent with having an accreted layer of material mixed to an average depth of about 0.2\\ms\\ (or alternatively having an accreted layer of 0.2\\ms\\ that remains unmixed).  Neither of these scenarios is consistent with the extensive mixing found by \\citet{2007A&A...464L..57S}.  A possible source for reduced thermohaline mixing efficiency has been suggested by \\citet{2008ApJ...677..556T}. These authors suggest that the action of gravitational settling will alter the composition gradient of the accreting star near its surface. Helium will settle from the surface, reducing the mean molecular weight at the surface but leading to an increase in the layers beneath. This produces a small region in which the mean molecular weight, $\\mu$, decreases outwards toward the stellar surface -- a situation which is stable to thermohaline mixing. This stabilising composition gradient (a so-called `$\\mu$-barrier') can inhibit the process of thermohaline mixing.  This paper extends the work of \\citet{2007A&A...464L..57S}, examining the effect of varying the composition of the accreted material (mimicking accretion from different masses of companion) and the amount of material that is accreted. We also investigate the effect that gravitational settling has on the extent to which material is mixed. ",
        "conclusions": "We have modelled the accretion of AGB material on to low-mass stars in order to determine what chemical signatures may be observed in CEMP stars. We have examined three specific cases: canonical evolution including only convective mixing, the inclusion of thermohaline mixing and the inclusion of both thermohaline mixing and gravitational settling. We find that thermohaline mixing can lead to accreted material being mixed with between 16 and 88\\% of the accreting star, depending on the mass and composition of the accreted material. When gravitational settling is included, thermohaline mixing is severely inhibited when only a small amount of material (around a few $10^{-3}$\\ms) is accreted because of the presence of a $\\mu$-barrier formed by the settling of helium. This barrier is less and less effective as the amount of accreted material is increased. It is also less effective for more massive companions as the barrier has had less time to be established before accretion on to the secondary occurs.  Models without thermohaline mixing produce turn-off objects with [C/H] values that are too high to match observations. Unless a substantial quantity of material has been accreted, these models also suffer from too much dilution at first dredge-up. Models with thermohaline mixing cannot reproduce the highest [C/H] values observed and predict the existence of low [C/H] objects which are not observed. In order to reproduce the highest [C/H] values, which seem to require the accretion of a large quantity of material from a companion, some alternative mechanism to suppress thermohaline convection is required. The inclusion of gravitational settling can solve the problem of the low [C/H] objects at low luminosity as the $\\mu$-barrier prevents mixing when the amount of accreted material is small. However, the inclusion of this physics presents another serious problem namely that carbon will settle from the surface during the main sequence, making it extremely difficult to form turn-off stars with [C/H] values of -1--0.  There also appears to be a problem with the abundance predictions of the AGB models. They do not predict substantial enhancements of both carbon {\\it and} nitrogen at the same time (except in a very narrow mass range). Further work needs to be done to find a mechanism capable of producing the observed abundance patterns."
    },
    "0807/0807.4382_arXiv.txt": {
        "abstract": "We re-analyse all of the archive observations of the Ophiuchus dark cloud L1688 that were carried out with the submillimetre common-user bolometer array (SCUBA) at the James Clerk Maxwell Telescope (JCMT). For the first time we put together all of the data that were taken of this cloud at different times to make a deeper map at 850$\\mu$m than has ever previously been published. Using this new, deeper map we extract the pre-stellar cores from the data. We use updated values for the distance to the cloud complex, and also for the internal temperatures of the pre-stellar cores to generate an updated core mass function (CMF). This updated CMF is consistent with previous results in so far as they went, but our deeper map gives an improved completeness limit of 0.1M$_{\\sun}$ (0.16 Jy), which enables us to show that a turnover exists in the low-mass regime of the CMF. The L1688 CMF shows the same form as the stellar IMF and can be mapped onto the stellar IMF, showing that the IMF is determined at the prestellar core stage. We compare L1688 with the Orion star-forming region and find that the turnover in the L1688 CMF occurs at a mass roughly a factor of two lower than the CMF turnover in Orion. This suggests that the position of the CMF turnover may be a function of environment. ",
        "introduction": "Star formation in molecular clouds occurs within prestellar cores, which are gravitationally bound cores within the clouds \\citep{1994MNRAS.268.276W,1996A&A.314.625A,1999MNRAS.305.143W, 2002Sci...295...76W, 2007prpl.conf...17D, 2007prpl.conf...33W}.  A number of observations have shown that the core mass function (CMF) of prestellar cores appears to mimic \\citep[Motte Andre \\& Neri 1998, hereafter MAN98,][]{1998ApJ...508L..91T, 2000ApJ...545..327J, 2001A&A...372L..41M, 2001ApJ...559..307J, 2002Sci...295...82K, 2002ApJ...575..950O, 2006ApJ...639..259J} the stellar initial mass function \\citep[IMF;][]{1955ApJ...121..161S}.  However, the comparison between the core mass function and stellar IMF has not often been accurately probed at lower masses. It is more difficult to study this part of the mass domain, but recent results have shown that the CMF exhibits a turnover at lower masses in a manner similar to the IMF \\citep{2007MNRAS.374.1413N}.  The Ophiuchus star-forming region is located at a distance of 139~pc \\citep{2008AN....329...10M} and is a site of low-mass star formation \\citep{1983ApJ...274..698W}. The region consists of two main clouds, L1688 and L1689, which both have extended streamers leading out to lengths of around 10 pc \\citep{1989ApJ...338..902L}. Specifically, it is the more massive of the two clouds, L1688, which is studied in this paper, and which is generally known as the Ophiuchus main cloud. Very high star formation rates have been measured here, with 14--40\\% of the molecular gas being converted into stars \\citep{1977AJ.....82..198V}.  The Ophiuchus cloud has been observed in many wavelengths from the visible to the submilimetre \\citep[e.g. MAN98, ][]{1983ApJ...269..182M, 1989ApJ...340..823W, 1992ApJ...401..667A, 1992ApJ...395..516G, 1997ApJS..112..109B, 2000ApJ...545..327J, 2001MNRAS.323.1025J, 2004ApJ...611L..45J, 2008ApJS..175..277D}. Because of this, the properties of the cloud are very well known and it is therefore a good place to probe low-mass star formation.  The Ophiuchus cloud is the nearest example of  `clustered' star formation (MAN98). This is important to study because most stars form in clustered environments \\citep{1993prpl.conf..429Z}. Ophiuchus may also be the nearest example of triggered star formation in action \\citep{1989ApJ...338..902L,2006MNRAS.368.1833N}, making it a prime candidate for study.  \\begin{figure*} \\includegraphics[angle=0,width=0.85\\textwidth]{./RhoOphWholeLabels2.eps} \\caption{Greyscale image and isophotal contour map of the SCUBA 850$\\mu$m continuum scan-map data of the Ophiuchus dark cloud L1688. Signal-to-noise contours at 5$\\sigma$ and 10$\\sigma$ are shown in black; 25$\\sigma$ and 100$\\sigma$ contours shown in white. 1$\\sigma$ noise levels vary from 15 to 40 mJy/beam (see text for details).} \\label{whole-map} \\end{figure*}  \\begin{figure} \\includegraphics[angle=0,width=0.40\\textwidth]{./RhoOphWholeJiggleLabels.eps} \\caption{Greyscale image and isophotal contour map of the SCUBA 850$\\mu$m continuum jiggle map data of the Ophiuchus dark cloud L1688. Contours at 0.25, 0.5 and 1.0 Jy/beam are shown in black; 2.0 and 4.0 Jy/beam contours are shown in white. 1$\\sigma$ noise levels vary from 10 to 180 mJy/beam.} \\label{jiggle-map} \\end{figure}  In this study we have combined all of the high signal-to-noise SCUBA wide-field scan-map data and narrow-field jiggle-map data taken of L1688, and re-reduced it to produce the deepest submillimetre map of this cloud ever made. L1688 is the region of the Ophiuchus cloud defined by \\citet{1989ApJ...338..902L} and outlined in their Figure~1a (it is marked out by a solid, 5K contour). Of the original regions covered at 1.3mm by MAN98 (Oph-A, -B, -C, -D, -E and -F) only one, Oph-D, is not included here, as it is not part of the central region of L1688. A newly discovered region, which we name Oph-J, is discussed. Two smaller regions, Oph-H and Oph-I are discussed by \\citet{2004ApJ...611L..45J} but are not included in this study. We produce a CMF and investigate the low-mass end of the CMF of the cloud. We compare this with previous findings and also with the Orion molecular cloud \\citep{2007MNRAS.374.1413N}. ",
        "conclusions": "In this paper we have re-analysed the SCUBA archive data for L1688, incorporating all available high signal-to-noise scan-map and jiggle-map data. An updated form of the CMF in the L1688 cloud complex has been presented using updated values for the distance to this region as well as new estimates for the temperatures of the cores.  We have shown that the CMF for L1688 is consistent with a three part power-law with slopes the same as seen in the stellar IMF. The higher-mass end of the CMF declines as a power law which is consistent with other studies of L1688 \\citep[MAN98;][]{2000ApJ...545..327J, 2006A&A...447..609S} as well as studies of Orion \\citep{2007MNRAS.374.1413N,2001A&A...372L..41M,2001ApJ...559..307J,2006ApJ...639..259J}.  Hence, the results are mostly in agreement with those found in earlier studies. However, our deeper maps have allowed the discovery of a turnover in the CMF at 0.7$M_{\\sun}$ which shows that the core mass function continues to mimic the stellar initial mass function to low masses. This agreement is indicative that the stellar IMF is determined at the prestellar core stage.  It has been shown that the relationship between the CMF and IMF is not necessarily a simple 1:1 translation in the mass axis. Consistency can also be achieved using a fully multiple star model."
    },
    "0807/0807.3794_arXiv.txt": {
        "abstract": "We have carried out a multi-site campaign to measure oscillations in the F5 star Procyon~A.  We obtained high-precision velocity observations over more than three weeks with eleven telescopes, with almost continuous coverage for the central ten days.  This represents the most extensive campaign so far organized on any solar-type oscillator.  We describe in detail the methods we used for processing and combining the data.  These involved calculating weights for the velocity time series from the measurement uncertainties and adjusting them in order to minimize the noise level of the combined data.  The time series of velocities for Procyon shows the clear signature of oscillations, with a plateau of excess power that is centred at 0.9\\,mHz and is broader than has been seen for other stars.  The mean amplitude of the radial modes is $38.1\\pm1.3$\\,\\cms{} (2.0 times solar), which is consistent with previous detections from the ground and by the WIRE spacecraft, and also with the upper limit set by the MOST spacecraft.  The variation of the amplitude during the observing campaign allows us to estimate the mode lifetime to be $1.5_{-0.8}^{+1.9}$\\,d.  We also find a slow variation in the radial velocity of Procyon, with good agreement between different telescopes. These variations are remarkably similar to those seen in the Sun, and we interpret them as being due to rotational modulation from active regions on the stellar surface.  The variations appear to have a period of about 10 days, which presumably equals the stellar rotation period or, perhaps, half of it.  The amount of power in these slow variations indicates that the fractional area of Procyon covered by active regions is slightly higher than for the Sun. ",
        "introduction": "Measuring solar-like oscillations in main-sequence and subgiant stars requires high-precision observations -- either with spectroscopy or photometry -- combined with coverage that is as continuous as possible. Most of the results have come from high-precision Doppler measurements using ground-based spectrographs, while measurements from spacecraft have also been reported (see \\citealt{B+K2007c} and \\citealt{AChDC2008} for recent summaries).  Procyon has long been a favourite target for oscillation searches.  At least eight separate velocity studies have reported an excess in the power spectrum, beginning with that by \\citet{BGN91}, which was the first report of a solar-like power excess in another star.  For the most recent examples, see \\citet{MLA2004}, \\citet{ECB2004}, \\citet{BMM2004} and \\citet{LKB2007}.  These studies agreed on the location of the excess power (around 0.5--1.5\\,mHz) but they disagreed on the individual oscillation frequencies.  However, a consensus has emerged that the large separation (the frequency separation between consecutive overtone modes of a given angular degree) is about 55\\,\\muHz. Evidence for this value was first given by \\citet{MMM98} and the first clear detection was made by \\citet{MSL99}.  Controversy was generated when photometric observations obtained with the MOST satellite failed to reveal evidence for oscillations \\citep{MKG2004,GKR2007,BAB2008}.  However, \\citet{BKB2005} argued that the MOST non-detection was consistent with the ground-based data.  Meanwhile, \\citet{R+RC2005} suggested that the signature of oscillations is indeed present in the MOST data at a low level (see also \\citealt{Mar08}).  Using space-based photometry with the WIRE satellite, \\citet{BKB2005b} extracted parameters for the stellar granulation and found evidence for an excess due to oscillations.  All published velocity observations of Procyon have been made from a single site, with the exception of two-site observations by \\citet{MLA2004}.  Here we describe a multi-site campaign on Procyon carried out in 2007 January, which was the most extensive velocity campaign so far organized on any solar-type oscillator.  The only other comparable effort to measure oscillations in this type of star was the multi-site photometric campaign of the open cluster M67 \\citep{GBK93}. ",
        "conclusions": "We have presented multi-site velocity observations of Procyon that we obtained with eleven telescopes over more than three weeks.  Combining data that spans a range of precisions and sampling rates presents a significant challenge.  When calculating the power spectrum, it is important to use weights that are based on the measurement uncertainties, otherwise the result is dominated by the noisiest data.  We have described in detail our methods for adjusting the weights in order to minimize the noise level in the final power spectrum.  Our velocity measurements show the clear signature of oscillations.  The power spectrum shows an excess in a plateau that is centred at 0.9\\,mHz and is broader than has been seen for other solar-type stars.  The mean amplitude of the radial modes is $38.1\\pm1.3$\\,\\cms{} ($2.04 \\pm 0.10$ times solar), which is consistent with previous detections from the ground and by the WIRE spacecraft, and also with the upper limit set by the MOST spacecraft.  The variation of the amplitude during the observing campaign allowed us to estimate the mode lifetime to be $1.5_{-0.8}^{+1.9}$\\,d.  We also found a slow variation in the radial velocity of Procyon, with good agreement between different telescopes.  These variations are remarkably similar to those seen in the Sun, and we interpret them as being due to rotational modulation from active regions on the stellar surface.  The variations appear to have a period of about 10 days, which presumably equals the stellar rotation period or, perhaps, half of it.  The amount of power in these slow variations indicates that the fractional area of Procyon covered by active regions is slightly higher than for the Sun.  The excellent coverage of the observations and the high signal-to-noise should allow us to produce a good set of oscillation frequencies for Procyon.  This analysis will be presented in subsequent papers."
    },
    "0807/0807.3331_arXiv.txt": {
        "abstract": "{The nature of Type Ia supernova progenitors is still unclear. The outstanding characteristic of the single-degenerate scenario is that it contains hydrogen in the binary companion of the exploding white dwarf star, which, if mixed into the ejecta of the supernova in large amounts may lead to conflicts with the observations thus ruling out the scenario.}{We investigate the effect of the impact of Type Ia supernova ejecta on a main sequence companion star of the progenitor system. With a series of simulations we investigate how different parameters of this system affect the amount of hydrogen stripped from the companion by the impact.}{The stellar evolution code GARSTEC is used to set up the structure of the companion stars mimicking the effect of a binary evolution phase. The impact itself is simulated with the smoothed particle hydrodynamics code GADGET2.}{We reproduce and confirm the results of earlier grid-based hydrodynamical simulation. Parameter studies of the progenitor system are extended to include the results of recent binary evolution studies. The more compact structure of the companion star found here significantly reduces the stripped hydrogen mass.}{The low hydrogen masses resulting from a more realistic companion structure are consistent with current observational constraints. Therefore, the single-degenerate scenario remains a valid possibility for Type Ia supernova progenitors. These new results are not a numerical effect, but the outcome of different initial conditions.} ",
        "introduction": "\\label{sec:introduction} While the progenitors for Type Ib/c and Type II supernovae are known, Type Ia supernovae (SNe~Ia) still elude an identification of their progenitor system. This is an unpleasant situation given the fact that these objects are one of the most important tools to determine cosmological parameters. By virtue of empirical calibration methods \\citep[e.g.,][]{phillips1993a} they can be used as  standardizable candles for distance measurements. This calls for an understanding of the mechanism of SNe~Ia; and indeed, some progress has been made in recent years in understanding the explosion mechanism in terms of thermonuclear explosions of white dwarf (WD) stars \\citep[e.g.,][]{reinecke2002d,gamezo2003a,roepke2005b, roepke2007b, mazzali2007a, roepke2007c}. In order to judge potential systematic errors in SN~Ia cosmology, a theoretical connection between the explosion characteristics and properties of the progenitor system would be desirable. Yet despite all efforts on both the theoretical and on the observational side, the nature of the progenitor system remains enigmatic.  The stabilization of WDs against gravity does not depend on a finite energy source such as the nuclear burning in normal stars. Due to the Fermi pressure of a degenerate electron gas, single WD stars are in principle eternally stable. Thus, some additional dynamics is required in order to reach an explosive state. The most likely possibility is a WD being part of a binary system and accreting matter from its companion.  Current progenitor models distinguish between the \\emph{single degenerate} and the \\emph{double degenerate} scenario. The former \\citep[proposed by][]{whelan1973a} assumes a ``normal'', non-degenerate star to be the binary companion -- either a main sequence (MS) star or a red giant (RG). In this case, the WD accretes mass from its MS or RG companion via Roche-lobe overflow  or by winds (symbiotic systems) until it approaches the Chandrasekhar mass. The densities reached in the core of the WD are then sufficiently high to trigger nuclear reactions which finally cause a thermonuclear explosion of the star \\citep[but note that an explosion before reaching the Chandrasekhar mass may also be possible, e.g.][]{fink2007a}. The double degenerate scenario \\citep{iben1984a, webbink1984a}, on the other hand, assumes a binary system of two WDs. Due to gravitational wave emission, the system becomes unstable at some point, the WDs merge, and may eventually explode in a SN Ia. A summary of arguments in favor of and against both scenarios can be found in \\cite{livio2000a}. In theoretical modeling, the single-degenerate Chandrasekhar-mass scenario has received most attention recently. By constraining the amount of fuel available in the thermonuclear explosion, it provides a natural explanation for the observed uniformity of SNe~Ia.  Observationally, its hard to distinguish between the progenitor scenarios. One fundamental difference is the complete absence of hydrogen in the double degenerate scenario as it assumes a merger of two carbon/oxygen WDs. In contrast, in the typical single degenerate scenario, hydrogen is the main constituent of the companion. An exception are helium-accretors in which a more evolved companion star has lost his hydrogen envelope. Thus, the WD accretes helium instead of hydrogen and hydrogen is missing in the system. \\citet{kato2003a} reported a possible detection of such an object. In the standard scenario, however, the companion star features a hydrogen envelope; and at least some part of it is expected to be carried away by the SN~Ia ejecta impacting the companion. This, in principle, causes a problem for the single-degenerate scenario, because the astronomical classification of SNe~Ia rests on the absence of hydrogen-features in the spectra of these events. The hydrogen stripped off from the companion will have rather low velocities. It may thus be detectable in nebular spectra, if abundant enough. For the single-degenerate scenario it is therefore of critical importance that the mass of stripped material is sufficiently low to be still consistent with the observations.  There have been a few attempts to search for hydrogen in nebular spectra of SN Ia. \\citet{mattila2005a} studied nebular spectra of SN 2001el. From modeling them, they derived an upper limit of $0.03\\, \\mathrm{M}_{\\odot}$ of solar abundance material at velocities lower than $1000\\, \\mathrm{km}\\, \\mathrm{s}^{-1}$. Recently, \\citet{leonard2007a} studied nebular spectra of SN~2005am and SN~2005cf. Based on the same model as \\citet{mattila2005a}, he estimated $\\lesssim$$0.01\\, \\mathrm{M}_{\\odot}$ of hydrogen material for both objects. Recently also hydrogen has been detected indirectly by \\citet{patat2007a} in circumstellar material of SN 2006X.  An alternative to this approach of observationally constraining the nature of the progenitor system is to directly search for the former companion star of the single-degenerate scenario in the remnants of historical galactic SNe~Ia. Such a search has been carried out in the remnant of Tycho Brahe's supernova of 1572 by \\citet{ruiz-lapuente2004a}, who claimed the identification of the binary companion. The star in question is a slightly evolved solar-type star, that moves with a radial velocity of $-108 \\, \\mathrm{km\\ s^{-1}}$ relative to the sun. It also has an atypical large tangential velocity of about $90 \\, \\mathrm{km\\ s^{-1}}$. A significantly larger velocity of the star compared to neighbours is expected as a result of the disappearing binary orbit. Other stars observed in the same area with similar distances move only with average radial velocities of about $-20$ to $-40 \\, \\mathrm{km\\ s^{-1}}$, with a velocity dispersion of about $20 \\, \\mathrm{km\\ s^{-1}}$.  On the theory side, \\citet{marietta2000a} presented two-dimensional hydrodynamical simulations of the impact of SNe~Ia on their companions. They found that $0.15 \\, \\mathrm{M}_{\\odot}$ were stripped from a Roche-lobe filling MS companion. This would rule out a MS-WD system for the SN~Ia analysed by \\citet{leonard2007a}, if it was representative. Recently, \\citet{meng2007a} pointed out that considering the effect of the mass transfer phase on the companion star may change the result significantly. They studied the impact of SNe~Ia on different companion stars analytically. In contrast to \\citet{marietta2000a}, who assumed the structure of single MS stars for the companion, \\cite{meng2007a} evolved it through the binary evolution phase before the explosion of the WD. They found at least $0.035\\, \\mathrm{M}_{\\odot}$ of stripped hydrogen for the companion. However, this result is only a lower limit, since they did not include mass loss by vaporization from the hot surface of the star. Thus, taken at face value, the currently available theoretical studies constitute a strong case against MS+WD progenitor systems for SNe~Ia.  The aim of the study presented here is to check and update the \\citet{marietta2000a} calculations with the results of recent detailed binary evolution models. \\citet{ivanova2004a} identified possible SN~Ia progenitors from a parameter study of  MS+WD binary evolution. There results are in agreement with other studies in this field \\citep[e.g.][]{langer2000a, han2004a}. Based on these results, we present an exploration of the effect of the impact of a SN Ia on different MS companions by 3D hydrodynamical simulations. Section \\ref{sec:code} summarizes the codes used. Section \\ref{sec:tests} demonstrates that our approach reproduces the results of \\cite{marietta2000a} and presents a resolution study. Section \\ref{sec:results} discusses an exploration of different progenitor systems and Section \\ref{sec:observ} derives observational implications of our results. A summary and an outlook conclude this work in Section \\ref{sec:conclusion}. ",
        "conclusions": "\\label{sec:conclusion} We studied the impact of the ejecta of SNe~Ia on main sequence companion stars in the context of the single-degenerate Chandrasekhar-mass scenario with hydrodynamical simulations. To this end, the (cosmological) GADGET2 smoothed particle hydrodynamics code was adapted to the stellar problem and employed in numerical simulations. It was shown that this SPH-based approach is capable of reproducing previous results obtained with a grid-based 2D scheme by \\cite{marietta2000a}. A resolution study indicated that with a few million particles the simulations yield numerically converged results.  We showed that the mass stripped from the companion star by the impact of the ejecta depends on the their kinetic energy and the binary separation. While the latter affects the mass of the stripped material significantly, the former has only a minor effect on it. This is due to the fact that the supernova explosion energy can only vary in a relatively narrow range given the restricted amount of fuel available for the nuclear energy generation.  The SPH approach was used to analyze the impact in a number of more realistic progenitor models than those employed in previous studies. For these, the companion stars were constructed with the stellar evolution code GARSTEC mimicking binary mass transfers with the parameters given by \\citet{ivanova2004a}. In the hydrodynamical impact simulations, we found about one order of magnitude less hydrogen material stripped off the companion by the impact of the supernova than predicted by previous studies. The main reason for this difference is a modified, more compact stellar structure of the companion star in combination with a resulting variation in the separation distance of the progenitor system. In particular the more compact state of the companion impedes the mass loss in the impact.  The reduced amount of hydrogen mixed into the ejecta of the supernova as predicted by our simulations leads to an agreement with observational studies of SN~Ia nebular spectra \\citep{mattila2005a,leonard2007a}. This removes the former disagreement between the available observations and simulations of the WD+MS progenitor system. Thus, to current knowledge, such a progenitor scenario is admissible in the context discussed here.  However, since the hydrogen masses predicted by our simulations are not far below the current observational upper limits, it may be possible in the near future to confirm or reject the studied progenitor scenario by either detecting hydrogen in SNe~Ia or lowering the limits by another order of magnitude. A stringent way of analysis would be to calculate synthetic spectra directly from the presented simulations and to compare the results with observations. This will be tackled in a forthcoming study."
    },
    "0807/0807.1334_arXiv.txt": {
        "abstract": "Investigating the link between supermassive black hole and galaxy evolution requires careful measurements of the properties of the host galaxies. We perform simulations to test the reliability of a two-dimensional image-fitting technique to decompose the host galaxy and the active galactic nucleus (AGN), especially on images obtained using cameras onboard the {\\it Hubble Space Telescope (HST)}, such as the Wide-Field Planetary Camera 2, the Advanced Camera for Surveys, and the Near-Infrared Camera and Multi-Object Spectrometer.  We quantify the relative importance of spatial, temporal, and color variations of the point-spread function (PSF).  To estimate uncertainties in AGN-to-host decompositions, we perform extensive simulations that span a wide range in AGN-to-host galaxy luminosity contrast, signal-to-noise ratio, and host galaxy properties (size, luminosity, central concentration).  We find that realistic PSF mismatches that typically afflict actual observations systematically lead to an overestimate of the flux of the host galaxy.  Part of the problem is caused by the fact that the {\\it HST}\\ PSFs are undersampled.  We demonstrate that this problem can be mitigated by broadening both the science and the PSF images to critical sampling without loss of information.  Other practical suggestions are given for optimal analysis of \\hst\\ images of AGN host galaxies. ",
        "introduction": "The mass of supermassive black holes (BHs) is strongly correlated with the luminosity (\\lbul; Kormendy \\& Richstone 1995; Magorrian et al. 1998) and the stellar velocity dispersion (\\vel; Gebhardt et al. 2000; Ferrarese \\& Merritt 2000) of the bulge of the host galaxy.  These scaling relations are often interpreted to be evidence that central BHs and their host galaxies are closely connected in their evolution (see reviews in Ho 2004).  The empirical correlations between BH mass and host galaxy properties can even be used as tools to track the progress of mass assembly during galaxy evolution (Peng et al. 2006a, 2006b; Woo et al. 2006; Ho 2007b). The central BH mass correlates most strongly with the bulge component of a galaxy rather than with its total mass or luminosity.  This has motivated detailed bulge-to-disk decompositions of the host galaxies to better quantify the intrinsic scatter of the \\mbh-\\mbul\\ and \\mbh-\\lbul\\ relations in the local Universe (Marconi \\& Hunt 2003; H\\\"{a}ring \\& Rix 2004).  To probe when the BH-host galaxy relations were established and how they evolved, it is of paramount importance to extend similar studies out to higher redshifts.  However, direct measurement of BH mass based on spatially resolved stellar or gas kinematics is unfeasible for all but the nearest galaxies with low levels of nuclear activity.  Accessing BHs and their host galaxies at cosmological distances requires a different approach---one that relies on active galactic nuclei (AGNs).  The masses of BHs in type~1 (unobscured, broad-line) AGNs can be readily estimated with reasonable accuracy using the virial technique with single-epoch optical or ultraviolet spectra (Ho 1999; Wandel et al. 1999; Kaspi et al.  2000; Greene \\& Ho 2005b; Peterson 2007). Quantitative measurements of the host galaxies with active nuclei, on the other hand, are less straightforward to obtain because the presence of the AGN introduces significant practical difficulties, as well as potential biases. This is especially problematic for the bulge component of the host, which is maximally affected by the bright AGN core.  A variety of techniques have been employed, both kinematical (e.g., Nelson et al. 2004; Onken et al. 2004; Barth et al. 2005; Greene \\& Ho 2005a, 2006; Ho 2007; Salviander et al. 2007; Ho et al. 2008; Shen et al. 2008) and photometric (e.g., McLure \\& Dunlop 2002; Peng et al. 2006b; Greene et al. 2008). \\begin{figure*} \\psfig{file=table1.ps,width=18.5cm,angle=0} \\vskip -0.4cm \\end{figure*}  The main challenge for photometric studies of active galaxies lies in separating the central AGN light from the host galaxy. While structural decomposition can be done by fitting analytic functions to the one-dimensional (1-D) light distribution, two-dimensional (2-D) analysis makes maximal use of all the spatial information available in galaxy images (Griffiths et al. 1994; Byun \\& Freeman 1995; de~Jong 1996; Wadadekar et al. 1999; Peng et al. 2002; Simard et al. 2002; de~Souza et al. 2004) and thus provides the most general and most robust method to decouple image subcomponents.  This flexibility proves to be especially important in the case of active galaxies, where the contrast between the central dominant point source and the underlying host can be very high.  In this regime, achieving a reliable decomposition requires knowing the point-spread function (PSF) to high accuracy.  This paper discusses the complications and systematic effects involved in photometric decomposition of AGN host galaxies, especially as it applies to images taken with the {\\it Hubble Space Telescope (HST)}.  Although numerous {\\it HST} studies of AGN hosts have been published, very few have explicitly investigated the systematic uncertainties or practical limitations of host galaxy decomposition.  We make use of an updated version of the 2-D image-fitting code GALFIT (Peng et al. 2002) to generate an extensive set of simulated images of active galaxies that realistically mimic actual {\\it HST} observations.  We then apply the code to fit the artificial images and quantify how well we can recover the input parameters for the AGN and for the host galaxy.  Our strategy resembles those employed in the quasar host galaxies studies by Jahnke et al. (2004) and S\\'{a}nchez et al. (2004), except that we have optimized our simulations to be applicable to nearby bright AGNs, a regime of most interest to us (Kim et al. 2008). Nearby bright AGN hosts observed using coarse pixels and low exposure time can be equally challenging to analyze as high-z AGNs observed with high resolution and signal-to-noise.  In other words, fundamentally, the difficulty of the analysis depends only on the following 3 relative conditions: angular resolution of the detector vs. the object being studied (i.e. object scale size in pixels), AGN-to-host contrast, and overall object signal-to-noise. Since we cover a wide range of parameter space the simulations described here should be applicable to most {\\it HST}\\ imaging studies of AGN hosts, including distant quasars. We hope that the results of this study can serve as a guide to other investigators working with similar data.  We describe the details of PSF variations in \\S{2}. In \\S{3}, we present the simulation procedure and our results from fitting artificial images using PSF models with varying degrees of realism.  A discussion and summary of the main results are given in \\S{4}. ",
        "conclusions": "In this study, we performed 2-D image-fitting simulations of AGN hosts to illustrate how systematics in the PSF mismatch may affect the deblending of the AGN and the host galaxy components.  Based on these simulations, we suggest practical strategies for how to observe and analyze host galaxies with bright active nuclei.  As seen in Figures 6--7, careful determination of the PSF is needed to extract very high-contrast subcomponents accurately.  This point was also underscored in the study of Hutchings et al. (2002).  We find that the three factors that most affect PSF structures are spatial distortion, temporal changes, and pixel undersampling.  The SED of the PSF matters to a much lesser degree, but it is worthwhile to match it when possible.  The spatial variation can be reduced by observing a stellar PSF at the same position as the science images. However, the temporal variance is trickier to avoid without observing PSFs concurrently with the science data, which is observationally expensive and rarely done in practice.  Even then, some temporal variability may happen even within a single orbit. Nevertheless, as there is evidence that PSF mismatch grows over time, and may not be completely periodic in nature, it is important to obtain a stellar PSF as close in time as possible to the science images. From our tests (\\S{2.1}), we recommend that stellar PSF images be taken within a month of the science images.  In the absence of a well-matched stellar PSF, synthetic PSFs generated with TinyTim are adequate substitutes (compare Fig. 10 to Fig. 7 and Fig. 9).  In all scenarios, whether using stellar or TinyTim PSFs, both the science image and the PSF image should be oversampled to reduce errors caused by pixel undersampling and subpixel shifting.  This can be accomplished by observing the science image and the PSF star using a four-point ``dither'' pattern to recover finer pixel sampling.  Or, if this option is not available, then simply broadening the images through convolution during the analysis stage is an acceptable alternative solution, and certainly better than no treatment at all (\\S{3.4} and Fig. 10).   \\begin{figure*} \\figurenum{15} \\psfig{file=f15.eps,width=18.5cm,angle=0} \\figcaption[fig15.ps] {Distribution of residuals for the magnitudes of the host for ({\\it left}) \\hnr $\\geq 0.2$ and ({\\it right}) \\hnr $< 0.2$.  After producing artificial images by varying $n$ between 2 and 6, we fit them with  $n$ fixed to 4.  This simulation gives an estimate of the measurement error for the bulge luminosity. \\label{fig15}} \\end{figure*}  When the host galaxies are faint and difficult to fit, it may be useful to hold the \\ser\\ index fixed to a constant value of either $n=1$ or $n=4$.  The decision about what value to use might be based on visual morphology, or by comparing $\\chi^2$ values of the bimodal priors.  A similar conclusion was reached by McLure et al. (1999), Jahnke et al.  (2004), and S\\'{a}nchez et al.  (2004) in their analysis of quasar host galaxies.  By varying the \\ser\\ index between $n$ = 2 and $n$ = 6, we find that the scatter is $\\sigma \\approx 0.3 $ mag for \\hnr $\\geq 0.2$ and \\nsn\\ $\\approx 1000$ (Fig. 15).  At higher contrast, \\hnr $\\leq 0.2$, the scatter increases to $\\sigma \\approx 0.4 $ mag.  Lastly, we briefly conducted a three-component bulge/disk/AGN decomposition to characterize uncertainties in estimating the bulge component.  We find that when the $S/N$ is high, the contrast is sufficiently low, and the bulges are sufficiently well-resolved, then a B/D/A decomposition can yield reliable bulge luminosity measurements.  These simulations are fully compatible with images of $z\\lesssim 0.3$ quasar host galaxies (e.g., McLure \\& Dunlop 1999). Our tests suggest that, when three-component decomposition is required, the uncertainty in the bulge luminosity increases by an additional $\\sim$0.1 mag compared to fits without a disk component.  Finally, we note in passing that our simulations allow the sky value to vary during the fitting.  For the current application, this choice makes little difference because the images have a large sky area and the profiles are idealized.  However, in real science images, the situation will be different.  From prior experience with actual \\hst\\ data, keeping the sky value fixed prevents the \\ser\\ index from going up to extremely high values in situations where the light profile of the galaxy is not well represented by the \\ser\\ function.  Thus, we recommend keeping the sky parameter fixed to a well-determined value whenever possible."
    },
    "0807/0807.1116_arXiv.txt": {
        "abstract": "The outer regions of disk galaxies show a drop-off in optical and H$\\alpha$ emission, suggesting a star formation threshold radius, assumed to owe to a critical surface density below which star formation does not take place. Signs of filamentary star formation beyond this threshold radius have been observed in individual galaxies in the H$\\alpha$ and recent GALEX surveys have discovered that 30\\% of disk galaxies show UV emission out to  2-3 times the optical radius of the galaxy.    We run smooth particle hydrodynamics simulations of disk galaxies with constant density extended gas disks to test whether over-densities owing to spiral structure in the outer disk can reproduce the observed star formation.  We indeed find that spiral density waves from the inner disk propagate into the outer gas disk and raise local gas regions above the star formation density threshold, yielding features similar to those observed. Because the amount of star formation is low, we expect to see little optical emission in outer disks, as observed.  Our results indicate that XUV disks can be simulated simply by adding an extended gas disk with a surface density near the threshold density to an isolated galaxy and evolving it with fiducial star formation parameters. ",
        "introduction": "The most basic description of galactic-scale star formation is the Schmidt relation, which assumes that the star-formation rate (SFR) is a power law function of the gas density \\citep{Schmidt-1959}. \\citet{Kennicutt-1998} showed that this relation accurately describes star formation in galaxies over a range of $10^{5}$ in gas surface density and $10^{6}$ in SFR per unit area.  However, beyond a critical galactocentric radius, SFRs determined from H$\\alpha$ emission are truncated relative to the expectation from the Kennicutt-Schmidt relation despite the availability of large reserves of low density cold gas \\citep{Kennicutt-1989, Martin-Kennicutt-2001, Wong-Blitz-2002}. The lack of star formation at large radii has primarily been attributed to dynamical disk stability, i.e., \\begin{equation} Q(r) = \\frac{\\Sigma_{c}(r)}{\\Sigma_{gas}(r)} = \\frac{\\kappa(r) c_{s}}{\\pi G \\Sigma_{gas}(r)} \\gsim 1 \\, , \\end{equation} \\noindent where $\\kappa$ is the epicyclic frequency at a given radius, $\\Sigma_{c}(r)$ is the critical surface density and $c_{s}$ is sound speed of the gas, proportional the velocity dispersion of the gas ( $\\sigma=c_{s}\\gamma^{-1/2}$, where $\\gamma$ is the ratio of specific heats) \\citep{Toomre-1964, Spitzer-1968, Quirk-1972}. The effective velocity dispersion of gas in spirals is roughly constant across the disk \\citep[e.g.,][]{vanderKruit-Shostak-1984, Anderson-et-al-2006} leaving $\\kappa$ as the only radially dependent quantity in the critical surface density. If velocity curves of spiral galaxies are approximately flat, $\\kappa$, and therefore $\\Sigma_{c}$, fall with $r$, implying that at some radius the disk becomes gravitationally stable, which may inhibit star formation.  In support of this notion, several studies have shown correlations between the radius at which star formation ceases and the radius at which $Q\\approx 2$ \\citep[][see \\citet{Schaye-2004} for a comparison of these results]{Kennicutt-1989, Martin-Kennicutt-2001, Wong-Blitz-2002}. More sophisticated models including the formation of cold molecular gas  have also been successful at reproducing observed SFR truncations \\citep{ Elmegreen-Parravano-1994, Schaye-2004} .    Observations indicate the presence of star formation beyond the threshold. Faint, isolated H\\,II regions distributed in spiral structures out to two optical radii have been found in face on galaxies \\citep{Ferguson-et-al-1998} and in cold disks in edge-on galaxies \\citep{Christlein-Zaritsky-2008}. Similar star formation has been seen in the Milky Way \\citep{Brand-et-al-2001} and dwarf galaxies with gas densities below the threshold \\citep{vanZee-et-al-1997}. A systematic survey by the Galaxy Evolution Explorer (GALEX: \\citet{Martin-et-al-2005}) recently discovered that 30\\% of disk galaxies show extended UV emission beyond the optical radius, and presumed star formation threshold, of the galaxy \\citep{Thilker-et-al-2007, Gildepaz-et-al-2005, Thilker-et-al-2005, Zaritsky-Christlein-2007}. In 2/3 of these cases, this emission is structured, lying in spiral or filamentary patterns (Type I XUV disks). In the other 1/3, the emission takes the form of a large zone in the outer regions of the galaxy with an enhanced UV/optical ratio relative to the inner disk (Type II XUV disk) \\citep[][hereafter T07]{Thilker-et-al-2007}.   To reconcile outer disk star formation with disk stability, it has been proposed that local densities above $\\Sigma_{c}(r)$ in the outer disk allow star formation beyond the truncation radii \\citep{Kennicutt-1989, Martin-Kennicutt-2001, Schaye-2004, Elmegreen-Hunter-2006, Gildepaz-et-al-2007}. The low SFR implies a low number of O stars producing H$\\alpha$ emission at a given time, or none at all, explaining why this star formation was undetected in the H$\\alpha$ in some galaxies \\citep{Gildepaz-et-al-2007}. Most of these disks have large amounts of H\\,I in their outer regions and UV and H$\\alpha$ complexes are often coincident with local H\\,I over-densities \\citep[][T07]{Ferguson-et-al-1998} supporting this picture. Analysis of the star formation profiles of these XUV disk galaxies indicates that the azimuthally averaged Kennicutt-Schmidt relation could hold beyond the dynamical threshold radius \\citep[][Hereafter B07]{Boissier-et-al-2007}. However, studies of star forming clumps in the XUV disk of M\\,83 indicate that traditional star formation laws, including the threshold radius, hold locally \\citep{Dong-et-al-2008}.   Theoretically, \\citet{Elmegreen-Hunter-2006} have shown that XUV disks may not be inconsistent with a star formation density threshold.  In their analytical model, gas clumping triggered by spiral density waves, radial variations in the interstellar medium (ISM) turbulence, and gas phase transitions lead to local regions of star formation in the extended gas disk. In this paper, we investigate XUV disks using smooth particle hydrodynamic (SPH) simulations of an isolated disk galaxy with an extended gas disk and a simple, commonly employed, star formation prescription based on the Kennicutt-Schmidt relation that includes a threshold density for star formation. We demonstrate that, even with these fiducial star formation laws, spiral structure propagates from the inner disk to the extended H\\,I disk and produces local regions of enhanced of gas density which trigger star formation similar to that seen in observed Type I XUV disks. ",
        "conclusions": "\\label{sec:conc}  We show that features similar to those seen in Type I XUV disks can be produced with fiducial star formation prescriptions, including a star formation threshold density, if there is gas present at large radii near the threshold density.  This is a clear illustration of the fact that a star formation threshold density does not imply the existence of a star formation threshold radius beyond which star formation ceases. As proposed in earlier studies and shown by \\citet{Elmegreen-Hunter-2006}, local over-densities, caused by spiral density waves in our simulation, allow star formation in gas that, when circularly averaged, is below or near the classical star formation threshold density.   It is possible that star formation will differ between inner and outer disks.  This could arise owing to different properties of the ISM in the inner and outer regions of disks, from, for example, a different ionized fraction owing to various levels of self-shielding against an ionizing UV background. Currently, our simulations do not have the resolution to probe the detailed physics of star formation. Here, we simply show that any simulation whose star formation prescription follows a Kennicutt-Schmidt relation with a star formation threshold density and has a high enough gas density for local over-densities to exceed the star formation threshold density in the outer disk will exhibit UV emission in the outer disk.  Understanding the mechanism for generating XUV disks has promising implications for probing galaxy structure in the outer parts of galaxies.  If the distribution of H\\,I can be determined based on the observed UV emission, this can be used to infer the extended H\\,I properties of disks. Outer disk H\\,I and its star formation history could also constrain recent gas accretion in galaxies and the formation theory of disks. For these reasons, the generation of XUV disks warrants further exploration. Simulations having different extended gas profiles (e.g. exponential, $1/R$) at various gas densities are needed to test whether we can create the range of properties seen in XUV disks, including Type II XUV disks. Upcoming extragalactic H\\,I surveys such as THINGs \\citep{Walter-et-al-2005} and ALFALFA \\citep{Giovanelli-et-al-2005} will provide better statistics on the types of H\\,I profiles seen on disk galaxies and how they correlate with XUV disk behavior. The role of interactions, particularly flybys, in exciting spiral structure and affecting the morphology of XUV disks also needs to be explored.  Finally, the application of a radiative transfer code to our SPH results to compare colors of these clumps to samples such as T07 and \\citet{Zaritsky-Christlein-2007} would provide interesting constraints on star formation such as the initial mass function in outer disks."
    },
    "0807/0807.1927_arXiv.txt": {
        "abstract": " ",
        "introduction": "\\paragraph{}     Inflation~\\cite{inflation} has emerged as the standard paradigm describing physics of the very early universe. Besides addressing several fine-tuning issues in big bang cosmology such as the flatness and horizon problems, it provides a framework to explain the origin of structure and the cosmic microwave background (CMB) anisotropy~\\cite{books}. While there is a plethora of effective field theory based models of inflation~\\cite{Lyth:1998xn}, many outstanding questions in inflationary cosmology {\\it require} a fundamental microscopic description. Conversely, recent observations of the CMB and large scale structure~\\cite{observation} lead us to increasingly precise measurements of the inflationary parameters. These measurements provide us with an exciting window to probe physics at ultra-high energies~\\cite{reconstruction}, far higher than what current and upcoming terrestrial accelerators can reach. Thus inflationary cosmology has become the perfect arena for fundamental theory to meet experiment.  String theory is currently our leading candidate for a quantum theory of gravity. Thus it is worthwhile to explore explicit realizations of inflation within this framework. In this paper, we will focus on one of the most well developed inflationary scenarios in string theory, i.e. $D$ brane inflation~\\cite{DvaliTye} (see also Refs.~\\cite{Dvali:2001fw,Burgess:2001fx}; for reviews, see Ref.~\\cite{reviews} and references therein), where the inflaton field is identified with the position of a space-filling mobile $D$ brane, usually a $D3$ brane, in a warped six dimensional manifold~\\cite{KKLMMT}. In the original scenario of Refs.~\\cite{DvaliTye,Dvali:2001fw,Burgess:2001fx,reviews,KKLMMT}, an additional $\\overline{D3}$ was introduced to drive inflation. The $\\overline{D3}$ brane is localized by the RR fluxes at the tip of a warped throat, thus inflation  proceeds as the mobile $D3$ is attracted by a weak $D3$-$\\overline{D3}$ Coulombic force to move slowly along the warped direction. However, it was also noted in Ref.~\\cite{KKLMMT} that because the volume modulus of the compactification couples non-trivially to the canonical inflaton, its stabilized value gives additional Hubble scale correction to the inflaton mass, causing the well known $\\eta$ problem~\\cite{Copeland:1994vg}.   An important step towards addressing the $\\eta$ problem explicitly in this concrete setting was recently made in Refs.~\\cite{Baumann0, Baumann1} (see also Refs.~\\cite{Panda:2007ie, McGill1, Krause:2007jk, Pajer:2008uy}). The key ingredient in the construction was the one loop threshold correction to the non-perturbative superpotential obtained in Ref.~\\cite{Gaugino} (see also Refs.~\\cite{Ganor,BHK}). In Ref.~\\cite{KKLT} and other stabilized compactifications, non-perturbative effects are often introduced to stabilize moduli. In the context of Ref.~\\cite{KKLT}, such effects come from instantons on a stack of $D7$ branes (or Euclidean $D3$ branes). Interestingly, the non-perturbative moduli stabilizing force also turns out to give the dominant contribution to the inflaton potential\\footnote{Although the Coulombic force is subdominant in comparison to the moduli stabilizing force, a $\\overline{D3}$ brane was still introduced to end inflation.}. This contribution arises because the mobile $D3$ brane backreacts on the moduli stabilizing $D7$ branes. The correction depends on the holomorphic four cycles within the conifold on which the $D7$ branes wrap. The embedding of the $D7$ branes breaks the isometry of the deformed conifold, and thus the inflationary trajectory depends sensitively on the choice of the embedding function {defining the loci of the $D7$ branes}. As a result, explicit slow-roll models have been constructed by a ``delicate'' tuning of the microscopic compactification parameters.   While the broken angular isometry directions are stabilized by the coordinate dependent non-perturbative superpotential, for a given $D7$ brane embedding, there are typically residual isometries preserved by the resultant scalar potential. The potential for the fields associated with these isometries remain flat during the inflationary epoch and so they can take arbitrary values without affecting the inflationary trajectory. Being almost massless, their quantum fluctuations give rise to a nearly scale invariant isocurvature perturbation spectrum. As argued by Lyth and collaborators in Refs.~\\cite{Lyth,Lyth1}, these isocurvature perturbations can be converted to the curvature perturbations at the end of inflation. In the context of $D$ brane inflation, inflation ends when the open string tachyon condenses between the mobile $D3$ and $\\overline{D3}$. The critical value of the canonical inflaton at which inflation ends $\\phi_\\tip$ depends on the residual symmetries as they enter into the tachyon potential. Since $\\phi_\\tip$ picks up spatial dependence through the quantum fluctuations of the light residual symmetries, inflation can end on a spatial slice of non-uniform energy density. {As we will see, this is the case for instance when the inflaton potential is dominated by the moduli stabilizing force towards the end of inflation. Thus,} one could in principle expect potentially significant contribution to the power spectrum and non-Gaussianities at the end of inflation.   In this paper, we study these multi-field effects at the end of brane inflation, and outline the necessary conditions for them to be significant. We then perform a case study for the setup considered in Ref.~\\cite{Baumann1}, by explicitly calculating the canonical inflaton potential near the tip of the deformed conifold, and demonstrate that inflation can persist in this region provided that the $D3$-$\\overline{D3}$ Coulombic attraction becomes subdominant. We also discuss various mechanisms to uplift the vacuum energy which results in a subdominant Coulombic potential all the way to the tip of a warped throat. We also show that the angular stable inflationary trajectory for the specific $D7$ brane embedding \\cite{Kuper} used in Ref.~\\cite{Baumann1} can be extended to the entire deformed conifold. However, along the specific trajectory considered in Ref.~\\cite{Baumann1}, we will see explicitly that the corresponding residual angular isometries have vanishing proper separations at the tip. Thus, for this specific $D7$ embedding, {no significant contribution to the curvature perturbation is generated at the end of inflation}. This implies that while multi-field effects can in principle be significant in brane inflation, they can only happen with other $D7$ embeddings, or with more than one stacks of $D7$ branes present.   This paper is organized as follows. In Section~\\ref{sec_D3brane}, we review the basic setup of flux compactification and brane inflation, in order to set up our notation. Readers who are familiar with the above topics can skip this section. In Section~\\ref{Lytheffect}, we recast the mechanism proposed in Ref.~\\cite{Lyth} in the context of brane inflation in a warped throat, and outline the necessary conditions for it to take place. In Section~\\ref{sec_scenario}, we discuss various possible uplifting mechanisms in a warp throat and propose a natural scenario for an uplifted potential to realize the effect of Ref.~\\cite{Lyth}. In Section~\\ref{sec_endofinf}, we explicitly calculate the canonical inflaton potential near the tip of the deformed conifold and the resulting slow-roll parameters. The degeneracy of the residual isometries  will also be shown. We end with some discussions in Section~\\ref{Discussion}. We relegate most of the calculational details in a number of Appendices. ",
        "conclusions": "\\label{Discussion} \\setcounter{equation}{0}     In this paper, we studied the systematics of multi-field effects at the end of warped $D$ brane inflation. We discussed the necessary criteria for the isocurvature perturbations generated by the angular motion of a mobile $D3$ brane to be converted into the curvature perturbations usually associated with its radial motion in this scenario. We found that the significance of the end of inflation effects considered in Ref.~\\cite{Lyth} depends on the specific mechanism for uplifting the vacuum energy. If the uplifting is due to some distant $\\overline{D3}$ branes or a $D$-term potential, the Coulombic potential can easily become subdominant even towards the end of inflation, and the effects described in Ref.~\\cite{Lyth} can in principle be significant. However, in the most explicit $D$ brane inflation constructed to date \\cite{Baumann0,Baumann1}, the $D7$ brane embedding chosen \\cite{Kuper} does not yield such effects, { regardless of the uplifting mechanism.} This latter result is specific to the embedding of the moduli stabilizing branes as well as the infrared geometry of the throat. % Along the stable trajectory for the embedding considered in Ref.~\\cite{Kuper}, the proper distance for the residual isometry direction vanishes in the entire throat, the moduli space vanishes at the tip. % It would be interesting to examine other $D7$ brane embeddings and/or other warped throats which leave a moduli space of vacua at the tip. Examples of such embeddings for the deformed conifold appeared in Ref.~\\cite{D3vacua}, where the residual isometry directions reside on the finite size $S^3$. However, finding an angular stable trajectory in these examples may remain challenging. Nevertheless, our results underscore the importance of multi-field effects in string inflation, as noted also in the context of DBI inflation recently in Ref.~\\cite{Langlois} (see also earlier discussions in Refs.~\\cite{Huang:2007hh,Easson:2007dh}).   As discussed in Section~\\ref{Lytheffect}, {the strength of the Lyth effect} depends on the ratio $\\varepsilon_k/\\varepsilon_e$. Since the flat region of the inflaton potential considered in Refs.~\\cite{Baumann0,Baumann1} is an inflection point, $\\varepsilon_k$ depends sensitively on where around the inflection point corresponds to the CMB scale. Given a $D$ brane inflation model which can yield {the Lyth effect} considered here, a precise determination of the amplitude of such effects would require the use of the full KS metric~\\cite{KS}. {This is yet another context in which details of the warped geometries in the infrared can have significant effects on the CMB observations~\\cite{inflation_tip}.} {Furthermore, regardless of the {Lyth effect} studied here,} a detailed comparison of the WMAP data with microscopic parameters of $D$ brane inflation requires identifying the relevant part of the inflaton potential which generates the observed CMB {anisotropy}, and the full KS metric is essential. Work along these lines is underway.   Finally, one may hope to also realize the curvaton mechanism~\\cite{curvaton} using these light fields. In the setup {we discussed}, however, inflation ends as $D3$ and $\\overline{D3}$ annihilate and thus the would-be curvaton fields themselves disappear. For the same reason, any multi-field effect~\\cite{multi} after inflation will not be present as long as they are associated with $D3$ or $\\overline{D3}$ branes. Nevertheless it would be interesting to implement the curvaton scenario in a different setup satisfying a number of constraints~\\cite{Gong:2006hf}.     \\subsection*{Acknowledgement} \\paragraph{}   We are indebted to Bret Underwood for numerous valuable insights and discussions, and for collaboration at the initial stage of this project. We are also grateful to Misao Sasaki for discussions and comments on the manuscript. We thank Daniel Baumann, Chong-Sun Chu, Min-Xin Huang, David Lyth, Fernando Marchesano, Liam McAllister, Peter Ouyang, Sudhakar Panda, and Fernando Quevedo for helpful discussions. HYC would like to thank KITP at USCB and the Sixth Simons Workshop at SUNY Stony Brook for their hospitalities where part of the work was being carried out. JG is grateful to the Santa Fe 08 Cosmology Summer Workshop for hospitality where this work was being finished. The work of HYC and GS is supported in part by NSF Career Award No. PHY-0348093, DOE grant DE-FG-02-95ER40896, a Research Innovation Award and a Cottrell Scholar Award from Research Corporation, and a Vilas Associate Award from the University of Wisconsin. JG is partly supported by the Korea Research Foundation Grant KRF-2007-357-C00014 funded by the Korean Government.     \\appendix"
    },
    "0807/0807.3439_arXiv.txt": {
        "abstract": "{We discuss the effect of atmospheric dispersion on the performance of a mid-infrared adaptive optics assisted instrument on an extremely large telescope (ELT). Dispersion and atmospheric chromaticity is generally considered to be negligible in this wavelength regime. It is shown here, however, that with the much-reduced diffraction limit size on an ELT and the need for diffraction-limited performance, refractivity phenomena should be carefully considered in the design and operation of such an instrument. We include an overview of the theory of refractivity, and the influence of infrared resonances caused by the presence of water vapour and other constituents in the atmosphere. `Traditional' atmospheric dispersion is likely to cause a loss of Strehl only at the shortest wavelengths (L-band). A more likely source of error is the difference in wavelengths at which the wavefront is sensed and corrected, leading to pointing offsets between wavefront sensor and science instrument that evolve with time over a long exposure. Infrared radiation is also subject to additional turbulence caused by the presence of water vapour in the atmosphere not seen by visible wavefront sensors, whose effect is poorly understood. We make use of information obtained at radio wavelengths to make a first-order estimate of its effect on the performance of a mid-IR ground-based instrument. The calculations in this paper are performed using parameters from two different sites, one `standard good site' and one `high and dry site' to illustrate the importance of the choice of site for an ELT.} ",
        "introduction": " ",
        "conclusions": "Atmospheric dispersion is a well known source of error in astronomical observations at visible and near-infrared wavelengths. With the advent of ELTs and their integrated AO facilities, atmospheric refractivity becomes a relevant quantity even in the mid-infrared. This is particularly the case if the traditional refractivity formulations are recalculated to include the effects of IR water and CO$_2$ resonances, which causes the refractivity to deviate substantially from the NIR values extrapolated to longer wavelengths.  The calculations in this paper, though in many places approximate in nature, show the likely magnitude of refractive effects on the AO correction for a mid-IR ELT instrument. While `simple' dispersion will only cause a significant drop in Strehl at the short-wavelength end of the region, differential refraction effects between visible/NIR WFS and mid-IR science instrument may cause a significant image degradation over long exposures, even at moderate zenith distances. This problem in particular will be investigated more closely, taking into account the effect of multiple laser guide stars which can only be sensed in the visisble.  Water vapour turbulence has recently been of interest in the interferometry and submillimetre communities, however, a thorough understanding of the phenomenon and experimental data are lacking (and non-existent in the framework of an ELT-sized filled aperture). Our first order estimates show that the Strehl loss can be significant, if not large, depending on the size of the outer scale. However, given the poor understanding of this type of turbulence both in the spatial and temporal domain, and of its correlation with dry air turbulence, these results show be regarded with caution. Further experimental results, specifically designed for ELT-type observations in the mid-IR, are required to help assess the problem.  The comparison of results between the a Paranal-like site and a `high and dry' location like Cerro Macon is also very important, and reinforces the notion that a high and dry location is more suitable for mid-IR astronomy. However, we expect some of these advantages to apply to NIR also, and would invite NIR colleagues to examine site-related effects more closely.  Correcting atmospheric dispersion on an ELT is a challenging task, even in today's telescopes. Solutions have been suggested for ELTs in recent years, but none have the capability of providing correction over a wide wavelength range (from visible to infrared). The problems will have to be addressed by a combination of technology (e.g. K-band wavefront sensing, NIR laser guide stars) and careful calibration. The first step, however, is a detailed theoretical understanding, and we intend to continue this work to make this possible."
    },
    "0807/0807.1266_arXiv.txt": {
        "abstract": "Large-scale diffuse radio emission is observed in some clusters of galaxies. There is ample of evidence that the emission has its origin in synchrotron losses of relativistic electrons, accelerated in the course of clusters mergers. In a cosmological simulation we locate the structure formation shocks and estimate their radio emission. We proceed as follows: Introducing a novel approach to identify strong shock fronts in an SPH simulation, we determine the Mach number as well as the downstream density and temperature in the \\MNUniverse\\ simulation which has $2 \\times 1024^3$ particles in a $500\\:h^{-1}\\,{\\rm Mpc}$ box and was carried out with non-radiative physics. Then, we estimate the radio emission using the formalism derived in \\citet{hoeft:07} and produce artificial radio maps of massive clusters. Several of our clusters show radio objects with similar morphology to large-scale radio relics found in the sky, whereas about half of the clusters show only very little radio emission. In agreement with observational findings, the maximum diffuse radio emission of our clusters depends strongly on their X-ray temperature. We find that the so-called accretion shocks cause only very little radio emission. We conclude that a moderate efficiency of shock acceleration, namely $\\xi_{\\rm e} \\lesssim 0.005$, and moderate magnetic fields in the region of the relics, namely 0.07 to $0.8\\,{\\rm \\mu G}$ are sufficient to reproduce the number density and luminosity of radio relics. ",
        "introduction": "\\label{sec-intro}  The large-scale structure of the Universe, composed of clusters, superclusters, filaments, and sheets of galaxies, is still in the process of formation. Overdense regions such as clusters and filaments keep accreting matter. Gas streams out of cosmic voids onto the sheets and filaments. When the newly accreted gas collides with the denser gas within these structures, shock fronts arise, dissipating the kinetic energy. In sheets and filaments, the gas follows the gravitational potential towards the clusters of galaxies. Eventually, the gas collides with the intra-cluster medium (ICM) with a few $1000\\:{\\rm km\\,s^{-1}}$ an gets heated to temperatures of $10^{7}$ to $10^{8}\\,{\\rm K}$. \\\\   The flow of gas is not as steady as the above picture suggests. A significant fraction of the gas accretion onto clusters is in the form of groups and clusters. The mergers of rich clusters are -- to our knowledge -- the most energetic events after the Big Bang. Kinetic energies of the order of $10^{64}$ erg are dissipated in giant shock waves. Only in recent years, X-ray telescopes have reached the necessary spatial and spectral resolution to detect the signatures of such shock waves in a few massive clusters. \\\\   A number of diffuse, steep-spectrum radio sources without optical identification have been observed in galaxy clusters. These sources have complex morphologies and show diffuse and irregular emission \\citep{kempner:01, slee:01, bacchi:03, feretti:05, giovannini:06}. They are usually subdivided into two classes, denoted as `radio halos' and `radio relics'. Cluster radio halos are unpolarised and have diffuse morphologies that are similar to those of the thermal X-ray emission of the cluster gas \\citep{giovannini:06}. Examples for clusters with radio halos are the Coma cluster \\citep{kim:89,deiss:97}, the galaxy cluster 1E\\,0657-56 \\citep{liang:00}, the X-ray luminous cluster A2163, and distant clusters such as CL\\,0016+16. The cluster A520 shows a halo with a low surface brightness with a clumpy structure \\citep{govoni:01}. Other examples can be found in \\citet{giovannini:99}. In general, radio halos are found in clusters with significant substructure and rich clusters with high X-ray luminosities and temperatures. The radio power correlates strongly with the cluster luminosity \\citep{feretti:05}. \\\\   Unlike halos, radio relics are typically located near the periphery of the cluster. They often exhibit sharp emission edges and many of them show strong radio polarisation \\citep{giovannini:04}. The sizes of relics and the distance to the cluster centre vary significantly. Examples for radio relic with sizes of one Mpc or even larger have been observed in Coma and A2256, which contain both a relic and a halo (as do A225, A1300, A2744 and A754). The cluster A3667 \\citep{roettgering:97} contains two very luminous, almost symmetric relics with a separation of more than five Mpc. The cluster A3376 shows an almost ring-like radio emission \\citep{bagchi:06}. The clusters A115 and A1664 show relics only at one side of the elongated X-ray distribution \\citep{govoni:01}. The relic source 0917+75 is particularly puzzling as it is located at 5 to $8\\,{\\rm Mpc}$ from the most nearby clusters. Other clusters show rather small relics as for example A85 \\citep{slee:01}. \\\\   The spectra of the diffuse radio sources indicate that their origin lies in synchrotron losses of relativistic electrons. The cooling time of the electrons which cause observable emission is of the order of one hundred Myr \\citep[see \\eg\\ the models in][]{slee:01}. The origin of the relativistic electron population which causes the emission is still not clear. There are essentially two classes of models that explain the presence of relativistic electrons. Either they are injected into the ICM via AGN activity or stellar feedback or they obtain their energy from particle acceleration at large-scale shock fronts in galaxy clusters. As discussed above, structure formation in the universe causes a variety of shock fronts in the intergalactic medium and the ICM \\citep{miniati:00, ryu:03, ryu:08}. These shock fronts are expected to be collisionless and capable of accelerating protons and electrons to relativistic energies. The correlation between the presence of diffuse radio emission in galaxy clusters and signs for a recent merger supports the scenario in which merger shocks generate the necessary relativistic electrons \\citep{feretti:06}. A3367 may serve as another piece of evidence: The radio relic is located where the merger induced bow shock is expected \\citep{roettiger:99}. Three mechanisms have been proposed for the action of the shock wave: (i) in-situ diffusive shock acceleration of electrons by the Fermi\\,I process \\citep[primary electrons,][]{ensslin:98, roettiger:99, miniati:01}, (ii) re-acceleration of electrons by compression of existing cocoons of radio plasma \\citep{ensslin:01, ensslin:02, hoeft:04}, or (iii) in-situ acceleration of protons and the production of relativistic electrons and positrons by ineleastic p-p collisions (secondary electrons). In the first two cases, the diffuse radio emission is roughly confined to the region of the shock fronts. In contrast, in the latter scenario the relativistic protons have long life times and can travel a large distance from their source before they release their energy. Hence, secondary electrons may be lead to radio halos. \\\\   In principle, observations of non-thermal cluster phenomena could provide an independent and complementary way of studying the growth of structure in our Universe and could shed light on the existence and the properties of the warm-hot intergalactic medium (WHIM), provided the underlying processes are understood. Sheets and filaments are predicted to host this WHIM with temperatures in the range $10^5$ to $10^7\\:{\\rm K}$ whose evolution is primarily driven by shock heating from gravitational perturbations breaking on mildly non-linear, non-equilibrium structures \\citep{cen:99}. Low-frequency aperture arrays such as {\\sc Lofar} are ideally suited to detect many diffuse steep-spectrum sources. In the next two years, the first {\\sc Lofar} survey is expected to chart a million galaxies and may discover hundreds of cluster radio halos \\citep{roettgering:06}. It is thus timely to study the distribution of diffuse radio sources in a cosmological context. \\\\    There have been efforts to simulate the non-thermal emission from galaxy clusters by modelling discretised cosmic ray (CR) energy spectra on top of Eulerian grid-based cosmological simulations \\citep{miniati:01b, miniati:04, miniati:07}. Recently, a series of papers explored the dynamical impact of CR protons on hydrodynamics in a cosmological SPH simulation \\citep{jubelgas:08, pfrommer:06, ensslin:07}. \\citet{skillman:08} presented a new method for identifying shock fronts in adaptive mesh refinement simulations and they computed the production of CRs in a cosmological volume adopting a nonlinear diffusive shock acceleration model. They found that CRs are dynamically important in galaxy clusters. \\citet{pfrommer:08} used {\\sc Gadget} simulations of a sample of galaxy clusters and implemented a formalism for CR physics on top of radiative hydrodynamics. They modelled relativistic electrons that are (i) accelerated at cosmological structure formation shocks and (ii) produced in hadronic interactions of CRs with protons of the ICM. \\citet{pfrommer:08} approximated both the CR spectrum and that of relativistic electrons locally by single power-laws with free parameters for the slope, the normalisation, and the low energy cut-off. Energy and momentum conservation, including source and sink terms, result in evolution equations for the spectra. They found that the radio emission in galaxy clusters is dominated by secondary electrons. Only at the location of strong shocks the contribution of primary electrons may dominate. In the cluster centres they found a radio emission of about $10^2 \\, h^3 \\: {\\rm mJy \\, arcmin^{-2}}$, while in the periphery the azimuthal average the emission is $10^{-3} \\, h^3 \\: {\\rm mJy \\, arcmin^{-2}}$. \\\\   Little is still known about the structure of shock fronts in the ICM. With high resolution imaging one can constrain the width of the transition and a deconvolution of the images gives an estimate for the density and pressure jump. Since the mean free path of protons in the cluster environment is of the order of Mpc, shock fronts in the ICM are collisionless. As studies of the best investigated collisionless shock front, namely the Earth bow shock, have revealed, the dissipation of the upstream kinetic energy is a complex process and depends on several shock properties (see \\eg\\ \\citet{burgess:07} for an introduction). For instance, the angle between the upstream magnetic field and the shock normal determines if an ion can gyrate between the upstream and downstream region, and the Mach number determines if instabilities in the shock region operate efficiently. However, shock fronts in the intra-cluster medium may differ significantly from the bow shock of the Earth as, for instance, the upstream plasma in the ICM is much less magnetised. Unfortunately, it is beyond the scope of current computer resources to study collisionless shock from the first principles, since scales from the electron gyro radius up to the large-scale structure of the shocks are involved. In hybrid simulations an effective small-scale response of electrons and protons is assumed. Using such a method, \\citet{kang:07} found that upstream CRs excite Alfv{\\'e}n waves and thereby amplify the magnetic field. \\\\   In this paper we combine a large cosmological simulation with a simple model for the radio emission of shock accelerated electrons. Our aim is to apply the emission model to the whole range of shock fronts generated during cosmic structure formation. A representative shock front sample is obtained from the \\MNUniverse\\ simulation which has been carried out with TreeSPH code {\\sc Gadget}-2. We have developed a novel approach for locating the shock fronts and to estimate their Mach number. For computing the radio emission we follow Hoeft \\& Br{\\\"uggen} (2007, HB07). They assumed that electrons are accelerated by diffusive shock acceleration and cool subsequently by synchrotron and inverse Compton losses. As a result the radio emission can be expressed as a function of downstream plasma properties, Mach number, and surface area of the shock front. Applying this radio emission model to the \\MNUniverse\\ simulation leads to radio-loud shock fronts with complex morphologies. We visualize these shock fronts to show where the radio emission is generated and compute artificial radio maps. As the \\MNUniverse\\ simulation provides a cosmologically representative cluster sample, we also investigate the relation between radio luminosity and X-ray temperature. \\\\   This paper is organised as follows: In \\Sec{sec-mn-universe} we briefly summarise the main characteristics of the \\MNUniverse\\ simulation. In \\Sec{sec-finding-shocks} we describe our approach for locating strong shock fronts in a SPH simulation and for estimating the Mach number. The radio emission model as worked out in HB07 is outlined and adopted for the SPH simulation in \\Sec{sec-radio-model}. Finally, in \\Sec{sec-results} we show our results for the shock fronts in the \\MNUniverse\\ simulation and for the radio properties of structure formation shocks. ",
        "conclusions": "\\label{sec-discussion}   We have presented a novel method for estimating the radio emission of strong shocks that occur during the process of structure formation in the universe. Our approach is based on an estimate for the shock surface area and we compute the radio emission per surface element using the Mach number of the shock and the downstream plasma properties. The advantage of our method is that we can accurately determine the emission of the downstream regime without using an approximate, discretised electron spectrum. The disadvantage is clearly that we do not include any evolution of the shock front and the downstream medium. However, our main goal here is to compute the radio emission from all shock fronts generated during cosmic structure formation. Our main assumption is that radio emission is caused by primary electrons accelerated at the shock front. To obtain a model suitable for the application in the framework of a cosmological simulation we have to neglect the complexity of collisonless shocks, details of the electron acceleration mechanisms, the amplification of magnetic fields by upstream cosmic rays, and so forth. The model reflects rather an expectation for the average radio emission entailed by structure formation shocks. \\\\   Our results show that in a cosmological simulation one indeed finds textbook examples for large-scale, ring-like radio relics as observed in A3667 and A3376. Our example {\\em A} shows two large half-shell shaped shock fronts and no radio emission in the centre of the merging cluster. However, cluster {\\em A} is not very massive, hence the radio luminosity of the two half-shells is low compared to A3667. Another nice example is cluster {\\em B} which shows a radio relic only one side of the cluster, similar to A115. Beside those spectacular relics in the periphery of galaxy clusters, we find that a lot of clusters show complex, radio-loud shock structures close to the cluster centre. This suggests that part of the central diffuse radio emission in galaxy clusters, \\ie the radio halo phenomena, can be attributed to synchrotron emission of primary electrons. \\\\   We have modelled the radio emission with conservative assumptions for the efficiency of shock acceleration, namely $\\xi_e = 0.005$, and for the strength of magnetic fields in the downstream region, namely of the order of 0.07 to $0.8\\,{\\rm \\mu G}$. Still, we find that the radio luminosity in the cluster region is on average significantly above observed values, see \\Fig{fig-cum-synchro}. Even lower values for the acceleration efficiency or the magnetic field strength would suffice to explain the abundance of diffuse radio objects on the sky. In conclusion, our findings support the scenario in which radio relics are generated by primary electrons. \\\\    In summary, we have analysed one of the largest hydrodynamical simulations of cosmic structure formation with respect to shock fronts in clusters of galaxies. We have developed a method to identify the shock fronts, to estimate their Mach numbers and their orientation. We have applied the method described in \\citet{hoeft:07} to determine the radio emission as a function of the shock surface area and the downstream plasma properties. Our analysis led to the following results:  \\begin{itemize}  \\item By evaluating the entropy and the velocity field we are able to identify strong shocks in the simulation and to determine their properties.  \\item Using conservative values for the efficiency of strong shocks to accelerate electrons and for strengths of magnetic fields, we are able to reproduce the number density and the luminosity of large-scale radio relics. Only a very few of the most massive clusters show such luminous, extended sources.   \\item Our model reproduces various spectacular sources of diffuse radio emission at the periphery of galaxy clusters. Other clusters show emission with complex morphology close to the cluster centre. These sources may be classified as part of a central radio halo, following the observational distinction between relics and halos.   \\item Our results reproduce the strong correlation between radio luminosity and cluster temperature. The highest radio luminosities occur in dense and hot environments such as massive clusters. In addition we find that a large number of galaxy clusters show only little diffuse radio emission.   \\item We find that the abundance of radio relics can be explained with efficiency of diffusive shock acceleration for electrons lower than $\\xi_e = 0.005$ and a strength of magnetic fields in the relic region lower than 0.07 to $0.8 \\, { \\rm \\mu G}$.   \\item All of the luminous radio relics belong to internal shock fronts. The accretion shocks are located at larger distances from the cluster centre. Since density and temperature are low at this location, their luminosity is too small to reach a similar flux as internal shock fronts.   \\end{itemize}    \\noindent {\\sc Acknowledgment }  MH acknowledges DLR funding under the grant 50 OX 0201. MB acknowledges the support by the DFG grant BR 2026/3. The \\MNUniverse\\ simulations have been done at the Barcelona Supercomputing Center (Spain) and analysed at NIC J\\\"ulich (Germany). GY thanks MEC (Spain) for financial support under project numbers FPA2006-01105 and AYA2006-15492-C03. Our collaboration has been supported by the European Science Foundation (ESF) for the activity entitled `Computational Astrophysics and Cosmology' ({\\sc AstroSim}).    \\newcommand{\\aap  }{A\\&A} \\newcommand{\\araa }{ARA\\&A} \\newcommand{\\apj  }{ApJ} \\newcommand{\\apjs }{ApJS} \\newcommand{\\apjl }{ApJL} \\newcommand{\\apss }{ApSS} \\newcommand{\\aapr }{A\\&A~Rev.} \\newcommand{\\aj   }{AJ} \\newcommand{\\jgr  }{JGR} \\newcommand{\\mnras}{MNRAS} \\newcommand{\\nat   }{Nature} \\newcommand{\\physrep  }{PhysRep}"
    },
    "0807/0807.0110_arXiv.txt": {
        "abstract": "We present a comprehensive study of the nature of the SDSS galaxies divided into various classes based on their morphology, colour, and spectral features. The SDSS galaxies are classified into early-type and late-type; red and blue; passive, H{\\protect\\scriptsize II}, Seyfert, and LINER, returning a total of 16 fine classes of galaxies. We examine the luminosity dependence of seven physical parameters of galaxies in each class. We find that more than half of red early-type galaxies (REGs) have star formation or AGN activity, and that these active REGs have smaller axis ratio and bluer outside compared to the passive REGs. Blue early-type galaxies (BEGs) show structural features similar to those of REGs, but their centres are bluer than REGs. H{\\protect\\scriptsize II} BEGs are found to have bluer centres than passive BEGs, but H{\\protect\\scriptsize II} REGs have bluer outside than passive REGs. Bulge-dominated late-type galaxies have red colours. Passive red late-types are similar to REGs in several aspects. Most blue late-type galaxies (BLGs) have forming stars, but a small fraction of BLGs do not show evidence for current star formation activity.  Differences of other physical parameters among different classes are inspected, and their implication on galaxy evolution is discussed. ",
        "introduction": "One of the fundamental issues of the observational cosmology is the evolutionary connection between various classes of galaxies. Today, galaxies are classified with various criteria. The most classical classification of galaxies is the Hubble sequence: elliptical galaxies, lenticular galaxies, spiral galaxies, barred spiral galaxies, and irregular galaxies \\citep{hub36,san61,dev74}. The main criterion of this classification is the morphology of galaxies; that is, the existence, size ratio, and appearance of spiral arms, disc, bulge, and bar. The Hubble sequence was established mainly based on the galaxies without nuclear activity, and such galaxies are often called \\emph{normal} galaxies. On the other hand, more and more galaxies showing active nuclei with broad line emission are being found; they are often called \\emph{abnormal} galaxies. According to the features of the line emission, active galactic nuclei (AGN) host galaxies are classified into several sub-classes: Seyfert 1 galaxies, Seyfert 2 galaxies, broad-line radio galaxies, narrow-line radio galaxies, low ionisation nuclear emission regions (LINERs), and so on. Luminosity is another criterion to classify galaxies. \\citet{san84} defined the galaxies with M$_{B} > -18$ as \\emph{dwarf} galaxies. In addition to those general classes of galaxies, there are some unusual galaxy classes with interesting properties: E+A galaxies with the spectral features of both very old stellar populations and very young stellar populations \\citep{dre92,yan06,yan08}; ultra-luminous infrared galaxies (ULIRGs) that are very bright in the mid- and far-infrared bands \\citep{san96,hwa07}; blue compact galaxies that show very compact morphology and high surface brightness with blue colour \\citep{thu97}; extremely red objects (EROs) whose optical $-$ infrared colour is extremely red \\citep{els88}; and so on. Recently, to deal with a large amount of survey data, many astronomers have classified galaxies simply into red sequence galaxies and blue sequence galaxies according to their colour-magnitude relation \\citep[e.g.][]{mar07}.  For a long time,  the individual properties of galaxies in those various classes have been investigated, and many efforts have been made to answer fundamental questions about galaxy evolution and connections between galaxy classes. Some examples of such fundamental questions are as follows: Why is there a conspicuous bimodality in the colour distribution of galaxies? Why does the colour bimodality not exactly agree with the morphological segregation? How do the environments affect the properties of galaxies? How did active galactic nuclei (AGNs) form? Is there any transition among the different classes of galaxies? Recent studies using large survey data discovered several interesting aspects of galaxy evolution, providing some answers to those fundamental questions.  For example, \\citet{bal06} studied the bivariate luminosity functions with galaxy type classification, and found that there is a clear morphological bimodality supporting the idea that merger and accretion are associated with bulges and discs, respectively. \\citet{ber05} showed that both luminosity and colour of early-type galaxies are correlated with stellar velocity dispersion, and that velocity dispersion may be also closely correlated with the age and metal abundance of early-type galaxies. \\citet{cho07} found that late-type galaxies show wider dispersion in several physical quantities than early-type galaxies, and that those physical quantities manifest different behaviours across $M_{\\star}\\pm1$ mag. \\citet{par07} investigated the environmental effects on various physical properties of galaxies, finding that a key constraint on galaxy formation models is the morphology-density-luminosity relation. \\citet{mat06} analysed the colour, 4000{\\AA} break, and age in various spectral classes of galaxies, suggesting that the median light-weighted stellar age of galaxies is directly responsible for the colour bimodality in the galaxy population.  Those previous studies, however, are not yet enough to explain the origins of the various galaxy classes and the evolutionary connections between the classes. One of the major limitations is that the galaxy classifications in most previous studies were often limited to only one or two properties. For instance, automatically-classified morphology \\citep[e.g.][]{par05,bal06}, galaxy colour \\citep[e.g.][]{mar07}, or spectral line features \\citep[e.g.][]{mat06} are used frequently in recent studies based on large survey data, but a multilateral classification study using all these criteria at the same time has not yet been seen. Such simplifications may be often very useful to understand the various and complicated phenomena in galaxy evolution. However, a simple classification can not distinguish detailed aspects of galaxy evolution. For example, early-type galaxies are often regarded as red and passive galaxies in most studies, but that is not always true. Some kinds of galaxies with early-type morphology were found to have blue colours possibly originating from young stellar populations or active nuclei \\citep{abr99,im01,fer05,lee06}. It is difficult to understand these kinds of unusual classes in studies with simple classifications.  Closely related to galaxy classification, a couple of studies searching for principal components of galaxy properties have been conducted recently. \\citet{ell05} examined the distribution of photometric, spectroscopic and structural parameters for 350 nearby galaxies using the Millennium Galaxy Catalogue \\citep{lis03}, arguing that most properties show a clear distinction between early-type galaxies and late-type galaxies. Later, \\citet{con06} carried out principal-component analyses of the properties of 22000 galaxies at $z<0.05$ using the Third Reference Catalogue of Bright Galaxies \\citep[RC3;][]{dev91}, finding that the three parameters determining a galaxy's physical state may be mass, star formation and interactions/mergers. These studies presented how useful to consider various parameters in galaxy classification, although some important components like AGNs were not considered in their analyses. \\citet{con06} showed that multiple independent components, rather than any single component, may determine the various properties of galaxies.  Therefore, to understand the nature of galaxies more comprehensively and to give stronger constraints on galaxy evolution models to \\emph{explain all kinds of galaxies}, it is necessary to investigate the nature of galaxies in various and finely-divided classes, and to find out the evolutionary connections between the classes. We have been doing a comprehensive study on a set of fine galaxy classes in the Sloan Digital Sky Survey \\citep[SDSS;][]{yor00}, based on their morphology, colour and spectral features. In this paper, the first in the series, we present the optical properties of galaxies in various fine classes. The outline of this paper is as follows. Section 2 shows the data set we used, and \\S3 describes the methods to classify the SDSS galaxies and to select volume-limited samples. We present the statistics of selected optical parameters and their luminosity dependence in \\S4. Based on the optical properties, we discuss the nature of galaxies in the fine classes in \\S5. Finally, the conclusions in this paper are given in \\S6. Throughout this paper, we adopt the cosmological parameters $h=0.7$, $\\Omega_{\\Lambda}=0.7$, and $\\Omega_{M}=0.3$. ",
        "conclusions": "We conducted a comprehensive study of the nature of the SDSS galaxies in various classes based on their morphology, colour and spectral features. Using three criteria, we classified the SDSS galaxies into early-type and late-type; red and blue; passive, H{\\protect\\scriptsize II}, Seyfert and LINER, resulting in 16 fine classes of galaxies in total. We estimated the luminosity dependence of seven physical quantities in each class, and compared the properties among classes, using a sub-sample with the same distribution of the velocity dispersion. From the analysis, we found that each galaxy class has its own distinguishable features. This shows that an analysis based on a simple classification may have a risk of mixing up different kinds of objects with different natures.  The red early-type galaxies include well-known typical elliptical galaxies ($p$REGs), but some REGs show evidence for additional star formation in their outer regions ($h$REGs). Some other REGs with AGNs ($s$REGs, $l$REGs) have structural properties showing the existence of larger disc components than $h$REGs, indicating the relationship between AGN activity and gas accretion. The blue early-type galaxies may be in the process of bulge formation. The structural similarity between $p$REGs, $p$BEGs and $h$BEGs, supports an evolutionary sequence of $h$BEGs $\\rightarrow$ $p$BEGs $\\rightarrow$ $p$REGs. Seyfert early-type galaxies have a close relationship with the outer disc components of early-type galaxies, and the disc components in LINER early-type galaxies ($l$REGs, $l$BEGs) are smaller than those in Seyfert early-type galaxies, on average. The blue late-type galaxies have properties in agreement with typical spiral galaxies. Most of them are star-forming ($h$BLGs), but a very small fraction of BLGs do not show any evidence of current star formation ($p$BLGs). Some of BLGs have an AGN ($s$BLGs, $l$BLGs), which are less detected at large inclination, and $s$BLGs show particularly bright centres on average. The median axis ratio in each class shows that some intrinsic BLGs with large inclination may often be classified as red late-type galaxies, due to strong extinction by dust in their disc. In addition to dust extinction, a large bulge-to-disc ratio may make a late-type galaxy red. Many $p$RLGs seem to be bulge-dominated late-type galaxies with small inclination, in which line emission is not detected due to the limited size of the spectroscopic fibre aperture. Like in BLGs, AGN activity is detected in some RLGs ($s$RLGs, $l$RLGs), which have small inclination.   This paper is the first in the series of comprehensive studies on the nature of the SDSS galaxies in finely-divided classes. In the following papers, we will inspect various aspects of galaxies in these classes, focusing on their multi-wavelength properties and environmental effects."
    },
    "0807/0807.3219_arXiv.txt": {
        "abstract": "{} {This study has three principal aims: (i) to increase the number of detected pulsation modes of 44 Tau, especially outside the previously known frequency ranges, (ii) to study the amplitude variability and its systematics, and (iii) to examine the combination frequencies.} {During the 2004/5 and 2005/6 observing seasons, high-precision photometry was obtained with the Vienna Automatic Photoelectric Telescope in Arizona during 52 nights. Together with previous campaigns, a data base from 2000 to 2006 was available for multifrequency analyses.} {Forty-nine pulsation frequencies  are detected, of which 15 are independent pulsation modes and 34 combination frequencies or harmonics. The newly found gravity mode at 5.30 cd$^{-1}$ extends the known frequency range of instability. Strong amplitude variability from year to year is found for the $\\ell$ = 1 modes, while the two radial modes have essentially constant amplitudes. Possible origins of the amplitude variability of the  $\\ell$ = 1 modes, such as precession of the pulsation axis, beating and resonance effects are considered. The amplitudes of the combination frequencies, $f_i + f_j$,  mirror the variations in the parent modes. The combination parameter, which relates the amplitudes of the combination frequencies to those of the parent modes, is found to be different for different parents. } {} ",
        "introduction": "The asteroseismic comparison between measurements and theoretical models of the radial and nonradial pulsation of stars on and near the main sequence involves the exact matching of the values of the pulsation frequencies, as well as reproducing the frequency range in which instability occurs. However, it is observationally difficult to determine the important borders of the frequency range in which a particular star is unstable. Due to the small amplitude growth rates of modes on the edges of the frequency instability range, only modes with small photometric amplitudes are expected to appear in these frequency regions.  If we consider the $\\delta$ Scuti star FG Vir as an example (Breger et al. 2005), all photometric amplitudes in the 30 - 45 cd$^{-1}$ region had amplitudes less than 1.0 millimag. In fact, only the addition of 218 additional nights of photometry to the previous data allowed the detection of the pulsation modes near 40 cd$^{-1}$. In addition to the problem of small amplitudes in specific frequency regions, it is also essential to separate the `real' modes from the combination frequencies, $f_i \\pm f_j$, which occur in these high-frequency and low-frequency regions. This separation requires very high frequency resolution so that the chance of accidental agreements between the expected and measured frequency values is essentially reduced to zero. In practice this requires campaigns covering several years.  The requirement of multiyear campaigns is strengthened by the amplitudes of many pulsators showing variability on timescales of months and, especially, years. To determine what types of modes show amplitude variability and to search for the astrophysical origin of this amplitude (and frequency) variability requires extensive observing campaigns. Finally, successful asteroseismology also demands the identification of individual pulsation modes. Both spectroscopic and photometric techniques have been applied: however, the large amounts of observing time required are easier to obtain on smaller telescopes with photometers.  The $\\delta$ Scuti star 44 Tau is ideal for a detailed study because of its extremely low rotational velocity of 3 $\\pm$ 2 km s$^{-1}$ (Zima et al. 2007), for which second-order effects of rotation can be neglected, and because its amplitudes place the star between the high-amplitude pulsators (HADS) with dominant radial modes and the average low-amplitude $\\delta$ Scuti variable with mostly nonradial modes. Zima et al. (2007) present high-resolution spectroscopy of 44 Tau, which was used together with radial velocity data to derive the photospheric element abundances and provide $m$ values for the pulsation mode identifications. Six axisymmetric and two prograde modes were identified.  In their asteroseismic study of 44~Tau, Lenz et al. (2008) present spherical mode degree identifications for the dominant modes of 44 Tau. Photometric amplitude ratios and phase differences suggest two radial, four $\\ell$ = 1, and three $\\ell = 2$ modes. Due to the measured $\\log g$ value of 3.6 $\\pm$ 0.1, both main sequence and post-main sequence models were examined. The identified modes can be fitted well in both evolutionary stages. The predicted frequency ranges for unstable modes are in good agreement with the observed ranges adopted by Lenz et al. (2008). Due to the different envelope structure of the main sequence and post-main sequence model, the borders of instability differ slightly. Consequently, new independent frequencies in the low-frequency and high-frequency regions possess important information for improving the models.  We have previously presented 90 nights of photometric observations obtained during the three observing seasons 2000/2001, 2001/2002, and 2002/3 (Antoci et al. 2007: this paper also lists a detailed observational history of the star). The data from three years led to detecting 29 frequencies, of which 13 were independent pulsation modes. The two comparison stars used showed low-frequency variability at the millimag level, which is, regrettably, quite common. Otherwise, these measurements are very accurate. Consequently, the publication omitted the 0 to 5 cd$^{-1}$ region awaiting future measurements with different comparison stars, as reported in the present paper. ",
        "conclusions": "Two additional seasons of high-precision photometric data of the $\\delta$ Scuti star 44 Tau were gathered with the Vienna Automatic Photoelectric Telescope in Arizona. A frequency analysis of the total data set including photometry from 2000 to 2006 led to detecting 49 pulsation frequencies, of which 15 are independent pulsation modes and 34 combination frequencies or harmonics.  We find that the amplitudes of the combination frequencies, $f_i + f_j$,  mirror the variations of the parent modes. The combination parameter, $\\mu$, which relates the amplitudes of the combination frequencies to those of the parent modes, is found to be different for different radial and nonradial parents. The combination frequencies involving the radial fundamental mode have combination parameters, $\\mu$, of 0.003, while those involving the first radial overtone can be as high as $\\approx 0.009$.  The time-base of 5 years allows us to study amplitude variability in more detail. Strong amplitude variability from year to year was found for the $\\ell$ = 1 modes, while the two radial modes have essentially constant amplitudes. Moreover, the amplitudes of all of the $\\ell = 1$ modes decrease during most of the observed period. We examined several possible reasons for the amplitude variability of the $\\ell$ = 1 modes:  (i) Precession of the pulsation axis: This causes variable amplitudes because the mode visibility depends on the inclination angle. While the amplitudes of the radial modes remain constant, the amplitudes of $\\ell$ = 1 modes will change. However, both axisymmetric and non-axisymmetric modes decrease at the same time; this observation cannot be explained in this scenario because the visibility of $m = 0$ modes behaves opposite to their non-axisymmetric counterparts.  (ii) Beating of close frequencies: Due to the slow rotation of 44 Tau, beating between the components of a rotationally split $\\ell=1$ triplet has been examined within the limits of the measured rotation rate. Generally, the predicted beating timescales are too short compared to the observed amplitude modulation. However, beating with hitherto unresolved, additional close frequencies (which are not a component of the same rotational triplet) could produce amplitude modulation at the observed timescales.  (iii) Resonance effects: The expected timescales for amplitude modulation in a 44 Tau model were computed. We find that only for a fraction of the variable $\\ell = 1$ modes could the predicted variability timescales match the observed modulation periods.  None of these three explanations can explain the observed amplitude variability completely. However, a combination of all these effects cannot be excluded. Even more data would be required to clarify this point.  The newly found gravity mode at 5.30 cd$^{-1}$ extends the previously known frequency range to lower frequencies. This detection agrees better with the predicted instability ranges for main sequence models. However, our main sequence models cannot fit the measured fundamental parameters such as effective temperature. Our post-main sequence models provide a good fit of the fundamental parameters but predict stability at 5.30 cd$^{-1}$. The problem with the present models underestimating the observed $\\delta$ Scuti star instability is common to many $\\delta$ Scuti stars and may be related to still incomplete opacity data."
    },
    "0807/0807.3443_arXiv.txt": {
        "abstract": "We report on an ongoing investigation into a seismic calibration of solar models designed for estimating the main-sequence age and a measure of the chemical abundances of the Sun. Only modes of low degree are employed, so that with appropriate modification the procedure could be applied to other stars. We have found that, as has been anticipated, a separation of the contributions to the seismic frequencies arising from the relatively smooth, glitch-free, background structure of the star and from glitches produced by helium ionization and the abrupt gradient change at the base of the convection zone renders the procedure more robust than earlier calibrations that fitted only raw frequencies to glitch-free asymptotics. As in the past, we use asymptotic analysis to design seismic signatures that are, to the best of our ability, contaminated as little as possible by those uncertain properties of the star that are not directly associated with age and chemical composition. The calibration itself, however, employs only numerically computed eigenfrequencies. It is based on a linear perturbation from a reference model. Two reference models have been used, one somewhat younger, the other somewhat older than the Sun. The two calibrations, which use BiSON data, are more-or-less consistent, and yield a main-sequence age $t_\\odot=4.68\\pm0.02\\,$Gy, coupled with a formal initial heavy-element abundance $Z=0.0169\\pm0.0005$. The error analysis has not yet been completed, so the estimated precision must be taken with a pinch of salt. ",
        "introduction": "The only way by which the age of the Sun can be estimated to a useful degree of precision is by accepting the basic tenets of solar-evolution theory and measuring those aspects of the structure of the Sun that are predicted by the theory to be indicators of age. The structure measurements must be carried out seismologically, and evidently one expects greatest reliability of the results when all the available helioseismic data are employed. However, the most relevant modes are those of lowest degree, because it is they that penetrate most deeply into the energy-generating core where the relic helium-abundance variation records the integrated history of nuclear transmutation. Moreover, it is also only they that can be measured in other stars. Therefore, there has been some interest in calibrating theoretical stellar models using only low-degree modes. The prospect was first discussed in detail by \\cite[Christensen-Dalsgaard (1984, 88)]{jcd84, jcd88}, \\cite{ulrich86} and \\cite{dog87}, although prior to that it had already been pointed out that the helioseismic frequency data that were available at the time indicated that either the initial helium abundance $Y_0$ or the age $t_\\odot$, or both, are somewhat greater than the generally accepted values \\cite[(Gough 1983)]{dog83}, an inference which is consistent with our present findings. Subsequent, more careful, calibrations were carried out by \\cite{guenther89}, \\cite{dog-nov90}, \\cite{guenther-demarque97}, \\cite{weiss-schlattl98}, \\cite{wd99}, \\cite{dog01} and Bonanno, Schlattl \\& Patern\\`o (2002). Not all of them addressed the influence of uncertainties in $Y_0$ on the determination of $t_\\odot$.  As a main-sequence star ages, helium is produced in the core, increasing the mean molecular mass $\\mu$, preferentially at the centre, and thereby reducing the sound speed. The resulting functional form of the sound speed $c(r)$ depends not only on age $t_\\odot$ but also on the relative augmentation of $\\mu(r)$, which itself depends on the initial absolute value of $\\mu$, and hence on $Y_0$. \\cite{dog01} tried to separate these two effects using the degree dependence of the small separation $d_{n,l}=3(2l+3)^{-1}(\\nu_{n,l}-\\nu_{n-1,l+2})$ of cyclic frequencies $\\nu_{n,l}$, where $n$ is order and $l$ is degree. This is possible, in principle, because modes of different degree and similar frequency sample the core differently. However, that difference is subtle, and the sensitivity to the relatively fine distinction between the effects of $t_\\odot$ and $Y_0$ on the functional form of $c(r)$ in the core is low. Consequently the error in the calibration produced by errors in the observed frequency data is uncomfortably high.  This lack of sensitivity can be overcome by using, in addition to core-sensitive seismic signatures, the relatively small oscillatory component of the eigenfrequencies induced by the sound-speed glitch associated with helium ionization \\cite[(Gough 2002)]{dog02}, whose amplitude is close to being proportional to helium abundance $Y$ \\cite[(Houdek \\& Gough 2007a)]{hgdog07a}. The neglect of that component in the previously employed asymptotic signature $d_{n,l}$ had not only omitted an important diagnostic of $Y$ but had also imprinted an oscillatory contamination in the calibration as the limits $(k_1, k_2)$, where $k=n+\\frac{1}{2}l$, of the adopted mode range was varied \\cite[(Gough 2001)]{dog01}. It therefore behoves us to decontaminate the core signature from glitch contributions produced in the outer layers of the star (from both helium ionization and the abrupt variation at the base of the convection zone, and also from hydrogen ionization and the superadiabatic convective boundary layer immediately beneath the photosphere). To this end a helioseismic glitch signature has been developed by \\cite{hgdog07a}, from which frequency contributions $\\delta\\nu_{n,l}$ can be computed and subtracted from the raw frequencies $\\nu_{n,l}$ to produce effective glitch-free frequencies $\\nu_{{\\rm s}n,l}$ to which a glitch-free asymptotic formula (\\ref{eq:asymp}) can be fitted. The solar calibration is then accomplished by interpolating the theoretical seismic signatures computed on a grid of solar models to the observations, using a standard grid to compute derivatives with respect to $t_\\odot$ and $Y_0$, and a carefully computed reference solar model designed to be close to the Sun. The result of the first preliminary calibration by this method, using BiSON data, has been reported by \\cite{hgdog07b}. Here we enlarge on our discussion of the analysis, and we augment our results with a calibration based on a second reference solar model. \\vspace{-5mm} \\begin{figure} \\centering \\includegraphics[height=.30\\textheight]{hg_fig1.eps} \\caption{Top left: The symbols (with error bars obtained under the assumption that the raw frequency errors are independent) represent second differences, $\\Delta_2\\nu$, of low-degree solar frequencies from BiSON. Top right: The symbols are second differences $\\Delta_2\\nu$ of adiabatic pulsation eigenfrequencies of solar Model S of \\cite{jcd96}. The solid curves in both panels are the diagnostics\\, (\\ref{eq:delnu}) -- (\\ref{eq:secdiff}), whose eleven parameters have been adjusted to fit the data optimally. Bottom: The symbols denote contributions $\\delta\\nu$ to the frequencies produced by the acoustic glitches of the Sun (left panel) and Model S (right panel).} \\end{figure} ",
        "conclusions": ""
    },
    "0807/0807.0440_arXiv.txt": {
        "abstract": "We present high spatial resolution 21~cm HI observations of EA01A and EA01B, a pair of interacting post-starburst, or E+A, galaxies at $z = 0.0746$. Based on optical HST/WFPC2 images, both galaxies are known to display disturbed morphologies. They also appear to be linked by a bridge of stars. Previous HI observations \\citep{chang01} had already uncovered sizable quantities of neutral gas in or near these galaxies but they lacked the spatial resolution to locate the gas with any precision within this galactic binary system. We have analysed deep, high resolution archival VLA observations of the couple. We find evidence for three gaseous tidal tails; one connected to EA01A and two emanating from EA01B. These findings confirm, independently from the optical imaging, that {\\em (i)} EA01A and EA01B are actively interacting, and that, as a consequence, the starbursts that occurred in these galaxies were most likely triggered by this interaction, and that {\\em (ii)} $6.6\\pm 0.9\\times10^9$~{\\Msun} of neutral gas are still present in the immediate vicinity of the optical bodies of both galaxies. The \\HI\\ column density is lowest at the optical positions of the galaxies, suggesting that most of the neutral gas that is visible in our maps is associated with the tidal arms and not with the galaxies themselves. This might provide an explanation for the apparent lack of ongoing star formation in these galaxies. ",
        "introduction": "Post-starburst galaxies (or PSGs, or E+A galaxies, or k+a/a+k galaxies) are characterized by their optical spectra, independent of any morphological or photometric considerations. PSGs contain a sizable population of very young stars, producing very prominent Balmer lines, but lack ongoing star formation, hence the absence of detectable emission lines. This suggests that PSGs are indeed observed very shortly, within less than 1~Gyr, after the end of a vigorous, abruptly truncated starburst \\citep{dres83,dres99,pog99}. Today, they constitute only a small fraction of the galaxy cluster population ($<1$\\%, \\citet{fab91}); at intermediate-redshifts, however, they constituted a substantial cluster population \\citep{bel95}. Although the trigger and the abrupt end of the starburst is still not fully understood, photometry of PSGs shows evidence for disturbed morphologies, e.g. tidal tails, suggesting that in many cases the trigger is most likely associated to a galaxy-galaxy merger or interaction \\citep{zab96,yang04,blake04,tran03,goto05,yang08}. Using numerical simulations, \\citet{bekki05} has shown that PSGs can be formed via a major merger of two gas-rich spiral galaxies. In these merger simulations, a starburst is triggered that consumes the available gas within a timespan of roughly 1~Gyr, ending the starburst abruptly. Recent \\HI\\ observations have uncovered large quantities of neutral hydrogen gas in a significant fraction of PSGs \\citep{chang01,buyle06,helmboldt07,buyle08}. This opens up the possibility of investigating the hypothesized merger origin of many PSGs.    \\begin{figure*} \\includegraphics[width=17cm]{f1.eps} \\caption{Channel maps overplotted on an optical HST/WFPC2 F702W image of EA01A/B. The contour levels are 2$\\sigma$, $3\\sigma$,\\ldots with $\\sigma$=0.2~mJy~beam$^{-1}$. The synthesized beam is depicted in the bottom-right corner of the first map (top left). \\label{channels}} \\end{figure*}  Indeed, high-resolution interferometric radio observations of the neutral gas would provide valuable additional information. On the one hand, \\HI\\ morphologies and velocity fields, if they are distorted, could corroborate the merger hypothesis, and, moreover, could reveal whether at least part of the neutral gas is still gravitationally bound and will remain available for future star formation. We selected the PSG binary EA01A/B as an interesting target for our study. Both galaxies have a PSG-type optical spectrum and have the same recession velocity. EA01A is the bluest galaxy of the PSG sample of \\cite{zab96}, suggesting it is also the youngest sample member. HST/WFPC2 images in the F435W and F702W bands \\citep{yang04} provide strong evidence that the couple is interacting. EA01A, positioned about 11$''$ east of EA01B, contains many stellar clusters that are very conspicuous compared with the overall rather low surface brightness of this galaxy. This diffuse appearance suggests it is on the verge of disintegrating due to the injected orbital energy. EA01B, on the other hand, is a bulge dominated early-type spiral galaxy; it sports strongly asymmetric stellar arms. Both galaxies appear to be connected by a stellar bridge. Previous HI observations \\citep{chang01} had already uncovered sizable quantities of neutral gas in or near these galaxies but they lacked the spatial resolution to locate the gas with any precision within this galactic binary system.  In section~\\ref{obs}, we describe our data reduction and analysis of deep archival VLA \\HI\\ observations of the EA01A/B galaxy binary. The results are presented and discussed in section~\\ref{discussion} and summarized in section~\\ref{conclusions}. ",
        "conclusions": "EA01A/B is a close pair of post-starburst (E+A) galaxies, surrounded by some $7 \\times 10^9$~{\\Msun} of neutral gas. Most of this gas resides in what appear to be three tidal arms, two of which are connected with EA01B. Together with optical HST/WFPC2 images, these observations show that EA01A and EA01B are actively interacting. The galaxies themselves show up as minima in the \\HI\\ column density. This lack of galaxy-bound dense neutral gas is most likely connected with the fact that these galaxies are in a post-starburst phase."
    },
    "0807/0807.2029_arXiv.txt": {
        "abstract": "We present 3.5m Apache Point Observatory\\footnote{The Apache Point Observatory 3.5-meter telescope is owned and operated by the Astrophysical Research Consortium.} second-epoch spectra of four low-metallicity emission-line dwarf galaxies discovered serendipitously in the Data Release 5 of the Sloan Digital Sky Survey (SDSS) to have extraordinary large broad H$\\alpha$ luminosities, ranging from 3 $\\times$ 10$^{41}$ to 2 $\\times$ 10$^{42}$ erg s$^{-1}$. The oxygen abundance in these galaxies is very low, varying in the range 12 + log O/H = 7.36 -- 7.99. Such extraordinarily high broad H$\\alpha$ luminosities cannot be accounted for by massive stars at different stages of their evolution. By comparing with the first-epoch SDSS spectra, we find that the broad H$\\alpha$ luminosities have remained constant over a period of 3 -- 7 years, which probably excludes type IIn supernovae as a possible mechanism of broad emission. The emission most likely comes from accretion disks around intermediate-mass black holes with lower mass limits in the range $\\sim$ 5$\\times$10$^5$ $M_\\odot$ -- 3$\\times$10$^6$ $M_\\odot$. If so, these four objects form a new class of very low-metallicity AGN that have been elusive until now. The absence of the strong high-ionization lines [Ne {\\sc v}] $\\lambda$3426 and He {\\sc ii} $\\lambda$4686 can be understood if the nonthermal radiation contributes less than $\\sim$ 10\\% of the total ionizing radiation. ",
        "introduction": "Active galactic nuclei (AGN) are thought to be powered by massive black holes at the centers of galaxies, accreting gas from their surroundings. Observations of AGN show that they generally possess a high metallicity, varying from solar to supersolar metallicities \\citep{S98,H02}. While the derived metallicities do depend on the detailed model assumptions, this appears to be a solid conclusion. Gas metallicity is known to be strongly correlated with the stellar mass of the host galaxy \\citep{T04}. Since AGN are usually found in massive, bulge-dominated galaxies that have converted most of their gas into stars by the present epoch, their gas metallicities are generally high. A question then arises: do low-metallicity AGN exist? If so, can we find them in low-mass galaxies?  To address these questions, \\citet{Gr06} have searched the Sloan Digital Sky Survey (SDSS) Data Release 4 (DR4) spectroscopic galaxy sample of over 500,000 objects to select out $\\sim$170,000 emission-line galaxies with high S/N spectra. They then use diagnostic line ratios to select out 23,000 Seyfert 2s galaxies.  Imposing an upper mass limit of 10$^{10}$ $M_\\odot$ to restrict themselves to low-mass galaxies, they are left with a sample of only $\\sim$ 40 AGN, which they found appear to have metallicities around half that of typical AGN, i.e. having solar or slightly subsolar values. The same high metallicity range is found in the sample of low-mass AGN of \\citet{G07}. Assessing their findings, \\citet{Gr06} are led to another question: ``Why are there no AGN with even lower metallicities?'' In this paper, we suggest that these low-metallicity AGN do exist although they are extremely rare.  In the course of a long-range program to search for extremely metal-deficient emission-line dwarf galaxies, \\citet{I07} have used the SDSS DR5 database of 675,000 spectra to assemble a large sample of emission-line galaxies. Two criteria were applied: 1) the [O {\\sc iii}] $\\lambda$4363 line must be detected to allow for a direct determination of element abundances; and 2) obvious high-metallicity AGN spectra are excluded. Thus, contrary to \\citet{Gr06} and \\citet{G07}, \\citet{I07} were not specifically looking for AGN. These criteria resulted in a sample of $\\sim$10,000 emission-line galaxies (ELG). While studying that sample to look for ELGs with broad components in their strong emission lines, \\citet{I07} came across four galaxies with very unusual spectra. The general characteristics of the four galaxies are given in Table \\ref{tab1}. Their absolute magnitudes are typical of dwarf galaxies. Because of their relatively large distance ( $z$ $\\sim$ 0.1-0.3) and relatively small angular sizes ($\\sim$ 1\\arcsec-2\\arcsec, only slightly larger than the seeing disk), their SDSS images (Fig. \\ref{images}) do not show much details. They possess a compact structure. Two galaxies, J1025+1402 and J1047+0739, have an approximately round shape, while the other two more distant galaxies, J0045+1339 and J1222+3602, have a distorted shape suggestive of mergers. Their colors are not blue like the other ELGs, but vary from red to yellow to green. Their spectra shown in Fig. \\ref{spectra}, resemble those of moderately to very low-metallicity high-excitation H {\\sc ii} regions: their oxygen abundances are in the range 12$+$$\\log$ O/H $\\sim$7.4--7.9, i.e. their heavy element mass fractions vary from $Z_\\odot$/19 to $Z_\\odot$/5 if the solar calibration 12$+$$\\log$O/H$=$8.65 of \\citet{A05} is adopted. \\citet{I07} found that there is however a striking difference: the strong permitted emission lines, mainly the H$\\alpha$ $\\lambda$6863 line, show very prominent broad components.  These are characterized by somewhat unusual properties: 1) their H$\\alpha$ full widths at zero intensity $FWZI$ vary from 102 to 158 \\AA, corresponding to expansion velocities between 2200 and 3500 km s$^{-1}$; 2) the broad H$\\alpha$ luminosities $L_{br}$ are extraordinarily large, varying from 3$\\times$10$^{41}$ to 2$\\times$10$^{42}$ erg s$^{-1}$.  This is to be compared with the range 10$^{37}$--10$^{40}$ erg~s$^{-1}$ found by \\citet{I07} for the other ELGs with broad-line emission.  The ratio of H$\\alpha$ flux in the broad component to that in the narrow component varies from 0.4 to 3.4, as compared to 0.01--0.4 for the other galaxies; 3) the Balmer lines show a very steep decrement, suggesting collisional excitation and that the broad emission comes from very dense gas ($N_e$$\\gg$10$^{4}$ cm$^{-3}$). Evidently, these galaxies are exceedingly rare, since they constitute only 4/675,000 or 0.0006\\% of our original sample.  To account for the broad line emission in these four objects, \\citet{I07} have considered various physical mechanisms: a) Wolf-Rayet (WR) stars; b) stellar winds from Ofp or luminous blue variable stars; c) single or multiple Supernova (SN) remnants propagating in the interstellar medium; d) SN bubbles; e) shocks propagating in the circumstellar envelopes of \\hbox{type\\,IIn} SNe; and f) AGN. While mechanisms a-d can account for $L_{br}$ $\\sim$ 10$^{36}$ to 10$^{40}$ erg s$^{-1}$, they cannot provide for luminosities that are 30 to 200 times greater. These very large luminosities are more likely associated with SN shocks or AGN. \\citet{I07} have considered \\hbox{type\\,IIn} SNe because their H$\\alpha$ luminosities are larger ($\\sim$10$^{38}$--10$^{41}$ erg~s$^{-1}$) than those of the other SN types, IIp and IIl, and they decrease less rapidly.  To decide whether type IIn SNe or AGN are responsible for the broad emission in these galaxies, monitoring of their spectral features on the relatively long time scale of several years is necessary. If broad features are produced by IIn type SNe, then we would expect a decrease in the broad line luminosities. No significant temporal evolution would be expected in the case of an AGN. Additionally, higher signal-to-noise ratio spectra are necessary to put better constraints on the presence of the high-ionization [Ne {\\sc v}] $\\lambda$3426 and He {\\sc ii} $\\lambda$4686 emission lines, good indicators of a source of hard non-thermal radiation. In order to check for temporal evolution, we have obtained second-epoch spectra of the above four galaxies with broad emission, using the 3.5m Apache Point Observatory (APO) telescope\\footnote{In fact, \\citet{I07} have identified 5 objects as AGN candidates. We have not included here the fifth candidate, J2230--0006$\\equiv$PHL 293B, because its broad H$\\alpha$ luminosity is only 8.6$\\times$10$^{37}$ erg s$^{-1}$, some 10$^3$ - 10$^4$ times lower than those of the other four AGN candidates. The broad H$\\alpha$ luminosity of J2230--0006 has been erroneously given in Table 8 of \\citet{I07} as having 10 times its true value. Given this relatively low luminosity and the fact that a new 3.5m APO spectrum of J2230--0006 shows that its broad hydrogen lines have a P Cygni profile with a blue-shifted absorption, the broad emission probably originates from a stellar wind rather than from an accretion disk around a AGN. Most likely, the broad emission in J2230--0006 is caused by a strong outburst in a bright luminous blue variable (LBV) star. Similar broad hydrogen emission has been detected recently by \\citet{P08} in the extremely metal-deficient dwarf galaxy DDO 68.}. We describe the observations in \\S2. We discuss in \\S3 the main properties of the broad emission and show that they can be accounted for by low-metallicity intermediate-mass AGN. Our conclusions are summarized in \\S4. ",
        "conclusions": "  1. The steep Balmer decrements of the broad hydrogen lines and the very high luminosities of the broad H$\\alpha$ line in all four galaxies (3$\\times$10$^{41}$ to 2$\\times$10$^{42}$ erg s$^{-1}$) suggest that the broad emission arises from very dense and high luminosity regions such as those associated with supernovae (SNe) of type IIn or with accretion disks around black holes. However, the relative constancy of the broad H$\\alpha$ luminosities over a period of 3--7 years likely rules out the SN mechanism. Thus, the emission of broad hydrogen lines is most likely associated with accretion disks around black holes. If so, these four objects would harbor a new class of AGN that are extremely rare (in 0.0006\\% of all galaxies). These AGN would be intermediate-mass black holes residing in low-metallicity dwarf galaxies, with an oxygen abundance that is considerably lower than the solar or super-solar metallicity of a typical AGN.  2. There is no obvious spectroscopic evidence for the presence of a source of a non-thermal hard ionizing radiation in all four galaxies: high-ionization emission lines such as He {\\sc ii} $\\lambda$4686 and [Ne {\\sc v}] $\\lambda$3426 emission lines were not detected at the level $\\leq$ 1 -- 2 percent of the H$\\beta$ flux. We have calculated a series of CLOUDY models with ionizing spectra which include both thermal stellar and nonthermal power-law ionizing radiation in order to account for the absence of the high-ionization lines. We find that the predicted fluxes of the high-ionization lines are below the detectability level if the spectral energy distribution $f_\\nu$ $\\propto$ $\\nu^\\alpha$ of the ionizing nonthermal radiation has $\\alpha$ $\\sim$ --1 and the nonthermal ionizing radiation is significantly diluted by the thermal stellar ionizing radiation contributing $\\la$ 3 percent of the total ionizing radiation, or the ionizing spectrum is steeper ($\\alpha$ $\\sim$ --2), and the nonthermal ionizing radiation contributes $\\la$ 10 percent of the total ionizing radiation.  3. The lower limits of the masses of the central black holes $M_{\\rm BH}$ of $\\sim$ 5$\\times$10$^5$ $M_\\odot$ -- 3$\\times$10$^6$ $M_\\odot$ in our galaxies are among the lowest found thus far for AGN."
    },
    "0807/0807.2503_arXiv.txt": {
        "abstract": "Several $f(R)$ modified gravity models have been proposed which realize the correct cosmological evolution and satisfy solar system and laboratory tests. Although nonrelativistic stellar configurations can be constructed, we argue that relativistic stars cannot be present in such $f(R)$ theories. This problem appears due to the dynamics of the effective scalar degree of freedom in the strong gravity regime. Our claim thus raises doubts on the viability of $f(R)$ models. ",
        "introduction": "The current accelerated expansion of the Universe is one of the deepest mysteries in cosmology~\\cite{Acceleration}. This acceleration may be due to some unknown energy-momentum component having the equation of state $p/\\rho \\approx -1$, or may be due to a modification of general relativity. In this paper, we are interested in the latter possibility. The simplest phenomenological way of modifying gravity is to consider a gravitational action described by a function of the Ricci scalar, $f(R)$, instead of the Einstein-Hilbert action. An early attempt is found, e.g., in~\\cite{StInf}, where a modification like $f(R)= R+R^2/\\mu^2$ was used to explain the accelerated expansion in the {\\em early} Universe. More recently, $f(R)$ modified gravity theories are often considered as a possible origin of the {\\em current} acceleration of the Universe~\\cite{Rev}.  Any modified theories of gravity must account for the late-time cosmology which is well established by observations, and at the same time must be consistent with solar system and laboratory tests of gravity. However, since $f(R)$ gravity has the equivalent description in terms of the Brans-Dicke theory with the Brans-Dicke parameter $\\omega=0$~\\cite{BD}, naively constructed models would result in violation of the above requirements~\\cite{Chiba, ESK, CSE}. For example, the original proposal of~\\cite{CDTT} employs $f(R) = R-\\mu^4/R$, which admits an acceleratedly expanding solution even in the absence of a dark energy component. In order to accommodate a late time acceleration, however, one must introduce a very low energy scale, $\\mu\\sim H_0$ (the present Hubble scale), leading to a very light scalar field, which predicts the parametrized post-Newtonian (PPN) parameter $\\gamma=(1+\\omega)/(2+\\omega)=1/2$. Obviously, this result contradicts the observational constraint $|\\gamma-1|\\lesssim 10^{-4}$~\\cite{Will}   To circumvent this difficulty, it is important to notice that the presence of matter may affect the dynamics of the extra scalar degree of freedom. The key idea is essentially the same as that of the ``chameleon'' model~\\cite{Chameleon, Cham2, Cham3}, in which the effective mass of the scalar field depends on the local matter density. In particular, the scalar field is very light for the cosmological density and is heavy for the solar system density, though the actual mechanism is slightly more complicated. The most successful class of $f(R)$ models~\\cite{Cembranos:2005fi, Hall, NvA, Tegmark, Li, AT, St, Hu, AB} incorporates this chameleon mechanism to evade local gravity tests. The experimental and observational consequences of this kind of $f(R)$ models are found in Refs.~\\cite{Tsujikawa1, Tsujikawa2, Brax}. (See also~\\cite{NO}.)   In this paper, we consider the strong gravity regime of the carefully constructed models of~\\cite{St, Hu, AB}. The strong gravity aspects of $f(R)$ theories have not been explored so much before. Recently, Frolov suggested that such $f(R)$ models generically suffer from the problem of curvature singularities which can be easily accessed by the field dynamics {\\em in the presence of matter}~\\cite{Frolov}. In other words, a curvature singularity may be caused not by diverging gravitational potential depth, $|\\Phi|=\\infty$, but rather by a slightly large gravitational field, $|\\Phi|\\lesssim 1/2$. This motivates us to study relativistic stars in the context of $f(R)$ gravity. Spherically symmetric stars in $f(R)$ gravity have been investigated so far in~\\cite{Stars1, Stars2, MV, BB}. (We confine ourselves to a {\\em metric} theory of $f(R)$ gravity. Using the {\\em Palatini} formalism, polytropic stars have been studied  in~\\cite{Sot}.) We shall show, both analytically and numerically, that stellar solutions with relatively strong gravitational fields cannot be constructed. Using the specific example of relativistic stars, we clarify how the singularity problem arises in the strong gravity regime of $f(R)$ theories. The singularity problem was also identified in~\\cite{ABCos} in a cosmological setting.   This paper is organized as follows. In the next section we derive equations of motion for $f(R)$ modified gravity, and define the specific model we study. In Sec.~\\ref{sec:basic}, we reinterpret the problem of finding the desired stellar configuration as the problem of the particle motion in classical mechanics. We give some analytic arguments in Sec.~\\ref{sec:An} and then we present our numerical results in Sec.~\\ref{sec:Num}. Finally, we draw our conclusions in Sec.~\\ref{sec:Conc}. ",
        "conclusions": "\\label{sec:Conc}   In this paper, we have studied the strong gravity aspect of $f(R)$ modified gravity models that reproduce the conventional cosmological evolution and evade solar system and laboratory tests~\\cite{St, Hu, AB}. It is known that $f(R)$ theories can be recasted simply in the Brans-Dicke theory with $\\omega=0$, but the potential for the effective scalar degree of freedom may play a complicated and nontrivial role. Moreover, the presence of matter may affect dynamics of the scalar field, possibly mimicking the chameleon model~\\cite{Chameleon}.    We have explored uniform density, spherically symmetric stars and their exterior geometry in the $f(R)$ model of~\\cite{St}. The main result of the present paper is summarized as follows: given model parameters, there is a maximum value of the gravitational potential produced by a star, above which no asymptotically de Sitter stellar configurations can be constructed. We show this both analytically and numerically. For example, the model with $n=1$ and $\\lambda\\approx 2$ gives $\\Phi_{{\\rm max}}\\approx 0.1$. This raises a warning sign for a class of $f(R)$ theories, because neutron stars cannot be present in such gravity models.   The underlying mechanism that hinders strong gravitational fields around matter is explained essentially as follows~\\cite{Frolov}. Consider a static matter distribution. The Newtonian potential obeys the Poisson equation $\\nabla^2\\Phi \\sim G \\rho$, while the equation of motion for the scalar field implies $\\nabla^2\\chi \\sim G\\rho$. From this, one can evaluate the excitation of the scalar degree of freedom around the matter distribution as $\\delta\\chi\\sim{\\cal O}(\\Phi)$. If the de Sitter minimum is located very close to the point $\\chi=1$, which corresponds to $R=\\infty$ in the effective potential, a slightly strong gravitational field will cause the problem of appearance of a curvature singularity.   Bearing the above evaluation in mind, let us comment on the other specific models of $f(R)$ gravity. The model of Hu and Sawicki~\\cite{Hu} and Starobinsky's one share the same structure of $f(R)$ in the high-curvature regime, i.e., $f(R)\\approx R-2\\Lambda_{{\\rm eff}}+C/R^{\\alpha}$ with $\\alpha>0$. Therefore, we expect that the same problem arises in the Hu and Sawicki's model. The model by Appleby and Battye is characterized by~\\cite{AB} \\begin{eqnarray} f(R)=\\frac{R}{2}+\\frac{1}{2a}\\ln\\left[ \\cosh(aR)-\\tanh(b)\\sinh(aR) \\right], \\end{eqnarray} where $a$ and $b$ are parameters. Since $\\chi = df/dR = [1+\\tanh(aR-b)]/2$, a positive curvature singularity corresponds to $\\chi=1$. Taking, for example, $b=1.5$, we find that $\\chi_1\\approx 0.93$ at the de Sitter minimum, which is very close to the dangerous curvature singularity (the result is independent of $a$). Since also in this model the effective potential is finite at $R=+\\infty$, we anticipate the same singularity problem.   Our choice of the parameters in the present paper gave $\\Phi_{{\\rm max}}\\approx 0.1$, for which neutron stars are unlikely to exist. However, there still remains a possibility that more realistic stellar environments and matter profiles weaken the bound on the potential by a factor of 2 or 3, and at the same time make the chameleon mechanism work~\\footnote{ The neutron star-white dwarf system PSR J1141--6545 can put strong constrains on alternative theories of gravity around relativistic stars~\\cite{Will, E-F}. For this reason, we need to invoke the chameleon mechanism to describe such stars.}. It is technically much more difficult to construct stellar configurations with a realistic equation of state, realistic energy densities, and realistic stellar environments. Such an elaborated modeling of relativistic stars might allow for $\\Phi_{{\\rm max}}$ as large as, say, 0.3, but the parameter space of the theory will be very restricted. To conclude, $f(R)$ theories that reproduce the correct behavior of weak gravity in the solar vicinity do not admit neutron star solutions without special care."
    },
    "0807/0807.2818_arXiv.txt": {
        "abstract": "We performed, for the first time, the simulation of spiral-in of a star cluster formed close to the Galactic center (GC) using a fully self-consistent $N$-body model. In our model, the central super-massive black hole (SMBH) is surrounded by stars and the star cluster. Not only are the orbits of stars and the cluster stars integrated self-consistently, but the stellar evolution, collisions and merging of the cluster stars are also included. We found that an intermediate-mass black hole (IMBH) is formed in the star cluster and stars escaped from the cluster are captured into a 1:1 mean motion resonance with the IMBH.  These ``Trojan'' stars are brought close to the SMBH by the IMBH, which spirals into the GC due to the dynamical friction.  Our results show that, once the IMBH is formed, it brings the massive stars to the vicinity of the central SMBH even after the star cluster itself is disrupted.  Stars carried by the IMBH form a disk similar to the observed disks and the core of the cluster including the IMBH has properties similar to those of IRS13E, which is a compact assembly of several young stars. ",
        "introduction": "Young and massive stars have been found within one parsec from the Galactic center (GC) \\citep{Krabbe95,Paumard06}. Some of them are $\\sim$1000 AU from the central SMBH. How these massive stars were brought to the vicinity of the SMBH has been a mystery. One possible scenario is the following.  A star cluster formed at a few tens of pc from the GC, and then spiraled in due to the dynamical friction \\citep{Gerhard01}. Previous simulations \\citep{PZ03,KM03,Kim04,GR05} have shown that the timescale of spiral-in of the star cluster can be short enough. However, how close the stars can actually approach the SMBH is not clear.  Another possible scenario is the in-situ formation in an accretion disk \\citep{LB03,Nay07,HN08}. Giant molecular clouds fall into the GC and form massive gaseous disks around the central BH. Stars form in the disk if it becomes gravitationally unstable and results in fragmentation. However, accretion disks have difficulty producing stars with eccentric orbits and a compact assembly of stars like IRS 13E, which is located at $\\sim 0.13$ pc in projection from the GC and contains half a dozen young stars and probably an IMBH \\citep{HM03,Maillard04}. \\citet{Mapelli08} argued that if a gas cloud undergoes a very close encounter (pericenter distance of 0.01 pc) with the central SMBH, the tidal compression could trigger the star formation, resulting in stars in close, bound orbits. However, how a cloud can come that close to the GC is not clear. On the other hand, if the star cluster has an eccentric orbit, the orbits of stars escaped from the cluster are also eccentric. The remnant of the core looks like IRS 13E.  We performed a fully self-consistent $N$-body simulation in which the internal dynamics of the cluster, that of the parent galaxy, and interactions between cluster stars and galaxy stars are correctly handled. In previous simulations, when the internal dynamics was followed by an accurate $N$-body code, the parent galaxy had to be modeled as a fixed potential with some fitting formulae for the dynamical friction. This means that the orbital evolution is not accurate. \\citet{Fujii08} showed that the actual orbital decay of the cluster is faster than that of previous simulations and the main reason is stars with a mass grater than 90\\% of an initial cluster escaping from this cluster.   We describe the method of our $N$-body simulation in section 2. In section 3 we show the results of simulations. Section 4 is for summary. ",
        "conclusions": "Using $N$-body simulation, we showed that many young and massive stars are carried to the GC by a star cluster due to the 1:1 mean motion resonance with an IMBH which is formed in the cluster.  In addition, we found that slingshots in the star cluster throw stars into orbits which pass near the GC.  These orbits have very high inclinations and are sometimes retrograde orbits.  They are new channels which carry young stars to the central parsec. Our simulation demonstrated the existence of massive stars and we explained why they form a disk-like structure.  The possible existence of two counter-rotating disks might suggest that two clusters have spiraled in.  The interaction and resonance between stars and IMBHs from multiple clusters might be responsible for the existence of stars which are very close to the central SMBH."
    },
    "0807/0807.0730_arXiv.txt": {
        "abstract": "The observational characteristics of quasi-periodic oscillations (QPOs) from accreting neutron stars strongly indicate the oscillatory modes in the innermost regions of accretion disks as a likely source of the QPOs. The inner regions of accretion disks around neutron stars can harbor very high frequency modes related to the radial epicyclic frequency $\\kappa $. The degeneracy of $\\kappa $ with the orbital frequency $\\Omega $ is removed in a non-Keplerian boundary or transition zone near the magnetopause between the disk and the compact object. We show, by analyzing the global hydrodynamic modes of long wavelength in the boundary layers of viscous accretion disks, that the fastest growing mode frequencies are associated with frequency bands around $\\kappa $ and $\\kappa \\pm \\Omega $. The maximum growth rates are achieved near the radius where the orbital frequency $\\Omega $ is maximum. The global hydrodynamic parameters such as the surface density profile and the radial drift velocity determine which modes of free oscillations will grow at a given particular radius in the boundary layer. In accordance with the peak separation between kHz QPOs observed in neutron-star sources, the difference frequency between two consecutive bands of the fastest growing modes is always related to the spin frequency of the neutron star. This is a natural outcome of the boundary condition imposed by the rotating magnetosphere on the boundary region of the inner disk. ",
        "introduction": "}  Wave modes in the boundary region in the inner accretion disk are a likely source of the distinct narrow frequency bands of quasi-periodic oscillations (QPOs) (van der Klis 2000) from neutron star sources in low mass X-ray binaries (LMXBs). The dependence of QPO frequencies on accretion rate suggests that the observed QPOs are connected with accretion disk modes. Psaltis, Belloni \\& van der Klis (1999) showed the existence of correlations between the frequencies of different QPO bands extending over a wide span of frequencies. The same correlation encompasses black hole as well as white dwarf and neutron star sources. This strongly implies that the frequency bands are determined by the oscillation modes of the accretion disk. The nature of the compact object, whether it is a black hole, white dwarf or neutron star, probably plays a role in exciting the same disk modes through possibly different mechanisms in different types of compact sources. The QPO bands modifying the X-ray luminosity are likely to belong to the boundary region between the disk and the compact object. For magnetic neutron stars, this boundary region is shaped by the interaction of the disk with the magnetosphere. The region where the rotation deviates from Keplerian flow is not necessarily very narrow; this boundary zone may have a size as large as a few tenths of the inner disk radius. We shall therefore use the terms \\textquotedblleft boundary\\textquotedblright \\thinspace and \\textquotedblleft transition\\textquotedblright \\thinspace region interchangeably.  Initial explorations of disk modes underlying QPO frequency bands from LMXBs were provided by Alpar et al. (1992) and Alpar \\& Y\\i lmaz (1997). The characteristics and possible excitation mechanisms of thin-disk oscillations were reviewed and discussed by Kato (2001). In many current models of QPOs, especially those of neutron stars, their frequencies are identified with test-particle frequencies (Stella, Vietri \\& Morsink 1999; Abramowicz et al. 2003). Most parts of an accretion disk, except, significantly, the transition regions at the inner boundary of the disk, are characterized by Keplerian rotation rates to a very good approximation. At any radius beyond the transition region, acoustic, magneto-acoustic and viscous corrections to the test-particle Keplerian orbital frequency are negligible. In the outer disk regime for disks around neutron stars and white dwarfs (and far enough from a central black hole), test-particle frequencies for radial and vertical perturbations of the orbit are degenerate with the Keplerian orbital frequency. The degeneracy is removed if the effects of general relativity are important. A Newtonian field of tidal or higher multi-pole structure would also lead to a split in the degeneracy. However, in Newtonian gravity, distortions of stellar shape even for the most rapidly rotating neutron stars will not introduce a significant level of non-degeneracy between the frequencies of test-particle oscillations that is comparable to the difference between the observed QPO frequency bands. Models employing the Keplerian frequency have been relatively successful in interpreting QPO frequency correlations. The Kepler (test particle) frequency interpretation has also been employed to place constraints on the masses and spins of compact objects. Empirically, predictions based on test-particle frequencies applied to QPO frequency correlations deviate from observations at a level of about 10\\%.  The Keplerian frequency represents the basic imprint of rotation on all the dynamical responses of the disk. Albeit simple, the identification of QPOs with test-particle frequencies has a number of shortcomings. First, precisely in a boundary layer where QPOs are expected to be excited, the dynamical frequencies of oscillations are not the test-particle frequencies. In the boundary layer, viscous and magnetic forces lead to deviation of orbital frequencies from Keplerian test-particle frequencies (see, e.g., Erkut \\& Alpar 2004). The resulting band of non-Keplerian rotation frequencies entail viscous and acoustic response and couplings between adjacent rings of the disk fluid. Hydrodynamic effects are therefore essential for an understanding of the frequency bands characterizing the boundary layer. In particular, the degeneracy of test-particle frequencies in the non-relativistic regime is lifted in the hydrodynamic boundary layer, where the radial epicyclic frequency is in fact larger than the orbital frequency. This simple observation (Alpar \\& Psaltis 2005) has important consequences for the interpretation of kHz QPO frequencies, in particular for the constraints on the neutron star mass-radius relation derived from kHz QPOs.  Second, test-particle frequencies do not distinguish between azimuthal sidebands. While kinematic models of QPOs like the beat-frequency model involve one specific band of frequencies, say $\\omega $, a large number of sidebands, of frequencies $\\omega _{m}\\cong \\omega -m\\Omega $, where $\\Omega $ is the orbital frequency, are allowed by azimuthal symmetry. Arguments as to why only one or two QPO bands are excited must involve choices imposed by the symmetries of the interaction between the magnetosphere and the accretion disk boundary, leading to resonances with particular frequencies at one or more radial regions in the disk. Without resonances, test-particle frequency spectra present no distinction between the modes. The realistic hydrodynamic modes may, in some parameter ranges, distinguish between both fundamental modes and all their azimuthal sidebands through the different growth or decay rates of these modes, \\emph{even in the case of free oscillations}. A reduction of the number of relevant (easily excitable) modes of free oscillations is certainly an important task for an understanding of accretion disks around neutron stars, white dwarfs and black holes. In the case of black holes, this classification of free hydrodynamic modes in terms of their growth rates is even more important as resonant excitation by the black hole is not available.  In this paper, we study the global modes of free oscillations at some position $r$ in the inner disk-boundary or transition region by analyzing the perturbed dynamical equations of a hydrodynamic disk, including pressure gradients, viscous and magnetic stresses. We consider the case of a disk around a neutron star with a magnetosphere. The neutron star is taken to be a \\emph{slow rotator}: the star's rotation rate is less than the value of the Keplerian rotation rate at the inner radius of the disk; $\\Omega _{\\ast }<\\Omega _{\\mathrm{K}}(r_{\\mathrm{in}})$. The actual rotation rate of the disk is set by the boundary condition $\\Omega (r_{\\mathrm{in}})=\\Omega _{\\ast }$. Thus in the steady state solutions for the disk, the rotation rate at the inner edge of the disk and in the transition region beyond is sub-Keplerian. The innermost disk radius $r_{\\mathrm{in}}$ determines the disk magnetosphere interface which may be subject to magnetic Rayleigh-Taylor or interchange instabilities. This instability operates mainly at the disk-magnetosphere interface because of the sharp density contrast across the radius $r=r_{\\mathrm{in}}$. The inner disk region for $r\\gtrsim r_{\\mathrm{% in}}$ consists of a non-Keplerian boundary region that joins the outer Keplerian disk with a continuous density distribution. As we will see in \\S\\ \\ref{ehdfom}, the surface density profile throughout the non-Keplerian boundary region increases with decreasing radii. Thus, our boundary region is not subject to interchange-like instabilities.  The free oscillation modes we explore by perturbing a steady state solution embody information about the stability of the disk through their growth or decay rates. Growing modes will provide a clue as to the origin of the QPO frequency bands, which could be excited by resonant forcing of the disk by time dependent interactions with the star's magnetosphere. We find, as expected, that growth or decay rates are determined by the dynamical effects of viscosity, with an additional dependence on the sound speed and the density and rotation-rate profiles in the non-Keplerian boundary region. For some boundary region parameters, not all fundamental modes of free oscillations and not all azimuthal sidebands grow, and among the growing modes the growth rates differ. For other choices of boundary region parameters the growth rates of all sidebands are similar. The frequencies of the modes differ from test-particle frequencies by amounts that can be several times larger than the corresponding QPO width. Most importantly, for a reasonable set of accretion disk parameters, we can show that only a certain few of the hydrodynamic free oscillation modes will grow, and that the frequency bands of these oscillations can be associated with the observed frequency bands. The nature of free oscillation modes in the boundary region depends on whether the neutron star is rotating at a rate slower or faster than the rotation rates prevailing in the inner boundary of the disk. In this paper, we will consider the more common and straightforward case of \\emph{slow rotators}, the case when the neutron star rotation rate $\\Omega _{\\ast }$ is less than $\\Omega _{\\mathrm{K}}(r_{% \\mathrm{in}})$, the Kepler rotation rate at some representative inner disk radius $r_{\\mathrm{in}}$. An investigation of \\emph{forced} resonant excitation of these prevalent modes, and the comparison and association with observed QPO bands will follow in a subsequent paper.  \\S\\ 2 lays out the basic assumptions and equations, \\S\\ 3 displays the mode analysis for the free oscillations, including a discussion of hydrodynamic effects, frequency bands, and growth rates. In \\S\\ 4, we discuss the results and present our conclusions. ",
        "conclusions": "We have studied the global hydrodynamic modes of long wavelength in the boundary or transition region of viscous accretion disks as a possible source of kHz QPOs in neutron-star LMXBs. The stability of the eigenmodes strongly depends on the global disk structure imposed by the boundary conditions as expected from the mode analysis in the long wavelength regime. Our local treatment takes account of the local effects of the global disk parameters on the excitation of hydrodynamic free oscillations.  We find that the frequencies and growth rates of the modes are mainly determined by global disk parameters such as $\\kappa /\\Omega $, the ratio of the radial epicyclic frequency to the orbital frequency, the radial surface-density profile $\\beta $, $\\Omega _{\\nu }/\\Omega _{s}$, the ratio of the radial drift velocity to the sound speed, and $\\Omega _{s}/\\Omega $, the ratio of the sound speed to the azimuthal velocity. The local values of $% \\kappa /\\Omega $ and $\\beta $ directly follow from the global solution for the rotational dynamics of a boundary-transition region model. The parameters $\\Omega _{\\nu }/\\Omega _{s}$ and $\\Omega _{s}/\\Omega $ depend additionally on the particular viscosity prescription and the details of the structure of accretion flow in the inner disk.  The hydrodynamic effects on the frequencies and growth rates of the modes are reflected by the parameters $\\Omega _{\\nu }/\\Omega _{s}$ and $\\Omega _{s}/\\Omega $. We have found that the growth rates of different modes for a given azimuthal wavenumber $m$ differ significantly when the hydrodynamic corrections are sufficiently large, i.e. $\\Omega _{\\nu}/\\Omega _{s}\\lesssim 1 $ and $\\Omega _{s}/\\Omega \\lesssim 1$. In this regime, the growing mode frequencies significantly exceed test-particle frequencies for a plausible range of $\\Omega _{s}/\\Omega $ at any particular radius within the boundary region. This is because the effect of hydrodynamic corrections on the growth rates of the modes is expected to be strong in a magnetic boundary layer or transition zone as discussed in \\S\\ 3.2. In the limit of small hydrodynamic corrections, i.e. $\\Omega _{\\nu}/\\Omega _{s}\\ll 1$ and $\\Omega _{s}/\\Omega \\ll 1$, the eigenmodes have test-particle frequencies with negligible growth rates. Our analysis shows that taking account of hydrodynamic effects has an important outcome: the frequencies of the growing modes gain higher values above test-particle frequencies as their growth rates increase for larger hydrodynamic corrections.  Throughout the boundary region, we have identified the growing mode frequencies by the eigenvalues that approach the test-particle frequencies $% \\omega _{1}^{(0)} =\\kappa $, $\\omega _{1}^{(1)} =\\kappa +\\Omega $, and $% \\omega _{2}^{(1)} =\\kappa -\\Omega $ in the limit of small hydrodynamic corrections (see panel a of Fig.~2). Modes such as $\\omega _{3}^{(0)} =0$ and $\\omega _{3}^{(1)}\\simeq \\Omega $ are not excited for the range of $% \\Omega _{s}/\\Omega $ used in Fig.~2 (panel b) and Fig.~3 (panel b). The difference frequency between successive bands of growing modes is $\\Delta \\omega \\simeq \\Omega _{*}$ as observed for the relatively slowly rotating neutron stars in LMXBs that show kHz QPOs (\\S\\ 3.2). In a boundary layer or transition region, the orbital frequency of the disk matter is close to the rotation frequency of the star ($\\Omega \\sim \\Omega _{*}$). As $\\Delta \\omega \\simeq \\Omega$, the separation between consecutive mode frequencies is always related to the spin frequency of the neutron star. This is the result of the boundary condition imposed by the rotating magnetosphere on the innermost disk. The modes of the test-particle frequencies $\\omega _{1}^{(0)}\\cong \\kappa $ and $\\omega _{2}^{(1)}\\cong \\kappa -\\Omega $ were originally proposed by Alpar \\& Psaltis (2005) to be associated with the upper and lower kHz QPOs observed in neutron-star LMXBs. The frequency separation between the higher and lower-frequency kHz QPO peaks decreases by a few tens of Hz as both QPO frequencies increase by hundreds of Hz in a range of $200-1200$ Hz per source in all observed sources, namely, Sco~X-1, 4U~$1608-52$, 4U~$1728-34$ and 4U~$1735-44$ (van der Klis 2000). This pair of frequencies satisfies the observed behavior that the difference frequency $\\Delta \\omega $ between two kHz QPOs decreases as both frequencies increase. According to our present analysis, the pairs of observed kHz QPO bands could be either $\\omega _{1}^{(1)}\\simeq \\kappa +\\Omega $ and $\\omega _{1}^{(0)}\\simeq \\kappa $ or $\\omega _{1}^{(0)}\\simeq \\kappa $ and $\\omega _{2}^{(1)}\\simeq \\kappa -\\Omega $, respectively, since the identification of the observed QPOs with $\\omega _{1}^{(1)}\\simeq \\kappa +\\Omega $ and $\\omega _{1}^{(0)}\\simeq \\kappa $ also meets the observed condition that $\\Delta \\omega $ decreases when both QPO frequencies increase. The growth rates of free oscillations produce a wide spectrum of sidebands, $\\omega _{1,2}^{(m)}\\cong \\kappa \\pm \\vert m\\vert \\Omega $, some of which fall in the range of observed power spectra. The reduction to only two of kHz QPO bands as observed in neutron star sources must therefore be a product of forced oscillations, resonances, and boundary conditions. The role of forced oscillations and resonances in the excitation of particular modes will be discussed in future work.  We find that the fastest growing modes with frequency branches $\\kappa $ and $\\kappa \\pm \\Omega $ are excited near the disk radius $r\\lesssim r_{0}$, where the orbital frequency is maximum. In the boundary layer, the loss of centrifugal support near $r_{0}$ leads to some radial acceleration of the disk matter. The subsequent rise in the radial drift velocity of the disk matter is accompanied by a sudden drop in the surface density as the vertically integrated dynamical viscosity is minimized (see \\S\\ 3.2). This picture is common to all boundary layers or transition zones for which magnetic braking is efficient (see Erkut \\& Alpar 2004 and references therein). Thus, independent of the boundary conditions and the magnetic field configuration in the boundary region, there exists a specific disk radius where the hydrodynamical background quantities such as the surface density $\\Sigma _{0}$ and the radial drift velocity $v_{r0}$ change dramatically. The steepness in the change of the surface density in the radial direction is reflected through the parameter $\\beta $ (see equation % \\ref{bet}). The sign and magnitude of $\\beta $ (the surface density profile in the disk) control whether the free oscillation modes will grow or decay. We see that the modes of frequency bands $\\kappa $ and $\\kappa \\pm \\Omega $ can be excited with the largest growth rates only for the disk radii $% r\\lesssim r_{0}$, where $\\beta $ is minimum and $\\Omega $ is maximum (see \\S\\ 3.2 for the range of $\\beta $). The inverse timescale, $\\Omega _{\\nu} $, associated with radial accretion also obtains its highest value near the same radii yielding the maximum growth rates in the accretion disk.  The boundary layers with sub-Keplerian rotation rates are usually expected to be realized in the innermost regions of accretion disks around \\emph{slow rotators}. For a \\emph{slow rotator}, the Kepler rotation rate at the inner disk radius $\\Omega _{\\mathrm{K}}(r_{\\mathrm{in}})$ prevails over the stellar rotation rate $\\Omega _{*}$. As we have illustrated in \\S\\ 3.2, the difference frequency between two consecutive bands of growing modes in the boundary region is close to the spin frequency of the neutron star ($% \\Delta \\omega \\simeq \\Omega _{*}$). This result agrees with the kHz QPO observations of the relatively slow rotators with spin frequencies below $% \\sim 400$~Hz in neutron-star LMXBs. In these sources, for example, in 4U~$1728-34$, the lower kHz QPO frequencies increase from 600 to 900 Hz and the upper kHz QPO frequencies increase from 950 to 1200 Hz, while the difference between the two observed kHz QPO frequency bands decreases slightly, of the order of 10 Hz, while remaining still commensurate with the spin frequency of the neutron star (M\\'{e}ndez \\& van der Klis 1999). A second class of kHz QPO sources, such as the sources KS~$1731-260$, Aql~X-1, 4U~$1636-53$, which have spin frequencies above $\\sim 400$~Hz is usually depicted as \\emph{fast rotators} (see Wijnands et al. 2003). In this second class of sources, for example in 4U~$1636-53$, the lower kHz QPO frequencies increase from 900 to 950 Hz and the upper kHz QPO frequencies increase from 1150 to 1190 Hz, while the difference between the two observed kHz QPO frequency bands again decreases by amounts of the order of 10 Hz (see van der Klis 2000 and references therein). The frequency separation between the kHz QPOs is close to half the spin frequency of the neutron star if it is a \\emph{fast rotator}. The emergence of two seemingly different classes of neutron-star LMXBs could be related to whether the rotation rate of the neutron star is less or greater than that of the inner disk matter. It is probable that $\\Omega _{*}>\\Omega _{\\mathrm{K}}(r_{\\mathrm{in}})$ for the relatively fast rotating neutron star sources for which the difference frequency between twin QPO peaks is $\\Delta \\omega \\simeq \\Omega _{*}/2$. Recent work by M\\'{e}ndez \\& Belloni (2007) actually suggests a more continuous distribution between $\\Delta \\omega \\simeq \\Omega _{*}$ and $% \\Delta \\omega \\simeq \\Omega _{*}/2$. The rotational dynamics of accretion flow in the innermost regions of accretion disks around these \\emph{fast rotators} could be quite different from those of the transition regions we consider here. We plan to address this issue in a subsequent paper."
    },
    "0807/0807.3418_arXiv.txt": {
        "abstract": "We explore CP violation effects on the neutrino propagation in dense environments, such as in core-collapse supernovae, where the neutrino self-interaction induces non-linear evolution equations. We demonstrate that the  electron (anti)neutrino fluxes are not sensitive to the CP violating phase if the muon and tau neutrinos interact similarly with matter. On the other hand, we numerically show that new features arise, because of the non-linearity and the flux dependence of the evolution equations, when the muon and tau neutrinos have different fluxes at the neutrinosphere (due to loop corrections or of physics beyond the Standard Model). In particular, the electron (anti)neutrino probabilities and fluxes depend upon the CP violating phase. We also discuss the CP effects induced by radiative corrections to the neutrino refractive index. ",
        "introduction": "One of the major open issues in neutrino physics is the possible existence of CP violation. As with the recent crucial experimental discoveries on neutrino oscillations, the answer to this question has fundamental implications in high-energy physics, astrophysics and cosmology e.g. to understand the matter versus anti-matter asymmetry in the Universe. The observation that the weak interaction violates the CP symmetry in the quark sector was first established in 1964 \\cite{Christenson:1964fg}. Future strategies to search for CP violation in the lepton sector depend upon the actual value of one yet unknown neutrino oscillation parameter, i.e. $\\theta_{13}$ \\cite{Ardellier:2006mn}, and require long term accelerator projects producing very intense neutrino beams \\cite{Volpe:2006in}. It is therefore essential to explore alternative avenues to get clues on this fundamental question, such as indirect effects in dense environments like core-collapse supernovae.  Core-collapse supernovae emit about 10$^{53}~$erg as neutrinos of all flavours during their rapid gravitational collapse. Such neutrinos might play a role on the two major supernova unsolved problems, namely understanding how the explosion finally occurs and where the nucleosynthesis of the heavy elements, produced during the {\\it r}-process, takes place. While neutrinos from a massive star were first observed during the SN1987A explosion, future observations of (extra)galactic or relic supernova neutrinos will help unravelling supernova physics and/or unknown neutrino properties. For example, in \\cite{Schirato:2002tg,Tomas:2004gr,Fogli:2004ff,Kneller:2007kg} the imprint of the shock wave on the neutrino time signal is investigated; while avenues for extracting information on the third neutrino mixing angle are discussed in \\cite{Dighe:1999bi,Engel:2002hg}. This searches require advances in the modelling of supernova dynamics, of neutrino propagation in dense environments and of our knowledge on neutrinos.  Impressive developments are currently ongoing in our understanding of neutrino propagation in dense matter. While solar experiments \\cite{Davis:1964aa,Ahmad:2001an,Eguchi:2002dm,Arpesella:2007xf}  have beautifully confirmed the oscillation enhancement induced by the coupling with matter - the Mikheev-Wolfenstein-Smirnov or MSW effect  \\cite{Wolfenstein:1977ue,Mikheev:1986wj} -, recent theoretical investigations have shown that the inclusion of the neutrino self-interactions in dense environments introducing a non-diagonal neutrino refractive index \\cite{Pantaleone:1992eq} gives rise to a wealth of new phenomena, as first pointed out in \\cite{Samuel:1993uw}. Various regimes have been identified : the synchronized one \\cite{Pastor:2001iu,Duan:2005cp}, the bipolar oscillations \\cite{Duan:2005cp,Hannestad:2006nj} and the spectral split phenomenon \\cite{Raffelt:2007cb,Raffelt:2007xt}. Since numerical calculations become more involved, analytical treatments for the three flavour case are being proposed (see e.g. \\cite{Dasgupta:2007ws}). The importance of the loop corrections to the neutrino refractive index, the $V_{\\mu \\tau}$ potential \\cite{Botella:1986wy}, is underlined in \\cite{EstebanPretel:2007yq}. Moreover constraints on neutrino mixing from shock re-heating and the {\\it r}-process nucleosynthesis, including the neutrino-neutrino interaction, are investigated in \\cite{Qian:1994wh,Sigl:1994hc}. The impact of the neutrino-neutrino interaction on the electron fraction relevant for the {\\it r}-process is investigated in \\cite{Pastor:2002we,Balantekin:2004ug}.  In a previous work \\cite{Balantekin:2007es} we have investigated the CP effects on the neutrino fluxes, on the electron fraction in a supernova (relevant for the {\\it r}-process) as well as the possible impact on the supernova neutrino signal in an observatory on Earth. In particular we have shown analytically that no effects can be found on the electron (anti)neutrino fluxes, when muon and tau neutrino have the same fluxes at the neutrinosphere; while significant effects are obtained numerically on the fluxes when they differ. The calculations in  \\cite{Balantekin:2007es} are obtained considering interaction with matter at tree level only. In this work we explore for the first time CP violation effects on the neutrino propagation in dense environments including the standard MSW effect, the neutrino self-interactions and the {\\it V}$_{\\mu \\tau}$ refractive index. We first show analytically that if the muon and tau neutrinos interact similarly, the electron (anti)neutrinos are not sensitive to the CP violating phase, even in presence of neutrino self-interactions. This result is general and valid for any matter density profile and/or initial neutrino luminosity. We present numerical results, obtained within the three flavour formalism, on the neutrino oscillation probabilities and fluxes within the star. In particular we show that both the probabilities and the fluxes become sensitive to the CP violating phase if muon and tau neutrino interact differently with matter (e.g. because of loop corrections or physics beyond the Standard Model.) The paper is structured as follows. Section II presents the theoretical framework for describing the neutrino propagation including the coupling with matter as well as the neutrino-neutrino interaction. Section III gives the analytical and numerical results. Conclusions are drawn in Section IV. ",
        "conclusions": "We have investigated possible effects of the CP violating phase on the neutrino propagation in dense matter when interaction with matter without/with loop corrections and the neutrino self-interaction are included. Our analytical results demonstrate that, at tree level, even when the neutrino-neutrino interaction is included there are no CP violating effects on the electron (anti)neutrino fluxes in the star unless $\\nu_{\\mu}$ and $\\nu_{\\tau}$ fluxes differ at the neutrinosphere. If such condition is not satisfied, a totally new feature arise, namely that the electron (anti)neutrino oscillation probabilities (and fluxes) become sensitive to the CP violating phase $\\delta$. The latter is also true when the loop corrections to the refractive index are included. We find numerically that, in most cases studied, the modifications introduced by the CP violating phase are larger (smaller) at low (high) energies than in the case where the neutrino-neutrino interaction is not included, and at the level of a few percent. We also find numerically that, even assuming that the muon and tau neutrinos have the same fluxes at the neutrinosphere, the CP effects induced by $V_{\\mu\\tau}$ only are amplified by the neutrino self-interactions up to several percent.  \\begin{figure}[t] \\vspace{.6cm} \\centerline{\\includegraphics[scale=0.3,angle=0]{nuefluxratioTeqVmutau.eps}\\hspace{.2cm} \\includegraphics[scale=0.3,angle=0]{nuefluxratioTdifVmutau.eps}} \\caption{Ratios of the $\\nu_e$ fluxes for a CP violating phase $\\delta=180^{\\circ}$ over $\\delta= 0^{\\circ}$ as a function of neutrino energy. They correspond to inverted hierarchy and small $\\theta_{13}$ and different distances from the neutron star surface, i.e. 200 km (dotted), 500 km (dashed), 750 (dot-dashed), 1000 (solid). The results include the $\\nu$-$\\nu$ interaction and the $V_{\\mu \\tau}$ refractive index. They are obtained using equal $\\nu_{\\mu}$ and $\\nu_{\\tau}$ fluxes at the neutrinosphere (left) or taking $T_{\\nu_{\\mu}}=1.05~T_{\\nu_{\\tau}}$ (right). For the $\\bar{\\nu}_e$ fluxes deviations up to 10$\\%$ are found at energies lower than 10 MeV.} \\label{fig:nueIHSA} \\end{figure}  \\vspace*{0.5cm} The authors acknowledge the support from \"Non standard neutrino  properties and their impact in astrophysics and cosmology\", Project No. ANR-05-JCJC-0023.  \\textit{} \\vspace*{-0.25cm}"
    },
    "0807/0807.1301_arXiv.txt": {
        "abstract": "{We analyze the data of Kamiokande-II, IMB, Baksan using a parameterized description of the antineutrino emission, that includes an initial phase of intense luminosity. The luminosity curve, the average energy of $\\bar\\nu_e$ and the astrophysical parameters of the model, derived by fitting the observed events (energies, times and angles) are in reasonable agreement with the generic expectations of the delayed scenario for the explosion.}  \\normalsize\\baselineskip=15pt ",
        "introduction": "We begin with a rapid historical excursus, with emphasis on the issues that are relevant for data analysis.  \\begin{itemize}  \\item Colgate \\& White 1966\\cite{cg} propose the paradigm for the explanation of the core collapse supernovae, where neutrinos are the key agents.  \\item Nadyozhin 1978\\cite{nad} concludes a detailed calculation that demonstrates an initial phase of intense neutrino luminosity.  \\item Bethe \\& Wilson 1985\\cite{bw} suggest that the energy deposition on a scale of half a second can re-energize the stalled shock wave.  \\item 1987: Kamiokande-II\\cite{kii}, IMB\\cite{imb}, Baksan\\cite{bak} and LSD\\cite{lsd} observe several events in correlation with SN1987A.  \\item Several authors--e.g., Bahcall 1989\\cite{bah}--remark that non-LSD data generically meet the expectations. Then the main interest shifts on the relevance of oscillations.  \\item Lamb \\& Loredo 2002\\cite{ll} (LL) discuss whether SN1987A data indicate specific imprints of the delayed scenario\\footnote{With the term `delayed scenario' we refer here and in the following to the scenario for the explosion put forward by Bethe and Wilson, that incorporates the initial phase of intense neutrino luminosity discussed by Nadyozhin (this is also called `standard scenario' or `neutrino assisted explosion').} such as the initial phase of intense luminosity.  \\item Imshennik~\\&~Ryazhskaya~2004\\cite{ir} suggest a 2 stage scenario with essential role of rotation and possible explanation of LSD events detected 4.5 hours earlier.  \\end{itemize} Next we summarize the present status: although there is mounting evidence that the delayed scenario is correct, a conclusive proof is still missing; while most people is convinced that SN1987A neutrinos confirmed the general picture of the explosion, there are annoying doubts on whether SN1987A was a standard object or not; while we continue sharpening our theoretical tools, we hope intensely in a Galactic supernova event to progress in our understanding.  In short, one should be aware that until the theoretical picture will be definitively assessed, the interpretation of the results from SN1987A will continue to contain elements of uncertainty. But if we want to test the expectations, we are forced to answer the question: {\\em What are the generic expectations for the delayed scenario?} We summarize the features that are relevant for our analysis: \\begin{itemize} \\item In conventional detectors (scintillators and water Cherenkov) the main detection reaction is the inverse-beta decay of electron antineutrinos on free protons (IBD): $\\bar\\nu_e p\\to n e^+ $.  \\begin{figure}[t] $$\\includegraphics[width=0.5\\textwidth]{vissaniFig1.eps}$$ \\caption{\\it Sketch of $\\nu_e$ and $\\bar\\nu_e$ (higher and lower) luminosity curves. The excess of $\\nu_e$ in the first 50 ms (`neutronization' phase) gives a small contribution to the total number of observable events. \\label{fig1}} \\end{figure}   \\item The main part of the emission happens in two stages as shown in Fig.~\\ref{fig1}. 10-20\\%  of the energy is radiated in an early phase, here called {\\bf accretion}, that should last about half-a-second; 80-90\\% of the energy is radiated later, during the phase of neutron star {\\bf cooling}, namely in a quiet thermal phase.   \\item The main reactions for energy radiation during accretion are $e^- p\\to n \\nu_e$ and $e^+   n\\to p \\bar\\nu_e$, the second being the inverse of IBD. The presence around the nascent neutron star of an abundant amount of $e^+e^-$ plasma ensures a large flux of  $\\nu_e$ and $\\bar\\nu_e$, that are key ingredients to revive the stalled shock wave. [See Janka\\cite{j} for a review.]  \\item During the cooling phase, neutrinos of all species ($\\nu_e,\\nu_\\mu,\\nu_\\tau,\\bar\\nu_e,\\bar\\nu_\\mu,\\bar\\nu_\\tau$) are radiated with similar luminosities; this feature is occasionally called `equipartition'.   \\end{itemize} In this talk, based on a work in collaboration with M.L.~Costantini and A.~Ianni\\cite{paglia}, we present an analysis of SN1987A data along these lines. We will compare two models for neutrino emission, namely the conventional one, when the neutrino emission is described by a `one-component' (cooling) model, and a `two-component' (cooling+accretion) model which resembles more closely the expectations of the delayed scenario. Preliminary results have been documented in\\cite{mosca,prep,ifae,brasile}; see also\\cite{jcap}, in particular Sect.~3.2 there. ",
        "conclusions": "Our study confirms earlier results with the 1-component model. The refined treatment of background, cross section, description of the scattering angle, inclusion of Baksan data, {\\em etc.}\\ do not lead to important changes.  We confirm the results of Lamb \\& {}Loredo in particular the very strong evidence for accretion, when their 2-component model is adopted.  We discussed an improved 2-component model, where the average energy and luminosity curves are constrained to be continuous, cooling follows accretion, oscillations (not very important {\\em a posteriori}) are included.   The best fit of $\\tau_a$, $M_a$, {\\em etc.}\\ are close to expectations; the binding energy $2.2\\times 10^{53}$ erg is lower than for the 1-component model; the evidence of accretion is not as strong as for the Lamb \\& Loredo model, but still important."
    },
    "0807/0807.3074_arXiv.txt": {
        "abstract": "{Fourier transform (or lag) correlators in radio interferometers can serve as an efficient means of synthesising spectral channels. However aliasing corrupts the edge channels so they usually have to be excluded from the data set. In systems with around 10 channels, the loss in sensitivity can be significant. In addition, the low level of residual aliasing in the remaining channels may cause systematic errors. Moreover, delay errors have been widely reported in implementations of broadband analogue correlators and simulations have shown that delay errors exasperate the effects of aliasing.} {We describe a software-based approach that suppresses aliasing by oversampling the cross-correlation function. This method can be applied to interferometers with individually-tracking antennas equipped with a discrete path compensator system. It is based on the well-known property of interferometers where the drift scan response is the Fourier transform of the source's band-limited spectrum.} {In this paper, we simulate a single baseline interferometer, both for a real and a complex correlator. \\revd{Fringe-rotation usually compensates for the phase of the fringes to bring the phase centre in line with the tracking centre. Instead, a modified fringe-rotation is applied. This enables an oversampled cross-correlation function to be reconstructed by gathering successive time samples.}} {Simulations show that the oversampling method can synthesise the cross-power spectrum while avoiding aliasing and works robustly in the presence of noise. An important side benefit is that it naturally accounts for delay errors in the correlator and the resulting spectral channels are regularly gridded} {} ",
        "introduction": "\\label{sec:intro}  The observing band of radio interferometers frequently needs to be split into sub-bands, either for spectral line observations or to reduce the effects of chromatic aberration. Fourier transform correlators offer an efficient method for dividing the observation band. In this scheme, the signal from the two arms of the interferometer are correlated at discrete delay steps, making direct measurements of the cross-correlation function (see Fig.~\\ref{fig:ccf-spec}). Taking the Fourier transform of the cross-correlation function gives the complex cross-power spectrum. For a signal of bandwidth $\\Delta\\nu$, Nyquist sampling theorem requires the signal to be sampled at \\revd{time intervals of $1/(2\\Delta\\nu)$} to avoid aliasing. But the cross-correlation function of a band-limited signal extends over an infinite delay range. So the Nyquist sampling theorem holds true only if we sample the cross-correlation function over an infinite delay range. Clearly this is not practical so the cross-correlation function is sampled over a finite range. This results in a recovered signal spectrum with tapered band edges which will overlap with its images in the spectral domain. This overlap causes aliasing and will corrupt the recovered cross-power spectrum.  \\begin{figure}[hbtp] \\centering \\mbox{\\subfigure[Real correlator]{ \\includegraphics[width=3.5in,trim=0 0 0 65,clip=true]{pic/8117fig1a.eps} }} \\mbox{\\subfigure[Complex correlator]{ \\includegraphics[width=3.5in]{pic/8117fig1b.eps} }} \\caption[The real and complex CCF]{{\\bf a)} {\\it Top plot:} The cross-correlation function measured by a real correlator (black) and the envelope (grey). This is a graphical representation of the cross-correlation function in Eq.~(\\ref{eqn:corr-response}) with the geometrical delay $\\taug = 0$. The open circles indicate the position of the detectors in instrument delay $\\taui$. {\\it Bottom plots:} The cross-power spectrum recovered by applying the Fourier transform on the cross-correlation function. We have assumed the telescope parameters in Table~\\ref{tab:tel-param}. {\\bf b)} The in-phase (real) and quadrature (imaginary) cross-correlation functions sampled by a complex correlator. The cross-power spectrum recovered in this case is single-sided.}\\label{fig:ccf-spec} \\end{figure}   \\revd{As an illustration, the plots in Fig.~\\ref{fig:under-vs-oversamp} show snapshots of the cross-correlation functions and the recovered spectra.} The left-most plots \\revd{(case~1)} show the cross-correlation function when the source transits an east-west baseline (that is when the path difference between the two arms of the interferometer is zero). The circles represent the measurements at discrete delay steps of the correlator. The next plots to the right \\revd{(case~2)} show the cross-correlation function a short while later. The right-most plots show the spectrum calculated by taking the Fourier transform of the discrete measurements. The filled grey circles indicate the spectrum at transit \\revd{(case~1)} and the dark circles are for \\revd{case~2}. In the top row of Fig.~\\ref{fig:under-vs-oversamp}a, the cross-correlation function was \\revd{sampled at 16 delays. This critically samples the {$\\Dnu=6$\\GHz} signal bandwidth but undersamples the {0--12\\GHz} basebandsignal,} so the positive and negative halves of the {6--12\\GHz} bands lie side-by-side. In \\revd{case~1}, the cross-correlation measurements trace out a delta function and give the characteristic flat spectrum. But in \\revd{case~2}, the amplitudes are perturbed. Ideally, the amplitudes should \\revd{not change} so this points to a fundamental problem. This could be worked round by oversampling the \\revd{baseband signal} at 64 or more delays as illustrated by Fig.~\\ref{fig:under-vs-oversamp}b. Now the two sets of amplitudes at different times match perfectly. Oversampling introduces a buffer between the two halves of the passband and also between their spectral images. This suppresses aliasing by reducing the overlaps between the signal and image bands. But clearly, sampling at 64 delays in hardware is not practical.    \\revd{This is manifested in the Fourier Transformed data (the spectrum) as temporal modulations in both the amplitude and phase (see Fig.~\\ref{fig:alias-cycles}).} In aliased signals, the noise components between the channels will also be correlated so the individual channels cannot be strictly treated as independent measurements. These effects are particularly pronounced in the edge frequency channels as seen in Fig.~\\ref{fig:alias-cycles} so these channels are usually rejected. This may be acceptable in systems with tens to hundreds of channels. But some correlators, particularly broadband analogue Fourier transform correlators with channels of order 10 (for example, \\citealt[][ hereafter {\\paperone}]{li2004,roberts2007,holler2007}), the loss constitutes a significant portion of the total bandwidth.     In principle, there are three ways to suppress aliasing: \\begin{enumerate} \\item Increasing the delay range over which the cross-correlation function is sampled. The spectral channel width will be narrower and the edges of the passband will be sharper. So the overlap between the signal spectrum and its images will be narrower. The edge channels will still have to be rejected but it will be a smaller portion of the whole data. \\item Oversampling the cross-correlation function at finer delay steps as already illustrated with Fig.~\\ref{fig:under-vs-oversamp}. \\item Reducing the bandwidth so that it is not critically-sampled. \\end{enumerate} Implementing either of the first two modifications in hardware would be costly and the third approach would lose sensitivity. In Sect.~\\ref{sec:os-method}, we propose a software-based approach to oversample the cross-correlation function and avoid aliasing. This is based on the well-known notion that the spectrum of a source can be obtained from the cross-correlation function measured by a drift scan. We can reconstruct the cross-correlation function using a combination of source tracking, path compensation and fringe-rotation. In Sect.~\\ref{sec:simulations}, we will illustrate an application of this technique with simulations. The system requirements are: (1) Individually-tracking antennas and (2) discrete-delay path compensation. We have modelled an analogue correlator here but the principles could equally apply to digital correlators.  \\revd{A} number of groups recently reported broadband analogue Fourier transform correlators suffering from delay errors (for example, \\citealt{harris2001}; {\\paperone}; \\citealt{roberts2007}). Simulations by {\\paperone} showed that delay errors make the effects of aliasing worse. Although the delay errors can be calibrated a sample at a time (recently by \\citealt{harris2001}; {\\paperone}), the method described here \\revd{will be} a natural way of accounting for delay errors. This \\revd{side benefit} is perhaps as significant as the alias suppression aspect of the method. The resulting spectral channels are also regularly gridded at the desired frequencies.  However there are a number of issues that must be considered in a practical system and these are discussed in Sect.~\\ref{sec:practical-issues}. This paper is a follow up to {\\paperone} where we described the development of a broadband 6--12{\\GHz} Fourier transform correlator for the Arcminute Microkelvin Imager \\citep[AMI;][]{kaneko2006,zwart2008}. AMI is a new interferometer designed to survey for clusters of galaxies by exploiting the Sunyaev-Zel'dovich effect \\citep{sunyaev1972}. The effects of aliasing in the correlator used in this instrument are discussed further in {\\paperone}. Before going into a detailed discussion of the oversampling method, we will first give a brief overview of how the data is processed in a conventional interferometric system.   \\begin{figure*}[hbtp] \\centering \\mbox{\\subfigure[The \\revd{baseband signal ({0--12\\GHz}) under and critically-sampled. The spectra show signs of aliasing. The cross-correlation function was sampled at 16 and 32 delays.}]{ \\includegraphics[width=17cm]{pic/8117fig2a.eps}}} \\mbox{\\subfigure[\\revd{Oversampling the baseband signal by $2\\times$ and $4\\times$ critical-sampling reduces aliasing.}]{ \\includegraphics[width=17cm]{pic/8117fig2b.eps}}} \\caption[Critical and oversampling]{The effects of different oversampling rates on the recovered spectra. The left-most plots show the cross-correlation function \\revd{from Eq.~(\\ref{eqn:corr-response4}) as a function of the residual delay {\\Dtau}. This is a snapshot when the source transits an east-west baseline so the geometrical delay $\\taug = 0$ (case~1).} The next \\revd{set of} plots to the right \\revd{(case~2) show the cross-correlation function a little later when the centre of the envelope has drifted by a quarter of a delay spacing}. The measurements at transit \\revd{(case~1)} are indicated by the grey circles and the measurements later on \\revd{(case~2)} are indicated by the black circles. The horizontal axis is the delay in {\\Dtau}, which will be defined in Sect.~\\ref{sec:real-comp-corr}. The right-most plots are the amplitudes calculated from the delay measurements. The filled grey marks are at transit \\revd{(case~1)} and the dark circles are a while later \\revd{(case~2)}. The two sets of marks should match but it can be seen in {\\bf a)} that they do not. This mismatch shows up as the temporal modulation seen in Fig.~\\ref{fig:alias-cycles}. In the upper row of {\\bf a)}, \\revd{although the {6\\GHz} signal bandwidth} is critically-sampled at 16 delays, \\revd{the} {0--12\\GHz} \\revd{baseband signal is undersampled.} The two halves of the spectrum are side-by-side. The spectrum is aliased at both the high and low-frequency ends.  In the lower row of {\\bf a)}, the highest frequency component ({12\\GHz}) is critically-sampled with 32 delays. Now only the upper frequency range suffers from aliasing. In {\\bf b)}, \\revd{the baseband signal is oversampled by a factor of 2 and 4 (sampled at 64 and 128 delays respectively)}. Now the two sets of spectral amplitudes match. Oversampling ensures a buffer between the two halves of the signal spectrum as well as with the image spectra. This reduces the spectral overlap and suppresses aliasing. In this case, oversampling the \\revd{baseband frequency} ({0--12\\GHz}) by a factor of 2 is sufficient (\\revd{an equivalent} sampling frequency of {48\\GHz} at 64 delays). Oversampling above this does not reduce aliasing any further. As a minor technical detail, the forms of the spectra in {\\bf b)} are slightly different from the one in Fig.~\\ref{fig:ccf-spec}. Here (and also in Fig.~\\ref{fig:alias-cycles}), we shifted the spectra by half a sub-band using the shift theorem of Fourier transforms. This ensures that the sub-bands span the desired frequency range and is discussed further in \\cite{holler2007}. Subsequent simulations will assume channel-shifting.}\\label{fig:under-vs-oversamp} \\end{figure*} ",
        "conclusions": "We have described a software-based method for overcoming aliasing in critically-sampled Fourier transform correlators. The method reduces aliasing by reconstructing the oversampled cross-correlation function from successive samples. Data from complex correlators can be fringe-rotated directly and efficiently in software. If using this method on real correlators, the signal should be fringe-rotated in hardware to avoid systematic errors.  The edge spectral channels are usually discarded because they are the worst affected by aliasing. For an 8-channel complex correlator simulated here, the sensitivity could be improved by up to 15~percent by retaining the edge channels. The additional computational overhead and complexity are offset by an improvement in sensitivity and reduction in systematic errors from aliasing. This could be a significant benefit for Fourier transform correlators with less than 10 spectral channels. Some data with low fringe rates will have to be rejected, but for AMI-like telescopes with fractional bandwidths $\\Dnu / \\nuRF \\gtrsim 0.4$ and maximum baseline to dish ratios of $D/d \\lesssim 10$, the fraction of data lost is relatively small. We have also shown that we can naturally compensate for delay errors to recover regularly-gridded frequency channels at the design frequencies."
    },
    "0807/0807.2244_arXiv.txt": {
        "abstract": "{ We review what can (and cannot) be learned if dark matter is detected in one or more experiments, emphasizing the importance of combining LHC data with direct, astrophysical and cosmological probes of dark matter. We briefly review the conventional picture of a thermally produced WIMP relic density and its connection with theories of electroweak symmetry breaking. We then discuss both experimental and theoretical reasons why one might generically expect this picture to fail.  If this is the case, we argue that a combined effort bringing together all types of data -- combined with explicitly constructed theoretical models -- will be the only way to achieve a complete understanding of the dark matter in our universe and become confident that any candidate actually provides the relic density.}  \\preprint{}  \\begin{document} ",
        "introduction": "One of the most compelling hints for physics beyond the standard model is the cosmological observation that nearly a quarter of our universe consists of cold dark matter. In the next few years, the Large Hadron Collider (LHC) shows the promise of producing these elusive particles and possibly measuring their microscopic properties \\cite{Brhlik:2000dm,list}. This will be challenging and using LHC observations to reconstruct a complete theory of cosmological dark matter could prove even more challenging, if possible at all. New information will soon come from PAMELA, GLAST, and other experiments and the interpretation of the reported DAMA annual modulation effect may be clarified.  One reason is because before we can claim that we know what the dark matter actually is, it will be essential to calculate the relic density of each candidate and compare the total with the cosmological relic density determined from the dark matter's gravitational influence -- because it is impossible to measure the relic density directly. This calculation will require a detailed knowledge of {\\em all} dark matter particles that contribute to the relic density, their interactions, and a knowledge of the cosmological expansion history. For example, we already know that neutrinos make up roughly a percent of the dark matter, and most complete theories have in addition to massive neutrinos also Weakly Interacting Massive Particles (WIMPs) and axions, so naively there is no convincing reason to expect one form to entirely dominate, as first emphasized in \\cite{Brhlik:2000dm}. In addition, to calculate the relic density of dark matter candidates we must know how they annihilate as the universe cools, so we can calculate the number left today.  Which diagrams are important depends on the particle interactions.  For example, if the dark matter particles of interest are the lightest superpartner it could behave like the partners of the W boson (wino, $\\tilde{W}$), the Higgs boson (higgsino, $\\tilde{h}$), the photino $\\tilde{\\gamma}$, or the sneutrino $\\tilde{\\nu}$, etc.  These interact differently, and annihilate differently, so it is necessary to know the relevant composition and couplings, which means measuring them, perhaps in the context of a theory that allows some to be calculated.  Each direct, indirect and collider channel can contribute to this information. It is also necessary to know the cosmological history of the universe from the end of inflation until the relic density is no longer changing. Of course, as soon as there are candidates there will be acclaim for the discovery of dark matter, but the question is too important to ignore the need for a proper calculation of the relic density.  In the remainder of the paper we present a review of dark matter, emphasizing the issue of reconstructing a complete picture from both the particle physics and cosmological perspectives -- referring the reader to \\cite{Bertone:2004pz} for a more general review, particularly of candidates and assuming an entirely thermal history. In Section \\ref{section1} we review the cosmological evidence for dark matter and its expected distribution inside our galaxy.  In the next section, after a critical overview of the standard calculation of thermal relic density of dark matter -- noting the numerous assumptions that are involved -- we then discuss how this naturally connects the dark matter with new physics near the scale of electroweak symmetry breaking.  In Section 4 we discuss what we call the ``dark matter inverse problem''.  That is, if we were able to reconstruct signatures of dark matter at LHC (the possibility of which we review) is it possible to also reconstruct the dark matter relic density from this information alone and be convinced that we have obtained a complete understanding of dark matter?  The answer is most likely no, and we will give many reasons -- both phenomenological and theoretical-- why this is not only unlikely, but also not expected from complete and self consistent theories. Instead, as we discuss in Section 5, by combining LHC data with other direct and indirect detection experiments it may be possible to attain a complete understanding of dark matter. Indeed, this is a particularly exciting time, since a number of experiments such as GLAST and PAMELA should be reporting data soon, many other experiments will come online in the near future, and when these astrophysical probes are combined with collider data and cosmological probes a more complete picture of dark matter will emerge.  In the last section we conclude with  a summary of how one should proceed given a detection of a dark matter candidate at LHC, direct, or indirect detection experiment. ",
        "conclusions": "Both cosmological observations and theoretical consistencies of electroweak symmetry breaking seem to suggest the presence of WIMPs. The expectation of data coming from LHC makes this an exciting time to ask whether these two notions of dark matter are one and the same -- or if the story is perhaps more complicated. We have seen in this review that theory would suggest the latter, but that by combining data from colliders with that coming from cosmological and astrophysical probes of dark matter a complete understanding of dark matter can probably be achieved.  In fact, one is tempted to compare this situation with the concordance approach of modern precision cosmology where one experiment alone (such as WMAP) is not enough to overcome uncertainties, but when combined with other probes (such as supernovae, baryon acoustic oscillations, galaxy surveys, etc.) one can obtain a more constrained and precise fit to the cosmological parameters.  Thus, by combining the data from different experiments in this way -- along with explicit and complete theoretical models that not only predict specific dark matter candidates, but also the cosmological expansion history -- we can hope to obtain a complete understanding of dark matter. \\\\  \\noindent {\\bf How should one proceed once a dark matter candidate is established?} \\\\ One good path is to calculate the relic density in models that are sufficiently complete to have a cosmological history, and accommodate electroweak symmetry breaking.  Presumably several different models would emerge as consistent with dark matter and LHC data, but all of them would have other predictions and tests that would allow one class of them to eventually become convincing.  Another approach would be to calculate the relic density as if the history were thermal (using the experimentally inferred self annihilation cross section).  Then there are several possibilities (as discussed above): \\bi \\item The relic density could come out about right (i.e., agree with the cosmological abundance determined by e.g. WMAP), in which case it would belong to the previous paragraph. \\item The relic density could come out smaller than the WMAP number, in which case either: \\bi \\item There is another source of dark matter, e.g. axions or some heavy hidden sector matter. \\item And/Or, the dark matter was non-thermally produced yielding an increased relic density. \\item And/Or, at the time the dark matter was produced the cosmic expansion was not entirely radiation dominated, e.g. kination. \\ei \\item The relic density could come out larger than the WMAP number, in which case either: \\bi \\item  There was non-thermal production yielding less than the thermal relic density. \\item  And\\footnote{Note:  As we discussed above, at the time non-thermal production occurs (e.g. from decay of a heavy scalar) this will produce substantial entropy.  This will dilute the thermal abundance, which will also contribute to the total relic density.  One could imagine more complicated scenarios where all of these factors play a role and the relic density of dark matter gets both thermal and non-thermal contributions (as is expected for example in the case of axions \\cite{arxiv:0807.1726}).}/Or, there were additional sources of entropy after thermal freeze-out (e.g. decays of heavier particles, such as scalars (moduli) or gravitinos). \\item And/Or the cosmic expansion was modified after thermal freeze-out of the dark matter (e.g. thermal inflation). \\ei \\ei  In all cases, we will be learning, amazingly, about the cosmological history of the universe, from LHC, direct, and indirect detection of dark matter."
    },
    "0807/0807.1717_arXiv.txt": {
        "abstract": "Near Earth Objects (NEOs) are fragments of remnant primitive bodies that date from the era of Solar System formation. At present, the physical properties and origins of NEOs are poorly understood. We have measured thermal emission from three NEOs --- (6037) 1988~EG, 1993~GD, and 2005~GL --- with Spitzer's IRAC instrument at 3.6, 4.5, 5.8, and 8.0~$\\mu$m (the last object was detected only at 5.8~and 8.0~$\\mu$m). The diameters of these three objects are 400~m, 180~m, and 160~m, respectively, % with uncertainties of around 20\\% % (including both observational and systematic errors). For all three the geometric albedos are around~0.30, in agreement with previous results that most NEOs are S-class asteroids. For the two objects detected at 3.6~and 4.5~$\\mu$m, diameters and albedos based only on those data agree with the values based on modeling the data in all four bands. This agreement, and the high sensitivity of IRAC, show the promise of the Spitzer Warm Mission for determining the physical parameters for a large number of NEOs. ",
        "introduction": "Near Earth Objects (NEOs) are bodies whose orbits pass within a few tenths of an AU of the Earth's orbit. As of this writing, there are around 5000 NEOs known. The Pan-STARRS program is likely to increase the number of known NEOs to $\\sim$10,000 or more by 2013. These bodies are of critical interest to both the scientific community and the public. The NEO population is the source of potential Earth-impacting asteroids (hence a Congressional mandate to study these objects), and some may be easily reached by spacecraft, enabling our exploration of the nearby Solar System. Because NEOs have only recently been perturbed out of orbits in the main asteroid belt, and so are relatively primitive objects, they contain information that records the origin of our Solar System and that may offer insight into both the past (via delivery of organic material) and future (via impact-caused extinctions) of life on Earth. However, the physical characterization of these objects is by far outpaced by discoveries. The NEO size and albedo distributions, crucial inputs for Solar System studies as well as the assessment of the NEO Earth impact hazard, are only poorly constrained, especially at the smallest sizes \\citep[e.g.,][]{StuartBinzel2004}.   The diameter and albedo of asteroids can be determined from thermal-infrared observations together with appropriate thermal modeling \\citep[e.g.,][]{STM,LebofskySpencer1989,HarrisLagerros2002}, provided the absolute magnitude $H$ \\citep[optical magnitude at a standardized observing geometry; see][]{HG} is known. A suitable thermal model for NEOs is the Near-Earth Asteroid Thermal model \\citep[NEATM, ][]{neatm}, which allows for simultaneous fits of the asteroid diameter, albedo, and effective surface temperature (parametrized through the beaming parameter $\\eta$) \\citep[e.g.,][]{neatm,HarrisLagerros2002}.   We have measured thermal emission from three NEOs with the Spitzer Space Telescope \\citep{werner04} and present our data (\\S2) and results (\\S3) here. Using the NEATM, we derive albedos and diameters for all three objects (\\S3). In \\S4 we comment on the apparent thermal inertias for these objects and demonstrate that a study of NEOs could profitably be carried out with the Spitzer Warm Mission. ",
        "conclusions": "All three objects have diameters less than 500~meters, making them among the smallest NEOs with known albedos and diameters, and among the smallest individual objects studied with the Spitzer Space Telescope. All three objects also have albedos close to~0.3, in agreement with the idea that the NEO population is dominated by S-class asteroids \\citep[e.g.,][]{binzel04}. \\citet{binzel04} also found that the albedos for S-class (and related classes) NEOs rise from their main belt average value around~0.22 to greater than~0.3 for objects $\\lesssim$500~m. Our results appear to confirm this trend (Figure~\\ref{sed}), though with small numbers and not insignificant error bars. It is quite premature to discuss the reality of the potentially interesting downward turn at even smaller sizes.  The best-fit (floating) $\\eta$ values found for 6037 and 1993~GD are roughly consistent with empirical expectations \\citep{Delbo2003}, which were recently used by \\citet{Delbo2007} to determine the typical thermal inertia of $D\\sim 1$~km NEAs. Thermal inertia is indicative of the presence or absence of loose material (regolith) on the surface and is a key parameter for model calculations of the Yarkovsky effect, a non-gravitational force that severely influences the orbital dynamics of small asteroids. (Note that \\citet{Vokrouhlicky2005} list 6037 as a potential target for direct observations of the Yarkovsky effect.) Our results suggest that our targets have unremarkable thermal inertias and may be similar to the 320~meter diameter S-type NEO (25143) Itokawa \\citep{ThMueller,MuellerDiss}, the target of the Hayabusa mission. However, more work and a systematic, large survey are needed to determine the typical thermal inertia of sub-km NEAs.   For 6037 and 1993~GD the diameters and albedos we derive using only 3.6~and 4.5~$\\mu$m data are in agreement with our other model solutions, particularly for 6037, which is strongly detected (SNR$>$10) in both bands. This agreement has important implications for the Spitzer Warm Mission. After Spitzer's onboard cryogen is exhausted, observations in IRAC bands 1 and 2 (3.6 and 4.5~$\\mu$m) can still be made with essentially no loss of sensitivity. Our results show the promise of capitalizing on the superior sensitivity of IRAC to determine the physical properties of a large number of NEOS during the Spitzer Warm Mission."
    },
    "0807/0807.4304_arXiv.txt": {
        "abstract": "The critical issue in cosmology today lies in determining if the cosmological constant is the underlying ingredient of dark energy. Our profound lack of understanding of the physics of dark energy places severe constrains on our ability to say anything about its possible dynamical nature. Quoted errors on the equation of state, $w(z)$, are so heavily dependent on necessarily over-simplified parameterisations they are at risk of being rendered meaningless. Moreover, the existence of degeneracies between the reconstructed $w(z)$ and the matter and curvature densities weakens any conclusions still further. We propose consistency tests for the cosmological constant which provide a direct observational signal if $\\Lambda$ is wrong, regardless of the densities of matter and curvature. As an example of its utility, our flat case test can warn of a small transition from $w(z)=-1$ of $20\\%$ from SNAP quality data at 4-$\\sigma$, even when direct reconstruction techniques see virtually no evidence for deviation from~$\\Lambda$. It is shown to successfully rule out a wide range of non-$\\Lambda$ dark energy models with no reliance on knowledge of $\\Omega_m$ using SNAP-quaility data and a large range for using $10^5$ supernovae as forecasted for LSST. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0807/0807.4852.txt": {
        "abstract": "The increased sensitivity of future radio telescopes will result in requirements for higher dynamic range within the image as well as better resolution and immunity to interference. In this paper we propose a new matrix formulation of the imaging equation in the cases of non co-planar arrays and polarimetric measurements. Then we improve our parametric imaging techniques in terms of resolution and estimation accuracy. This is done by enhancing both the MVDR parametric imaging, introducing alternative dirty images and by introducing better power estimates based on least squares, with positive semi-definite constraints. We also discuss the use of robust Capon beamforming and semi-definite programming for solving the self-calibration problem. Additionally we provide statistical analysis of the bias of the MVDR beamformer for the case of moving array, which serves as a first step in analyzing iterative approaches such as CLEAN and the techniques proposed in this paper. Finally we demonstrate a full deconvolution process based on the parametric imaging techniques and show its improved resolution and sensitivity compared to the CLEAN method. \\\\ {\\bf Keywords}: Radio astronomy, synthesis imaging, parametric imaging, minimum variance, robust beamforming, convex optimization, CLEAN. ",
        "introduction": "The future of radio astronomical discoveries depends on achieving better resolution and sensitivity while maintaining immunity to terrestrial interference which is rapidly growing. To achieve the improved sensitivity and higher resolution new instruments are being designed. The Square Kilometer Array (SKA) \\footnote{For information on the SKA project the reader is referred to http:$//$www.skatelescope.org.} \\cite{hall2005} and the Low Frequency Array (LOFAR) \\cite{bregman2005} are two of these advanced instruments. Achieving higher sensitivity to observe faint objects results in  high dynamic range requirements within the image, where strong sources can affect the imaging of the very weak sources. On the other hand, Moore's law \\cite{moore65} together with recent advances in optimization theory \\cite{boyd2004} open the way to the application of more advanced computational techniques. In contrast to hardware implementation, these image formation algorithms, that are implemented in software can benefit from the continuing increase in computational power, even after the antennas and the correlators will be built. In this paper we extend the parametric deconvolution approach of \\cite{leshem2000a} to obtain better power estimation accuracy and higher robustness to interference and modeling uncertainty. The algorithms presented here can also be used in conjunction with real-time interference mitigation techniques as described in \\cite{leshem2000a} and \\cite{leshem2000b}.  We briefly describe the current status of radio astronomical imaging techniques. For a more extensive overview the reader is referred to \\cite{thompson86}, \\cite{yen85} or \\cite{taylor99}. A good historic perspective can be found in \\cite{kellerman2001}, whereas \\cite{sault2007} provides a very recent perspective.  The principle of radio interferometry has been used in radio astronomy since 1946 when Ryle and Vonberg constructed a radio interferometer using  dipole antenna arrays \\cite{ryle52}. During the 1950's several radio interferometers which use the synthetic aperture created by movable antennas have been constructed. In 1962 the principle of aperture synthesis using earth rotation has been proposed \\cite{ryle62}. The basic idea is to exploit the rotation of the earth to obtain denser coverage of the visibility domain (spatial Fourier domain). The first instrument to use this principle was the five kilometer Cambridge radio telescope.  During the 1970's new instruments with large aperture have been constructed. Among these we find the Westerbork Synthesis Radio Telescope (WSRT) in the Netherlands and the Very Large Array (VLA) in the USA. Recently, the Giant Microwave Telescopes (GMRT) has been constructed in India and the Allen Telescope Array (ATA) in the US. Even these instruments subsample the Fourier domain, so that unique reconstruction is not possible without some further processing known as deconvolution. The deconvolution process uses some a-priori knowledge about the image to remove the effect of ``dirty beam'' side-lobes.  Two principles dominate the astronomical imaging deconvolution. The first method was proposed by Hogbom \\cite{hogbom74} and is known as CLEAN. The CLEAN method is basically a sequential Least-Squares (LS) fitting procedure in which the brightest source location and power are estimated. The response of this source is removed from the image and then the process continues to find the next brightest source, until the residual image is noise-like.  During the years it has been partially analyzed \\cite{schwarz78}, \\cite{schwarz79} and \\cite{tan86}. However full analysis  of the method is still lacking due to its iterative nature. The CLEAN algorithm has many recent flavors which are capable of faster performance e.g., the Clark version \\cite{clark80} and the Cotton-Schwab \\cite{schwab84}. In these versions the dirty image is recomputed only after several point sources have been estimated. Furthermore the sources are subtracted from the ungridded visibility. This results in better suppression of the sources. Interestingly it is well known that when the noise model is non-white (as in radar clutter modeled by ARMA processes in SAR applications) algorithms such as RELAX \\cite{li96} outperform the CLEAN algorithm.  A second approach proposed by Jaynes \\cite{Jaynes57} is maximum entropy deconvolution (MEM). The basic idea behind MEM is the following.  Not all images which are consistent with the measured data and the noise distribution satisfy the positivity demand, i.e., the sky brightness is a positive function. Consider only those that satisfy the positivity demand. From these select the one that is most likely to have been created randomly. This idea has also been proposed in \\cite{frieden72} and applied to radio astronomical imaging in \\cite{gull78}. Other approaches based on the differential entropy have also been proposed \\cite{ables74}, \\cite{wernecke77}. An extensive collection of papers discussing the various methods and aspects of maximum entropy can be found in the various papers in \\cite{roberts84}. Briggs \\cite{briggs95} proposed a non-negative least squares approach (NNLS) which eliminates the need for iterative processing. However, the computational complexity is very large.  In this paper we use a reformulation of the image formation problem as a parameter estimation problem using a set of covariance matrices, measured at the various observation epochs \\cite{leshem2000a}. This yields a model where the array response is time varying. Previous research on time varying arrays and their application to direction-of arrival (DOA) estimation includes \\cite{zeira95}, \\cite{zeira96} and \\cite{sheinvald98}. In \\cite{leshem2000a} we proposed a simplified ML estimator. Lanterman \\cite{lanterman2000} developed a full EM algorithm for implementing the MLE proposed in \\cite{leshem2000a}. The algorithm performs quite well although it is quite complex compared to the solutions described in this paper.  The above mentioned algorithms assume perfect knowledge of the instrumental response (point spread function). Due to various internal and external effects this assumption holds only approximately. One way to overcome this problem is the use of calibrating sources. An unresolved source with known parameters is measured, and by relating the model errors to the array elements a set of calibration equations is solved. A much more appealing solution is to try to improve the fitting between the data and the sky model by adjusting the calibration parameters. Another possibility \\cite{noordam82} is to use the redundant structure of the array to solve for the calibration parameters (this is possible only for some arrays which have redundant baselines, such as the WSRT). A good overview of the various techniques is given in \\cite{pearson84}.  In this paper we extend the above methods in several directions. First, we extend the parametric formulation of \\cite{leshem2000a} to the non-coplanar array case and to polarimetric imaging. Then we propose relatively ``low'' complexity approaches for the image deconvolution problem, based on Minimum Variance Distortionless Response (MVDR) and its robust extensions. We call these extensions Least-Squares Minimum Variance Imaging (LS-MVI). We provide a new type of dirty image that has isotropic noise response, something desirable in imaging applications. This is done by generalizing the work of Borgiotti and Kaplan \\cite{borgiotti79} to the moving array case. We discuss acceleration techniques, related to the Clark \\cite{clark80} and the Cotton-Schwab \\cite{schwab84} approaches to CLEAN. However, in contrast to the classical CLEAN case, the accelerated algorithm involves semi-definite programming of low order, ensuring that the covariance matrices remain positive semi-definite after the subtraction. We provide analytic expressions for the asymptotic bias of the MVDR based imaging. We also relate the classical self-calibration technique to a novel extension of the robust Capon beamformer to the moving array case, showing that self-calibration can be cast in terms of semi-definite programming. This has advantage over previous approach to self-calibration since the covariance structure maintains its positive definite structure.  We also demonstrate full parametric deconvolution process based on the proposed technique and compare CLEAN and LS-MVI on simulated images. We will show in simulations that LS-MVI, indeed, significantly outperforms the classical CLEAN algorithm over a wide range of parameters, in terms of better resolution, higher sensitivity and dynamic range.  The structure of the paper is as follows.  In section \\ref{sec:astron} we describe the astronomical measurement equation. The measurement equation is subsequently rephrased in a more convenient matrix formulation both for non co-planar and co-planar arrays. It is then extended with the effect of noise and unknown calibration parameters. We also discuss extension to polarimetric measurements. In section \\ref{sec:LS_MVI} we discuss the new Least Squares Minimum Variance Imaging and extensions to self calibration using robust Capon techniques. In section \\ref{sec:MVDR_analysis} we present bias analysis of the MVDR imaging technique.  In \\ref{sec:simulation} we describe several computer simulations demonstrating the gain in parametric deconvolution relative to the CLEAN in terms of resolution, sensitivity and capability to model extended structures. We also provide example of the tightness of the statistical bias analysis of the MVDR DOA estimation with a moving array. We end up with some conclusions. ",
        "conclusions": "In this paper we extend the matrix formulation of \\cite{leshem2000a} to non co-planar arrays and polarimetric measurements. Then we propose a new parametric imaging technique that improves the resolution and sensitivity over the classical CLEAN algorithm. The method is based on several improvements: A new type of dirty image, LS estimation of the powers and semi-definite constraints. We show how the technique can be combined into self-calibration using semi-definite programming. Our semi-definite self-calibration algorithm also provides a new approach to robust beamforming with a moving array, extending the techniques of \\cite{li03}, \\cite{vorobyov03} and \\cite{lorenz05}. We provide statistical analysis of the location estimator. Simulated examples comparing full deconvolution using LS-MVI and comparing them to the CLEAN method are presented. These simulations demonstrate that the parametric approach has higher resolution, is more robust to source structures, and performs better in noisy situations. The great potential of the methods proposed in this paper is a first step, towards the development of more advanced imaging techniques, capable of providing higher dynamic range and interference immunity as required by the radio telescopes of the future. %"
    },
    "0807/0807.1467_arXiv.txt": {
        "abstract": "{} {We searched for the presence of extended emission-line regions (EELRs) around low-redshift QSOs.} {We observed a sample of 20 mainly radio-quiet low-redshift quasars ($z<0.3)$ by means of integral field spectroscopy. After decomposing the extended and nuclear emission components, we constructed \\Ox\\ $\\lambda5007$ narrow-band images of the EELR to measure the total flux. From the same data we obtained high S/N ($>$50) nuclear spectra to measure properties such as \\Ox/\\Hb\\ flux ratios, \\Fe\\ equivalent widths and H$\\beta$ line widths.} {A significant fraction of the quasars (8/20) show a luminous EELR, with detected linear sizes of several kpc. Whether or not a QSO has a luminous EELR is strongly related with nuclear properties, in the sense that an EELR was detected in objects with low \\Fe\\ equivalent width and large H$\\beta$ FWHM.  The EELRs were detected preferentially in QSOs with larger black hole masses. There is no discernible relation, however, between EELR detection and QSO luminosity and Eddington ratio.}  {}{} ",
        "introduction": "Quasars may show extended emission-line regions (EELRs) stretching over characteristic scales of ten to hundred kpc. Luminous EELRs have been found in particular around steep-spectrum radio-loud quasars \\citep[RLQs,][]{Boroson:1984,Heckman:1991,Crawford:2000,Stockton:2002,Fu:2006} and also radio galaxies \\citep[e.g.][]{Villar-Martin:2006,Villar-Martin:2007}. There are also indications that radio-quiet quasars (RQQs) can have EELRs \\citep[e.g.,][]{Boroson:1985,Stockton:1987,Bennert:2002}. Most existing studies, however, are based either on long-slit spectroscopy or on narrow-band imaging; while the former technique will generally capture only part of any extended emission, the latter does not provide any spectroscopic information.  Integral field spectroscopy is a relatively new and powerful tool to study AGN host galaxies. The combined diagnostic power of imaging and spectroscopy allows for a flexible treatmeant of the data, facilitating the construction of broad- or narrow-band images, of 2-dimensional velocity maps as traced by emission or absorption lines, or the spectral analysis of selected regions within the field of view.  Here we present preliminary results from a study of 20 low-redshift QSOs, most of which are RQQs, observed with an integral field unit (IFU).  After decomposing nuclear and extended emission we identified 8 quasars with clearly detectable EELRs. We focus on exploring the relation between the nuclear spectral properties of a QSO and the existence or non-existence of a luminous EELR. A full account of our IFU study will be given in a separate paper (Husemann et al., in preparation; hereafter Paper~II). ",
        "conclusions": "By means of integral field spectroscopy, we have revealed the presence of extended emission line regions (EELRs) in a sample of predominantly radio-quiet low-redshift QSOs.  The sizes and luminosities are much higher than those of the well-known EELRs in Seyfert galaxies, extending out to several kpc and therefore stretching over the entire host galaxy (and possibly beyond).  Some 40~\\% of the objects in our sample show a prominent EELR, while we failed to detect any EELR signature in the remaining 60~\\%. The detection or non-detection of an EELR can be linked to the \\emph{nuclear} spectral properties of the quasar itself, independent of their radio classification, in excellent agreement with what \\citet{Boroson:1984} and \\citet{Boroson:1985} found from their off-nuclear long-slit spectra. Our observations go further in that our integral field spectral data map out the entire nuclear-subtracted EELRs and derive their luminosities.  Correlations between various nuclear spectral properties were already studied by many authors \\citep[e.g.][]{Boroson:1992,Sulentic:2000,Netzer:2004}. Here we provide new evidence that properties of the nuclear spectra, depending on the sub-pc scale conditions in AGN, are linked to properties of the host galaxies extending over several kpc. It thus appears possible to predict with high confidence whether or not a QSO will show a very extended emission line region just from its nuclear spectrum: If it has broad \\Hb\\ and weak \\Fe, it will probably have not only strong nuclear, but also strong extended \\Ox\\ emission.  It is tempting to interpret this relation between sub-pc and super-kpc scales in the context of the current discussion about AGN feedback affecting the host galaxies. Since we can rule out the QSO luminosity as a main driver of the differences between presence or absence of an EELR, the remaining possibilities are that there could be differences in the structure of the nuclear region (obscuration in connection with possible inclination effects); or differences due to an outflow or jet, or the host galaxies of the two QSO classes are intrinsically different, in particular with respect to their overall distribution of warm ionised gas. We conclude that QSO host galaxies come in at least two flavours: with or without extended ionized gas. Whether this highlights different initial conditions or a sequence of evolutionary stages we cannot say at present. More and better quality data will be needed to elucidate these intriguing trends."
    },
    "0807/0807.1698_arXiv.txt": {
        "abstract": "We present optical time series spectroscopy of the pulsating white dwarf star G~29-38 taken at the Very Large Telescope (VLT). By measuring the variations in brightness, Doppler shift, and line shape of each spectrum, we explore the physics of pulsation and measure the spherical degree ($\\ell$) of each stellar pulsation mode.  We measure the physical motion of the g-modes correlated with the brightness variations for three of the eight pulsation modes in this data set.  The varying line shape reveals the spherical degree of the pulsations, an important quantity for properly modeling the interior of the star with asteroseismology. Performing fits to the H$\\beta$, H$\\gamma$, and H$\\delta$ lines, we quantify the changing shape of the line and compare them to models and previous time series spectroscopy of G~29-38. These VLT data confirm several $\\ell$ identifications and add four new values, including an additional $\\ell$=2 and a possible $\\ell$=4. In total from both sets of spectroscopy of G~29-38, eleven modes now have known spherical degrees. ",
        "introduction": "DAVs, or ZZ Cetis, are variable white dwarf stars with hydrogen atmospheres. They are the  coolest known white dwarf pulsators and reside in an instability strip near 12,000~K. Their pulsations are multi-periodic with amplitudes as high as 5~\\percent\\ and periods between 70~s and 1100~s. In the past few years the number of known DAVs has increased from a few dozen to more than a hundred \\cite[see][]{voss07,castanheira06,mullally05,mukadam04}. The pulsators at the hot end of the instability strip show lower amplitudes and more stable pulsations, while the others show more complex pulsations including variability of the excited modes, and an abundance of harmonic and combination modes.  The high gravity (log(g)$\\sim$8 cgs) of the DAVs favor gravity-mode type pulsations \\citep{rkn82}, where motion along the surface of the star produces regions with varying temperatures.  To describe the spatial distribution of these regions, each of the pulsations on the star are characterized by a spherical harmonic (Y$_{\\ell,m}$) and the radial order, $n$. The spherical degree, $\\ell$, and azimuthal order, $m$, describe the number of nodal lines across the surface of the star, and directly affect observations of the pulsations.  The pulsation's radial component allows asteroseismological models to determine the internal structure of the star. The DAVs have evaded unambiguous asteroseismological modeling in part because of the difficulty in measuring the spherical degree of their pulsation modes.  Usually a value of $\\ell$=1 is assumed in these models. However, it is possible to fit modes with higher values of $\\ell$ and this can significantly change the parameters of the model. For example, \\citet{bradley06} showed that using the identification of $\\ell$=4 \\citep{thompson04}, instead of $\\ell$=2 \\citep{cast04} for one mode, changes the modeled depth of the Hydrogen layer in G~185-32 by a factor of 100.  One method to determine $\\ell$ independent of asteroseismological modeling was introduced by \\citet{robinson95}. They showed that $\\ell$ can be obtained from time series optical and UV spectroscopy. The effects of limb darkening increase in the blue wavelengths and decrease through the spectral lines, changing the fraction of integrated light coming from the limb.  As such, the pulsation amplitude varies with both wavelength and the spherical degree of the mode. \\citet*{c00} (hereafter C00) measured chromatic amplitudes (pulsation amplitude at each wavelength bin) across the Hydrogen Balmer lines and demonstrated that spectral variations can indeed be used to identify spherical degree. The same method was applied to the DAVs HS~0507+0434B, G~117-B15A, and G~185-32 \\citep{kotak02a, kotak04, thompson04} and the DBV GD~358 \\citep{kotak03}.  Generally, the method worked well, although for one mode -- the 272~s mode in G~117-B15A -- the chromatic amplitudes are extremely puzzling \\citep{kotak04}.  Also, they found that the chromatic amplitudes provided unambiguous identifications only for the strongest modes; for weaker ones, noise obscured the expected visible differences.  In order to improve the precision, \\citet{kotak03} attempted to use the slope and curvature in selected wavelength ranges of the chromatic amplitudes as measures of spherical degree, with somewhat limited success.  A different approach was demonstrated by \\citet{thompson04} where they analysed fits to the spectral lines rather than the observed fluxes in each spectral bin. The chromatic amplitudes created from these fits did allow for mode identifications from lower signal-to-noise data.  G~29-38 is one of the brightest DAVs with complex pulsations and some of the largest amplitudes, making it a prime candidate for time series spectroscopy. It has an abundance of modes, but they are not all excited simultaneously and many observation runs are necessary to measure them all.  \\citet{kleinman98} compiled 10~years of photometric observations and found a possible series of $\\ell$=1 modes on G~29-38.  Time series spectroscopy taken at the Keck Observatory with the LRIS spectrograph provided chromatic amplitudes of G~29-38 that directly detected an $\\ell$=2 mode at 776~s (C00) and gave strong evidence for another at 920~s \\citep{kotak02a}. Using the same data, \\citet*{vk00} (hereafter VK00) measured the periodic line-of-sight velocities associated with the longitudinal motion of the gas. For the 776~s mode they found relatively large motion as compared to the measured flux amplitude, again indicating that the mode is $\\ell$=2 \\citep{dz82}.  The measurements of the pulsation velocities were confirmed by \\citet{thompson03} using a series of high resolution spectra of G~29-38 and provided observational evidence regarding the convective driving theory proposed for DAVs \\citep{bhill83, gw1}.  The velocities also hint at the nature of the many combination and harmonic modes present on these pulsators. These combinations and harmonics most likely result from nonlinear mixing at the surface of the star and are not real modes that probe the interior \\citep{wu01, montgomery05}. VK00 reinforced this idea by finding no velocity signature associated with the combinations and harmonics on G~29-38. However, \\citet{thompson03} reported one combination mode with a significant velocity detection.  In this paper we show the results of time series spectroscopy of G~29-38 taken at the VLT (Very Large Telescope) in 1999.  We present the light and velocity curves measured from the series of spectra.  We fit 13 modes in the light curve and find variations at the same period in the velocity curve for the largest amplitude pulsations. We then perform fits to the Balmer lines to quantify the changing line shape.  Since we are using fits to the observed spectral lines, we can present the line shape variations in terms of the fitted values: the normalized equivalent width variations of the Gaussian and Lorentzian functions. We perform the same analysis on the 1996 Keck spectral series presented by VK00 and C00 as a check of our method and recover their results of an $\\ell$=2 mode among a series of $\\ell$=1 modes. Application of the same method to our VLT data adds four additional mode identifications to the G~29-38 pulsation spectrum. ",
        "conclusions": ""
    },
    "0807/0807.2522_arXiv.txt": {
        "abstract": "The Guide Star Catalog II (GSC-II) is an all-sky database of objects derived from the uncompressed Digitized Sky Surveys that the Space Telescope Science Institute has created from the Palomar and UK Schmidt survey plates and made available to the community.  Like its predecessor (GSC-I), the GSC-II was primarily created to provide guide star information and observation planning support for {\\it Hubble Space Telescope}. This version, however, is already employed at some of the ground-based new-technology telescopes such as GEMINI, VLT, and TNG, and will also be used to provide support for the James Webb Space Telescope (JWST) and Gaia space missions as well as the Large Sky Area Multi-Object Fiber Spectroscopic Telescope, one of the major ongoing scientific projects in China.  Two catalogs have already been extracted from the GSC-II database and released to the astronomical community. A magnitude-limited ($R_F$ = 18.0) version, GSC2.2, was distributed soon after its production in 2001, while the GSC2.3 release has been available for general access since 2007.  The GSC2.3 catalog described in this paper contains astrometry, photometry, and classification for 945,592,683 objects down to the magnitude limit of the plates. Positions are tied to the International Celestial Reference System; for stellar sources, the all-sky average absolute error per coordinate ranges from $0.''2$ to $0.''28$ depending on magnitude. When dealing with extended objects, astrometric errors are ~20\\% worse in the case of galaxies and approximately a factor of 2 worse for blended images.  Stellar photometry is determined to 0.13-0.22 mag as a function of magnitude and photographic passbands ($R_F$, $B_J$, $I_N$). Outside of the galactic plane, stellar classification is reliable to at least 90\\% confidence for magnitudes brighter than $R_F=19.5$, and the catalog is complete to ~$R_F$=20. ",
        "introduction": "The Guide Star Catalog II (GSC-II) is an astronomical database constructed from the scanned images of 9541 Palomar and UK Schmidt photographic sky survey plates digitized at Space Telescope Science Institute (STScI). These same plate images are also known as the Digitized Sky Survey (DSS), and are accessible directly from the STScI eb site as well as from a number of major astronomical data centers around the world. A subset of the original images, based on the second-epoch Palomar surveys only (also known as DPOSS), is available separately from Caltech \\citep{2003BASBr..23Q.197D}. Whilst all the images represent over 8 terabytes of data that are archived at STScI, the use of the H-transform compression technique \\citep{1992doss.conf..167W,1994SPIE.2199..703W} has enabled a reduction in data volume to $\\approx 1$ terabyte for distribution to data centers.  Each individual uncompressed image was processed to detect the objects, and calibrations were obtained for each plate by polynomial modeling against photometric and astrometric reference catalogs. The average accuracy in position for stellar objects of intermediate magnitude ($V\\lesssim 18.5$) is $\\sim 0.3''$, whilst the corresponding astrometric precision, on a scale of approximately 0.5 degrees, is close to $0.2''$. Photometric calibrations, carried out in the natural plate passband, show typical errors of 0.1-0.2 mag and systematic offsets lower than 0.1 mag. Stars are correctly classified as such with at least 90\\% confidence for magnitudes brighter than $R_F=19.5$.  The plate catalogs were loaded into a database system known as Catalog of Objects and Measured Parameters from All Sky Surveys (COMPASS), \\citep[see][]{1998asal.confE...3L}. This database was built using \"Objectivity\", a commercial Object-Oriented Database Management System; it constitutes a repository for  $\\approx 5$ billion measurements, and is $\\approx2.5$ terabytes in size. Once the individual observations have been cross-matched to link different measurements of unique astronomical objects, a catalog can be exported from the database as a collection of FITS binary tables. Then, the export catalog is made available to the community via Web services and also provided to a number of data centers. Over the last few years there have been a number of interim releases from the GSC-II, primarily to consortium members for telescope operations. The latest public release, the GSC~2.3.2, contains 945,592,683 unique objects and is $\\approx 170$ GB in size. Details about the release content and access methods are available from the ST ScI and Osservatorio Astronomico di TOrino (OATo) Web sites.  In this paper we describe the construction, calibration, and overall quality of GSC~2.3.2 , the current catalog release (for simplicity, GSC~2.3 in the rest of this paper), which was investigated through internal tests and a number of external comparisons to suitable literature data.  \\subsection{Background} In the age of new-technology telescopes and space-based missions, it is useful to recall that an important part of observational astronomy has historically been the creation of catalogs containing the reference and (often) target objects required to support observing programs. From a historical viewpoint it is fascinating that astrometry---the oldest branch of astronomy---is so vital to the success of modern high-tech observatories. As the technological complexity (and cost) of building and operating telescopes has increased enormously over the last 20 years, so has the effort to provide the best scientific return for these investments. One of the important issues is the optimization of observing efficiency, which depends on the use of proper pointing and input catalogs as well as on the access to digitized versions of the optical sky surveys.  The need for a deeper, all-sky catalog was highlighted early by the creation of the first Guide Star Catalog (GSC-I) to satisfy the pointing requirements of the {\\it Hubble Space Telescope} ({\\it HST}). An in-depth description of the GSC-I catalog may be found in a set of three papers \\citep{1990AJ.....99.2019L,1990AJ.....99.2059R,1990AJ.....99.2082J}. The GSC-I was used for {\\it HST} observation planning as well as target acquisition and tracking and has proved to be very reliable for its intended purpose. In addition, since its publication on CD-ROM, this catalog had become widely used by many ground-based telescopes to speed up the process of finding guide stars.  Although the GSC-I has been used with great success operationally and scientifically, it became clear (even during its construction) that it was possible to improve its usefulness by addressing the known systematic calibration errors. Indeed, it was clear that an increase in scope to include multicolor, multi-epoch data would eventually lead to the requirements for a new much improved archive, the GSC-II. In addition to the risk of damaging the UV-sensitive MAMA detectors in the second generation of {\\it HST} instruments by O stars as faint as 19th magnitude, it was generally acknowledged that the proper motions of the guide stars during the proposed lifetime of {\\it HST} would pose a potential problem for the accurate pointing of the telescope during the latter years of its operation. It was also realized that a fainter archive with multicolors such as the GSC-II would address a number of other astronomical needs such as: supporting adaptive optics on the next generation of large telescopes, remote or queue scheduling capabilities, and improved detector instruments requiring color information of the target objects.  In order to be prepared for this, STScI negotiated access to the original plates of the Palomar Observatory Sky Survey (POSS-II) and, in partnership with the Anglo-Australian Observatory (AAO), undertook the Second Epoch Southern (SES) red survey with the UK Schmidt Telescope Unit (UKSTU).  These new surveys, when combined with the plate material used in the GSC-I and the earlier POSS-I and UKSTU surveys, provide the material to generate a GSC-II based on, at least, two epochs and three passbands.  \\begin{figure*}[t] \\centering \\includegraphics[scale=.70,angle=0]{gsc2paper_filter.ps} \\caption{Transmission curves of the photographic passbands $B_J, R_F,$ and $I_N$ for the Palomar (solid lines) and AAO (dashed line) Schmidt surveys, compared to the Johnson--Kron--Cousins  BVRI$_{\\rm c}$ filters. \\label{fig:photometry:filters}} \\end{figure*}   \\subsection{The GSC-II Project} STScI and OATo first began a collaboration to start developing GSC-II in 1989. STScI was primarily motivated by telescope operations, and OATo was interested in the scientific applications of this catalog for galactic structure.  Once the project was underway, additional resources were contributed by ESO and GEMINI (who wished to use GSC-II for VLT and GEMINI telescope operations respectively), ST-ECF (as part of the NASA-ESA {\\it HST} funding agreement), and the Astrophysics Division of ESA (for science projects). A significant fraction ($\\sim 30\\%$) of the plate processing was performed at ST-ECF using the same pipeline software system described later. ",
        "conclusions": "In Table \\ref{GSCI-IIreqs}, we compare the requirements provided in the GSC-II {\\it implementation plan}---originally foreseen to reach magnitude 18 in $V$---with the actual performance of the current version of the catalog down to the same magnitude limit. In all items we have met or exceeded the specifications except for the proper motions.  Table \\ref{GSCI-IIspecs} summarizes the main characteristics of GSC~2.3, which essentially superseeds the requirements in that it reaches 2-3 mag depeer. The astrometric and photometric errors reported in the table are derived from those of Tables \\ref{tab:GSC23_SDSS_astro} and \\ref{tab:GSC23_SDSS_F}, and positional uncertainties are obtained by summing in quadrature the error in each coordinate. A consideration of the global statistics shows that even with the inclusion of fainter objects we are able to meet the original specifications.  Whilst no further improvements are required (or funded) for its use in {\\it HST} operations, we intend to continue development to produce the best possible catalog for scientific as well as operational uses. Our future plans include an astrometric recalibration using the UCAC~3 and applying a magnitude-dependent correction; this should not only reduce the absolute positional errors but amend the systematic error that is affecting the computed proper motions. A photometric re-reduction of all plates will also benefit from the improved overall quality of the available calibrating sequences. Finally, we plan to reclassify the objects splitting the nonstar classification into galaxy and blend as well as better determine the magnitudes of galaxies using an algorithm that does not assume a stellar profile \\citep{2007ApJS..170...33P}.  As of cycle 15, GSC~2.3 is the guide catalog for {\\it HST}, and it is a reference catalog for the VLT and Gemini adaptive optics programs, where real guide stars are still preferred to laser ``guide stars''; it has also been provided to the {\\it Chandra}, {\\it Galex}, {\\it XMM-Newton} and {\\it Swift} missions. In the preparation for the Gaia mission, the current catalog version is used as a snapshot of what Gaia is expected to observe \\citep{2006MmSAI..77.1172D} and as a base for the Initial Gaia Source List being compiled by OATo for the Gaia data reduction. Finally, GSC-II future release(s) will be at the heart of the Astrometric Support System of the very ambitious Large Sky Area Multi-Object Fiber Spectroscopic Telescope (LAMOST) undertaken by the Chinese Academy of Science, and will be part of the guide system for JWST \\citep{Spagna2001,2003AAS...202.0410S}.  In addition to the operational uses of the catalog itself, the GSC-II database has also been mined for many scientific studies, such as planetary nebulae \\citep{2004A&A...420..207K}, young open clusters \\citep{2000A&A...357..460S}, halo white dwarfs \\citep{2006A&A...448..579C}, peculiar objects \\citep{2001Natur.413..139M,2002A&A...393L..45C},  X-ray pulsars \\citep{Panzera2003}, and the structure and kinematics of stellar populations of our Galaxy \\citep{2006A&A...451..125V,2007MNRAS.375.1381K}.   Finally, it is worth going back to the density map presented in Figure \\ref{fig:skymap}, as it does illustrate GSC~2.3: it provides accurate positions, magnitudes in three bands, and stellar classification for 4$\\pi$ steradians to levels of completeness commensurate with the underlying stellar density. Future versions extracted from the GSC-II II database will only improve on this situation, including proper motions and increasing the catalog accuracy as a result of improvements in reference catalogs and reduction procedures."
    },
    "0807/0807.0705_arXiv.txt": {
        "abstract": "Within the next five years, a number of direct-imaging planet search instruments, like the {\\it VLT SPHERE} instrument, will be coming online.  To successfully carry out their programs, these instruments will rely heavily on a-priori information on planet composition, atmosphere, and evolution.  Transiting planet surveys, while covering a different semi-major axis regime, have the potential to provide critical foundations for these next-generation surveys.  For example, improved information on planetary evolutionary tracks may significantly impact the insights that can be drawn from direct-imaging statistical data.  Other high-impact results from transiting planet science include information on mass-to-radius relationships as well as atmospheric absorption bands.  The marriage of transiting planet and direct-imaging results may eventually give us the first complete picture of planet migration, multiplicity, and general evolution. ",
        "introduction": "{\\underline{{\\bf Ground-Based}}}  {\\it -  HiCIAO}: An adaptive optics coronagraphic simultaneous-differential-imager for the Subaru 8.2 meter telescope (Tamura et al. 2006);  currently in commissioning phase.  {\\it -  Project 1640}:  An adaptive optics coronagraphic integral field spectrograph for the Palomar 5 meter Telescope (Hinkley et al. 2008);  currently in commissioning phase.  {\\it -  SPHERE}:  An adaptive optics coronagraphic simultaneous-differential-imager and integral field spectrograph for one of the VLT 8.2 meter telescopes (Beuzit et al. 2008);  commissioning expected 2011.  {\\it -  GPI}:  An adaptive optics coronagraphic integral field spectrograph for the Gemini South 8 meter telescope (Macintosh et al. 2006);  commissioning expected 2011.  {\\underline{{\\bf Space-Based}}}  {\\it Terrestrial Planet Finder Coronagraph\\footnote[1]{http://planetquest.jpl.nasa.gov/TPF-C}?}  {\\it Terrestrial Planet Finder Interferometer\\footnote[2]{http://planetquest.jpl.nasa.gov/TPF-I}/Darwin\\footnote[3]{http://www.esa.int/science/darwin}?}  \\vspace{4 mm} All of these systems will rely on results from transiting planet science to accurately interpret their observations. ",
        "conclusions": ""
    },
    "0807/0807.0475_arXiv.txt": {
        "abstract": "We present a method to simulate color, 3-dimensional images taken with a space-based observatory by building off of the established {\\em shapelets} pipeline.  The simulated galaxies exhibit complex morphologies, which are realistically correlated between, and include, known redshifts. The simulations are created using galaxies from the 4 optical and near-infrared bands (\\textit{B}, \\textit{V}, \\textit{i} and \\textit{z}) of the Hubble Ultra Deep Field (UDF) as a basis set to model morphologies and redshift.  We include observational effects such as sky noise and pixelization and can add astronomical signals of interest such as weak gravitational lensing.  The realism of the simulations is demonstrated by comparing their morphologies to the original UDF galaxies and by comparing their distribution of ellipticities as a function of redshift and magnitude to wider HST COSMOS data.  These simulations have already been useful for calibrating multicolor image analysis techniques and for better optimizing the design of proposed space telescopes. ",
        "introduction": "As astronomical surveys become deeper and wider, analysis techniques correspondingly become more complex and demanding. To calibrate these methods, extensive work has already been invested in the simulation of monochromatic astronomical imaging. Simulation packages have been developed to incorporate a semi-analytic model of galaxy number counts and evolution \\cite{EWBM}, or to mimic the properties of real observations \\cite{MRCB}. However, there are currently no packages able to create correlated images across several bands.  Multi-band image simulations are firstly useful to develop and calibrate analysis methods that use multicolor data. Many measurements in astronomy  (for example photometry, astrometry and shape measurement) are  ``inverse problems,'' where variation in a signal is easy to introduce but difficult to measure, usually due to complications involving observational seeing and noise. Simulated data provide the best way to calibrate such methods because these variables can be controlled. A known astronomical signal can be inserted into simulated data, and the accuracy of a method can be judged by examining any errors in its recovery.  One example is the measurement of weak gravitational lensing. In weak lensing, light from background galaxies is lensed by foreground matter distributions, causing a shear (distortion) of the background galaxies' shapes. The distortion is easy to add during the construction of simulated data. Although the lensing signal is achromatic, color simulations can be used to test sophisticated measurement methods that take advantage of \\begin{itemize} \\item the increased number of shear-measurable galaxies if some galaxies are only sufficiently bright in certain bands, \\item reduced {\\em noise} on shear measurement  (by $\\sqrt{N}$) if the intrinsic shapes of galaxies are uncorrelated between $N$ bands, and \\item reduced systematic {\\em bias} on shear measurement if the intrinsic shapes of galaxies are correlated between bands \\cite{JainJarvis}. \\end{itemize} One common challenge in weak lensing measurement is the deconvolution of galaxy shapes from the instrumental point-spread function (PSF). Since the PSF is different in each band, PSF-dependent biases will be averaged out by looking at multiple  bands. Conversely, biases inherent to a method will not be ameliorated. Developing multicolor analysis techniques to exploit these tricks requires multicolor simulations.  Multi-band image simulations are also useful to optimize the design and improve the science case for planned, multi-band imaging surveys such as SNAP \\cite{Aldering} or Euclid \\cite{DUNE}.  These surveys require multiple bands in order to observe different types of galaxies, observe objects typically obscured in other bands, observe objects out to different redshifts, and most importantly to obtain photometric redshifts for galaxies. Engineering requirements for the design of these instruments can be derived via image simulations by measuring the (often complex and subtle) effects of engineering parameters on scientific return \\cite{High}. Predictions for the scientific return of a given mission can be similarly estimated \\cite{WLII,WLIII}.  A full demonstration of the potential gains in multicolor shear measurement, or a full optimization of a future space-based lensing mission is beyond the scope of this paper. The purpose of this paper is to present a method for simulating deep, multi-color space-based images with correlated morphologies and redshifts. The simulation pipeline we present here will serve as a basis for performing the optimization of both shear measurement techniques and future space missions in future papers.  Our simulation pipeline generalizes the single-color method of \\cite{MRCB}, representing complex galaxy morphologies as ``shapelets'' \\cite{Cartesian,Polar}. Shapelet-based simulations are already widely used for weak lensing. The Shear TEsting Program (STEP) used similar simulated data to test and improve shape measurement and PSF correction methods \\cite{Mass2006}. Our generalization to multi-band, 3-dimensional simulations thus increases the realism and utility of a well established technique.  This paper is organized as follows.  In \\S 2 we give a brief review of shapelets and how they can be used to generate simulated images.  In \\S 3 we present the methodology by which we create multi-band, 3-dimensional simulations.  In \\S 4 we test the realism of our simulations through comparison to the real HST data. Lastly, in \\S 5, we discuss the conclusions and summarize our findings. ",
        "conclusions": "We presented a method to create an arbitrary amount of 3-dimensional, color, simulated, unique deep space images.  The simulations are created by perturbing a galaxy's polar shapelet coefficients in such a way as to create unique but realistic objects.  The previous simulation pipeline has been expanded to correlate morphologies across four wavelength bands and now also include a redshift distribution.  Though we currently use four wavelength bands, our simulation pipeline is flexible enough to include an arbitrary number of colors, should such a data set become available and useful.  Our simulations were tested by comparing them to the original UDF images.  They were found to have similar morphologies to the original galaxies.  Additionally, the weak lensing cosmological properties of the simulations were tested against COSMOS HST data.  Within reasonable error, our simulations were found to be consistent."
    },
    "0807/0807.2646_arXiv.txt": {
        "abstract": "{} {For the first time, the lunar occultation technique has been employed on a very large telescope in the near-IR with the aim of achieving systematically milliarcsecond resolution on stellar sources. } { We have demonstrated the burst mode of the ISAAC instrument, using a fast read-out on a small area of the detector to record many tens of seconds of data at a time on fields of few squared arcseconds. We have used the opportunity to record a large number of LO events during a passage of the Moon close to the Galactic Center in March 2006. We have developed and employed for the first time a data pipeline for the treatment of LO data, including the automated estimation of the main data analysis parameters using a wavelet-based method, and the preliminary fitting and plotting of all light curves. } { We recorded 51 LO events over about four hours. Of these, 30 resulted of sufficient quality to enable a detailed fitting.  We detected two binaries with subarcsecond projected separation and three stars with a marginally resolved angular diameter of about 2 milliarcseconds. Two more stars, which are cross-identified with SiO maser, were found to be resolved and in one case we could recover the brightness profile of the extended emission, which is well consistent with an optically thin shell. The remaining unresolved stars were used to characterize the performance of the method. } { The LO technique at a very large telescope is a powerful and efficient method to achieve angular resolution, sensitivity, and dynamic range that are among the best possible today with any technique. The selection of targets is naturally limited and LOs are fixed-time events, however each observation requires only a few minutes including overheads. As such, LOs are ideally suited to fill small gaps of idle time between standard observations. } ",
        "introduction": "The method of observing lunar occultations (LO) of background stars to derive their angular diameter, as well as other characteristics such as binarity, has been employed for decades and provides angular resolution at the milliarcsecond (mas) level independently of the diameter of the telescope used. This surpasses the  diffraction limit of even the largest single telescopes and rivals the resolution of long-baseline interferometry (LBI) even with baselines of hundreds of meters.  Basically, the LO technique relies on the lunar limb as a diffracting edge, rather than on the diameter of the telescope. The diffraction fringes are generated in space, and due to their relatively fast motion over the telescope they must be sampled at rates of order 0.1-1\\,kHz. Combined together, these factors greatly diminish the adverse effects of atmospheric turbulence, which are the main limit of other high angular resolution techniques. LO also permit not only a model-dependent derivation of source parameters such as the angular diameter and binary parameters, but also a model-independent reconstruction of the brightness profile of the source according to maximum-likelihood, or even a unique reconstruction by light curve inversion in special cases. The time required for observation (dominated by overheads, since an occultation lasts much less than 1\\,s) and for data analysis is significantly shorter than for most other high angular resolution methods.  Of course, LO suffer from a number of significant drawbacks, first among them that the sources cannot be chosen at will. The Moon covers only about 10\\% of the celestial sphere in its apparent orbit. LO are fixed time events, and as such subject to weather and instrumental downtimes. Finally, the lunar limb only provides a 1-D scan of the source, although observations from different sites, or at different epochs, can in principle be combined under favorable circumstances. Due mainly to the chromatic properties of the scattered light background around the Moon and of the diffraction fringes, the near-IR is ideally suited to observe LO. More details on the method, its performance, the data analysis and the results can be found in Richichi (\\cite{AR_thesis}, \\cite{richichi96}) and Fors (\\cite{thesis}). The CHARM2 catalogue alone (Richichi et al. \\cite{charm2}) lists 1815 LO results in the field of high angular resolution.  While the angular resolution of LO is set mainly by the lunar limb rather than by telescope size, this latter parameter obviously has a crucial role in determining the limiting sensitivity. Richichi (\\cite{1997IAUS..158...71R}) studied the performance of LO under a number of circumstances, including the then-futuristic use of IR array detectors on a 8-m class telescope. This has recently become reality: Richichi et al. (\\cite{Messenger2006}) and Fors et al. (\\cite{SEA2007}, \\cite{fors08}) reported in a preliminary form on the use of the ISAAC instrument in the so-called burst mode at the ESO VLT 8.2\\,m Antu telescope.  In the present paper we provide a detailed account of the observations carried out during a few hours in the night of March 21, 2006, when 51 LO events were recorded. A second batch of observations was carried out in August 2006: due to their large number and their different nature, these sources will be discussed in a separate paper. In Sect.~\\ref{data} we describe the observations and we provide details of the data reduction. In Sect.~\\ref{results} we discuss the results, which include new binaries, resolved stars and the near-IR counterparts of two radio maser sources. We also characterize the performance of this specialized observing mode, which combines the highest angular resolution and sensitivity possible today in the near-IR. In the conclusions we describe the integration of the LO technique in the service mode operations scheme in place at the Paranal observatory and our plans for routine LO observations at that site. ",
        "conclusions": "We have presented the first LO results obtained using the so-called burst mode of the ISAAC instrument at the ESO Antu VLT telescope. This mode permits to record data streams  on a small subarray at high temporal frequency (3.2\\,ms, although mostly 6.4\\,ms were used for the present work). We recorded 51 occultation events over about four hours during a close approach of the Moon to the Galactic Center. The events were reappearances from behind the dark limb. This, coupled with initial problems of focussing and image quality due mostly to the commissioning nature of the run, resulted in several sources being missed or without sufficient quality. Nevertheless, 30 of the events resulted in light curves of good, or even excellent quality, including some of the best SNR traces ever recorded with this technique. We have revealed two binaries, three stars with a marginally resolved angular diameter of the order of 2\\,mas, and two resolved masing AGB stars which appear to be in the foreground of the galactic bulge. For one of these, {\\object 2MASS 17453224-2833429}, we were able to recover the brightness profile of the extended emission. This has the typical signature of an optically thin shell, with inner radius of order 10-15\\,mas, corresponding to about 40\\,AU. Follow-up studies are required to obtain further information on these stars, given that they are all heavily reddened and most have no known counterpart or literature entries.  Our analysis, including the unresolved sources, has established an unprecedented performance of the LO technique in the near-IR using a very large telescope and a fast read-out mode. We find that the angular resolution varies between 4 and 1\\,mas, and even less in the case of very high SNR. Data quality is generally impressive, compared to smaller telescopes, mostly due to the scintillation reduction effect of the large mirror. The dynamic range is also significantly improved with respect to previous observations, and we have shown that in some cases it would have been possible to detect a companion fainter than the primary by about 8 magnitudes in broad K band on angular scales comparable to the Airy disk of the telescope. This result compares favorably with other AO-assisted high contrast imaging techniques.  The random selection of sources to be occulted and the fixed-time nature of the events remain the main downsides of the LO technique. However, thanks to the very short telescope time required for each observation, LO are an ideal filler program for times when the telescope is idle or other programs cannot be conveniently executed. The service mode observations implemented at the VLT provide the right context for this, and we have obtained approval for such a filler program during ESO observing periods 80 and 81 (from October 2007 through September 2008). Several hundreds of OBs have been prepared for each period, and are in the queue awaiting execution whenever five minutes of telescope time become available."
    },
    "0807/0807.1703_arXiv.txt": {
        "abstract": "{MHD turbulence is known to exist in shearing boxes with either zero or nonzero net magnetic flux. However, the way turbulence survives in the zero-net-flux case is not explained by linear theory and appears as a purely numerical result that is not well understood. This type of turbulence is also related to the possibility of having a dynamo action in accretion discs, which may help to generate the large-scale magnetic field required by ejection processes.} {We look for a nonlinear mechanism able to explain the persistence of MHD turbulence in shearing boxes with zero net magnetic flux, and potentially leading to large-scale dynamo action.} {Spectral nonlinear simulations of the magnetorotational instability are shown to exhibit a large-scale axisymmetric magnetic field, maintained for a few orbits. The generation process of this field is investigated using the results of the simulations and an inhomogeneous linear approach. We show that quasilinear nonaxisymmetric waves may provide a positive back-reaction on the large-scale field when a weak inhomogeneous azimuthal field is present, explaining the behaviour of the simulations. We finally reproduce the dynamo cycles using a simple closure model summarising our linear results.} {The mechanism by which turbulence is sustained in zero-net-flux shearing boxes is shown to be related to the existence of a large-scale azimuthal field, surviving for several orbits. In particular, it is shown that MHD turbulence in shearing boxes can be seen as a dynamo process coupled to a magnetorotational-type instability.} {} ",
        "introduction": "The problem of angular momentum transport is a central issue of accretion disc theory. Following \\cite{SS73}, angular momentum transport is often modelled assuming the disc is turbulent, using a kind of turbulent viscosity (the so-called $\\alpha$ model). However, the way discs may become turbulent is still a highly debated subject. A possible route to turbulence is the magnetorotational instability \\citep[MRI,][]{BH91a} which appears in discs sufficiently ionised to be coupled with the magnetic field lines \\citep{G96}. This instability has been extensively studied in its nonlinear regime using local \\citep{HGB95,SHGB96} and global \\citep{H00} numerical simulations. More recently, the effect of non-ideal MHD has been investigated numerically in shearing boxes with either zero \\citep{FPLH07} or nonzero \\citep{FSH00,LL07} net magnetic flux, showing a strong dependence on the magnetic Prandtl number $Pm$.  These new results bring back the question of the efficiency of MHD turbulence in accretion discs since $Pm$ is believed to vary by several orders of magnitude in real objects \\citep{BH08}. In particular, as turbulence seems to disappear in zero-net-flux simulations for $Pm\\le 1$, the existence of a turbulent flow in low-$Pm$ objects such as protoplanetary discs is questionable. However, as pointed out by \\cite{FPLH07}, today simulations can reach Reynolds numbers that are very small (up to $10^4$) compared to those of real discs ($\\sim 10^{10}$). Therefore, no clear conclusion for accretion discs can be drawn from these numerical results.  It has been pointed out by \\cite{PCP07} and \\cite{FP07} that numerical resolution can also play an important role in simulations. In particular, increasing resolution in zero-net-flux simulations \\emph{without} any physical dissipation leads to a weaker turbulence and transport. One might conclude from these results that turbulence should disappear in real astrophysical systems \\citep{PCP07}. We think however that this conclusion is questionable since the ideal MHD model does not hold in turbulent flows, mainly because a dissipation scale necessarily exists, at which non-ideal effects are \\emph{not} negligible. Of course, ideal MHD simulations must include some sort of numerical dissipation in their algorithm, which then implicitly defines a dissipation scale of the order of the numerical grid scale. However, this algorithm-dependent dissipation is quite different from physical dissipation processes \\citep{LB07}, leading potentially to numerical artifacts for large-scale properties, such as turbulent transport.  On the other hand, the way MHD turbulence is sustained in these simulations is not well understood. It is known that when a mean vertical field is applied, the magnetorotational instability destabilises the flow and leads to developed three-dimensional turbulence \\citep{BH91a,HGB95}. This picture is not directly applicable to zero-net-flux simulations since the magnetic field can be dissipated either by a finite resistivity, or by the turbulence itself (see \\cite{HGB96} for an extensive discussion of this point). Therefore, even if one assumes that the MRI appears locally because of a given magnetic field configuration, one has to regenerate this field with some sort of dynamo mechanism. The whole process sustaining turbulence in zero-net-flux simulations is therefore a \\emph{nonlinear} mechanism, potentially involving a kind of magnetorotational instability at some stage. As already mentioned by \\cite{BH92}, this means that this disc dynamo is intrinsically nonlinear, as it requires a non-negligible Lorentz force, and cannot be described as a kinematic dynamo. Following these ideas, \\cite{ROP07} obtained a steady nonlinear solution to the MHD equations in the zero-net-flux case, using rigid conducting walls as radial boundary conditions. This solution is clearly a first step toward the understanding of the mechanism working in shearing boxes, although the boundary conditions and Reynolds numbers are significantly different compared to turbulent simulations. Note also that \\cite{BNST95} studied this dynamo process in discs, using boundary conditions allowing for mean flux variations. Although a azimuthal field was generated in their simulations, no physical understanding of the underlying process was provided.  In this paper, we describe a possible mechanism able to sustain MHD turbulence in the shearing box with zero net flux. First, we recall the resistive MHD equations in the shearing box, and the numerical method used to solve them. Next, we investigate the temporal evolution of some zero-net-flux simulations, and we exhibit a long-timescale cycle for the large-scale magnetic field. The origin of this cycle is described using a spectral decomposition, and we demonstrate that its origin is related to some properties of the nonaxisymmetric structures of the flow. We then study a linear theory of nonaxisymmetric waves in the presence of a large-scale magnetic field similar to the one observed in simulations. We show that this linear theory predicts the right properties for the nonaxisymmetric structures and partially explains the nonlinear cycle. We then provide a phenomenological closure model for the evolution of the large-scale field, based on our previous findings. We show that this model reproduces the basic behaviour of the cycles, and provides a mechanism able to sustain turbulence in dissipative MHD shearing boxes. The existence of a long-timescale cycle and large-scale structures in these simulations is the most significant finding of our investigation, mainly because it shows that MHD turbulence is able to generate large scales, independent of the dissipative scales. This mechanism can also be seen as a way to generate large-scale magnetic fields, providing a dynamo mechanism specific to accretion disc turbulence. These findings are discussed in detail in the concluding section, along with the implications for real astrophysical systems.  \\begin{figure*} \\centering \\includegraphics[width=1.0\\linewidth]{0152fg01.eps} \\caption{Time-evolution of $B_x$, averaged in the $x$ and $y$ directions. Notice the long-lived ($T\\sim 50\\, S^{-1}$) vertical structures, mostly found with long vertical wavelengths.} \\label{bxtm}% \\end{figure*} ",
        "conclusions": "\\subsection{Summary} In this paper, we have investigated the behaviour of the large-scale magnetic field in zero-net-flux simulations of the magnetorotational instability. We have first shown that the large-scale azimuthal field $B_x(z)$ is subject to a long-timescale oscillation when the flow is turbulent. Studying the induction equation, we have seen that these cycles are primarily due to an oscillation of a large-scale radial magnetic field $B_y$ and an azimuthal EMF $\\mathcal{E}_x$, whereas the radial EMF $\\mathcal{E}_y$ acts like a turbulent resistivity on $B_x$. Studying the nonaxisymmetric modes of the system, we found that the axisymmetric EMF is due mostly to the largest nonaxisymmetric structures, showing that the cycles are essentially a large-scale process. To understand the generation of the EMF, we have investigated the behaviour of nonaxisymmetric linear waves in the presence of an inhomogeneous and constant in time azimuthal magnetic field $B_x(z)$. This linear analysis explains most of the properties of the EMF observed in the nonlinear simulations. In particular, we have shown that the large-scale axisymmetric $\\mathcal{E}_x$ can be reversed for large $B_x$, explaining the cycles in the simulations. To summarise these results, we have considered a simplified closure model, encapsulating the linear properties. This model is able to reproduce the main behaviours of the cycles despite its simplicity, showing that the physics involved in the cycles is well described by our linear analysis.  \\begin{figure}[h!] \\centering \\includegraphics[width=1.0\\linewidth]{0152fg38.eps} \\caption{Scheme of the mechanism responsible for the large-scale magnetic field cycles. The dashed lines represent the nonlinear interactions which are not discussed in the present work. These interactions are nonetheless required to obtain a self-sustained mechanism (see text).} \\label{mechanism}% \\end{figure}   To summarise our findings, we sketch the mechanism responsible for the cycles in Fig.~\\ref{mechanism}. In this sketch, we have assumed that the nonaxisymmetric excitation (seed) is related to the nonlinear coupling of amplified shearing waves (dashed line). This specific process has not been studied in the present work and is just given here as a way to ``close the loop''. Note however that a more complicated mechanism may be involved in this back-reaction without changing the global picture. Moreover, we have described the contribution of $\\mathcal{E}_x(z)$ as a generic ``dynamo'' effect, but one should remember that this term may be either correlated or anticorrelated with $B_x$, as a function of the strength of $B_x$. Finally, we have omitted viscous and resistive effects for simplicity, as they have a small impact on the mechanism itself.  \\subsection{Comparison with previous works}  We would like to stress that the model presented here does not constitute a full description of a sustaining mechanism for MHD turbulence in discs. Indeed, we have assumed in the linear analysis and in the closure model that the flow is able to generate continuously small-amplitude shearing waves. As mentioned previously, this process involves a nonlinear coupling of nonaxisymmetric waves, which is beyond the scope of this work. However, a fully consistent self-sustaining process would require a description of this effect. According to \\cite{FPLH07}, MHD turbulence appears in zero-net-flux shearing-box simulations only when $Pm>1$. However, most of our analysis does not depend on the magnetic Prandtl number, and the dynamo process described here may work for arbitrarily low $Pm$. This suggests that the shearing-wave excitation mechanism depends on the Prandtl number, and may be too weak in the $Pm<1$ case.  This confirms that the excitation should be related to the small turbulent scales, which are still poorly understood.  Interestingly, our results indicate that the large-scale field generation mechanism is not destroyed for large Reynolds numbers, provided that the background flow is turbulent. Therefore, the same kind of mechanism might be at work in real astrophysical discs, generating a large-scale field in a sufficiently ionised and turbulent plasma. One should note however that the configuration used in our simulations is not comparable to the real geometry of an accretion disc. In particular, we do not know how the aspect ratio will modify the mechanism of the cycles, nor how the vertical stratification may enter this picture. However, as discovered by \\cite{BNST95} in the case of a stratified flow with nonperiodic vertical boundary conditions, one observes azimuthal field cycles with a long timescale.  These cycles have a significantly longer period ($T\\sim 200\\, S^{-1}$) compared to ours and involve a modification of the net azimuthal magnetic flux. Moreover, it has been shown by \\cite{BD97} that in this case, the cyclic behaviour may be explained by a classical $\\alpha$ effect proportional to the vorticity\\footnote{Note that this $\\alpha$ effect might be related to the presence of vertical stratification and the breaking of the reflectional symmetry}. Because of these fundamental differences, we cannot conclude that the mechanism responsible for all these cycles is the same without further investigation. Nevertheless, it seems that a more systematic investigation of this kind of long-timescale behaviour is required in order to better understand the way MHD turbulence sustains magnetic fields in various configurations.  As mentioned in the introduction, \\cite{PCP07} suggested that turbulent transport should vanish for large enough resolution and Reynolds number. Therefore, the turbulent transport associated with this mechanism is another important issue. As one can check easily in the numerical simulation and in the toy model, the large-scale $B_x$ and $B_y$ are roughly phase-shifted by $\\pi/2$, leading to a very small Maxwell stress ($\\langle B_x B_y\\rangle\\sim 10^{-4}$). Therefore, in our picture, most of the transport must come from the shearing waves, which have been shown to have nonzero Maxwell and Reynolds stresses. Unfortunately, we cannot put a precise value on the transport due to the shearing waves since it is controlled by their initial excitation. Therefore, this mechanism cannot give a precise answer on the expected transport in a real accretion disc. From a more phenomenological point of view, we expect the large-scale magnetic field strength produced by our mechanism not to depend on the Reynolds numbers if they are sufficiently large (see section \\ref{numlin}). Since this large-scale field couples with all the other modes available, we expect the large-scale nonaxisymmetric structures to have roughly the same amplitude as the large-scale field, in the limit of a fully developed turbulence. If this picture is correct, we would then expect a minimum transport of the order of a few times  $10^{-3}$, independently of the Reynolds numbers. Note however that this conclusion relies on the assumption that the large-scale field does not depend on the dissipation coefficients. We have shown that this assumption is plausible, although a precise analytical description of our linear analysis would be required to ascertain this statement.  The present description of an accretion disc dynamo is related to the steady self-sustaining solutions obtained by \\cite{ROP07} by numerical continuation methods. Both mechanisms involve a large-scale azimuthal magnetic field with zero net flux, which is generated through the shearing of a radial field.  Both also involve a nonaxisymmetric magnetorotational-type instability that regenerates the radial field through its nonlinear feedback.  However, the present mechanism is intrinsically local and unsteady, and may work for arbitrarily large Reynolds numbers whereas the steady, highly symmetrical solutions of \\cite{ROP07} depend on the existence of walls and are apparently restricted to low Reynolds numbers. Nevertheless, one may also think that these steady solutions correspond to some fixed points in phase space, around which the shearing box solutions oscillate, creating the dynamo cycles we have described. If this picture is correct, it would be an elegant way to unify these different approaches, and possibly to find other mechanisms of the same kind.  This work might also be compared to the shear dynamo simulations of \\cite{YH08}. However, we emphasise that our analysis involves a nonlinear dynamo, whereas the shear dynamo described by \\cite{YH08} is a kinematic (linear) dynamo, in which turbulence is \\emph{forced}. Moreover, the shear dynamo occurs in a non-rotating system in which the MRI does not appear.  Applying our linear analysis to such a system does not produce a $\\gamma$ effect that could explain the shear dynamo in either linear or nonlinear regimes.  Finally, we would like to stress that the problem of a non-linear dynamo, involving several linear instabilities, has also been described by \\cite{CBC03} in the context of magnetically buoyant flows. Although the underlying instabilities are somewhat different, they also found a cyclic behaviour which, as in the present work, cannot be fully described by a classical $\\alpha$ effect. This may indicate that all these instabilities share some common properties yet to be discovered, allowing for the development of a nonlinear dynamo.  \\subsection{Future work} As discussed previously, the main issue raised by our findings is how the turbulence is able to excite shearing waves. Interestingly, the same kind of problem arises in the case of the shearing box with a mean azimuthal field \\citep{BH92}. Therefore, a simpler way to study this effect is to investigate the way turbulence is sustained in the presence of a mean azimuthal field. In particular, an interesting test would be to check that in this case, the presence of turbulence also depends on $Pm$, in a similar way as in \\cite{FPLH07}. This would be a good indication that the large-scale dynamo process \\emph{itself} does not depend on $Pm$ at large $Re$, even though the underlying turbulent regeneration mechanism process does. Moreover, this idea would provide a new way to understand the $Pm$-dependence observed in MRI simulations with a nonzero mean vertical field \\citep{LL07}, and potentially to describe the asymptotic transport in the limit of high and low $Pm$.  On the other hand, the dynamo process itself requires further investigation. The linear analysis provided here has been computed using essentially numerical tools, as the analytical description of this problem is somewhat technical. This analytical description is nevertheless required, as it will give a precise and formal constraint on the $\\gamma$ effect and its dependence on the azimuthal field strength. It may also be a way to have a more physical understanding of the underlying mechanism responsible for the resistive and $\\gamma$ effects, leading to a potential generalisation of this mechanism to a wider class of flow. Finally, we would be able to check, at least linearly, if this dynamo process depends on the non-ideal effects, and potentially to confirm our conjecture."
    },
    "0807/0807.1535_arXiv.txt": {
        "abstract": "We present new multi-wavelength observations of the dwarf Seyfert 1 galaxy POX~52 in order to investigate the properties of the host galaxy and the active nucleus, and to examine the mass of its black hole, previously estimated to be $\\sim 10^{5}$ \\msun. {\\it Hubble Space Telescope} ACS/HRC images show that the host galaxy has a dwarf elliptical morphology ($M_I= -18.4$ mag, \\sersic\\ index $n=4.3$) with no detected disk component or spiral structure, confirming previous results from ground-based imaging. X-ray observations from both \\chandra\\ and \\xmmn\\ show strong (factor of 2) variability over timescales as short as 500 s, as well as a dramatic decrease in the absorbing column density over a 9 month period. We attribute this change to a partial covering absorber, with a 94\\% covering fraction and $N_{H} = 58^{+8.4}_{-9.2} \\times 10^{21}$ cm$^{-2}$, that moved out of the line of sight in between the \\xmmn\\ and \\chandra\\ observations. Combining these data with observations from the VLA, {\\it Spitzer}, and archival data from 2MASS and {\\it GALEX}, we examine the spectral energy distribution (SED) of the active nucleus.  Its shape is broadly similar to typical radio-quiet quasar SEDs, despite the very low bolometric luminosity of $\\lbol\\ = 1.3 \\times 10^{43}$ \\ergs. Finally, we compare black hole mass estimators including methods based on X-ray variability, and optical scaling relations using the broad \\hbeta\\ line width and AGN continuum luminosity, finding a range of black hole mass from all methods to be $\\mbh\\ = (2.2-4.2) \\times\\ 10^{5}$ \\msun, with an Eddington ratio of $\\lbol/\\ledd \\approx 0.2-0.5$. ",
        "introduction": "\\begin{figure*} \\begin{center} \\scalebox{0.25}{\\includegraphics{f1.eps}} \\end{center} \\caption{\\hst\\ ACS/HRC images of POX 52 and GALFIT models. The top and bottom rows show the F814W and F435W images, respectively. In each row, the left panel is the ACS image, the middle panel is the full GALFIT model, and the right panel shows the fit residuals. For the F435W filter, the double-\\sersic\\ host galaxy model is shown.} \\label{image-panel} \\end{figure*}  Stellar and gas dynamical studies have proven to be the most secure method of measuring the masses of supermassive black holes in nearby galaxies, and black holes with masses of $\\sim10^6~$--$~10^9$ \\msun\\ have now been detected in a few dozen galaxies (see Ferrarese \\& Ford 2005 for a recent review). With just a few exceptions, however, stellar-dynamical and gas-dynamical techniques for black hole mass measurement cannot be applied to Seyfert galaxies and quasars, because most luminous active galaxies are too distant for their black hole's gravitational sphere of influence to be resolved. Most black hole mass estimates in broad-lined active galactic nuclei (AGNs) are based on indirect methods that rely on scaling relations resulting from reverberation-mapping data \\citep{Kaspi, bentz06}. Using these methods, it has recently become possible to search for active galaxies having extremely low-mass black holes, in order to extend our understanding of black hole demographics to objects having black hole masses of $\\mbh \\lesssim 10^6$ \\msun.  The best nearby example of a black hole in the $<10^6$ \\msun\\ mass range is that in the dwarf Seyfert 1 galaxy NGC~4395 \\citep{FilS89,fh03}. NGC~4395 is an essentially bulgeless, late-type dwarf galaxy with type 1 Seyfert characteristics such as broad emission lines \\citep{FilS89, FHS93} and rapid X-ray variability \\citep{Iwa00, Shih03, Moran05}. Located at a distance of only $\\sim4.3$~Mpc \\citep{thim04}, it is an excellent target for multiwavelength observations despite the tiny luminosity of its central engine ($\\lbol \\approx 6\\times10^{40}$ \\ergs; Moran \\etal\\ 1999). A recent ultraviolet (UV) reverberation-mapping measurement found \\mbh~$=~(3.6~\\pm~1.1)~\\times~10^5$~\\msun\\ \\citep{Pet05}. \\citet{GH04, GH07c} searched the Sloan Digital Sky Survey (SDSS) archives to find AGNs with low-mass black holes and have found numerous examples of Seyfert 1 galaxies with $\\mbh < 10^6$ \\msun, which we will refer to as the GH sample. However, most of the SDSS objects found by Greene \\& Ho are much more distant than NGC 4395, making detailed studies of their nuclei and host galaxies more difficult. Although the GH sample does contain examples of AGNs in low-mass disk galaxies \\citep[see also][]{Dong}, NGC 4395 remains an exceptional case due to its extreme late-type morphology as well as its proximity.  \\object{POX 52}, also known as G1200-2038 or PGC 038055 ($D~=~93$ Mpc for $H_0= 70$ \\kms\\ Mpc\\per), is an interesting counterpart to NGC~4395. It was originally discovered in an objective prism study by \\citet{KSB87}, who classified it as a dwarf disk galaxy with a Seyfert 2 nucleus, but noted the presence of a weak broad component of the \\hbeta\\ line. \\citet{bar04} re-observed POX 52, classifying it as a type 1 AGN based on clear detections of broad \\hal\\ and \\hbeta\\ and noting that its optical emission-line spectrum is nearly identical to that of NGC 4395. Both NGC 4395 and POX 52 can be classified as narrow-line Seyfert 1 (NLS1) galaxies, since they satisfy the defining criterion of having FWHM(\\hbeta)$<2000$ \\kms. NLS1 galaxies are often, but not always, characterized by strong \\ion{Fe}{2} emission and relatively weak [\\ion{O}{3}] emission when compared to \\hbeta\\, along with steep soft ($0.5 - 2.0$ keV) spectral slopes and rapid variability in the X-ray \\citep{OP85, Boller96}. NGC~4395 and POX~52 are therefore somewhat atypical NLS1 galaxies in that they have very high-excitation spectra with large [\\ion{O}{3}]/\\hbeta\\ ratios and very weak \\ion{Fe}{2} emission.  Using the virial scaling relation calibrated by \\citet{Kaspi}, Barth \\etal\\ estimated a black hole mass of $\\mbh\\approx1.6\\times10^{5}$ \\msun, making this one of the least massive black holes identified in any AGN. This mass is consistent with the mass expected based on the \\msigma\\ relation (extrapolated to small \\sigmastar) for the measured stellar velocity dispersion $\\sigmastar = 36 \\pm 5$ \\kms. They also estimated the bolometric luminosity (\\lbol\\ $\\approx 2 \\times 10^{43}$ \\ergs) and the Eddington ratio (\\lbol/\\ledd\\ $\\approx 0.5 - 1$) for POX~52. Thus, although POX~52 and NGC~4395 have very similar black hole masses and emission-line spectra that are nearly identical in appearance, the AGN in POX~52 is $\\sim 300$ times more luminous than that in NGC~4395, which has an estimated \\lbol/\\ledd\\ $\\approx 10^{-3}$ \\citep{Pet05}. \\citet{bar04} used ground-based images to study the morphology of POX~52 and found that the host galaxy was best fit with a \\sersic\\ model with an index of $n = 3.6 \\pm 0.2$, effective radius of $r_{e} = 1\\farcs{2}$, and absolute magnitude $M_{V} = -17.6$. Combined with \\sigmastar\\ measured from the spectrum, these results show POX~52 to fall near the dwarf elliptical sequence \\citep{Geha03} in the fundamental plane projections of \\citet{Burstein97}, making it the first example of a dwarf elliptical to host a Seyfert nucleus. Although the past few years have seen a large increase in the number of known AGNs with low-mass black holes, POX~52 remains a valuable target for further study, because of its relatively small distance and the fact that it is a less luminous and presumably less massive galaxy than the majority of the host galaxies in the GH SDSS sample.  Among nearby dwarf ellipticals, there is little evidence for the presence of central black holes.  The closest well-studied example is NGC 205 in the Local Group, and stellar-dynamical observations with the \\emph{Hubble Space Telescope (HST)} have set an upper limit to the black hole mass of $\\mbh < 3.8\\times10^4$ \\msun\\ \\citep{Val05}, which is below the mass expected from a simple extrapolation of the \\msigma\\ relation. For dwarf elliptical galaxies outside the Local Group, stellar-dynamical observations lack the spatial resolution needed to detect black holes, and AGN surveys are the best available alternative for black hole searches. A double nucleus in one Virgo dwarf elliptical galaxy has been interpreted as morphological evidence for a stellar disk in orbit around a central black hole \\citep{deb06}, but aside from NGC~205 there are essentially no direct observational constraints on the black hole occupation fraction in dwarf elliptical galaxies. POX~52 represents a unique opportunity to study, in depth, the unambiguous presence of a low-mass central black hole in a nearby dwarf elliptical galaxy.  Here, we present new multiwavelength observations of POX 52, including the first targeted \\hst\\ and X-ray observations of this unusual object. \\hst\\ images are used to examine the host galaxy morphology; our results confirm the previous conclusions that the galaxy is consistent with a dwarf elliptical morphology with a \\sersic\\ index of $\\sim 4$. We describe X-ray observations obtained with both the \\chandra\\ and \\xmmn\\ X-ray Observatories in order to study the temporal and spectral properties of the active nucleus, and to constrain its bolometric luminosity and black hole mass. We also include observations from the Very Large Array (VLA) and {\\it Spitzer} Infrared Spectrometer (IRS), as well as archival data from 2MASS and \\emph{GALEX}, in order to investigate the spectral energy distribution (SED) of POX 52 over a wide range in frequency. We use recently updated versions of the broad-line virial scaling relations to obtain new estimates of the mass of the black hole, and in combination with the multiwavelength SED we determine the bolometric luminosity and Eddington ratio for the AGN. In addition to the \\hst\\ data presented here, we were also awarded time as part of the same project to obtain ultraviolet and optical spectra of POX 52 with the Space Telescope Imaging Spectrograph (STIS), but STIS failed before the observations were scheduled to be taken. ",
        "conclusions": "New \\hst\\ observations have confirmed POX 52 to be a dwarf elliptical galaxy, with a \\sersic\\ index of $n \\approx 4.3$, $r_e \\approx 1.15$, and $M_I = -17.3$. These parameters, along with the previously measured stellar velocity dispersion from \\citet{bar04}, show POX 52 to reside near the dwarf elliptical family in the fundamental plane. POX 52 was the first AGN found to reside in a dwarf elliptical galaxy, but more recent studies of low-mass AGNs and their host galaxies have found a number of similar objects \\citep{GH08}.  We see strong variability over $500 - 10^{4}$ s timescales with both \\chandra\\ and \\xmmn\\ and a substantial change in the absorbing column density due to partial covering over a 9 month time period. Data from the Optical Monitor on \\xmmn\\ showed no significant variability in the UV over short $\\sim 10^3 - 10^5$ s timescales. We see no variability within the $2\\sigma$ errors in the 7 month period between the OM and {\\it GALEX} NUV observations, but this does allow variations of $10-35$\\% on the $1\\sigma$ uncertainty level. The {\\it Spitzer} spectrum shows evidence of possible star-formation due to the PAH features and red continuum, but that analysis is deferred to a later paper.  Including data from the VLA, {\\it Spitzer}, \\emph{GALEX}, and 2MASS, we compile an extensive SED that shows POX 52 to be consistent with the radio-quiet template of more massive quasars and estimate a bolometric luminosity of \\lbol\\ $= 1.3 \\times 10^{43}$~\\ergs\\ by incorporating all of the available data. Using a variety of mass estimators, we find a range of black hole mass \\mbh~$\\approx (1.8-4.2) \\times 10^{5}$ \\msun. We find that we are unable to determine the break frequency from the \\chandra\\ PSD, either due to a low signal-to-noise ratio or insufficient temporal sampling. If $\\nu_B \\approx 10^{-4}$ Hz, then our mass estimate from the X-ray excess variance of $\\mbh = (3.9 \\pm 1.7) \\times 10^5$ \\msun\\ is consistent with the optical mass estimates. Otherwise, new X-ray observations of POX 52 in an unabsorbed state with the larger effective area of \\xmmn\\ would be needed in order to determine $\\nu_B$ and estimate the black hole mass with this method.  We have shown POX 52 to have an SED consistent with a scaled-down version of more massive type 1 quasars and that it is in a different accretion state than NGC 4395, as seen in the two orders of magnitude disparity between Eddington ratios. The number of low-mass AGNs known is increasing with new results from large AGN surveys, such as \\citet{GH04, GH07c}, but POX 52 remains unique in that it is one of the nearest and best studied of the low-mass AGNs."
    },
    "0807/0807.2911_arXiv.txt": {
        "abstract": "We extend the Hubble diagram up to $z = 5.6$ using 63 gamma-ray bursts (GRBs) via peak energy-peak luminosity relation (so called Yonetoku relation), and obtain constraints on cosmological parameters including dynamical dark energy parametrized by $P/\\rho\\equiv w(z) = w_0 + w_a \\cdot z/(1+z)$. It is found that the current GRB data are consistent with the concordance model, ($\\Omega_m = 0.28, \\Omega_{\\Lambda} = 0.72, w_0 = -1, w_a = 0$), within two sigma level. Although constraints from GRBs themselves are not so strong, they can improve the conventional constraints from SNeIa because GRBs have much higher redshifts. Further we estimate the constraints on the dark-energy parameters expected by future observations with GLAST (Gamma-ray Large Area Space Telescope) and \\swift by Monte-Carlo simulation. Constraints would improve substantially with another 150 GRBs. ",
        "introduction": "\\label{sec:intro}  In our previous paper \\citep{Kodama2008}, we calibrated the peak energy-peak luminosity relation of GRBs with 33 nearby events ($z < 1.62$) whose luminosity distances were estimated from those of large amount of SNeIa \\citep{Riess2007,Wood-Vasey2007,Davis2007}. This calibrated Yonetoku relation, derived without assuming any cosmological models, can be used as a new cosmic distance ladder toward higher redshifts. Then we determined the luminosity distances of 30~GRBs in $1.8 < z < 5.6$ using the calibrated relation and calculated the likelihood varying $(\\Omega_m, \\Omega_{\\Lambda})$. We obtained $(\\Omega_m, \\Omega_{\\Lambda}) = (0.37^{+0.14}_{-0.11}, 0.63^{+0.11}_{-0.14})$ for a flat universe, which is consistent with the concordance cosmological model within one sigma level. Our logic to obtain a new distance ladder is similar to that for SNeIa, that is, we calibrate a new distance indicator (Yonetoku relation) at low redshifts ($z < 1.62$) using the well established indicators (SNeIa). Then we assume the new relation holds at high redshifts ($z > 1.62$), although more detailed analysis is needed for possible selection bias and evolution effects in the relation \\citep{Oguri2006}.  Currently the number of GRBs with $z > 1.62$ is relatively small ($\\sim 30$) and the statistical and systematic errors of the Yonetoku relation are not so small compared with SNeIa. Even then, GRBs are still effective to probe dark energy, especially for dynamical dark energy whose energy density becomes large at high redshifts, because the mean redshift of GRBs is higher than that of SNeIa \\citep{Amati2008,Liang2008,Basilakos2008, Schaefer2007,Ghirlanda2006,Firmani2006a}. There would be many ways to characterize the time variation of dark energy. Here we adopt a simple phenomenological model as \\citep{Chevallier2001,Linder2003} $P/\\rho \\equiv w(z) = w_0 + w_a \\cdot z/(1+z) = w_0 + w_a (1-a)$, where $w_0$ and $w_a$ are constants and $a$ is the scale factor of the universe. For this model, GRBs would give strong constraints on $w_a$ which represents the time dependence of dark energy.  In this Letter, we present constraints on cosmological parameters such as $w_0$ and $w_a$ using both SNeIa with $z < 1.8$ and GRBs with $z < 5.6$. In \\S~2, we briefly review the main results of the previous paper \\citep{Kodama2008} and the theoretical and observational basis of the analysis. In \\S~3-1, we assume cosmological constant ($w = -1$) and obtain constraints on ($\\Omega_m, \\Omega_{\\Lambda}$). In \\S~3-2, we assume a flat universe and non-dynamical ($w_a = 0$) dark energy and obtain constraints on $(w_0, \\Omega_m)$. In \\S~3-3, we fix $\\Omega_m = 0.28$ for simplicity, and obtain the plausible values of $w_0$ and $w_a$. In \\S~4, we discuss how these constraints will be improved in future observations of high redshift GRBs by such as GLAST. Throughout the paper, we fix the current Hubble parameter as $H_0 = 66~{\\rm km~s^{-1}Mpc^{-1}}$. ",
        "conclusions": ""
    },
    "0807/0807.2250_arXiv.txt": {
        "abstract": "Inelastic dark matter, in which WIMP-nucleus scatterings occur through a transition to an excited WIMP state $\\sim$ 100 keV above the ground state, provides a compelling explanation of the DAMA annual modulation signal.  We demonstrate that the relative sensitivities of various dark matter direct detection experiments are modified such that the DAMA annual modulation signal can be reconciled with the absence of a reported signal at CDMS-Soudan, XENON10, ZEPLIN, CRESST, and KIMS for inelastic WIMPs with masses $O(100 \\;\\gev)$.  We review the status of these experiments, and make predictions for upcoming ones. In particular, we note that inelastic dark matter leads to highly suppressed signals at low energy, with most events typically occurring between 20 to 45 keV (unquenched) at xenon and iodine experiments, and generally no events at low ($\\sim 10$~keV) energies. Suppressing the background in this high energy region is essential to testing this scenario. The recent CRESST data suggest seven observed tungsten events, which is consistent with expectations from this model.  If the tungsten signal persists at future CRESST runs, it would provide compelling evidence for inelastic dark matter, while its absence should exclude it. ",
        "introduction": "The DAMA collaboration has observed an annual modulation in the nuclear recoil rates of their experiment with a confidence level that now exceeds $8\\sigma$ \\cite{Bernabei:2008yi}. This observation is consistent with weakly interacting massive particles (WIMPs) in the halo striking the target nuclei slightly less and more often as the Earth moves with and against the WIMP wind. The WIMP interpretation has been sharply criticized due to its apparent inconsistency with the null results of several other direct detection experiments, including CDMS and XENON. In this paper, we demonstrate that inelastic dark matter \\cite{TuckerSmith:2001hy} continues to offer a compelling explanation of both the annual modulation signal as well as the lack of a reported signal at CDMS, XENON, ZEPLIN, KIMS, CRESST  and other experiments.   Inelastic dark matter (iDM) \\cite{TuckerSmith:2001hy} was originally proposed to reconcile the DAMA annual modulation observation and null results from CDMS. The basic model is a simple extension of the standard WIMP model. An inelastic WIMP has two basic properties: \\begin{itemize} \\item{In addition to the dark matter particle $\\chi$, there exists an excited state $\\chi^*$, with a mass $m_{\\chi^*}-m_{\\chi} = \\delta \\approx  \\beta^2 m_\\chi \\sim 100$~keV heavier than the dark matter particle.} \\item{Elastic scatterings off of the nucleus, i.e., $\\chi N \\rightarrow \\chi N$ are suppressed, compared with the inelastic scatterings $\\chi N \\rightarrow \\chi^* N$}. \\end{itemize} The splitting $\\delta$ is comparable to the kinetic energy of a WIMP in the halo. This causes the kinematics of the scattering process to be significantly modified compared with a WIMP that scatters elastically.  The altered scattering kinematics leads to a fundamental difference in how an inelastic WIMP shows up in direct detection experiments:  only those with sufficient kinetic energy to upscatter into the heavier state will scatter off nuclei. The minimum velocity to scatter with a deposited energy $E_R$ is \\begin{eqnarray} \\beta_{\\rm min} &=& \\sqrt{\\frac{1}{2 m_N E_R}} \\left( \\frac{m_N E_R}{\\mu} + \\delta \\right), \\label{eq:betamin} \\end{eqnarray} where $m_N$ is the mass of the target nucleus and $\\mu$ is the reduced mass of the WIMP/target nucleus system.  This minimum velocity requirement  means that experiments tend to probe only the higher velocity region of the WIMP halo velocity distribution.  This simple modification leads to three key  features that change the relative sensitivities of dark matter experiments. Specifically \\cite{TuckerSmith:2001hy,TuckerSmith:2004jv}, \\begin{itemize} \\item{ {\\it Heavier targets are favored over lighter targets} - For a given $E_R$ and $\\delta$, a range of velocities in the halo will be accessible to a given dark matter experiment.  \\Eref{eq:betamin} shows that the range available for  heavier target nuclei will be greater than that for lighter targets. Because the WIMP velocity distribution  is expected to fall rapidly above the peak velocity of $\\sim 220\\; {\\rm km/s}$, targets with heavier nuclei tend to have greater sensitivity than those with lighter nuclei.} \\item{{\\it Modulation of the signal is significantly enhanced} - For conventional WIMPs, the modulation is typically of the order of a few percent. For inelastic WIMPs, because one is generally sampling a higher velocity component of the WIMP velocity distribution, where numbers are changing rapidly as a function of velocity, the modulation can be much higher. In the extreme limit, it can be such that there are particles available to scatter at DAMA/LIBRA in the summer, but not in the winter, yielding a $100 \\%$ modulation of the signal.% } \\item{{\\it The spectrum of events is dramatically changed, suppressing or eliminating low-energy events} - The spectrum of conventional WIMPs rises exponentially at low energy. As a consequence, significant sensitivity gains can be achieved by pushing the energy threshold of the experiment to ever lower energy. With inelastic DM, however, the kinematical changes cause the signal to be suppressed or even eliminated at sufficiently low recoil energies. The result is a spectrum which can peak at 20 keVr and above.} \\end{itemize} In this paper we demonstrate that together, these effects can reconcile DAMA/LIBRA with the null results of other experiments. The CRESST experiment, with its tungsten target, has a robust sensitivity. We shall see that the seven tungsten events seen there are consistent with expectations from this scenario.  Before proceeding to quantitative studies of the different experiments, we consider qualitatively how the above differences from conventional WIMPs reconcile the various experimental results. As a first example, we consider the limits of CDMS (Ge) on DAMA. As germanium is significantly lighter than iodine, the sensitivity of CDMS is considerably reduced. Indeed, for some regions of parameter space, no WIMPs  in the halo are capable of inelastic scattering at CDMS, while some are  capable of scattering at DAMA/LIBRA. In practice, there is some residual sensitivity in CDMS to inelastic WIMPs at the high-velocity edge of the halo distribution.  This suggests that bounds on inelastic WIMPs are rather sensitive to the velocity cutoff; we treat this carefully below.  In contrast to CDMS, the XENON experiment should be sensitive to inelastic nuclear recoils in a similar way to the DAMA experiment due to the close proximity of the respective target nucleus masses (Xe and I).  However, here there are two crucial differences between the two experiments' reported results:  (1)  the XENON experiment's energy threshold is far smaller than DAMA's (unquenched) energy threshold, and (2) XENON took their data between October 6 to February 14, which happens to be roughly during the time when the Earth travels \\emph{with} the galactic WIMP wind, thus reducing average WIMP velocity.  Due to a combination of effects from kinematics and nuclear form factors, conventional WIMPs have a rapidly falling spectrum as a function of energy. This has led to a significant push to lower the energy threshold of experimental searches, in order to optimize sensitivity. XENON's bound on the cross section, for instance, is especially strong because of the dearth of events from 4.5 - 14.5 keV.  By comparison, interpreted as unchanneled iodine scatters, the unquenched energy range for the DAMA annual modulation signal is roughly 22-66 keV.  Hypothesis-testing approaches for elastic WIMPs can achieve very stringent limits, even in the presence of background events at high energy, if the low energy region is under good control.  In contrast, an inelastic WIMP that leads to an annual modulation in DAMA often recoils off nuclei with energies in a range that is nearly \\emph{orthogonal} to the range used by XENON to set their bounds on elastic scattering WIMPs. This is particularly interesting as XENON's blind search identified nine candidate events between 14.5 keV to 26.9 keV, and (outside their analysis range) another fourteen candidate events between 26.9 keV to 45 keV.  This spectrum is certainly inconsistent with the distribution of nuclear recoils expected from an elastic scattering WIMP.  Yet, we demonstrate that inelastic WIMPs with masses $O$(100 GeV) can be made consistent with the DAMA annual modulation, leaving no trace in the low energy bins of XENON, while contributing between a few to tens of nuclear recoil events in the XENON high energy bins.  Thus, we find that XENON does not rule out inelastic dark matter precisely because they have twenty-three events in the signal region where inelastic WIMPs would be expected to be found.  Finally, the heightened sensitivity to the high-velocity part of the WIMP halo velocity distribution causes the modulation effect to be significantly enhanced. This makes the time of year of an experiment's run especially important. If the modulation is a few percent, the precise dates of data taking are of little consequence, but in the presence of $O$(1) modulation, it can clearly change the limits. In the extreme case, where at some times of the year there are no dark matter particles capable of scattering, it can remove all sensitivity of (in this case) XENON to DAMA signals.  These qualitative arguments are borne out by more careful analyses. The layout of the paper is as follows: in section \\ref{sec:indep} we discuss, without reference to halo models, the relative implications of unmodulated signals to the DAMA modulation signal, and show that consistency of the DAMA signal arising from quenched iodine scatters requires a model with low-energy events suppressed, and modulation enhanced well beyond the standard $\\sim 2\\%$ level expected for a standard WIMP. In sections \\ref{sec:models} and \\ref{sec:kine}, we discuss models for iDM and the kinematical implications of iDM at various experiments, which include such low-energy suppressions and enhanced modulation. In section \\ref{sec:limits} we discuss the limits arising from various experiments, and find the allowed parameter space for these models. We find significant regions of parameter space open, with masses $\\sim 50-200\\; \\gev$, which are otherwise closed in the absence of inelasticity. Finally, in section \\ref{sec:discussion}, we consider the implications of the limits for current and future experiments, and conclude. ",
        "conclusions": "\\label{sec:discussion} In light of the recent results from DAMA/LIBRA, it is important to continue to consider whether simple, viable models of dark matter exist which can reconcile the presence of the annual modulation signal with the lack of a reported signal from the many other experiments. Most recent analyses \\cite{Foot:2008nw,Feng:2008dz,Petriello:2008jj,Bottino:2008mf} have focused on the light ($\\sim$ few GeV) regions of parameter space, and have exploited the possibility of channeling \\cite{Bernabei:2007hw} in opening low mass windows.  In contrast, the mass range for inelastic dark matter is more typical of a standard WIMP. The inelastic scattering can arise quite simply from particle physics models with approximate symmetries, and so remains an appealing possibility, quite distinct from the low-mass window.  Of particular interest is the spectral data from DAMA/LIBRA. The low ($2-2.5\\; \\kev$) bin shows significantly lower modulation than what usually occurs for most WIMPs, where the lowest bin often (but not exclusively) shows the highest modulation. This spectral feature arises automatically with inelastic dark matter, and drives our numerical fits to large values of $\\delta$, where modulation is enhanced, and low energy rates at other experiments are highly suppressed.  We have seen that the three basic features of iDM, namely, preference for heavy targets, enhanced modulation, and suppression of low-energy events, allow the DAMA modulation to be consistent with other experiments. CDMS is suppressed due to its light target, while XENON is evaded by the suppression of events between 4.5 and 14.5 keV.  Interestingly, experiments such as ZEPLIN and KIMS, which are significantly weaker than XENON and CDMS for conventional WIMPs, are competitive in studies of inelastic dark matter. This is because they have target masses similar to iodine, and their energy range focuses on the range most relevant for studying the DAMA signal. In addition to studies similar to those presented here, studies of modulation in the presence of background events may prove quite effective, as the modulated fraction can be O(10\\%) with inelastic dark matter.  Of particular importance in this vein is CRESST. Aside from a xenon-target signal shifted to a higher energy range, the second robust prediction of iDM \\cite{TuckerSmith:2004jv} was the inevitable signal from tungsten scattering events. The seven events arising in the CRESST experiment are consistent with the rate expected from iDM. Should they persist in future exposure, this would be strong evidence for the inelastic nature of dark matter. Conversely, the absence of such events would exclude the inelastic explanation of the DAMA modulated signal.  .  Ultimately, this scenario makes two clear predictions: a signal rate in the 20-50 keVr range for iodine and xenon targets, and significant rates on tungsten targets, whose signal would naturally be peaked at 20-30 keV, and extending up to possibly 80 keV. Germanium targets still have tremendous reach in their next rounds, and such a signal would be peaked at high ($\\sim$ 70 keV) energies.  The nature of dark matter remains one of the most important questions in physics. We have seen here that the inelastic dark matter scenario continues to provide an explanation of the DAMA modulation, consistent with the  results of other experiments. The robust predictions of this scenario make it exceedingly testable, and consequently, the next generation of experiments should determine if the dark matter is inelastic.  \\vskip 0.15in \\noindent{{\\em Note added:} As this work was being prepared, \\cite{Petriello:2008jj} appeared, which also considered the inelastic scenario and concluded that the 50-200 GeV window for inelastic dark matter was excluded by CDMS and XENON. We disagree with this conclusion.  We believe the different results obtained in \\cite{Petriello:2008jj} may have arisen from the approximations made there for the halo escape velocity and for the form of the modulation.} \\vskip 0.2in \\noindent {\\bf Acknowledgments} The authors would like to thank Aaron Pierce for extensive and useful discussions. The authors would like to thank Kaixuan Ni for very insightful discussions on XENON, and providing us the information on nuclear recoils above 27 keV. The authors would further like to thank Sunkee Kim for consultation and information on the KIMS experiment. The authors would like to thank Pierluigi Belli for important information on the DAMA/LIBRA results, and Rafael Lang for discussions on CRESST. NW and GDK acknowledge the support of the KITP Santa Barbara, where some of this work was performed, and the support in part by the National Science Foundation under Grant No. PHY05-51164.  NW thanks the Aspen Center for Physics for their hospitality, where this work was completed. GDK was supported by the Department of Energy under contract DE-FG02-96ER40969. The work of DTS was supported by NSF grant 0555421. SC and NW were supported by NSF CAREER grant PHY-0449818 and DOE grant \\# DE-FG02-06ER41417.   \\vskip 0.15in"
    },
    "0807/0807.4476_arXiv.txt": {
        "abstract": "{ New \\HI observations of the nearby radio-loud galaxies Centaurus~A and B2~0258+35 show broad absorption ($\\Delta v_{\\rm{absorp}}\\sim 400$~\\kms ) against the unresolved nuclei. Both sources belong to the cases where blue- and redshifted absorption is observed at the same time. In previous Cen~A observations only a relative narrow range of redshifted absorption was detected. We show that the data suggest in both cases the existence of a circumnuclear disk. For Cen~A the nuclear absorption might be the atomic counterpart of the molecular circumnuclear disk that is seen in CO and H$_2$. Higher resolution observations are now needed to locate the absorption and to further investigate the structure and kinematics of the central region of the AGN and the way the AGN are fueled. ",
        "introduction": "The central region of AGN and the way AGN are fueled are important topics in current extra-galactic astronomy. Mergers have always been considered important for the triggering of AGN by providing the mechanism that could bring gas to the central regions and fuel the black hole \\citep[e.g.][]{hibbard96,mihos96,barnes02}. While this seems to be likely in powerful radio galaxies, recent observations show that other AGN - e.g. low-power radio galaxies - exist where there is no evidence for accretion through mergers. The activity in these galaxies appears to be associated with the slow accretion of gas from the ISM/IGM \\citep{best05,balmaverde08}.\\\\ A powerful tool to learn more about the structure of the central regions of AGN and the fueling mechanism is the study of the distribution and kinematics of the gas. In order to perform detailed studies of single objects we have selected two \\Hi -rich radio-loud galaxies, the merger remnant Centaurus~A and B2~0258+35 (NGC~1167). The later is unlikely to be a recent merger but has a young radio source suggesting that ``cold accretion'' from the IGM plays the dominant role by providing the fuel for the AGN activity (see Sect.~3).\\\\ Neutral hydrogen is an ideal tracer of accretion, interaction and merging events. \\HI in absorption has been often detected in radio sources \\citep[see e.g.][]{conway97,vermeulen03,morganti05}. These absorption structures have shown a variety of characteristics. \\HI absorption profiles only redshifted (relative to the systemic velocity), were initially found in a number of radio galaxies and interpreted as gas clouds that are falling into the nucleus, possibly indicating accretion of gas which deliver the fuel for the AGN \\citep[e.g.][]{vangorkom89}. However, more sensitive and broader band observations have shown that the picture is more complicated and that blueshifted absorption occurs even more often than redshifted absorption \\citep[e.g.][]{vermeulen03,morganti05}.\\\\ In other cases, the \\HI absorption is centered on the systemic velocity of the galaxy and these cases have been often interpreted as \\HI associated with a circumnuclear disk (or torus), see e.g. \\citet{conway97}, \\citet{peck01}. Support for the picture of a circumnuclear disk comes from theoretical work that indeed predict the existence of such circumnuclear \\HI disks around (active) black holes \\citep[see e.g.][]{maloney96,loeb08}. Thus, the picture is more complicated and the questions concerning the central structure and fueling remains. \\begin{figure*}[t!] \\includegraphics[height=4.6cm]{cena.eps} \\caption{\\footnotesize Centaurus~A: Left panel: \\HI absorption against the unresolved nucleus as derived from the ATCA observations \\citep{morganti08}. The total absorption width ranges from 400 to 800~\\kms . The arrow indicates the previously known (redshifted) absorption. The vertical line gives the systemic velocity. Right panel: Position-velocity plot of the \\HI (grey-scale and thin contours) and superimposed the CO emission (thick contours; from \\citet{liszt01}, taken along position angle 139\\degr . Note that the CO observations do not extend beyond a radius of about 1 arcmin. Figure taken from \\citet{morganti08}.} \\label{cena} \\end{figure*} ",
        "conclusions": "Using new, broad band observations of Cen~A and B2~0258+35 we have shown the presence of blue- and redshifted \\HI absorption against the unresolved nuclei. Both galaxies belong to the group of radio-loud galaxies where blue- and redshifted absorption is observed at the same time. Previously, for Cen~A only a narrow range of redshifted absorption was detected.\\\\ We have shown \\citep{morganti08} that the \\HI absorption in Cen~A is not evidence of gas infall into the AGN, but instead is due to a cold, circumnuclear rotating disk. The absorption profile of the WSRT data suggests the same interpretation for B2~0258+35. Both observations are in agreement with theoretical work which predicts the existence of such circumnuclear disks. However, these results leave the fueling question open.\\\\ Sensitive, higher resolution observations (VLBI for Cen~A and VLA A-array for B2~0258+35) are now needed (and proposed) to detect and locate the absorption in order to further explore the characteristics of the nuclear \\HI and to test the hypotheses of circumnuclear disks.\\\\ Both sources are unique objects. Cen~A is the closest AGN and therefore the linear scale is much better than for any other radio-loud source. Hence, the inner structure and fueling mechanism can be studied in great detail. Because the background structure of B2~0258+35 is extended around all directions of the black hole, higher spatial resolution data will allow to investigate the complete 2-D kinematics in the vicinity of the AGN. For both sources it will be interesting to investigate the presence of velocity gradients (which might be evidence for circumnuclear disks) as well as radial motions (infall/outflow?)."
    },
    "0807/0807.3060_arXiv.txt": {
        "abstract": "If there is a fifth force in the dark sector and dark sector particles interact non-gravitationally with ordinary matter, quantum corrections generically lead to a fifth force in the visible sector. We show how the strong experimental limits on fifth forces in the visible sector constrain the direct detection cross section, and the strength of the fifth force in the dark sector. If the latter is comparable to gravity, the spin-independent direct detection cross section must typically be $\\lesssim 10^{-55}$ cm$^2$.  The anomalous acceleration of ordinary matter falling towards dark matter is also constrained: $\\eta_{\\rm OM-DM} \\lesssim 10^{-8}$. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0807/0807.1912_arXiv.txt": {
        "abstract": "The interferometric technique known as peeling addresses many of the challenges faced when observing with low-frequency radio arrays, and is a promising tool for the associated calibration systems. We investigate a real-time peeling implementation for next-generation radio interferometers such as the Murchison Widefield Array (MWA). The MWA is being built in Australia and will observe the radio sky between 80 and 300 MHz. The data rate produced by the correlator is just over 19 GB/s (a few Peta-Bytes/day). It is impractical to store data generated at this rate, and software is currently being developed to calibrate and form images in real time. The software will run on-site on a high-throughput real-time computing cluster at several tera-flops, and a complete cycle of calibration and imaging will be completed every 8 seconds. Various properties of the implementation are investigated using simulated data. The algorithm is seen to work in the presence of strong galactic emission and with various ionospheric conditions. It is also shown to scale well as the number of antennas increases, which is essential for many upcoming instruments. Lessons from MWA pipeline development and processing of simulated data may be applied to future low-frequency fixed dipole arrays. ",
        "introduction": "\\label{Introduction}  The Murchison Widefield Array (MWA) is an 80-300 MHz synthesis array that is being built in Western Australia, with construction to be completed in 2010. The shire of Murchison has a quiet radio environment, making it an excellent site for this and other radio facilities \\cite{Bowman2007a}. Each of the 512 antennas will be a $4\\times4$ tile of dipoles. An analogue beamformer at each antenna combines the signals from the 16 dipoles, producing an electronically steerable primary beam with a width of approximately 25$^\\circ$ at 150 MHz. When the signals from all antennas are combined, the array will have a synthesized beam with a width of approximately 4.5$^\\prime$ at 150 MHz. The main science goals of the MWA are the detection of redshifted 21cm emission from the Epoch of Reionization (EoR) \\cite{Bowman2006}, transient detection (for example \\cite{Bhat2007}), and remote heliospheric sensing \\cite{Salah2005}. A schematic of two MWA antenna tiles is shown in Fig. \\ref{plot:schematic}.  To make a map of the sky using radio interferometry, one typically builds up an estimate of the 2D Fourier transform of the sky, then applies a Fourier transform to obtain the image. This is known as synthesis imaging, and a good overview of the subject is given in \\cite{Thompson-etal.2001}. The measured data that are used to build up the Fourier interference pattern are spatial cross-correlations -- or visibilities -- that are obtained by correlating voltage streams collected by many pairs of antennas. The visibilities lie in the uvw coordinate frame, where $w$ is the component of the antenna separation vector (or \\emph{baseline} vector) in the direction of the field center (in units of wavelengths), and $u$ and $v$ are orthogonal coordinates in the plane normal to $w$ (aligned with the corresponding image coordinate axes, $l$ and $m$). This situation is not naturally a 2D Fourier transform. For small images, $w$ is multiplied by a term that is approximately zero and the 2D nature holds. For large fields this is not the case, but the problem can still be reduced to 2D transforms (a good overview is given in \\cite{Cornwell2005}).\\footnote{The MWA will produce \\emph{snapshot} images, which means each antenna pair contributes a single visibility to each image. The MWA visibilities will be approximately coplanar, so a 2D Fourier relationship will hold even for wide-field images.} Post processing typically involves calibrating the visibilities, gridding them onto the uv-plane to form a regularly sampled interference pattern, and then applying a 2D FFT to form an image. Techniques such as self-calibration can then be used in an attempt to improve the calibration by iterating back and forth between the visibilities and the image.  For a number of reasons, MWA visibilities cannot be processed in this way. Many of these effects are common to all low frequency arrays, and are described in detail in \\cite{Erickson1999} and \\cite{Thompson-etal.2001}. Each antenna has a different direction-dependent response over the field of view, which cannot simply be divided out. These response patterns may also change significantly over the course of an observation. Furthermore, the ionosphere causes direction-dependent phase shifts that effectively change the position and polarization state of sources during an observation. These effects mean that we cannot make a fully calibrated interference pattern in the standard way. What we can do is use the measured visibilities to iteratively fit ionospheric phase shifts and antenna gains towards many bright catalogue sources, and store these fits to aid deconvolution and resampling processes after the images have been made. These measurements are the focus of this paper, but before they are discussed some of these challenges will be reviewed more closely.  \\begin{figure}[ht] \\centering \\includegraphics[width=7.8cm]{fig1.png}  \\caption{Schematic diagram of an MWA interferometer baseline and sky. Each MWA receiving element is a phased tile comprising a $4\\times4$ grid of crossed dipoles ({\\it vertical planar bowtie structures}). The main response lobes are steered electronically by beam-formers ({\\it BF}) to establish the instrument field of view. BF output signals are sampled at baseband and filtered digitally in receiver electronics ({\\it RX}) and correlated to provide cross-power spectra. In addition to compact objects (e.g., quasars and pulsars), the MWA sky includes foreground emission from the galaxy, which is chiefly synchrotron emission that is linearly polarized and traces turbulent and ordered magnetic fields in the interstellar medium ($\\vec{B}$). Magnetized plasma from solar coronal mass ejections or other activity move outward through the heliosphere. This medium and the Earth ionosphere induce primarily refractive shifts in source positions (higher order distortions in the observing passband are small due to the limited extent of the array) and time-variable Faraday rotation of polarization along different lines of sight. The calibration scheme described here enables polarization-sensitive imaging of the diffuse galactic and compact source populations and assembly of a sky model. Knowledge of the compact background sources is ``confusion limited'' (i.e., constrained by multiplicity of sources within one instrument resolution element and contamination of each image pixel by sidelobes of distant sources). Confusion increases with image sensitivity and the size of a resolution element..}  \\label{plot:schematic} \\end{figure}  Unlike the radio sky at higher frequencies, which appears sparsely populated at the sensitivity levels of modern instruments, the sky to be observed by the MWA is full of sources. The high density of sources and large angular resolution of the array will result in images that are confusion limited, with significant flux coming from background galaxies in each synthesized beam, as well as from the sidelobes of other sources in the primary beam of each antenna. The sky is also quite complex. There is significant emission on many angular size scales, from compact extragalactic sources and pulsars to diffuse galactic synchrotron radiation \\cite{Erickson1999}. The latter experiences Faraday rotation and depolarization in the ionized interstellar medium, and as a result exhibits a strong linearly polarized component that is highly dependent on position and frequency \\cite{Wieringa1993}. This polarized all-sky signal is at least a few orders of magnitude brighter than the unpolarized EoR signal, and a fully-polarized calibration formalism, such as the Hamaker-Bregman-Sault measurement equation \\cite{Sault-etal.1996}, is required to reduce contamination from the spatially structured, linearly polarized emission from the galaxy.  If the raw interferometer data are stored for offline processing, iterative calibration and deconvolution algorithms can be used to address many of the problems described below (see, for example, \\cite{Bhatnagar2008} and \\cite{Nijboer2007}). However, it is impractical to store the 19 GB s$^{-1}$ data stream coming from the MWA correlator,\\footnote{The $N_a=512$ antennas lead to $N_a(N_a-1)/2=130816$ different cross-correlation measurements. They are made in a correlator for 4 different polarization products and up to 3072 different frequency channels every 0.5 seconds. Each visibility is represented by 3 real bytes and 3 imaginary bytes, which leads to $\\sim19$ GB/s.} and the MWA will store images. This means that much of the calibration must take place in real time, before or during the imaging process. At the heart of the calibration system for the MWA is the calibrator measurement loop (CML), which measures apparent angular offsets induced by the ionosphere and the system gain toward known compact astronomical sources across the sky. These measurements are used to fit models of the ionosphere and instrument response, and support subtraction of strong sources that limits sidelobe contamination during calibration. As in \\cite{Noordam2004} and \\cite{vanderTol2007}, measuring and subtracting the contribution of each source is carried out sequentially, so that the stronger sources are removed before measurements of weaker sources are made. In this paper we do not consider multivariate fits of parameters for all of the calibrator sources simultaneously, as described in \\cite{vanderTol2007}, but we will discuss it briefly in section \\ref{Discussion}.  Estimation of calibration parameters is greatly over constrained due to the large number of antenna pairs ($1.31\\times10^5$). On the other hand, the wide-field nature of the instrument and real-time computing requirement pose challenges, the most important of which are listed below.  \\begin{enumerate}  \\item \\emph{Direction-dependent gain and polarization response}. Each MWA receiving element is a $4\\times4$ array of fixed crossed dipoles (Fig. \\ref{plot:schematic}). The phased beams are steered with an analogue beamformer, which will typically be updated every 5 to 10 minutes to compensate for rotation of the Earth. The common approach of assuming that the polarized receptors are orthogonal over the field of view with a small amount of direction-dependent leakage cannot be used. The direction-dependent instrumental polarization of the antenna beams will be significant, and it will be measured along with the direction-dependent gain using many calibrator sources spread over the entire sky. These measurements will be repeated as the field of interest moves across the antenna beams.  \\item \\emph{Confusion}. Since the MWA's primary beams cover such a large section of the sky, each field mapped by the MWA will contain hundreds of relatively bright sources. To calibrate the array, we require accurate flux density measurements of known sources. Such measurements can be corrupted by faint sources within the synthesized beam of the array (``confusion'') as well as the sidelobes of brighter sources outside the synthesized beam (``sidelobe contamination'').  Confusion and sidelobe contamination can arise from both compact and large-scale sources such as extragalactic radio galaxies and galactic synchrotron emission respectively. The Galactic synchrotron, in particular, has a polarized component and structure on many spatial scales. Since the interferometer baselines will respond to large-scale structure differently depending on their length and orientation, calibration of data that includes bright resolved sources such as the Sun or Galactic plane must be performed carefully.  \\item \\emph{Ionosphere}. The 3-dimensional ionosphere can significantly perturb the waves coming from celestial sources. The maximum antenna separation of the MWA is short enough that, for a given source and during normal ionospheric conditions, all of the antennas have approximately the same line of sight through the ionosphere. This assumption will be used throughout. Under these conditions there will be no defocusing and the ionosphere can be described by a two-dimensional phase screen that makes sources appear to move away from their true positions. (Faraday rotation of incident polarization is a second ramification, but this will be considered in a future paper. We limit consideration here to ionospheric calibration using unpolarized sources, which cannot be used to calibrate polarization position angles \\cite{Sault-etal.1996}, \\cite{Hamaker2000a}.)  \\item \\emph{Real-time data reduction}. The real-time nature of the MWA means that compute-intensive processes need to avoided where possible. Unfortunately, this means that many of the promising techniques currently being investigated, such as iterative self-calibration and deconvolution algorithms \\cite{Bhatnagar2008}, and some wide-field imaging algorithms \\cite{Cornwell2005}, cannot currently be implemented in real time. They either cannot be used at all or need to be approximated.  \\end{enumerate}  Calibration occurs in a back-end known as the real-time system (RTS), which consists of a visibility integrator (time and frequency), the CML, and an imaging pipeline. These tasks run sequentially, and as mentioned later the processing load is split over frequency. The imaging pipeline incorporates gridding, imaging FFTs, correction for ionospheric and wide-field distortion of the sky, Stokes conversion of images,\\footnote{Stokes parameters describe the polarization state of a signal as an unpolarized component, I, two linearly polarized components, Q and U, and a circularly polarized component, V. They are used extensively in radio astronomy, see, for example, \\cite{BornWolf1999} and \\cite{Thompson-etal.2001}.} and astronomical coordinate conversion. MWA primary science drivers require that the time and frequency resolution are sufficient to ensure that stationary signals from sources throughout the antenna field of view are coherent for the highest frequencies and longest antenna separations, where interference fringe phases vary most rapidly with time.  Most of the key elements of the real-time calibration system have been coded and are regularly tested with simulated data, as described in section \\ref{Examples}. To date, the tests have focused on unpolarized cosmic signals, but the response of the instrument polarizes the signals during reception, and the processing employs the fully-polarized description discussed in the ensuing chapters. Apart from tolerance testing -- to determine optimal bandwidths, number of calibrators, etc. -- and algorithm development for real-time operation, the main piece of outstanding work is the incorporation of polarized calibrators into the system to fully constrain primary beam models and the Faraday rotation component of the ionosphere.  While the discussion and examples given below focus on the MWA, the techniques are applicable to other low-frequency array projects, such as the SKA Molonglo Prototype (SKAMP), the Long Wavelength Array (LWA), and the Low Frequency Array (LOFAR). That said, the instantaneous synthesized beam of the MWA does make it particularly well suited to this type of processing. We will not go into specific details on how the techniques can be optimized, which at any rate will be different for the different arrays, and refer to papers such as \\cite{Lonsdale2004} for an overview of the power of large-$N$ radio arrays.  After outlining the assumptions and mathematical model in the next section, the steps in the CML are discussed in more detail in section \\ref{The Calibrator Measurement Loop}, followed by a discussion on algorithm convergence and performance. We then finish with an analysis of peeling simulated MWA data in section \\ref{Examples}. ",
        "conclusions": "\\label{Discussion}  \\subsection*{Performance}  The relatively short cadence time of the RTS means we need to make compromises when designing real-time peeling process. One of these compromises will be the number of sources peeled every 8 seconds. Several methods for calibrating antenna gains are compared in \\cite{Boonstra2003}. These are alternatives to the technique discussed in section \\ref{Instrumental Gain Measurements}. The number of complex multiplications required by most of the techniques scales as $N_a^3$, where $N_a$ is the number of antennas (or, more generally, the number of receivers being correlated to form visibilities). The technique that required the fewest computations was the logarithmic least squares (LOGLS) algorithm, which scaled as $N_a^2$. The numbers given in \\cite{Boonstra2003} for a single polarization version are $2N_a^2$ multiplications with an additional $16N_a^2$ for weighting.  How does this compare to the algorithm discussed in section \\ref{Instrumental Gain Measurements}? In our application we measure two polarizations, so if we consider each antenna as 2 polarized receptors the number of multiplications for the LOGLS algorithm is $2(2N_a)^2$. There will probably be another factor of 2 since correlations between the receptors on the same antenna will most likely need to be considered.\\footnote{The authors of \\cite{Boonstra2003} note that LOGLS is not easily generalized to a dual-polarized telescope array, but, for the sake of comparison, suppose that such a generalization exists.} The number of complex multiplications used to determine the calibration solutions in section \\ref{Instrumental Gain Measurements} is $O(24N_a^2)$, with an additional $O(12N_a^2)$ for weighting. This efficiency appears to be about as high as one might reasonably expect to achieve.  We are not just dealing with a single source, however. The CML needs to pre-peel all of the calibrator sources, and then, for each source, unpeel, rotate all of the visibilities, solve for the ionospheric offset, solve for the antenna gains, and peel. Table \\ref{tbl:FLOPS} shows the approximate number of floating-point operations used in each of these steps. These numbers were obtained by listing the main operations in the inner loops of the routines and multiplying each by our best estimate of the associated floating-point operations. These numbers are not exact; they are provided as a rough indication of where processing time will be spent. Also, the number of sources processed by each of the routines need not be the same. For example, we might peel and make gain measurements for 50 sources, but make ionospheric measurements on a few hundred more (which requires only the third and fourth rows).  \\begin{table}[htb] \\centering \\caption{Approximate floating-point operations required for each source in the CML (with 512 antenna tiles and a single frequency channel).} \\begin{tabular}{ c c } {\\bf Routine} & {\\bf Floating-point operations (millions)} \\\\ \\hline peeling (applied 3 times)       &  31 \\\\ rotate and accumulate           &  26 \\\\ ionospheric sums and rerotation &  21.5 \\\\ measure tile gains              &  37 \\\\ \\hline total                           & $O$(180) \\\\ \\end{tabular} \\label{tbl:FLOPS} \\end{table}  While the algorithm scales well, there is still on the order of 180 million floating-point operations required during each 8 second calibration cycle. For the whole array, some of the rows in the table need to be multiplied by the number of frequency channels (768), while others need to be multiplied by the number of frequency sub-bands ($\\sim50$). They also need to be multiplied by the number of sources, as discussed in the previous paragraph. We anticipate a few trillion floating-point operations for the CML over the 8 seconds. This does not include various overheads such as memory access that will increase the number of operations by a factor of a few.  \\subsection*{Options for Reducing the Load}  Each ionospheric refraction measurement can use data from every baseline, polarization and frequency. That is over 1.5 gigasamples for a single 2D offset. For sources that are only used for ionospheric measurements, all of the processing listed in table \\ref{tbl:FLOPS} will be reduced if we only use a subset of the visibilities. Most of the baselines are short, and many of those will be redundant. Furthermore, long baselines do not see as much of the extended galactic structure as short baselines, and they measure the apparent offset with higher angular resolution. Initial investigations suggest that we may be able to ignore more than 99\\% of the short baselines for the stronger sources.  Since the tile gains are changing slowly, we do not necessarily need to make measurements for every source every 8 seconds. This can be exploited by making measurements for the strongest sources every 8 seconds, and cycling through subsets of the other sources. This way we still have gain measurements distributed across the sky every few minutes when we make fits for the tile beams, but we only run the full peel algorithm on 20 or 30 sources at a time.  Rather than solving for direction-dependent parameters sequentially, one could fit for all of the calibrator sources simultaneously. In fact, the peeling algorithm is really just a robust and efficient method for reducing the number of unknowns and finding the multivariate solutions. One can also reduce the degrees of freedom by changing the fit parameters to quantities that do not change (or change slowly) with time or frequency, and solve using multiple snapshots, as discussed in detail in \\cite{vanderTol2007}. For the MWA, we hope to be able to use slowly varying direction-dependent dipole gains and phases to describe the primary beam of each tile. Once we have high quality measurements of our tile beams in the field, we will look more closely at fitting these dipole parameters directly. However, source subtraction will then be based on these dipole fits, not on direct tile gain measurements.  \\subsection*{Foreground Subtraction}  One of the primary challenges in the search for a signature from the EoR is that of foreground subtraction. At best the signal will be several orders of magnitude weaker than the galactic foreground and it is important that we understand the nature of the residuals from the calibration and peeling process. One of the drawbacks of peeling strong sources in real time and then averaging the resulting images together is that any residuals are also averaged into the mix. All of the calibration data will be stored in a database for use in offline processing, and peeling errors can be assessed and reduced at that stage. However, if there is any concern that residuals from the peeling process might mimic the EoR signal then peeling can be used for calibration only, and images formed from unpeeled visibilities.  There is significant effort going into techniques for removing foregrounds during offline processing. These include techniques that exploit spectral differences in the foregrounds and the EoR signature (see for example \\cite{Morales2006} and the discussion and references in \\cite{Furlanetto2006}), and the rotation measure synthesis techniques discussed in \\cite{Burn1966} \\& \\cite{Brentjens2005}. Direction-dependent deconvolution techniques such as the one described in \\cite{Bhatnagar2008} will also be essential for imaging.  \\subsection*{Radio Frequency Interference}  Of concern for any telescope operating at MWA frequencies is radio frequency interference (RFI). Even though the Murchison site is extremely radio quiet \\cite{Bowman2007a}, the array will still have to deal with some RFI. This includes communication and military satellite signals, reflections of FM radio broadcasts, and natural interference such as lightning. Due to the low spectral occupancy of the RFI, the high quality polyphase filter banks used to isolate frequencies, and the campaign-mode operation of the array, we will adopt the traditional strategy of flagging and ignoring contaminated data before imaging. Missing frequency or time samples can be accounted for in the weighting of the various least-squares algorithms.    \\label{Summary}  We have described a general approach for making measurements of strong point sources that can be used in the calibration process of wide-field, low-frequency radio arrays. This approach has been adopted for the MWA, and there is an ongoing effort to develop the required software, as well as to understand the benefits and limitations of the approach.  We have used simulated visibility data to show that the peeling algorithm works well in situations that are of major concern for future radio telescopes: crowded fields, strong galactic emission, and ionospheric refraction. The algorithm exhibits fast convergence, which is important since sources will be moving in and out of antenna sidelobes and the algorithm needs to be able to keep up with the antenna gain and phase changes, as well as changes in the ionosphere. Critical parts of the process are shown to be computationally efficient, and parts of the system lend themselves to significant levels of optimization.  The next step is a detailed analysis of the convergence properties of the algorithms, and a series of tolerance tests to investigate how the algorithms will behave in the various conditions we expect to encounter. This includes observations of weak emission that is masked by significant polarized diffuse foregrounds, high dynamic range observations close to the sun, and observations in the presence of severe ionospheric conditions, such as during activation and recombination of the ionosphere. As data from the initial deployment of antennas become available in late 2008 and early 2009, we will get a clearer picture of how harmful phenomena such as source variability (due to ionospheric scintillation, for instance) and dipole mutual coupling (which will affect our tile beam models) can be, and these can be worked into the tolerance tests."
    },
    "0807/0807.4530_arXiv.txt": {
        "abstract": "We study the impact of inhomogeneous hydrogen reionization on the thermal evolution of the intergalactic medium (IGM) using hydrodynamic + radiative transfer simulations where reionization is completed either early ($z\\sim9$) or late ($z\\sim6$). In general, we find that low-density gas near large-scale overdensities is ionized and heated earlier than gas in the large-scale, underdense voids. Furthermore, at a later time the IGM temperature is inversely related to the reionization redshift because gas that is heated earlier has more time to cool through adiabatic expansion and Compton scattering. Thus, at the end of reionization the median temperature-density relation is an inverted power-law with slope $\\gamma-1\\sim-0.2$, in both models. However, at fixed density, there is up to order unity scatter in the temperature due to the distribution of reionization redshifts. Because of the complex equation-of-state, the evolved IGM temperature-density relations for the redshift range $4\\lesssim z\\lesssim6$ can still have significant curvature and scatter. These features must be taken into account when interpreting the Ly$\\alpha$ absorption in high redshift quasar spectra. ",
        "introduction": "The cosmic reionization of hydrogen fundamentally changes the thermal and ionization conditions for the high redshift ($z\\gtrsim6$) intergalactic medium (IGM), converting a cold and neutral gas into a warm and highly ionized one \\citep[see][for a review]{2008Loeb}. During this inhomogeneous process, the radiative transfer (RT) of the UV field, primarily generated by stellar sources, will proceed such that large-scale, overdense regions near sources are generally ionized and heated earlier than large-scale, underdense regions far from sources \\citep[e.g.][]{2004BarkanaLoeb, 2004WyitheLoeb, 2004FZH, 2008Lee}. The photo-ionization of hydrogen (HI $\\rightarrow$ HII) and helium (HeI $\\rightarrow$ HeII) is expected to heat the IGM to temperatures of $\\gtrsim10^4$ K, with higher temperatures obtained for harder radiation spectra \\citep[e.g.][]{1990MiraldaEscude, 1994MiraldaEscude}. Thereafter, the thermal evolution will be governed by adiabatic cooling/heating for underdense/overdense gas, Compton scattering on the Cosmic Microwave Background (CMB), atomic line cooling, shock-heating, and additional photo-heating only where recombination is efficient.  In principle, the temperature of the IGM can be quantified through the thermal broadening of the Lyman alpha ($\\lya$) forest lines \\citep[e.g.][]{1999Schaye, 2000Schaye, 2000Ricotti, 2001McDonald} and the Jeans smoothing of the $\\lya$ forest flux power spectrum \\citep[e.g.][]{2000Theuns, 2001Zaldarriaga, 2004Viel}. Measurements of the temperature can then be used to put constraints on when the reionization of hydrogen and helium occurred \\citep[e.g.][]{2002Theuns, 2003HuiHaiman}. Note that the full reionization of helium (HeII $\\rightarrow$ HeIII) is expected to occur at lower redshifts ($z\\sim3$), triggered by quasars with harder spectra \\citep[see][for recent modelling]{2008Furlanetto, 2008bBolton, 2008McQuinn}. The current observational constraints have large uncertainties, although better measurements will come in time and improved models will be required to interpret them.  In this {\\it Letter} we focus on the impact of inhomogeneous hydrogen reionization on the thermal evolution of the IGM. In particular, we study the equation-of-state and measure the temperature-density relation and its scatter using hydrodynamic + RT simulations. \\citet{1997HuiGnedin} have studied the effects of uniform reionization and found that the temperature-density relation is well approximated by a positive power-law slope for the low-density IGM. However, inhomogeneous reionization can give rise to a much different equation-of-state, especially right after completion. Specifically, the large-scale, underdense voids are initially hotter than higher density gas near sources, that having been ionized and heated earlier have also had more time to cool. This gives rise to an equation-of-state characterized by an inverted (negative slope) temperature-density relation. At later times, the underdense gas cools adiabatically faster and the temperature contrast decreases and eventually reverses back to the positive slope. ",
        "conclusions": "We have studied the photo-ionization and photo-heating of the IGM from hydrogen reionization using two hydrodynamic + RT simulations. We considered two basic models in which reionization is completed early ($z\\sim9$) and late ($z\\sim6$) and found for both models that the temperature of a low-density region at a later time is inversely related to the redshift of reionization of that region. low-density gas near large-scale overdensities is ionized and heated earlier than gas in the large-scale, underdense voids. As a result, at the end of reionization the median temperature-density relation is an inverted power-law with slope $\\gamma-1\\sim-0.2$. There is up to order unity scatter in the temperature at fixed density due to the distribution of reionization redshifts. We expect that both the slope and scatter will depend on the reionization history, especially its duration. Furthermore, it is known that radiative transfer effects from Lyman limit systems can modify the temperature distribution \\citep[e.g.][]{1999AbelHaehnelt}, although more for helium rather than hydrogen reionization because for the latter, the UV spectrum from stellar sources is relatively soft and only weak spectral filtering can occur.  We conclude that at the high redshift range $4\\la z\\la 6$, the equation-of-state of the IGM is more complicated than the commonly assumed forms (isothermal or a tight power-law relation). It is important to keep this result in mind when interpreting the Ly$\\alpha$ absorption in quasar spectra \\citep[e.g.][]{2002Fan, 2006Fan, 2006Lidz, 2008Gallerani}. In fact, a recent analysis by \\citet{2007Becker} has already suggested that adopting an inverted temperature-density relation instead of an isothermal model may have profound implications on the interpretation of the reionization process based on the high redshift Ly$\\alpha$ forest. We will study the observational signatures of inhomogeneous hydrogen reionization in an upcoming paper."
    },
    "0807/0807.4750_arXiv.txt": {
        "abstract": "With the TRIS experiment we have performed absolute measurements of the sky brightness in a sky circle at $\\delta = +42^{\\circ}$ at the frequencies $\\nu =$ 0.60, 0.82 and 2.5 GHz. In this paper we discuss the techniques used to separate the different contributions to the sky emission and give an evaluation of the absolute temperature of the Cosmic Microwave Background. For the black-body temperature of the CMB we get: $T_{cmb}^{th}=(2.837 \\pm 0.129 \\pm 0.066)K$ at $\\nu=0.60$ GHz; $T_{cmb}^{th}=(2.803 \\pm 0.051 \\ ^{+0.430} _{-0.300})K$ at $\\nu=0.82$ GHz; $T_{cmb}^{th}=(2.516 \\pm 0.139 \\pm 0.284)K$ at $\\nu=2.5$ GHz. The first error bar is statistic (1$\\sigma$) while the second one is systematic. These results represent a significant improvement with respect to the previous measurements. We have also set new limits to the free-free distortions, $ -6.3 \\times 10^{-6} < Y_{ff} < 12.6 \\times 10^{-6}$, and slightly improved the Bose-Einstein upper limit, $|\\mu| < 6 \\times 10^{-5}$, both at 95\\% confidence level. ",
        "introduction": "At decimetric wavelenghts the sky brightness temperature ($T_{sky}$) can be written as the sum of different contributions: cosmic microwave background ($T_{cmb}$), galactic emission ($T_{gal}$), unresolved extragalactic radio sources ($T_{uers}$).  The temperature of the cosmic microwave background (CMB), $T_{cmb}$, can be considered, for our purposes, as isotropic. The largest anisotropy component is the dipole term $\\Delta T_{dipole}(\\alpha,\\delta) = T_d \\ \\cos \\theta$, with $ T_d = 3.381 \\pm 0.007$ mK, and $\\theta$ angle between the direction of observation and the maximum of the dipole at ($\\alpha = 11^h 12^m.2 \\pm 0^m.8$; $\\delta = -7^{\\circ}.06 \\pm 0^{\\circ}.16$) (\\cite[]{Bennet_96}; \\cite[]{Fixen_96} and \\cite[]{Fixen_02}). It gives a level of anisotropy well inside the error budget of our absolute measurements. Therefore we will consider in the present paper only the possible dependence of the CMB brightness temperature on the frequency ($T_{cmb}(\\nu)$). The frequency dependence could be related to the spectral distortion of the CMB.  The most accurate measurement of the CMB  temperature so far made is the one obtained by the COBE-FIRAS team (\\cite[]{Mather_90}; \\cite[]{Mather_94}; \\cite[]{Fixen_96}; \\cite[]{Mather_99}; \\cite[]{Fixen_02}). Between 60 to 600 GHz they found a spectrum fully compatible with a black body emitting at the thermodynamic temperature $T_{cmb}^{th}=2.725 \\pm 0.001 \\ K$ (see \\cite[]{Fixen_02}). Major spectral distortions in this range are excluded by COBE-FIRAS and the parameters of both the Bose-Einstein and Compton distortions are largely constrained: $|\\mu| < 9 \\times 10^{-5}$ and  $|y| < 15 \\times 10^{-6}$ (both at 95\\% CL, see \\cite[]{Fixen_96}).  Although very accurate the FIRAS result covers only part of the CMB frequency spectrum. In particular it does not cover the low frequency region where TRIS measurements were made and where distortions can be expected. Table \\ref{tab1} collects the measurements of the CMB temperature made at frequencies below few GHz starting from the 80's. At these low frequencies the error bars are much larger than at FIRAS frequencies, essentially because of the large uncertainties associated to the absolute calibration procedures and to the presence of important foregrounds.  To improve the low frequency situation it has been proposed in the past to carry on measurements of the CMB temperature from space: LOBO (\\cite[]{Sironi_95}; \\cite[]{Sironi_97}) for which TRIS is a pathfinder and DIMES \\cite[]{Kogut_96}. DIMES has been then transformed in a balloon program (ARCADE), whose results have been published recently (see \\cite[]{Kogut_04}; \\cite[]{Fixen_04}; \\cite[]{Singal_06}).  The black-body spectrum of the CMB is the result of the complete thermalization of the Universe in the pre-recombination era. At that time the strong interactions among the different components of the cosmic plasma were efficient enough to re-establish immediately the thermodynamic equilibrium after an energy injection. Nevertheless small deviations from the black-body distribution can survive if the energy injection occurred at a red-shift $z\\lesssim 10^6$. In the pre-recombination era ($z\\gtrsim 10^4 - 10^3$) a kinetic equilibrium is re-established through the Compton scattering, Double Compton scattering and the Bremsstrahlung mechanism. The resulting spectrum assumes a typical shape described by the Bose-Einstein or by the Compton distribution. For a review of these mechanisms see \\cite[]{Zeldovich_69}, \\cite[]{Sunyaev_70}, \\cite[]{Zeldovich_72}, \\cite[]{Illarionov_75a}, \\\\ \\cite[]{Illarionov_75b}, \\cite[]{Sunyaev_80}, while detailed calculations have been performed by \\cite{Burigana_91a}, \\cite{Burigana_91b}, \\cite{Burigana_95} and \\cite{Daly_91}. These types of distortions are described by the chemical potential $\\mu$ and the Comptonization parameter $y$, whose possible values have been already limited by COBE-FIRAS.  Energy injections after the recombination era are responsible of spectral distortions without any possibility of re-thermalization. Among these processes there is the photon injection through free-free mechanism in a re-ionized universe (see \\cite[]{Bartlett_91}). This effect depends on the square of the wavelength and therefore is dominant at very low frequencies and is not constrained by the COBE-FIRAS measurements. Other possible distortions are related to the formation of primordial molecules and to the formation of structures at large scale. Both these mechanisms are difficult to be modeled but, at least in the case of the formation of primordial molecules, a not negligible effect could be expected at decimetric wavelengths (see \\cite[]{Varsha_77} and \\cite[]{Dubrovich_95}).  The galactic emission ($T_{gal}(\\nu,\\alpha,\\delta)$) is anisotropic and has a power law frequency spectrum: $T_{gal}(\\nu) = T_{gal}(\\nu_0) (\\nu/\\nu_0)^{\\beta}$. Its importance and properties are analyzed in the companion paper \\cite[]{TRIS-III} (hereafter Paper III) from which we take the spectral index $\\beta$ we will use in the following.  The contribution of the unresolved extragalactic radio sources (UERS), $T_{uers}$, has also a power law frequency spectrum, but for a low resolution experiment it can be considered isotropic. In the following we will take it from the model developed by \\cite{Gervasi_08} on the basis of the source number counts measurements available in literature.  After a brief summary of TRIS experiment (Section \\ref{TRIS}), in Section \\ref{Calculation} we discuss the techniques used to separate the components of $T_{sky}$ and the results obtained for $T_{cmb}$. The implications of these results in terms of spectral distortions are presented in Section \\ref{Discussion}. ",
        "conclusions": "Starting from the absolute measurements of the sky brightness temperature, performed by the TRIS experiment, and presented in Paper I \\cite[]{TRIS-I}, we have evaluated the absolute temperature of the Cosmic Microwave Background at $\\nu =$ 0.60, 0.82 and 2.5 GHz.  The thermodynamic temperatures of the CMB we get are: $T_{cmb}^{th}=(2.837 \\pm 0.129 \\pm 0.066)K$ at $\\nu=0.60$ GHz; $T_{cmb}^{th}=(2.803 \\pm 0.051 \\ ^{+0.430} _{-0.300})K$ at $\\nu=0.82$ GHz; $T_{cmb}^{th}=(2.516 \\pm 0.139 \\pm 0.284)K$ at $\\nu=2.5$ GHz. The first error bar is 1$\\sigma$ statistics, while the second one is the systematic on the zero level assessment.  Thanks to improvements of the absolute calibration system and in the foregrounds separation technique TRIS succeded in reducing previous uncertainties by a factor $\\sim 9$ at $\\nu=0.60$ GHz and by a factor $\\sim 7$ at $\\nu=0.82$ GHz. At 2.5 GHz TRIS results are in agreement with the previous measurements.  These results, used to look for CMB spectral distortions, give an upper limit to the chemical potential $|\\mu| < 6 \\times 10^{-5}$ (95\\% CL) used to describe the \\emph{BE} distortions. We have also constrained the free-free distortions to: $ -6.3 \\times 10^{-6} < Y_{ff} < 12.6 \\times 10^{-6}$ (95\\% CL), approaching the values suggested by observations of the Lyman-$\\alpha$ forest \\cite[]{Haiman_97}."
    },
    "0807/0807.2176_arXiv.txt": {
        "abstract": "The flight calibration of the spectral response of CCD instruments below 1.5 keV is difficult in general because of the  lack of strong lines in the on-board calibration sources typically available. We have been using \\name\\, the brightest supernova remnant in the Small Magellanic Cloud, to evaluate the response models of the ACIS CCDs on the Chandra X-ray Observatory (CXO), the EPIC CCDs on the XMM-Newton Observatory, the XIS CCDs on the {\\em Suzaku} Observatory, and the XRT CCD on the {\\em Swift} Observatory.  E0102 has strong lines of O, Ne, and Mg below 1.5 keV and little or no Fe emission to complicate the spectrum. The spectrum of E0102 has been well characterized using high-resolution grating instruments, namely the XMM-Newton RGS and the CXO HETG, through which a consistent spectral model has been developed that can then be used to fit the lower-resolution CCD spectra. Fits with this model are sensitive to any problems with the gain calibration and the spectral redistribution model of the CCD instruments. We have also used the measured intensities of the lines to investigate the consistency of the effective area models for the various instruments around the bright O ($\\sim570$~eV and $\\sim654$~eV) and Ne ($\\sim910$~eV and $\\sim1022$~eV) lines. We find that the measured fluxes of the O~{\\small VII}~triplet, the O~{\\small VIII}~Ly~$\\alpha$ line, the Ne~{\\small IX}~triplet, and the Ne~{\\small X}~Ly~$\\alpha$ line generally agree to within $\\pm10\\%$ for all instruments, with 28 of our 32 fitted normalizations within $\\pm10\\%$ of the RGS-determined value.  The maximum discrepancies, computed as the percentage difference between the lowest and highest normalization for any instrument pair, are 23\\% for the O~{\\small VII}~triplet,  24\\% for the O~{\\small VIII}~Ly~$\\alpha$ line, 13\\% for the Ne~{\\small IX}~triplet, and 19\\% for the Ne~{\\small X}~Ly~$\\alpha$ line. If only the CXO and XMM are compared, the maximum discrepancies are 22\\% for the O~{\\small VII}~triplet,  16\\% for the O~{\\small VIII}~Ly~$\\alpha$ line, 4\\% for the Ne~{\\small IX}~triplet, and 12\\% for the Ne~{\\small X}~Ly~$\\alpha$ line. ",
        "introduction": "\\label{sect:intro}  %  This paper reports the progress of a working group within the {\\em International Astronomical Consortium for High Energy Calibration} (IACHEC) to develop a calibration standard for X-ray astronomy in the bandpass from 0.3 to 2.5 keV.  A brief introduction to the IACHEC organization, its objectives and meetings, may be found at the web page {\\tt http://www.iachec.org/}.  Our working group was tasked with selecting celestial sources with line-rich spectra in the 0.3-2.5~keV bandpass which would be suitable cross-calibration targets for the current generation of X-ray observatories.  The desire for strong lines in this bandpass stems from the fact that the quantum efficiency and spectral resolution of the current CCD-based instruments is changing rapidly from 0.3 to 1.5~keV and also significantly around the Si~K~edge, but the on-board calibration sources currently in use typically have strong lines at only two energies, 1.5~keV (Al~K$\\alpha$) and 5.9~keV (Mn~K$\\alpha$). The only option available to the current generation of flight instruments to calibrate any time variable response is to use celestial sources.  The missions which have been represented in this work are the {\\em Chandra X-ray Observatory}\\cite{weiss00,weiss02} (CXO), the {\\em X-ray Multimirror Mission}\\cite{jansen2001} (XMM-Newton), the {\\em ASTRO-E2 Observatory} ({\\em Suzaku}), and  the {\\em Swift} Gamma-ray Burst Mission.  Data from the following instruments have been included in this analysis: the {\\em High-Energy Transmission Grating} (HETG) and the {\\em Advanced CCD Imaging Spectrometer}\\cite{bautz98,garmire03,garmire92} (ACIS) on the CXO, the  {\\em Reflection Gratings Spectrometers}\\cite{denherder2001} (RGS), the {\\em European Photon Imaging Camera} (EPIC) {\\em Metal-Oxide Semiconductor}\\cite{turner2001} (MOS) CCDs and the EPIC p-n junction\\cite{strueder2001} (pn) CCDs on XMM-Newton, the {\\em X-ray Imaging Spectrometer} (XIS) on {\\em Suzaku}, and the {\\em X-ray Telescope}\\cite{burrows2005,godet2007} (XRT) on {\\em Swift}.   Suitable calibration targets would need to possess the following qualities.  The source would need to be constant in time, to have a simple spectrum defined by a few bright lines with a minimum of line-blending, and to be extended so that ``pileup'' effects in the CCDs are minimized but not so extended that the off-axis response of the telescope dominates the uncertainties in the response.  Our working group focused on supernova remnants (SNRs) with thermal spectra and without a central source such as a pulsar, as the class of source which had the greatest likelihood of satisfying these criteria.  We narrowed our list to the Galactic SNR Cas-A, the Large Magellanic Cloud remnant N132D and the Small Magellanic Cloud remnant \\name\\/.  We discarded Cas-A since it is a relatively young ($\\sim350$~yr) SNR with significant brightness fluctuations in the X-ray, radio, and optical over the past three decades, it contains a faint (but apparently constant) central source, and it is relatively large (radius $\\sim3.5$~arcminutes).  We discarded N132D because it has a complicated, irregular morphology in the X-rays and its spectrum shows strong, complex Fe emission.  We therefore settled on \\name\\ as the most suitable source given its relatively uniform morphology, small size (radius $\\sim0.4$~arcminutes), and comparatively simple X-ray spectrum.  \\begin{figure}[h] \\begin{center} \\begin{tabular}{c} \\includegraphics[width=3.0in,angle=0]{4panel.ps} \\includegraphics[width=3.0in,angle=0]{xis_panel.ps} \\end{tabular} \\end{center} \\vspace{-0.15in} \\caption[image] { \\label{fig:image} Images of E0102 from ACIS S3 (top left), MOS(bottom left), XRT(middle top), pn(middle bottom), and XIS(right). The black circles indicate the extraction regions used for the spectral analysis. The high angular resolution of the CXO's mirrors are evident in the structure resolved in the SNR and the small extraction region.} \\end{figure}    The SNR \\name\\ (hereafter E0102) was discovered by the {\\em Einstein Observatory}\\cite{seward1981}.  It is the brightest SNR in X-rays in the {\\em Small Magellanic Cloud} (SMC). E0102 has been extensively imaged by CXO\\cite{gaetz2000,hughes2000} and XMM-Newton\\cite{sasaki2001}. Figure~\\ref{fig:image} shows an image of E0102 with the relevant spectral extraction region for each of the instruments included in this analysis.  E0102 is classified as an ``O-rich'' SNR and has an estimated age of $\\sim1,000$~yr. The source diameter is small enough such that a high resolution spectrum may be acquired with the HETG on the CXO and the RGS on XMM-Newton. The HETG spectrum\\cite{flanagan2004} and the RGS spectrum\\cite{rasmussen2001} both show strong lines of O, Ne, and Mg with little or no Fe. E0102's spectrum is relatively simple compared to a typical SNR spectrum. Figure~\\ref{fig:rgs_spec_lin} displays the RGS spectrum from E0102. The strong, well-separated lines in the energy range 0.5 to 1.5~keV make this source a useful target for calibration observations. The source is extended enough to reduce the effects of photon pileup, which distorts a spectrum. Although moderate pileup is expected in all the non-grating instruments when observed in modes with relatively long frame times. The source is also bright enough to provide a large number of counts in a relatively short observation.  \\begin{figure}[h] \\begin{center} \\begin{tabular}{c} \\includegraphics[width=4.9in,angle=270]{rgspn_mod_tbabs_tbvarabs_2apec_line_ratios_jd_v1.9_rgs_lin_ratio_v2.ps} \\end{tabular} \\end{center} \\vspace{-0.15in} \\caption[image] { \\label{fig:rgs_spec_lin}  RGS1/2 spectrum of E0102 from a combination of 23 observations. Note the bright lines of O, Ne, and Mg. } \\end{figure} ",
        "conclusions": "\\label{sect:conclusions}  We have use the line-dominated spectrum of the SNR E0102 to test the response models of the ACIS S3, MOS, pn, XIS, and XRT CCDs below 1.5~keV. We have fitted the spectra with the same model in which the continuum and absorption components and the weak lines are held fixed while allowing only the normalizations of four bright lines/line complexes to vary.  We have compared the fitted line normalizations of the O~{\\small VII}~For line, the O~{\\small VIII}~Ly~$\\alpha$ line, the Ne~{\\small IX}~Res line, and Ne~{\\small X}~Ly~$\\alpha$ line to examine the consistency of the effective area models for the various instruments in the energy ranges around 570~eV, 654~eV, 915~eV, and 1022~eV. We find that the instruments are in general agreement with 28 of the 32 scaled normalizations within $\\pm10\\%$ of the RGS determined values. We find that the scaled line normalizations agree to within 23\\%, 24\\%, 13\\%, \\& 19\\% for O~{\\small VII}, O~{\\small VIII}, Ne~{\\small IX}, \\& Ne~{\\small X} when all instruments are considered.  When only CXO and XMM-Newton are considered, we find that the fitted line normalizations agree to within 22\\%, 16\\%, 4\\%, \\& 12\\% for O~{\\small VII}, O~{\\small VIII}, Ne~{\\small IX}, \\& Ne~{\\small X}.    \\begin{figure}[tbh] \\begin{center} \\begin{tabular}{c} \\includegraphics[width=3.5in,angle=270]{mos1-spie.ps} \\end{tabular} \\end{center} \\vspace{-0.15in} \\caption[MOS1] { \\label{fig:mos1_spec} MOS1 spectrum from OBSIDs 0123110201 and 0135720601. Note the excellent spectral resolution of the MOS data. } \\end{figure}  \\begin{figure}[bth] \\begin{center} \\begin{tabular}{c} \\includegraphics[width=3.5in,angle=270]{rgspn_mod_tbabs_tbvarabs_2apec_line_ratios_jd_v1.9_P0412980301PNS001_log_lin_ratio.ps} \\end{tabular} \\end{center} \\vspace{-0.15in} \\caption[pn] { \\label{fig:pn_spec} pn spectrum from OBSID0412980301 . The second (lower) curve shows the same data but with a linear axis which has been shifted downwards for clarity. Note the high count rate and the pattern in the residuals which might indicate an issue with the spectral redistribution function. } \\end{figure}  \\begin{figure}[tbh] \\begin{center} \\begin{tabular}{c} \\includegraphics[width=3.5in,angle=270]{iachec08c_xis1_fittedspec.ps} \\end{tabular} \\end{center} \\vspace{-0.15in} \\caption[pn] { \\label{fig:xis_spec} XIS1 spectrum from OBSID 100044010 showing the model for the XRB RX~J0103. Note the residuals below 0.5~keV and above 1.5~keV.} \\end{figure}  \\begin{figure}[bth] \\begin{center} \\begin{tabular}{c} \\includegraphics[width=3.5in,angle=270]{e0102_xrt_pc_g0to12.ps} \\end{tabular} \\end{center} \\vspace{-0.15in} \\caption[pn] { \\label{fig:xrt_spec} Combined {\\em Swift} XRT spectrum from four observations.  Note the excess from 1.5 to 2.5 keV, this is a signature of pileup. } \\end{figure}      \\vspace*{-0.1in}"
    },
    "0807/0807.2340_arXiv.txt": {
        "abstract": "{Recent observations of the Galactic center revealed a nuclear disk of young OB stars near the massive black hole (MBH), in addition to many similar outlying stars with higher eccentricities and/or high inclinations relative to the disk (some of them possibly belonging to a second disk). In addition, observations show the existence of young B stars (the 'S-cluster') in an isotropic distribution in the close vicinity of the MBH ($<0.04$ pc). We use extended N-body simulations to probe the dynamical evolution of these two populations. We show that the stellar disk could have evolved to its currently observed state from a thin disk of stars formed in a gaseous disk, and that the dominant component in its evolution is the interaction with stars in the cusp around the MBH. We also show that the currently observed distribution of the S-stars could be consistent with a capture origin through 3-body binary-MBH interactions. In this scenario the stars are captured at highly eccentric orbits, but scattering by stellar black holes could change their eccentricity distribution to be consistent with current observations. ",
        "introduction": "High resolution observations have revealed the existence of many young OB stars in the galactic center (GC). Accurate measurements of the orbital paramters of these stars give strong evidence for the existence of a massive black hole (MBH) which govern the dynamics in the GC \\citep{sch+02b,ghe+03a}. Most of the young stars are observed in the central 0.5 pc around the MBH. The young stars population in the inner 0.04 pc (the 'S-stars' or the 'S-cluster') contain only young B-stars, in apparently isotropic distribution around the MBH, with relatively high eccentricities ($0.3\\le e\\le0.95$) \\citep{ghe+03a,eis+05}. The young stars outside this region contain many O-stars residing in a stellar disk moving clockwise in the gravitational potential of the MBH \\citep{lev+03,gen+03a,lu+06,pau+06,tan+06}. The orbits of the stars in this disk have average eccentricity of $\\sim0.35$ and the opening of the disk is $h/R\\sim0.1$, where $h$ is the disk height and $R$ is its radius. \\citet{pau+06} and \\citet{tan+06} suggested the existence of another stellar disk rotating counter clockwise and almost perpendicular to the CW disk. This disk is currently debated as many of the young stars are have intermediate inclinations, and are possibly just outliers that do not form a coherent disk structure \\citep{lu+06}.  Here we briefly report on our study of the dynamical evolution of the young stars in the GC, both in the stellar disk and in the S-cluster. We use extensive N-body simulations with realistic number of stars  ($10^{3}-10^{5}$) using the  {\\tt gravitySimulator}, a 32-node cluster at the Rochester Institute of Technology that incorporates GRAPE accelerator boards in each of the nodes \\citep{har+07}. Thus we are able to probe the dynamics of the stars near the MBH and their stellar environment. We study two basic issues: (1) the long term evolution of the S-stars up to their lifetime of a few $10^{7}$ yrs, including their dynamical interaction with stars in the vicinity of the MBH; (2) The evolution of a realistic stellar disk, taking into account both the effects of non-equal mass stars, as studied earlier, and more importantly the effect of the interactions of disk stars with the stellar cusp around the MBH. As we show, the latter component proves to be more important than the other components discussed in previous studies. A detailed report of our complete set of out simulations (not shown here), in addition to analytic calculations will be presented in upcoming papers (Perets et al., in preparation) ",
        "conclusions": "\\label{sec:summary}  The dynamical evolution of the young stars in the GC both in the stellar disk(s) and in the inner S-cluster is not yet understood. We used N-body simulations to study the dynamics and origin of these stars. We found that the S-stars close to the MBH in the GC could be stars that were captured following a binary disruption by the MBH, and later on dynamically evolved due to scattering by other stars, or stellar black holes, to obtain their currently observed orbits. We also show the the young stellar disk could have formed as a cold (thin) circular disk and evolve to its currently observed thick (hot) disk, mostly due to scattering by cusp stars, whereas self relaxation of the disk plays a more minor role, especially in regard to the more massive stars seen in observations."
    },
    "0807/0807.0729_arXiv.txt": {
        "abstract": "% This first Subaru international conference has highlighted the remarkably diverse and significant contributions made using the 8.2m Subaru telescope by both Japanese astronomers and the international community. As such, it serves as a satisfying tribute to the pioneering efforts of Professors Keiichi Kodaira and Sadanori Okamura whose insight and dedication is richly rewarded. Here I try to summarize the recent impact of wide field science in extragalactic astronomy and cosmology and take a look forward to the key questions we will address in the near future. ",
        "introduction": "%  It's a pleasure to summarize the first Subaru international conference. The conference has been striking in many ways. Before I arrived in Hayama, I was impressed by one of the two conference posters: the woodblock print (Ukiyo-e) `The Great Wave off Kanagawa\" by the artist Katsushika Hokusai (1760-1849). This portrays desperate fishermen struggling with a tsunami against the backdrop of Mt Fuji. I thought this seemed an alarming image to use to encourage overseas visitors to come to the meeting, particularly given the proximity of the conference venue to the ocean and the location depicted! Fortunately, I'm glad to report our views of Mt Fuji and the ocean have been peaceful and serene throughout our stay.  The second striking thing about the meeting is the cause for celebrating the achievements of the Subaru telescope. Although we've heard panoramic results from many observatories and facilities, it's clear much of the focus has been with Subaru. We're congratulating Professor Keiichi Kodaira for his leadership in bringing this remarkable facility to fruition, and Professor Sadanori Okamura and his colleagues for their insight in developing the remarkable prime focus camera, Suprime-Cam. Subaru remains the only fully-steerable 8 meter class telescope with a prime focus imager; we've seen ample evidence of the wisdom of this decision. With further instruments such as MOIRCS, FMOS and hopefully eventually WFMOS, the Subaru telescope is destined to remain an unique facility in wide-field astronomy.  In an interesting presentation, Iye (2007) examined a recent article by Grothkopf et al (2007) which analyzed the Hirsch (2005) $h$-index\\footnote{The $h$-index is the number $n$ of refereed articles each of which has more than $n$ citations} of the four major 8-10m telescopes. Subaru ranks third in this analysis (behind Keck and the VLT) with a $h$-index of 40. However, as citations are cumulative, older facilities naturally garner more citations. The $m$-index takes this into account by dividing the number of citations by the years of routine operation. Grothkopf et al find that when this is done, the newer VLT is having a comparable impact to Keck, both have $m\\simeq$10, but  Iye was dismayed to find Subaru remains third with $m\\simeq$6.  Not one to be defeated easily, Iye invented the $i$-index which is the $m$-index divided by the number of telescopes in the observatory. As the VLT comprises 4 independent telescopes, and Keck and Gemini comprise two each, Subaru is clearly disadvantaged by this factor. When this correction is made, as a single telescope, Iye was pleased to find Subaru leads the pack!  Even if citations are never a fully-accurate guide of the discovery rate of a telescope, I think we'd all agree that Subaru has, through its unique wide field capabilities, contributed enormously to the growth and international standing of Japanese astronomy. ",
        "conclusions": ""
    },
    "0807/0807.0203_arXiv.txt": {
        "abstract": "{  We have acquired a deep i-band image of the BL Lacertae object \\object{S5 0716+714} while the target was in an low optical state. Due to the faintness of the nucleus, we were able to detect the underlying host galaxy. The host galaxy is measured to have an I-band magnitude of 17.5 $\\pm$ 0.5 and an effective radius of (2.7 $\\pm$ 0.8) arcsec. Using the host galaxy as a ``standard candle'', we derive z = $0.31 \\pm 0.08$ (1$\\sigma$ error) for the host galaxy of \\object{S5 0716+714}.  This redshift is consistent with the redshift z = 0.26 determined by spectroscopy for 3 galaxies close to \\object{S5 0716+714}.  The effective radius at z = 0.31 would be 12 $\\pm$ 4 kpc, which is consistent with values obtained for BL Lac host galaxies.  An optical spectrum acquired during the same epoch shows no identifiable spectral lines.  } ",
        "introduction": "BL Lacertae objects (BL Lacs) are active galactic nuclei (AGN) characterized by a featureless nonthermal continuum, high optical and radio polarization, and variability across the entire electromagnetic spectrum.  These properties arise from a relativistic jet pointed almost towards the observer \\citep{1979ApJ...232...34B}. The strong continuum emission from the jet in the optical sometimes presents a problem when attempting to properly characterize the underlying host galaxy or measure its redshift, especially since by definition the emission lines are very weak in BL Lacs. Hence, it is unsurprising that even after very deep searches \\citep[e.g.][]{2006AJ....132....1S} many BL Lacs remain without a spectroscopically determined redshift.  The redshift is a fundamental property of any extragalactic object, without which e.g. the luminosity or intrinsic variability properties of the target cannot be determined.  Furthermore, many BL Lacs have been recently detected at TeV gamma-rays using ground-based Cherenkov telescopes MAGIC, HESS, and VERITAS \\citep[e.g.][]{2007arXiv0712.3352H}.  Since VHE gamma-rays can be absorbed by the interaction with low energy photons of the EBL via pair production, this opens the possibility of measuring the amount of extragalactic background light \\citep[EBL, ][]{nikishov1962,1992ApJ...390L..49S} in the optical through to the far infrared, which provides important information about the galaxy and star formation history.  The absorption depends strongly on the distance of the source and the energy of the gamma-rays. If the redshift of the source is known, the VHE gamma-ray spectrum can be used to derive limits on the EBL \\citep[e.g.][]{2006Natur.440.1018A,2007A&A...471..439M}.  On the other hand, if the redshift is unknown, the VHE gamma-ray spectra can be used to set an upper limit to the redshift of the source \\citep[e.g.][]{2007ApJ...655L..13M}.  After its discovery as a bright (S$_{\\rm 5 Ghz}$ $>$ 1 Jy) flat-spectrum ($\\alpha \\leq -0.5$, S$_{\\nu} \\propto \\nu^{\\alpha}$) radio source, the BL Lac object \\object{S5 0716+714} has been studied intensively at all frequency bands. The object is highly variable with rapid variations observed from radio to X-ray bands \\citep{1996AJ....111.2187W,2006A&A...451..797O}. The nucleus of \\object{S5 0716+714} is typically bright in the optical \\citep[see e.g.][]{2005AJ....130.1466N}, thus previous attempts to either characterize its host galaxy \\citep{1993A&AS...98..393S, 2000ApJ...532..816U, 2002A&A...381..810P} or to determine its redshift spectroscopically \\citep{1993A&AS...98..393S, 1996MNRAS.281..425M} have not been successful. Thus, in spite of numerous variability studies of \\object{S5 0716+714}, it has never been possible to determine reliably e.g. the linear dimensions and luminosities of the varying components.  \\object{S5 0716+714} has also been detected in TeV gamma-rays \\citep{atel1500}. Given that previous estimates \\citep[e.g.][]{1996AJ....111.2187W} place \\object{S5 0716+714} at z $>$ 0.3, this would make the target one of the most distant TeV sources detected so far. Given the obvious importance of \\object{S5 0716+714} to EBL studies, it is be important to secure an accurate spectroscopic redshift for \\object{S5 0716+714} or at least constrain the redshift by some other means.  \\object{S5 0716+714} is regularly monitored at Tuorla Observatory as part of the blazar monitoring program\\footnote{http://users.utu.fi/$\\sim$kani/1m/index.html}. On December 17, 2007 we observed \\object{S5 0716+714} to go into a fairly deep optical minimum (R $\\sim$ 14.8) and initiated prompt i-band imaging at the Nordic Optical Telescope (NOT) to detect its host galaxy. The imaging was performed five days after the minimum and the results are reported in this Letter.  Throughout this Letter, we use the cosmology $H_0 = 70$ km s$^{-1}$ Mpc$^{-1}$, $\\Omega_{M}$ = 0.3 and $\\Omega_{\\Lambda}$ = 0.7. ",
        "conclusions": "A deep i-band image of the BL Lacertae object \\object{S5 0716+714} acquired at the Nordic Optical Telescope while the target was in an low optical state has enabled us to detect the underlying host galaxy. The host galaxy has an I-band magnitude of 17.5 $\\pm$ 0.5 and an effective radius of (2.7 $\\pm$ 0.8) arcsec. Using the host galaxy as a ``standard candle'' and a method proposed by \\cite{2005ApJ...635..173S}, we derive a redshift of z = $0.31 \\pm 0.08$ (1$\\sigma$ error) for the host galaxy of \\object{S5 0716+714}. This redshift is consistent with the redshift z = 0.26 determined by spectroscopy for 3 galaxies close to \\object{S5 0716+714}.  The corresponding effective radius of the host galaxy at z = 0.31 would be $12 \\pm 4$ kpc, also consistent with the values obtained for BL Lac host galaxies in other studies. An optical spectrum obtained at the same epoch shows a featureless continuum with no identifiable spectral lines."
    },
    "0807/0807.4036_arXiv.txt": {
        "abstract": "The Dark UNiverse Explorer (DUNE) is a wide-field imaging mission concept whose primary goal is the study of dark energy and dark matter with unprecedented precision. To this end, DUNE is optimised for weak gravitational lensing, and also uses complementary cosmological probes, such as baryonic oscillations, the integrated Sachs-Wolf effect, and cluster counts. Immediate additional goals concern the evolution of galaxies, to be studied with groundbreaking statistics, the detailed structure of the Milky Way and nearby galaxies, and the demographics of Earth-mass planets. DUNE is a medium class mission consisting of a 1.2m telescope designed to carry out an all-sky survey in one visible and three NIR bands (1deg$^2$ field-of-view) which will form a unique legacy for astronomy.  DUNE has been selected jointly with SPACE for an ESA Assessment phase which has led to the Euclid merged mission concept. ",
        "introduction": "\\label{intro} Dark energy and dark matter comprise the bulk of the mass-energy budget of the Universe and pose some of the most fundamental questions in physics.  The Dark UNiverse Explorer\\footnote{{\\tt http://www.dune-mission.net}} (DUNE) is a wide field mission concept designed to study these dark cosmic components with unprecedented precision. To do so, DUNE will use weak gravitational lensing along with other cosmological probes. In these proceedings, we give a brief summary of the DUNE mission concept which was recently proposed to ESA's Cosmic Vision programme. A description of the focal plane instrumentation can be found in an adjoining paper\\cite{jeffmark} and a more detailed description of the ESA proposal is provided in a previous publication\\cite{ref08}. An earlier and simpler version of DUNE was described in previous SPIE Proceedings\\cite{gr06,ref06}.  Gravitational deflection of light by intervening dark matter concentrations causes the images of background galaxies to acquire an additional ellipticity of order of a percent, which is correlated over scales of tens of arcminutes.  Utilisation of this cosmological probe relies on the measurement of image shapes and redshifts for several billion galaxies, both requiring space observations for PSF stablity and photometric measurements over a wide wavelength range in the visible and near-IR (NIR).   Furthermore, in order to break as many degeneracies as possible, and to provide independent constraints, complementary approaches should be used. DUNE has thus been designed to provide three additional cosmological probes : Baryon Acoustic Oscillations (BAO), the Integrated Sachs-Wolfe effect (ISW), and galaxy Cluster Counts (CC). It is therefore a unique mission to probe the dark Universe in different independent ways.  DUNE will tackle the following questions: What are the dynamics of dark energy?  What are the physical characteristics of the dark matter?  What are the seeds of structure formation and how did structure grow?  Is Einstein's theory of General Relativity the correct theory of gravity?  DUNE will combine its unique space-borne observations with existing and planned ground-based surveys, and hence greatly increases the science return of the mission while limiting costs and risks. The panoramic visible and NIR surveys required by DUNE's primary science goals will afford unequalled sensitivity and survey area for further studies.  Additional surveys at low galactic latitudes and in deep patches of the sky will open new scientific windows. DUNE will explore the nature of Dark Matter by measuring precisely the sum of the neutrino masses and by testing the Cold Dark Matter paradigm. It will test the validity of Einstein's theory of gravity. In addition, DUNE will investigate how galaxies form, survey all Milky-Way-like galaxies in the 2$\\pi$ extra-galactic sky out to $z \\sim 1$ and detect thousands of galaxies and AGN at $6<z<12$. It will provide a detailed visible/NIR map of the Milky Way and nearby galaxies and provide a statistical census of exoplanets with masses above 0.1 Earth mass and orbits greater than 0.5 AU.  Furthermore, DUNE surveys will provide a unique all-sky map in the visible and NIR at unprecedented spatial resolution for such large scale of 0.25 and 0.3 arcesec respectively. It will thus complement other space missions.   The DUNE mission will allow the investigation of a broad range of astrophysics and fundamental physics. These and the corresponding surveys are described in present and following section.  The last sections describe the mission profile, payload instrument and the current status of the mission concept.        \\subsection{The Dark Universe} \\subsubsection{The nature of Dark Energy}  A variety of independent observations overwhelmingly indicate that the cosmological expansion began to accelerate when the Universe was around half of its present age. Presuming the correctness of General Relativity this requires a new energy component known as Dark Energy. The simplest case would be Einstein Cosmological Constant ($\\Lambda$), in which the dark energy density would be exactly homogeneous and independent of time. However, such an interpretation conflicts with predictions of vacuum energy from Particle Physics by 120 orders of magnitude. For this reason, cosmologists are motivated to consider models of a dynamical dark energy, or even to contemplate modifications to General Relativity.   DUNE will deduce the expansion history using the distance-redshift relation ($D(z)$) and growth of structure. It can thus probe the nature and properties evolution of dark energy in two independent ways.  A single accurate technique can rule out many of the suggested members of the family of dark energy models, but it cannot test the fundamental assumptions about gravity theory. If General Relativity is correct, then either $D(z)$ or the growth of structure can determine the expansion history. In more radical models that violate General Relativity, however, this equivalence between $D(z)$ and growth of structure does not apply.  For this purpose, DUNE will use a combination of the following cosmological probes.  {\\it Weak Lensing - A Dark Universe Probe:} As light from galaxies travels towards us, its path is deflected by the intervening mass density distribution, causing the shapes of these galaxies to appear distorted by a few percent. The weak lensing method measures this distortion by correlating the shapes of background galaxies to probe the density field of the Universe. By dividing galaxies into redshift (or distance) bins, we can examine the growth of structure (as a function of redshift) and make three-dimensional maps of the dark matter. As the evolution of the growth of structure is different in different scenarios of Dark Energy and dark matter, we have access to the nature of these dark components.  An accurate lensing survey, therefore, requires precise measurements of the shapes of galaxies as well as information about their redshifts. High-resolution images of large portions of the sky are required, with low levels of systematic errors that can be achieved only via observations from a thermally stable satellite in space. Analyses of the dark energy require precise measurements of both the cosmic expansion history and the growth of structure. Weak lensing stands apart from all other available methods because it is able to make accurate measurements of both effects\\cite{ama07,ama07b}.  Given this, the optimal dark energy mission (and dark sector mission) is one that is centred on weak gravitational lensing and is complemented by other dark energy probes.  {\\it Baryon Acoustic Oscillations (BAO) -- An Expansion History Probe:} The scale of the acoustic oscillations caused by the coupling between radiation and baryons in the early Universe can be used as a 'standard ruler' to determine the distance-redshift relation. Using DUNE, we can perform BAO measurements using photometric redshifts yielding the three-dimensional positions of a large sample of galaxies. All-sky coverage in the NIR enabled by DUNE, but impossible from the ground, is essential to reach the necessary photometric redshift accuracy for this BAO survey.  {\\it Cluster Counts (CC) -- A Growth of Structure Probe:} Counts of the abundance of galaxy clusters (the most massive bound objects in the Universe) as a function of redshift are a powerful probe of the growth of structure. There are several ways to exploit the DUNE large-area survey for cluster detection: weak lensing, strong lensing, optical richness, and cross-correlation with X-ray or Sunyaev-Zeldovich surveys (eg. Planck, E-Rosita, etc).  {\\it Integrated Sachs-Wolfe (ISW) Effect -- A Higher Redshift Probe:} The ISW effect is the change in CMB photon energy as it passes through a changing potential well.  Its presence indicates either space curvature, a dark energy component or a modification to gravity. The ISW effect is measured by cross-correlating the CMB with a foreground density field covering the entire extra-galactic sky, as measured by DUNE. The presence and shape of the spectrum of this cross-correlation will thus inform us about the nature of dark energy. Because it is a local probe of structure growth, ISW will place complementary constraints on dark energy, at higher redshifts, relative to the other probes\\cite{dous08}.  The combination of all these probes will provide strong constraints on the properties and evolution of dark energy (see Fig.~\\ref{figfom}).  \\begin{figure} \\begin{center} \\begin{tabular}{c} \\includegraphics[width=0.4\\textwidth,angle=0]{Errors2_new.ps} \\end{tabular} \\end{center} \\caption{Expected errors on the dark energy equation of state parameters (in a scenario where $w(a) = w_0+(1-a) w_a$) for the four probes used by DUNE. The light blue band indicates the expected errors from Planck. By combining BAO, CC and ISW we begin to reach similar accuracies to that of the lensing survey. This combination thus gives DUNE two independent and comparable measures of dark energy. } \\label{figfom} \\end{figure}   \\subsubsection{Properties of Dark Matter}  Besides dark energy, one major component of the concordance model of cosmology is dark matter ($\\sim90$\\% of the matter in the Universe, and $\\sim 25$\\% of the total energy). The standard assumption is that the dark matter particle(s) is cold and non-collisional (CDM). Its nature \\emph{may} well be revealed by experiments such as the Large Hadron Collider (LHC) at CERN, but its physical properties may prove to be harder to pin down without astronomical input. One way of testing this is to study the amount of substructure in dark matter halos on scales 1-100'', which can be done using high order galaxy shape measurements and strong lensing with DUNE. Weak lensing measures with DUNE can constrain the total neutrino mass and number of neutrino species through observations of damping of the matter power spectrum on small scales. Combining DUNE measurements with Planck data would reduce the uncertainty on the sum of neutrino masses to 0.04eV, and should therefore measure the neutrino mass\\cite{kit08}.   \\subsubsection{The Seeds of Structure Formation}  It is widely believed that cosmic structures originated during inflation from vacuum fluctuations in primordial quantum fields stretched to cosmic scales. In the most basic inflation models, the power spectrum of these fluctuations is predicted to be close to scale-invariant.  As the Universe evolved, these initial fluctuations grew.  CMB measurements probe their imprint on the radiation density at $z \\sim 1100$. Density fluctuations continued to grow into the structures we see today. Combined with Planck, the DUNE weak lensing observations and auto-correlation function of galaxies  will provide a measurement of the shape of the primordial spectrum with unprecedent accuracy ($1\\%$ on $\\sigma_8$ and $n_s$).  \\subsubsection{Probing Einstein's Gravity}  Various modifications to gravity on large scales (e.g. by extra dimensions, superstrings, non-minimal couplings or additional fields) have been suggested to avoid the need to invoke dark matter and dark energy. The weak lensing measurements of DUNE will be used to test the underlying theory of gravity, using the fact that modified gravity theories typically alter the relation between geometrical measure and the growth of structure.  DUNE will be able to measure the growth factor exponent $\\gamma$ (signature of the deviation from Einstein gravity\\cite{amend08}) with a precision of 2\\%.  \\subsection{Tracking the Formation of Galaxies and AGNs} In order to disentangle the complex processes involved in galaxy formation, we need high-resolution near-IR space-based images to study galaxy morphology and large area for crucial, rare events, such as the merger rate of very massive galaxies; DUNE will provide this to $z \\sim 1.5$.  Using DUNE's weak lensing maps, we will study the relationship between galaxy mass and light, the bias, by mapping the total mass density and stellar mass, luminosity, morphological type of glaxies.  While at present only a few massive clusters at $z>1$ are known, DUNE will find hundreds of Virgo-cluster-mass objects at $z>2$, and several thousand clusters of M=$1-2 \\times 10^{13}$Mo; the latter are the likely environments in which the peak of QSO activity at $z\\sim2$ takes place, and will hold the empirical key to understanding the peak period of QSO activity.  Using the Lyman-dropout technique in the near-IR, a deep survey (DUNE-MD) will be able to detect the most luminous objects in the early Universe ($z>6$): $\\sim 10^4$ star-forming galaxies at $z\\sim8$ and up to $10^3$ at $z\\sim10$, for SFRs $>30-100$Mo/yr. It will also be able to detect significant numbers of high-$z$ quasars: up to $10^4$ at $z\\sim7$, and $10^3$ at $z\\sim9$. By applying the Gunn-Peterson test to this large statistically relevant sample of objects we will put stringent constraints on the end of the period of reionisation of the Universe.  DUNE will also detect a very large number of strong lensing systems: about $10^5$ galaxy-galaxy lenses, $10^3$ galaxy-quasar lenses and 5000 strong lensing arcs in clusters\\cite{meneg07}. It is also estimated that several tens of galaxy-galaxy lenses will be double Einstein rings\\cite{gav08}, which are powerful probes of the cosmological model as they simultaneously probe several redshifts.  In addition, during the course of the DUNE-MD survey (over 6 months), we expect to detect $\\sim 3000$ Type Ia Supernovae with redshifts up to $z\\sim0.6$. This will yield a measurement of the SN rate with unprecedented accuracy, thus providing information on their progenitors. This survey will also discover a comparable number of Core Collapse SNe (Types II and Ib/c), out to $z\\sim0.3$, whose rate provides an independent measurement of the star formation history.   \\subsection{Studying the Milky Way}  DUNE will also leas to breakthroughs in Galactic astronomy. The extragalactic survey, DASS-EX, complemented by the shallower survey of the Galactic plane (DASS-G with $|b|<30$) will provide all-sky high resolution (0.23'' in the wide red band, and 0.4'' in YJH) deep imaging of the stellar content of the Galaxy, allowing the deepest detailed structural studies of the thin and thick disk components, the bulge/bar, and the Galactic halo (including halo stars in nearby galaxies such as M31 and M33) in bands which are relatively insensitive to dust in the Milky Way.  DUNE will be little affected by extinction and will supersede all of the ongoing surveys in terms of angular resolution and sensitivity (photometric depth and low background).  DUNE will thus enable the most comprehensive stellar census of late-type dwarfs and giants, brown dwarfs, He-rich white dwarfs, along with detailed structural studies, tidal streams and merger fragments.  DUNE's sensitivity will also open up a new discovery space for rare stellar and low-temperature objects via its H-band imaging.  Studying the Galactic disk components requires the combination of spatial resolution (crowding) and dust-penetration (H-band) that only DUNE can deliver. It will also yield the most detailed and sensitive survey of structure and substructure in nearby galaxies especially of their outer boundaries, thus constraining merger and accretion histories.  \\subsection{Search for Extra-Solar Planets}  Using the microlensing effect, DUNE can provide a statistical census of exoplanets in the Galaxy with masses over $0.1 M_\\oplus$ from orbits of 0.5 AU from their parent star to free-floating objects. This includes analogues to all the solar system's planets except for Mercury, as well as most planets predicted by planet formation theory. Microlensing is the temporary magnification of a galactic bulge source star by the gravitational potential of an intervening lens star passing near the line of sight. A planet orbiting the lens star, will have an altered magnification, showing a brief flash or a dip in the observed light curve.  Because of atmospheric seeing, and poor duty cycle even using networks, ground-based microlensing surveys are only able to detect a few to 15 $M_\\oplus$ planets in the vicinity of the Einstein ring radius (2-3 AU). A dedicated survey (DUNE-ML), using the high angular resolution of DUNE, and the uninterrupted visibility and NIR sensitivity afforded by space observations will provide detections of microlensing events using as sources G and K bulge dwarfs stars and therefore can detect planets down to $0.1-1 M_\\odot$ from orbits of 0.5 AU.  Moreover, there will be a very large number of transiting hot Jupiters detected towards the galactic bulge as a free ancillary science. ",
        "conclusions": "\\begin{table} \\begin{center} \\caption{DUNE Baseline summary} \\label{baseline} \\begin{tabular}{|l|l|} \\hline Science objectives & Cosmology and Dark Energy\\\\ &  Galaxy formation,   Extra-solar planets\\\\ \\hline Surveys &  20,000 deg$^2$ extragalactic 20,000 deg$^2$  galactic\\\\ & 100 deg$^2$ medium-deep, 4 deg$^2$ planet hunting\\\\ \\hline Requirements & 1 visible band (R+I+J) for high-precision\\\\ & shape measurements\\\\ &  3 NIR bands (Y, J, H) for photometry\\\\ \\hline Payload & 1.2m telescope, Visible \\& NIR cameras \\\\ & with 0.5 deg$^2$ FOV each\\\\ \\hline Service module & Mars/Venus express, Gaia heritage \\\\ \\hline Spacecraft & 2013kg launch mass\\\\ \\hline Orbit & Geosynchronous\\\\ \\hline Launch & Soyuz S-T Fregat\\\\ \\hline Operations & 4 year mission\\\\ \\hline \\end{tabular}  \\end{center} \\end{table}  The DUNE mission concept can be seen as the next step in precision cosmology. ESA's Planck mission will bring unprecedented precision to the measurement of the high redshift Universe. This will leave the dark energy dominated low redshift Universe as the next frontier in high precision cosmology. Constraints from the radiation perturbation in the high redshift CMB, probed by Planck, combined with density perturbations at low redshifts, probed by DUNE, will form a complete set for testing all sectors of the cosmological model. In this respect, a DUNE+Planck programme can be seen as the next significant step in testing, and thus challenging, the standard model of cosmology. DUNE will offer high potential for ground-breaking discoveries of new physics, from dark energy to dark matter, initial conditions and the law of gravity.  DUNE will (i) measure both effects of dark energy by using weak lensing as the central probe; (ii) place this high precision measurement of dark energy within a broader framework of high precision cosmology by constraining all sectors of the standard cosmology model (dark matter, initial conditions and Einstein gravity); (iii) through a collection of unique legacy surveys be able to push the frontiers of galaxy evolution and the physics of the local group; and finally (iv) be able to obtain information on extrasolar planets, including Earth analogues.  The DUNE concept has been recently proposed to ESA's Cosmic Vision programme and has been selected jointly with SPACE\\cite{space} for an ongoing ESA Assesment Phase which has led to the merged {\\it Euclid} mission concept."
    },
    "0807/0807.1049_arXiv.txt": {
        "abstract": "{When low-mass stars form, the collapsing cloud of gas and dust goes through several stages which are usually characterized by the shape of their spectral energy distributions. Such classification is based on the cloud morphology only and does not address the dynamical state of the object.} {In this paper we investigate the initial cloud collapse and subsequent disk formation through the dynamical behavior as reflected in the sub-millimeter spectral emission line profiles. If a young stellar object is to be characterized by its dynamical structure it is important to know how accurately information about the velocity field can be extracted and which observables provide the best description of the kinematics. Of particular interest is the transition from infalling envelope to rotating disk, because this provides the initial conditions for the protoplanetary disk, such as mass and size.} {We use a hydrodynamical model, describing the collapse of a core and formation of a disk, to produce synthetic observables which we compare to calculated line profiles of a simple parameterized model. Because we know the velocity field from the hydrodynamical simulation we can determine in a quantitative way how well our best-fit parameterized velocity field reproduces the original. We use a molecular line excitation and radiation transfer code to produce spectra of both our hydrodynamical simulation as well as our parameterized model.} {We find that information about the velocity field can reasonably well be derived by fitting a simple model to either single-dish (15$''$ resolution) lines or interferometric data (1$''$ resolution), but preferentially by using a combination of the two. The method does not rely on a specific set of tracers, but we show that some tracers work better than others. Our result shows that it is possible to establish relative ages of a sample of young stellar objects using this method, independently of the details of the hydrodynamical model.} {} ",
        "introduction": "Low-mass young stellar objects (YSOs) are well-studied phenomena. Observationally, these objects have been measured in many different wavelength regimes in order to determine intrinsic properties such as their mass, luminosity, mass accretion rate, etc. YSOs are traditionally classified by the shape of their spectral energy distribution (SED)~\\citep{lada1984}, which reflects their general morphology. Subsequently, the envelope and disk densities, temperatures and chemistry can be studied using millimeter continuum and molecular line measurements. All of these derived quantities have been used to piece together a consistent chronology of the process of star formation and to establish a reference from which the age of a YSO can be reliably obtained~\\citet{mardones1997,gregersen2000}.  A theoretical evolution scenario of YSOs has been established for many years \\citep[e.g.,][]{ulrich1976,shu1977,cassen1981,terebey1984,bodenheimer1990, basu1998,yorke1999}. In these models (with the exception of the Shu model, which is spherically symmetric) a circumstellar disk is formed, after the collapse of the initial cloud sets in, due to conservation of angular momentum. It is therefore natural to use the kinematic properties of the models to characterize the evolutionary stage of the object, because the distribution of angular momentum changes in a monotonous way as the cloud contracts and matter spins up. Unfortunately, measuring the velocity field observationally is not such a simple matter. Kinematic information can only be derived from spectral emission lines and there is no way to directly solve for the velocity field from measurements of such lines. The reason for this is that the spectral profile depends on a range of physical parameters, including the local density, temperature, molecular abundance, and turbulence in the region where the transition in consideration is excited. All of this can be taken into consideration and self-consistent models, including a general parameterized velocity field, can be built and fitted to the observed spectra. The question is, however, how reliable the derived velocity field is given the input model, and to what extent the derived velocity fields are consistent with the prediction of the models.  The velocity field of YSOs has been well studied observationally in the last 20 years. At first, mainly infall were observed \\citep[e.g.,][]{calvet1992, vanlangevelde1994}, but soon afterward, objects showing a mix of infall and rotation was discovered \\citep{Saito1996,Ohashi1997,Hogerheijde2001, belloche2002}. More recently, protoplanetary disks have been studied in very high resolution using sub-millimeter interferometry and many of these have velocity fields with rotation only \\citep[e.g.,][]{Lommen2008}. The common interpretation of these observations is that the objects showing infall only are early-type embedded objects, the rotating disks are the end product of the collapse, and the objects showing both infall and rotation are in some kind of transition. Initial modeling work using a parameterization for the velocity field has been done by \\citet{myers1996} for low-mass star formation and by \\citet{keto1990} for the formation of high-mass stars. However, it has never been attempted to quantify the degree of the transition, how far the object has passed from the embedded phase to the protoplanetary disk phase.  In this paper we will address this question, by comparing synthetic spectra, calculated from a hydrodynamical simulation of a collapsing cloud, to a model with a simple parameterized velocity field. The comparison is done in spectral space in the same way as one would compare a model to real observations. By considering different transitions of different molecules in two different resolution settings, we investigate how to obtain the most reliable characterization of the velocity field. The importance of being able to reliably model the observations of YSOs using a simple generic model becomes evident when considering the capabilities of the upcoming ALMA (Atacama Large Millimeter Array). While it is possible to build sophisticated custom models, or even directly use the hydrodynamical simulation as a template model to describe single objects, this is a very time-consuming process. ALMA will produce very high resolution observations ($\\sim$ 0.01$''$) in snapshot mode and in order to deal with such vast samples of YSOs, simple yet reliable models are needed.  The layout of this paper is as follows: In Sect.~\\ref{model} we present the hydrodynamical simulation of~\\citet{yorke1999} that we use to make synthetic observations and describes our model and the way we fit this to the synthetic spectra. Section~\\ref{lines} presents the synthetic observables and also a few analytical examples of idealized cases for comparison. Section~\\ref{results} shows our results of the spectral comparison and the derived velocity fields, followed by a discussion and conclusion in Sects. \\ref{discussion} and \\ref{conclusion}, respectively. ",
        "conclusions": "We have explored the possibility of placing Young Stellar Objects in an evolutionary ordering based on their kinematical configurations by fitting a simple parameterized model to sub-millimeter observations. We have tested the feasibility of this by fitting the simple velocity model to a time series of synthetic observations with known velocity fields, generated by a hydrodynamical simulation. We find that the model reproduces the synthetic spectra reasonably well and that the best fit parameters which describe the velocity field are in agreement with the velocity field in the hydrodynamical simulation. We therefore conclude that it is possible to extract reliable information on astronomical objects if we replace the synthetic spectra with real observations.  We find that it is difficult, though not impossible, using single-dish data alone, due to the shallowness of the $\\chi^2$-space, but feasible if interferometric data are used, especially when combined with one or more single-dish lines. Using the latter option, we find that the central dynamical mass can be determined to within 20\\%. However, we have not been able to establish a single ``best way'' to obtain the parameters throughout the collapse. We find that the best result is obtained by trying several different methods to constrain the fits, and then pick the result of the one which shows the most peaked $\\chi^2$-space. That said, most of our best constrained points were obtained using a combination of low-resolution and high-resolution and this is the preferred starting point.  We further conclude that the molecular species used are not crucial for the result, as long as the line wings are clearly defined and uncontaminated in the case of the single-dish lines, and also that the transition is optically thick. For the high-resolution data, signal-to-noise is the main concern, so strong lines are preferred (CO, HCO$^+$, etc.).  Due to the uncertainty of the absolute time scale of the collapse calculated by our relatively simple hydrodynamical simulation, we cannot yet establish an absolute age of a given object, but relative ages among a sample of objects can be obtained, provided that the velocity field of a collapsing cloud evolves similar to that of the simulation. Indeed, if an age calibration could be made, for example using the chemical properties of one or more objects, this could prove a powerful method to place a large number of young stars in an evolutionary sequence, because of the relatively straight-forward observations needed to perform the analysis. The hydrodynamical simulation presented in this paper is not required in order to analyze real objects using our method. However, with an improved and more realistic hydrodynamical scheme, it might be possible to calibrate the evolution of the velocity field so that an absolute time scale of star formation can be established. \\\\  \\noindent \\emph"
    },
    "0807/0807.0423_arXiv.txt": {
        "abstract": "We demonstrate the feasibility of uncovering supermassive black holes in late-type, ``quiescent'' spiral galaxies by detecting signs of very low-level nuclear activity. We use a combination of x-ray selection and multi-wavelength follow-up. Here, we apply this technique to NGC 3184 and NGC 5457, both of type Scd, and show that strong arguments can be made that both host AGNs. ",
        "introduction": "As discussed in Mathur et al.\\ (these Proceedings) and in \\cite{gmff08}, the only way to detect low-mass supermassive black holes (SMBHs) may be by their accretion activity. We choose to use x-ray selection to identify candidate AGNs for the following reasons: First, x-ray emission is a hallmark of AGNs. Second, x-rays can penetrate obscuring material which may be hiding the line emitting regions. Third, there are fewer sources of x-rays in a galaxy than there are of optical and UV emission and so dilution of the AGN signature by host galaxy light is less of a problem, even if the AGN is moderately obscured. Fourth, even if, as expected in some theories \\citep{n00,l03} AGNs that have luminosities or accretion rates below a cut-off value do not have broad-line regions, they should still be detectable in x-rays. The disadvantage is that x-ray observations by themselves cannot always distinguish between AGNs and other x-ray sources, such as x-ray binaries (XRBs) and ultraluminous x-ray sources (ULXs). Multi-wavelength data are needed to determine the type of source.  As examples, we present here two late-type spiral galaxies, NGC 3184 and NGC 5457 (M101). Their nuclear optical spectra show no signs of AGNs \\cite{hfs97-3}. But using archival \\cxo\\ data and multi-wavelength data from the literature, we show that the sum total of evidence strongly suggests that the galaxies host AGNs. ",
        "conclusions": "There are compelling, though not conclusive, arguments that NGC 3184 and NGC 5457 actually host AGNs even though neither shows signs of AGNs in their optical spectra. These AGNs are excellent candidates for being low-mass SMBHs, as they reside in low-mass bulges. A similar approach can be used to uncover more candidate low-mass SMBHs.   \\begin{theacknowledgments} We are grateful to D.~A.~Dale for kindly providing \\textit{Spitzer} fluxes for the nucleus of NGC 3184 prior to publication. Support for this work was provided by the National Aeronautics and Space Administration through Chandra Award Number GO7-8111X 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.  \\end{theacknowledgments}"
    },
    "0807/0807.0615_arXiv.txt": {
        "abstract": "Systematic survey for multiperiodicity in the LMC Cepheids (Moskalik, Ko{\\l}aczkowski \\& Mizerski \\cite{MKM04}, \\cite{MKM06}) has led to discovery of several new forms of pulsational behaviour. One of them is periodic amplitude and phase modulation observed in many first/second overtone (FO/SO) double mode Cepheids. In the current paper we present detailed discussion of this newly discovered phenomenon, based on a combined OGLE+MACHO sample of double mode pulsators. ",
        "introduction": "We searched for additional signal in the data using a standard consecutive pre\\-white\\-ning technique. To that effect, we first fitted the data with the double frequency Fourier sum representing pulsations in two radial modes:  \\vskip 0pt  $$m(t) = \\langle m\\rangle + \\sum_{j,k} {\\rm A}_{jk} \\sin [2\\pi(j{\\rm f}_1 + k{\\rm f}_2)t + \\phi_{jk}].\\eqno(1)$$  \\smallskip  \\noindent The frequencies of the modes, ${\\rm f}_1$ and ${\\rm f}_2$, were also optimized. The residuals of the fit were then searched for additional periodicities. This was done with the Fourier transform, calculated over the range of $0-5$\\thinspace c/d. In the next step, a new Fourier fit with all frequencies discovered so far was performed and the fit residuals were searched for additional frequencies again. The process was stopped when no new frequencies were detected. ",
        "conclusions": ""
    },
    "0807/0807.0092_arXiv.txt": {
        "abstract": "{The discovery of true solar analogues is fundamental for a better understanding of the Sun and of the solar system. Despite a number of efforts, this search has brought only to limited results among field stars. The open cluster M67 offers a unique opportunity to search for solar analogues because its chemical composition and age are very similar to those of the Sun.} {We analyze FLAMES spectra of a large number of M67 main sequence stars to identify solar analogues in this cluster.} {We first determine cluster members which are likely not binaries, by combining proper motions and radial velocity measurements. We concentrate our analysis on the determination of stellar effective temperature, using analyses of line-depth ratios and H$\\alpha$ wings, making a direct comparison with the solar spectrum obtained with the same instrument. We also compute the lithium abundance for all the stars.} {Ten stars have both the temperature derived by line-depth ratios and H$\\alpha$ wings within 100 K from the Sun. From these stars we derive, assuming a cluster reddening $E(B-V)=0.041$, the solar colour $(B-V)_\\odot=0.649\\pm0.016$ and a cluster distance modulus of 9.63. Five stars are most similar (within 60 K) to the Sun and candidates to be true solar twins. These stars have also a low Li content, comparable to the photospheric abundance of the Sun, likely indicating a similar mixing evolution.} {We find several candidates for the best solar analogues ever. These stars are amenable to further spectroscopic investigations and planet search. The solar colours are determined with rather high accuracy with an independent method, as well as the cluster distance modulus.} ",
        "introduction": "\\label{sec:Intro}  The specificity of the Sun and of our solar system have been the subject of active investigation over the last 5 decades. How typical is the Sun for a star of its age, mass, and chemical composition? How typical is that solar-type stars host planetary systems? Are they similar at all to ours?  The quest to find stellar analogues to the Sun has been going on for a long time (for an extensive review see, e.g., \\citealt{Cayrel1996}), and it stems from the poor knowledge we have of the Sun when seen `as a star' and from how typical the Sun is for a G2 type star, for its age, chemical composition, population. It is, however, after the discovery of the first exo-planets (\\citealt{MayQue1995}) that this quest became even more compelling, because to find stars similar to our own would allow us to answer to fundamental questions related to the origin of the solar system, the frequency of planetary systems similar to ours, and eventually the formation of life in other exo-planetary systems (\\citealt{Cayrel1996}). The need to identify in the night sky solar proxies to be used for spectroscopic comparison is also diffuse, in particular for the analysis of small solar system bodies (B\\\"ohnhardt, private communication).  Among the most recent results in this research, \\cite{Melendez2006} used high resolution, high signal-to-noise ratio Keck spectra to show that HD 98618 is a very close solar twin, and \\cite{King2005} proposed HD 143436 after analyzing 4 stars pre-selected from literature. These stars seem to compare well with the best known solar twin, HR 6060, first analyzed by \\cite{PorSil1997}, and subsequently confirmed by \\cite{SoubTri2004}, who made a comparative study of several hundreds of ELODIE spectra. Finally, \\cite{Melendez2007} have shown HIP~56948 to be the best solar twin known to date both in stellar parameters and in chemical composition, including a low lithium abundance.  The open cluster M67 is a perfect target to search for solar analogues. Recent chemical analyses (\\citealt{Tau2000, Randich2006, Pace2008}), show that this cluster has a chemical composition (not only Fe, but also all the other elements) extremely similar to the solar one, as close as allowed by the high precision of the measurements. The analysis resulted in [Fe/H]=$-$0.03$\\pm$0.03 for \\cite{Tau2000}, [Fe/H]=0.03$\\pm$0.01 for \\cite{Randich2006}, and [Fe/H]=0.03$\\pm$0.03 for \\cite{Pace2008}.  There are other two additional characteristics which make M67 strategical. The first one is that all the determinations of age give for this cluster an age encompassing that of the Sun (3.5-4.8 Gyr; \\citealt{Yadav2008}), while the age determination for field stars is always uncertain. The second characteristic is that M67 is among the very few clusters showing Li depleted G stars (\\citealt{Pasquini1997}). This is an important point because, as pointed out by \\cite{Cayrel1996}, even if many stars appear to have most characteristics similar to the Sun, their Li abundance is usually 10 times higher than in our star. Since Li is likely an indicator of the complex interaction taking place in the past between the stellar external layers and the hotter interior, the choice of stars which also share the same Li abundance with the Sun is an additional property to pinpoint the true analogues.  In our opinion, the search of analogues to the Sun and to the solar system can be well performed in open clusters (OCs), which show a homogeneous age and chemical composition, common birth and early dynamical environment. As a consequence, they provide an excellent laboratory for investigating the physics of solar stars and of planetary system evolution, besides being excellent probes of the structure and evolution of the Galactic disk.  M67 is a rich cluster, therefore it provides us with the opportunity to find many stars candidates sharing similar characteristics, and not only one. This is fundamental to obtain some meaningful statistics, and the cluster hosts many main sequence (MS) stars of mass around the solar mass, which form a continuous distribution (Fig.~\\ref{fig:cmd_M67}).  Finding several solar analogues in M67 will also help in providing an independent estimate of the solar colors, a quantity which still suffers of some relevant uncertainty (see, e.g., \\citealt{Holmberg2006}), as well as an independent estimate of the distance modulus of the cluster.  The present paper is the culmination of a work, which involved the chemical determination of this cluster (\\citealt{Randich2006,Pace2008}), photometry and astrometry (\\citealt{Yadav2008}) to obtain membership, and FLAMES/GIRAFFE high resolution spectroscopy to clean this sample from binaries, and to look for the best solar analogues using the line-depth ratios method (\\citealt{GrayJoha1991,Biazzo2007}) and the wings of the H$\\alpha$ line (\\citealt{CayBen1989}) to determine accurate temperatures with respect to the Sun. In addition, the Li line is used to separate Li-rich from Li-poor stars.  \\begin{figure} \\centering \\includegraphics[width=9cm]{cmd_M67.ps} \\caption{Portion of the colour-magnitude diagram of M67. Our selected targets encompass the solar colour, and are high probability proper motion (\\citealt{Yadav2008}) and radial velocity single members. The red points refer to the stars observed with FLAMES/GIRAFFE in three nights.} \\label{fig:cmd_M67} \\end{figure} ",
        "conclusions": "By using selected observations with FLAMES/GIRAFFE at the VLT, we have made a convincing case that the open cluster M67 hosts a number of interesting potential solar twins, and we have identified them. We have computed spectroscopic accurate effective temperatures for all the stars with two methods. The color-temperature relationships we derive can be used to determine temperatures for MS solar-metallicity stars.  By computing the average solar twin colours, we have obtained a precise estimate of the solar $(B-V)$: $(B-V)_{\\odot}=0.649\\pm0.016$.  By averaging the magnitude of the solar twins, we have determined an accurate distance modulus for M67: 9.63$\\pm0.06_{\\rm stat}\\pm0.05_{\\rm sys}$, which is in excellent agreement with the most recent estimates, which were based on different, independent methods and data sets.  We have determined for all the stars Li abundances, confirming the presence of a large Li spread among the solar stars of this cluster, but showing for the first time, that the Li extra-depletion appears only in stars cooler than 6000 K. The candidate solar twins have Li abundance similar to that of our star, indicating that they also share with the Sun a similar mixing history."
    },
    "0807/0807.4678_arXiv.txt": {
        "abstract": "We construct a new brane-world model composed of a bulk --with a dilatonic field--, plus a brane --with brane tension coupled to the dilaton, cold dark matter and an induced gravity term. It is possible to show that depending on the nature of the coupling between the brane tension and the dilaton this model can describe the late-time acceleration of the brane expansion (for the normal branch) as it moves within the bulk. The acceleration is produced together with a mimicry of the crossing of the cosmological constant line ($w=-1$) on the brane, although this crossing of the phantom divide is obtained without invoking any phantom matter neither on the brane nor in the bulk. The role of dark energy is played by the brane tension, which reaches a maximum positive value along the cosmological expansion of the brane. It is precisely at that maximum that the crossing of the phantom divide takes place. We also show that these results remain valid when the induced gravity term on the brane is switched off. ",
        "introduction": "\\label{sec1}  This is a very exciting time for cosmology with the overwhelming amount of new observational data that theorists try to explain. One of the biggest puzzles that the theoretical community faces is to explain the recent speed up in the universe rate of expansion discovered  first through observations from distant type Ia supernova a decade ago \\cite{Riess:1998cb,Perlmutter:1998np}. This late-time acceleration of the universe has been latter on confirmed by other independent observational probes  based for example on  measurement of  the cosmic microwave background radiation, the clustering of galaxies on very large scales and the baryon acoustic oscillations \\cite{Tegmark:2003ud,Rapetti:2004aa,Spergel:2006hy,Percival:2007yw,Giannantonio:2008zi,Komatsu:2008hk}.  A plethora of different theoretical models have been so far proposed to explain this phenomenon \\cite{Copeland}, although unfortunately none of the models advanced so far is both completely convincing and well motivated. A cosmological constant corresponding to roughly two thirds of the total energy density of the universe is perhaps the simplest  way  to explain the late-time speed up of the universe --and in addition  would match rather well the observational data. However,  the expected theoretical value of the cosmological constant is about 120 orders of magnitude larger than the value needed to fit the data \\cite{Durrer:2007re}.  An alternative approach to explain the late-time acceleration is to invoke an infrared modification of general relativity  on large scales which, by weakening the gravitational interaction on those scales, allows the recent speed up of the universal expansion. This approach is also motivated by the fact that we only have precisions observations of gravity from sub-millimiter scales up to solar system scales while the Hubble radius, which is the scale relevant for the cosmic acceleration, is many orders of magnitude larger.  A quite promising scheme in this approach is the Dvali, Gabadadze and Porrati (DGP) model \\cite{Dvali:2000hr} which  corresponds to a 5-dimensional (5D) induced gravity brane-world model \\cite{Deffayet,IG, Sahni:2002dx,LDGP2,Bouhmadi-Lopez:2004ys}, where a low-energy modification occurs with respect to general relativity; i.e. an infrared  effect takes place, leading to two  branches of solutions: (i) the self-accelerating branch and (ii)  the normal branch.  As it name would suggest the self-accelerating branch solution gives rise to a  late-time accelerating brane universe. The acceleration of the brane expansion arises naturally, even without invoking the presence of any dark energy on the brane to produce the speed-up. Not surprisingly therefore this self-accelerated feature of the DGP model has lead to considerable  research activity \\cite{reviewDGP}. Furthermore, it has been recently shown  that by embedding the DGP model in a higher dimensional space-time the ghost issue present in the original model \\cite{Koyama:2007za} may be curred  \\cite{deRham:2007xp} while preserving the existence of a self-accelerating solution \\cite{Minamitsuji:2008fz}.  The normal branch also constitutes in itself a very interesting result of the DGP model however, as it can mimic a phantom-like behaviour on the brane by means of the $\\Lambda$DGP scenario \\cite{Sahni:2002dx}. At this respect we remind the reader that observational data do not seem incompatible with a phantom-like behaviour \\cite{Percival:2007yw} and therefore we should keep an open mind about what is producing the recent inflationary era of our universe. On the other hand, this phantom-like behaviour may well be a property acquired only recently by dark energy. This leads to an interest in building models that exhibit the so called crossing of the phantom divide line $w=-1$; for example in the context of the brane-world scenario \\cite{LDGP2,crossing}. The most important aspect of the $\\Lambda$DGP model is that the phantom-like mimicry is obtained without invoking any real phantom-matter  \\cite{phantom} which is known to violate the null energy condition and induce quantum instabilities\\footnote{We are referring here to a phantom energy component described through a minimally coupled scalar field with the wrong kinetic term.} \\cite{Cline:2003gs}. In the  DGP  scenario it is also possible to obtain a mimicry of the crossing of the phantom divide at the cost of invoking a dynamical dark energy on the brane \\cite{LDGP2}, for example modelled by a quiessence fluid or a (generalised)  Chaplygin gas.  In this paper we will show that there is an alternative form of mimicking the crossing the cosmological constant line $w=-1$ in the brane-world scenario. More precisely, we consider a 5D dilatonic bulk with a brane endowed with (or without)  an induced gravity term, a brane matter content corresponding   to cold dark matter, and a brane tension $\\lambda$ that depends on the minimally coupled bulk scalar field. We will show that in this set-up the normal branch expands in an accelerated way due to $\\lambda$ playing the role of \\textit{dark energy} --through its dependence on the bulk scalar field. In addition, it turns out that $\\lambda$ grows with the brane scale factor until it reaches a maximum positive value and then starts decreasing. Therefore, in our model the brane tension mimics crossing the phantom divide. Most importantly no matter violating  the null energy density is invoked in our model.   The paper can be outlined as follows: in section \\ref{sec2} we describe the general bulk plus brane scenario; in Section \\ref{sec3} we then analyse the vacuum (i.e., $\\rho_{m}=0$) solution to prepare the ground for the analysis carried out in section \\ref{sec4}. There we show that, under some assumptions on the nature of the coupling parameters between $\\lambda$ and $\\phi$, $1+w_{\\rm{eff}}$ changes sign as the brane evolves, with $w_{\\rm{eff}}$ the effective equation of state for the brane tension. In section \\ref{sec5} we show that the the presence of the induced gravity term is not a necessary ingredient in order to mimic the crossing of the phantom divide line  in this context. Our conclusions are presented in section \\ref{sec6}. Finally, in the appendix \\ref{appendix} we present an analytical  proof for the mimicry of the crossing of the phantom divide in the model presented in section~\\ref{sec2}. ",
        "conclusions": "\\label{sec6}   In this paper we have shown the existence of a mechanism that mimics the crossing of the cosmological constant line $w=-1$ in the brane-world scenario, and which is different from the one introduced in Refs.~\\cite{Sahni:2002dx,LDGP2}. More precisely, we have shown that if we have a 5D dilatonic bulk with or without an induced gravity term on the brane (normal branch), a brane tension $\\lambda$ which depends on the minimally coupled bulk scalar field, and a  brane matter content corresponding only to cold dark matter, then under certain conditions the brane tension grows with the brane scale factor until it reaches a maximum positive value at which it mimics crossing the phantom divide, and then starts decreasing. Most importantly no matter violating  the null energy condition is invoked in our model. Despite of the transitory phantom-like behaviour of the brane tension no big rip singularity is hited along the brane evolution.  In the model with an induced gravity term on the brane (normal branch or non-self-accelerating branch), the constraint equation fulfilled by the brane tension is too complicated to be solved analytically (see Eqs.~(\\ref{Friedmann}) and (\\ref{constraint})). However, we have shown that under certain physical and mathematical conditions -cold dark matter dominates at higher redshifts and it dilutes a bite faster than dust during the brane expansion as well as a mathematical condition that guarantees the non-existence of a local minimum of the brane tension- it is possible for the brane tension to cross the cosmological constant line. The analytical proof (see the appendix) has been confirmed by numerical solutions.  Furthermore, we have shown that for some values of the parameters the normal branch inflates eternally to the future due to the brane tension $\\lambda$ playing the role of dark energy through its dependence on the bulk scalar field.  On the other hand, in the model without an induced gravity term things are much easier to analyse as it is possible to get an analytical solution for the brane tension \\cite{Kunze:2001ji}. For the analytical solution and under the same physical and mathematical conditions assumed in the model with an induced gravity effect, we have shown that the brane tension grows until it reaches its maximum positive value and then it starts decreasing driving the late-time acceleration of the brane. It is precisely at that maximum value that the brane tension mimics the crossing of the phantom divide as its effective equation of state parameter $w_{\\rm{eff}}$ is such that $1+w_{\\rm{eff}}$ changes its sign in a smooth way. We have also imposed observational bounds -from the dark matter equation of state  \\cite{Muller:2004yb} and  the big bang nucleosynthesis \\cite{Kunze:2001ji} that constraints the modified Friedmann equation and the effective gravitational constant of the brane- to constraint the parameters of the model.  In summary, in the model presented here the mimicry of the phantom divide crossing is based on the interaction between the brane and the bulk through the brane tension that depends explicitly on the scalar field that lives in the bulk. We have also shown that the brane undergoes a late-time acceleration epoch."
    },
    "0807/0807.3744_arXiv.txt": {
        "abstract": "We present a photometric analysis of the star clusters Lindsay\\,1, Kron\\,3, NGC\\,339, NGC\\,416, Lindsay\\,38, and NGC\\,419 in the Small Magellanic Cloud (SMC), observed with the \\textit{Hubble Space Telescope} Advanced Camera for Surveys (ACS) in the F555W and F814W filters. Our color magnitude diagrams (CMDs) extend $\\sim$3.5 mag deeper than the main-sequence turnoff points, deeper than any previous data. Cluster ages were derived using three different isochrone models: Padova, Teramo, and Dartmouth, which are all available in the ACS photometric system. Fitting observed ridgelines for each cluster, we provide a homogeneous and unique set of low-metallicity, single-age fiducial isochrones. The cluster CMDs are best approximated by the Dartmouth isochrones for all clusters, except for NGC\\,419 where the Padova isochrones provided the best fit. Using Dartmouth isochrones we derive ages of $7.5 \\pm 0.5$~Gyr (Lindsay\\,1), $6.5 \\pm 0.5$~Gyr (Kron\\,3), $6 \\pm 0.5$~Gyr (NGC\\,339), $6 \\pm 0.5$~Gyr (NGC\\,416), and $6.5 \\pm 0.5$~Gyr (Lindsay\\,38). The CMD of NGC\\,419 shows several main-sequence turn-offs, which belong to the cluster and to the SMC field. We thus derive an age range of 1.2-1.6~Gyr for NGC\\,419. We confirm that the SMC contains several intermediate-age populous star clusters with ages unlike those of the Large Magellanic Cloud (LMC) and the Milky Way (MW). Interestingly, our intermediate-age star clusters have a metallicity spread of $\\sim$0.6~dex, which demonstrates that the SMC does not have a smooth, monotonic age-metallicity relation. We find an indication for centrally concentrated blue straggler star candidates in NGC\\,416, while for the other clusters these are not present. Using the red clump magnitudes, we find that the closest cluster, NGC\\,419 ($\\sim$50~kpc), and the farthest cluster, Lindsay\\,38 ($\\sim$67~kpc), have a relative distance of $\\sim$17~kpc, which confirms the large depth of the SMC. The three oldest SMC clusters (NGC\\,121, Lindsay\\,1, Kron\\,3) lie in the north-western part of the SMC, while the youngest (NGC\\,419) is located near the SMC main body. ",
        "introduction": "Star clusters are powerful tools for probing the star-formation history and the associated chemical evolution of a galaxy. As one of the closest star forming galaxies with star clusters covering a wide range of ages, the Small Magellanic Cloud (SMC) is a preferred location for detailed studies of this class of objects. The SMC is the only dwarf galaxy in the Local Group containing populous intermediate-age star clusters of all ages. The SMC appears to be part of a triple system together with the Large Magellanic Cloud (LMC) and the Milky Way (MW). Its star formation activity may be triggered by interactions with its companions \\citep[e.g.][]{yoshi03}. The proximity of the SMC allows us to resolve individual stars in compact and massive star clusters of intermediate and old age, down to the sub-solar stellar mass regime.  The globular cluster (GC) system of the MW exhibits a range of ages between $\\sim$10.5 and 14~Gyr \\citep[e.g.,][]{deang05} with the oldest populations belonging to the most ancient surviving stellar systems. In the Galactic halo, a ''young'' group of star clusters is found with Pal\\,1 being the youngest with an age of $8 \\pm 2$~Gyr \\citep{Rosenberg98}. Theories explaining the origin of these so-called young halo clusters, consider them to have been captured by the MW \\citep{buon95}, to have been formed during interactions between the MW and the Magellanic Clouds \\citep{fupe95}, or to have been accreted from destroyed and/or merged dwarf satellites \\citep[e.g.,][]{zinn93,mack04}.  The star formation history of the LMC shows pronounced peaks that coincide with the times of possible past close encounters between the LMC, SMC and MW, indicative of interaction-triggered cluster formation \\citep[e.g.,][]{gir95}. In the LMC, two epochs of cluster formation have been observed that are separated by an ''age gap'' of about 4-9~Gyr \\citep[e.g.,][]{holtz99, john99, Harris01}. In the early epoch a well-established population of metal-poor ($\\langle[Fe/H]\\rangle\\sim -2$) star clusters with comparable properties to Galactic halo clusters \\citep{sunt92, olsen98, dutra99} was formed. These clusters are as old as the oldest globular clusters in the MW and in the Galactic dwarf spheroidal companions \\citep{greb04}. In a second epoch, a large population of intermediate-age clusters with ages less than 3-4~Gyr have developed.  In contrast, the SMC contains only one old GC, NGC\\,121, which is 2-3~Gyr younger than the oldest GC in the LMC and MW \\citep{glatt08} (Paper~I). The second oldest SMC star cluster, Lindsay\\,1, has an age of $7.5 \\pm 0.5$~Gyr, and since then compact populous star clusters have formed fairly continuously until the present day \\citep[e.g.,][]{daco02}. Furthermore, the intermediate-age clusters in the SMC might survive for a Hubble time, due to their high mass and the structure of the SMC (no bulge or disk to be passed) \\citep{hunter03,lamers05,gieles07}.  The existing age determinations to this point have often been associated with large uncertainties. Stellar crowding, field star contamination and faintness of the main-sequence turnoffs made the measurement of precise ages difficult. These problems affect in particular ground-based data. Another difficulty is the large depth extent of the SMC which exacerbates the distance modulus-reddening degeneracy for each cluster. These uncertainties can affect the age determination considerably.  The capabilities of the Advanced Camera for Surveys (ACS) aboard the \\textit{Hubble Space Telescope} (HST) provide an improvement both in sensitivity (depth) as well as angular resolution, which is essential for reliable photometric age determinations in dense clusters. We present improved cluster ages and distance determinations for Lindsay\\,1, Kron\\,3, NGC\\,416, NGC\\,339, Lindsay\\,38, and NGC\\,419. This is part of a ground-based and space-based program to uncover the age-metallicity evolution of the SMC. Our space-based imaging data were obtained with HST/ACS and our ground-based spectroscopy was obtained with \\textit{Very Large Telescope} (VLT). We combine our photometric results with spectroscopic metallicity determinations to obtain a well-sampled age-metallicity relation.  The age-metallicity relation determined so far indicated that SMC clusters of similar age may differ by several tenths of dex in metallicity \\citep{daco98}. Previous studies provided ages and metallicities of SMC star clusters using a variety of techniques and telescopes (see $\\S$~\\ref{sec:ana}). Combining all published cluster ages for e.g. Kron\\,3 (5-10~Gyr) \\citep{gasc66,alc96,migh98,udal98,rich00}, we find a wide range of ages for some key star clusters depending on the method used for the determination.  \\begin{deluxetable*}{cccccc} \\tabletypesize{\\scriptsize} \\tablecolumns{6} \\tablewidth{0pc} \\tablecaption{Journal of Observations} \\tablehead{ \\colhead{Cluster} & \\colhead{Date} & \\colhead{} & \\colhead{Total Exposure Time} & \\colhead{R.A.} & \\colhead{Dec.} \\\\ \\colhead{} & \\colhead{yy/mm/dd} & \\colhead{Filter} & \\colhead{(s)} & \\colhead{} & \\colhead{} } \\startdata Lindsay\\,1     &2005$/$08$/$21& F555W & 40.0   & $0^h03^m53.19^s$ & $-73\\arcdeg28'15.74''$ \\\\ &&       & 1984.0 & $0^h03^m52.66^s$ & $-73\\arcdeg28'16.47''$\\\\ && F814W & 20.0   & $0^h03^m53.19^s$ & $-73\\arcdeg28'15.74''$\\\\ &&\t& 1896.0 & $0^h03^m52.66^s$ & $-73\\arcdeg28'16.47''$\\\\ Kron\\,3        &2006$/$01$/$17& F555W & 40.0   & $0^h24^m41.64^s$ & $-72\\arcdeg47'47.49''$\\\\ &&\t& 1984.0 & $0^h24^m41.92^s$ & $-72\\arcdeg47'45.49''$\\\\ && F814W & 20.0   & $0^h24^m41.64^s$ & $-72\\arcdeg47'47.49''$\\\\ &&\t& 1896.0 & $0^h24^m41.92^s$ & $-72\\arcdeg47'45.49''$\\\\ NGC\\,339       &2005$/$11$/$28& F555W & 40.0   & $0^h57^m47.40^s$ & $-74\\arcdeg28'26.25''$\\\\ &&\t& 1984.0 & $0^h57^m47.13^s$ & $-74\\arcdeg28'24.16''$\\\\ && F814W & 20.0   & $0^h57^m47.40^s$ & $-74\\arcdeg28'26.25''$\\\\ &&\t& 1896.0 & $0^h57^m47.13^s$ & $-74\\arcdeg28'24.16''$\\\\ NGC\\,416       &2006$/$03$/$08& F555W & 40.0   & $1^h07^m53.59^s$ & $-72\\arcdeg21'02.47''$\\\\ (WFC)\t       &&\t& 1984.0 & $1^h07^m54.09^s$ & $-72\\arcdeg21'01.79''$\\\\ && F814W & 20.0   & $1^h07^m53.59^s$ & $-72\\arcdeg21'02.47''$\\\\ &&\t& 1896.0 & $1^h07^m54.09^s$ & $-72\\arcdeg21'01.79''$\\\\ NGC\\,416       &2006$/$08$/$12& F555W & 70.0   & $1^h07^m58.76^s$ & $-72\\arcdeg21'19.70''$\\\\ (HRC)\t       &&\t& 1200.0 & $1^h07^m58.96^s$ & $-72\\arcdeg21'19.30''$\\\\ && F814W & 40.0   & $1^h07^m58.76^s$ & $-72\\arcdeg21'19.70''$\\\\ &&\t& 1036.0 & $1^h07^m58.96^s$ & $-72\\arcdeg21'19.30''$\\\\ Lindsay\\,38    &2005$/$08$/$18& F555W & 40.0   & $0^h48^m57.14^s$ & $-69\\arcdeg52'01.77''$\\\\ &&\t& 1940.0 & $0^h48^m56.76^s$ & $-69\\arcdeg52'03.07''$\\\\ && F814W & 20.0   & $0^h48^m57.14^s$ & $-69\\arcdeg52'01.76''$\\\\ &&\t& 1852.0 & $0^h48^m56.76^s$ & $-69\\arcdeg52'03.07''$\\\\ NGC\\,419       &2006$/$01$/$05& F555W & 40.0   & $1^h08^m12.53^s$ & $-72\\arcdeg53'17.72''$\\\\ (WFC)\t       &&\t& 1984.0 & $1^h08^m12.71^s$ & $-72\\arcdeg53'15.49''$\\\\ && F814W & 20.0   & $1^h08^m12.53^s$ & $-72\\arcdeg53'17.72''$\\\\ &&       & 1896.0 & $1^h08^m12.71^s$ & $-72\\arcdeg53'15.49''$\\\\ NGC\\,419       &2006$/$04$/$26& F555W & 70.0   & $1^h08^m17.93^s$ & $-72\\arcdeg53'03.60''$\\\\ (HRC)\t       &&\t& 1200.0 & $1^h08^m17.78^s$ & $-72\\arcdeg53'02.80''$\\\\ && F814W & 40.0   & $1^h08^m17.93^s$ & $-72\\arcdeg53'03.60''$\\\\ &&       & 1036.0 & $1^h08^m17.78^s$ & $-72\\arcdeg53'08.80''$\\\\ \\enddata \\label{tab:journalobs} \\end{deluxetable*}  Here we present the deepest available photometry with HST/ACS, which allows us to carry out the most accurate age measurements obtained so far. We determine the ages of these clusters utilizing three different isochrone models, which also yields distances. In the next Section we describe the data reduction procedure. In $\\S$~\\ref{sec:method} we present the color-magnitude diagrams (CMD) of the clusters and discuss their main features. In $\\S$~\\ref{sec:age} we describe our age derivation method and present our results. We give an estimate of the distances of our clusters long the line-of-sight in $\\S$~\\ref{sec:dist} and present a discussion and a summary in Sections $\\S$~\\ref{sec:ana} and $\\S$~\\ref{sec:sum}, respectively.  \\begin{figure*} \\epsscale{1.2} \\plotone{f1.eps} \\caption{Spatial Distribution in 2D of our cluster sample (red circles). The location of eight additional SMC clusters, for which reliable ages from the literature are available, is shown (blue crosses). We obtain a complete sample of all intermediate-age and old SMC star clusters (see $\\S$~\\ref{sec:age}), which we will discuss in $\\S$~\\ref{sec:dist} and~\\ref{sec:ana}. One of the clusters, Lindsay\\,116, lies outside the coordinate boundaries of the Figure. The cluster locations are shown superimposed on a star map of the SMC generated using the point source catalog of the Small Magellanic Cloud Photometric Survey \\citep{zar02} for stars with V$<$16.5~mag.} \\label{fig:clusters} \\end{figure*} ",
        "conclusions": "\\label{sec:ana}  \\subsection{Comparison of our age determination with previous studies} \\label{sec:comp}  Previous studies done by several different authors provided ages and metallicities of SMC star clusters using a variety of techniques and telescopes. Therefore, if we combine all published cluster ages, we find a wide range of ages and metallicities for a given cluster, depending on the method used for the determination: Lindsay\\,1 has an age range from 7.3-10~Gyr \\citep{gasc66,gasc80,ole87,sara95,migh98,udal98,alc03}, Kron\\,3 from 5-10~Gyr \\citep{gasc66,rich84,alc96,migh98,udal98,rich00}, and NGC\\,339 from 5-7.9~Gyr \\citep{migh98,udal98,rich00}.  No other cluster has such a wide range of different age determinations as NGC\\,416, reaching from 2.5-11.2~Gyr \\citep{dur84, els85, bica86, migh98, udal98, rich00}. The cluster is located close to the SMC main body where a large interstellar extinction is expected. The separation of field stars from the real cluster members was a major problem in the age determination process, among uncertain values for metallicity, reddening and distance. Using photometry obtained with the Wide Field Planetary Camera 2 (WFPC2) aboard HST, \\citet{migh98} found an absolute age of $6.6 \\pm 0.5$~Gyr for NGC\\,416, while \\citet{rich00} derived an age of 7.1 to 11.2~Gyr using the same data set.  The only available CMD of Lindsay\\,38 is provided by \\citet{pia01}. The observation was carried out with the Cerro Tololo Inter-American Observatory (CTIO) 0.9~m telescope using the Tektronix 2K \\#~3 CCD. They presented the first age determination of Lindsay\\,38 with $6 \\pm 0.6$~Gyr. For NGC\\,419, the latest CMD was published by \\citet{rich00} based on WFPC2 data. \\citet{udal98} published an age of 3.3~Gyr and \\citet{rich00} give an age range of $1.0-1.8$~Gyr.  For Lindsay\\,1, Kron\\,3, NGC\\,339, NGC\\,416, and NGC\\,419, the latest and deepest available CMD was provided with WFPC2 \\citep{migh98,rich00}, while for Lindsay\\,38 only ground-based data existed.  \\begin{deluxetable*}{ccccccc} \\tablecolumns{7} \\tablewidth{0pc} \\tablecaption{Age Comparison} \\tablenote{Comparison of our ages derived with the Dartmouth isochrones. For NGC\\,419 Padova isochones provided the best fit. } \\tablehead{ \\colhead{Ref.Source} & \\colhead{Lindsay\\,1} & \\colhead{Kron\\,3} & \\colhead{NGC\\,339} & \\colhead{NGC\\,416} & \\colhead{Lindsay\\,38} & \\colhead{NGC\\,419} \\\\ \\colhead{} & \\colhead{Gyr} & \\colhead{Gyr} & \\colhead{Gyr} & \\colhead{Gyr} & \\colhead{Gyr} & \\colhead{Gyr}} \\startdata This paper\t\t   & $7.5 \\pm 0.5$ & $6.5 \\pm 0.5$ & $6 \\pm 0.5$   & $6 \\pm 0.5$   & $6.5 \\pm 0.5$ & $1.2-1.6$     \\\\ \\citet{rich00}  \t   &\t   -\t   &   $5.6-7.9$   &  $5.0-7.9$    &   $4.0-7.1$   &\t  -\t   & $1.0-1.8$     \\\\ \\citet{migh98}  \t   & $7.7 \\pm 0.4$ & $4.7 \\pm 0.7$ & $5.0 \\pm 0.6$ & $5.6 \\pm 0.6$ &\t-\t   &\t-\t   \\\\ \\citet{udal98}  \t   &   $9.0$\t   &  $7.5$\t   &\t$4.0$\t   &  $6.6$\t   &\t  -\t   &  $3.3$\t   \\\\ \\citet{sara95}  \t   & $7.3 \\pm 0.6$ &\t -\t   &\t -\t   &\t   -\t   &\t   -\t   &\t-\t   \\\\ Alcaino et al. (1996, 2003) & $9-10$\t   &  $8$\t   &\t -\t   &\t   -\t   &\t   -\t   &\t-\t   \\\\ \\citet{pia01}\t\t   &\t   -\t   &\t  -\t   &\t   -\t   &\t    -\t   & $6.0 \\pm 0.6$ &\t-\t   \\\\ \\enddata \\label{tab:age_com} \\end{deluxetable*}  In Table~\\ref{tab:age_com} we compare our ages using the best-fitting Dartmouth isochrones (except NGC\\,419, for which the Padova isochones provided the best fit) with results published in the most recent studies based on HST/WFPC2 photometry. The data reach $\\sim$2~mag below the turnoff points, while our ACS data have a depth of 3.5~mag below turnoffs. We can see that \\citet{migh98} derived a similar age for Lindsay\\,1, while for the remaining clusters in the overlapping sample they found younger ages than the ones derived here. \\citet{rich00}, who used the same WFPC2 ''snapshots'' as \\citet{migh98}, gave age ranges for certain metallicities for the clusters in their sample, which cover the ages determined in this paper.  The ages published by \\citet{udal98} using OGLE (Optical Gravitational Lensing Experiment) data, do not exhibit a general trend to older or younger ages as compared to our results, and the age difference varies for each cluster. The OGLE survey is a shallow ground-based survey with a limiting magnitude of $\\sim$21~mag. \\citet{sara95} used the B-V color difference between the HB and the RGB for star clusters with red HB morphologies for their age determination. The CMDs were obtained using data from the 2048 RCA prime-focus CCD on the CTIO 4~m telescope \\citep{ole87} and the photometry reaches V$\\sim$23~mag. The age found for Lindsay\\,1 is in excellent agreement with our result.  Alcaino et al. (1996, 2003) used photometry for Lindsay\\,1 obtained with the 1.3~m Warsaw telescope, Las Campanas Observatory and reaches V$\\sim$22~mag. The age was determined by using the so-called vertical method, based on the difference between the luminosity of the MSTO and the HB level. For Kron\\,3, the photometry was taken with the EFOSC-2 CCD camera at the 2.2~m Max-Planck-Institute telescope of ESO , La Silla, and reaches V$\\sim$23~mag \\citep{alc96}. The age was determined using isochrones. In both studies, the resulting ages are higher than our values. \\citet{pia01} were the first to publish an age for Lindsay\\,38, which is in excellent agreement with the age derived in this paper.  Most CMDs published in previous studies do not go deep enough to show a clearly outlined MSTO, which is an essential feature for most age determination techniques. \\citet{migh98} determined their cluster ages relative to the age of Lindsay\\,1, measuring the difference between the RC and the RGB and found similar ages as in this paper. Kron\\,3 is an exception for which the authors derived a younger age due to large error associated with the MS photometry. \\citet{rich00} fitted isochrones to the red clump and also calculated the difference between the MSTO and the RC ($\\Delta V^{RC}_{TO}$) in combination with the calibration of \\citet{walker92}. The cluster ages found in this paper are within the age ranges given by \\citet{rich00}. CMDs using ground-based photometry reach $\\sim$V=20~mag, which is not deep enough to show the SGBs or the MSTOs, which can lead to large age differences.   \\subsection{Age range and spatial distribution} \\label{sec:agerange}  The intermediate-age SMC star clusters Lindsay\\,1, Kron\\,3, NGC\\,339, NGC\\,416, and Lindsay\\,38 form a continuous age sequence from 6 to 7.5~Gyr. The SMC is the only dwarf galaxy of the Local Group known to contain populous star clusters in this age range. The only ''true'' globular cluster in the SMC, NGC\\,121, has an age of 10.5-11.5~Gyr (Paper~I), but is still 2-3~Gyr younger than the oldest LMC and MW globular clusters \\citep[e.g., ][]{Ole96,olsen98,john99,mack04}. Between NGC\\,121, and the second oldest cluster, Lindsay\\,1, there is a small age-gap ($\\sim$3~Gyr), in which no surviving star cluster has been formed.  In our sample, we have four clusters with ages between 6-6.5~Gyr, and one that is significantly younger (1.2-1.6~Gyr). Good quality ages are available from ground-based and space-based observations for ten additional intermediate-age SMC star clusters. Combining them with our star clusters, we obtain a complete sample of all intermediate-age and old SMC star clusters: Kron\\,28, Kron\\,44, Lindsay\\,116, Lindsay\\,32, Lindsay\\,11, NGC\\,152, NGC\\,361, NGC\\,411, Lindsay\\,113, and BS90 (Table~\\ref{tab:lit_ages}).  \\begin{deluxetable*}{ccccc} \\tablecolumns{5} \\tablewidth{0pc} \\tablecaption{Literature cluster ages} \\tablenote{Ages for nine additional intermediate-age clusters from the literature. } \\tablehead{ \\colhead{Cluster} & \\colhead{Age} & \\colhead{Data} & \\colhead{Method} & \\colhead{Ref.Source}  \\\\ \\colhead{} & \\colhead{Gyr} & \\colhead{} & \\colhead{} & \\colhead{}} \\startdata Kron\\,28     &  $2.1 \\pm 0.5$ & CTIO 0.9~m telescope / Tektronix 2K \\#~3 CCD& $\\Delta V_{MSTO}^{RC,HB}$  & \\citet{pia01}   \\\\ Kron\\,44     &  $3.1 \\pm 0.8$ & CTIO 0.9~m telescope / Tektronix 2K \\#~3 CCD& $\\Delta V_{MSTO}^{RC,HB}$  & \\citet{pia01}   \\\\ Lindsay\\,116 &  $2.8 \\pm 1.0$ & CTIO 0.9~m telescope / Tektronix 2K \\#~3 CCD& $\\Delta V_{MSTO}^{RC,HB}$  & \\citet{pia01}   \\\\ Lindsay\\,32  &  $4.8 \\pm 0.5$ & CTIO 0.9~m telescope / Tektronix 2K \\#~3 CCD& $\\Delta V_{MSTO}^{RC,HB}$  & \\citet{pia01}   \\\\ Lindsay\\,11  &  $3.5 \\pm 1.0$ & CTIO 4.0~m telescope / RCA CCD  \t    & Isochrones \t\t& \\citet{mould92} \\\\ NGC\\,152     &  $1.4 \\pm 0.2$ & HST/WFPC2\t\t\t\t    & Isochrones \t\t& \\citet{crowl01} \\\\ NGC\\,361     &  $8.1 \\pm 1.2$ & HST/WFPC2\t\t\t\t    & Isochrones \t\t& \\citet{migh98}  \\\\ NGC\\,411     &  $1.2 \\pm 0.2$ & HST/WFPC2\t\t\t\t    & Isochrones \t\t& \\citet{alsa99}  \\\\ Lindsay\\,113 &  $4.0 \\pm 0.7$ & HST/WFPC2\t\t\t\t    & $d_{B-V}$\\tablenotemark{b} & Mighell et al. (1998a) \\\\ BS90\t    &  $4.3 \\pm 0.1$ & HST/ACS  \t\t\t\t    & Isochrones \t\t& \\citet{sabbi07} \\\\ \\enddata \\tablenotetext{b}{The method used by Mighell et al. (1998a) is defined by the (B-V) color difference between the mean color of the red clump and the RGB at the level of the RGB. This value then was compared with Lindsay\\,1, NGC\\,416, and Lindsay\\,113.} \\label{tab:lit_ages} \\end{deluxetable*}  For none of these clusters deep HST photometry is available, thus their ages should be considered with some caution. For NGC\\,152, NGC\\,361, NGC\\,411, Lindsay\\,113, and BS90 ''snapshots'' are available taken with WFPC2 (reaching V$\\sim$23~mag), and ACS (BS90, reaching V$\\sim$26~mag).  Looking at Figure~\\ref{fig:clusters}, we clearly see that the youngest clusters are located near the SMC main body, while the clusters with ages higher than $\\sim$4~Gyr lie in the outer parts. NGC\\,361 seems to be an exception, but the cluster age is still uncertain, and the literature age of 8.1~Gyr probably is too high. \\citet{crowl01} determined a distance of $51.7 \\pm 1.8$~kpc for N361 whereby the cluster lies $\\sim$ 7.5~kpc ahead of the SMC center. Another exception is BS90 that lies near the SMC main body, even though the cluster has an age of $\\sim$4.3~Gyr. The three oldest SMC clusters (NGC\\,121, Lindsay\\,1, Kron\\,3) are located in the north-western part of the SMC. We note that Lindsay\\,116 cannot be seen in Figure~\\ref{fig:clusters}, because it is located $6\\arcdeg .1$ south-east of the bar and lies therefore outside the displayed area.   The closest cluster in our sample, NGC\\,419, and the farthest cluster, Lindsay\\,38, have a relative radial distance of 17~kpc from each other. We can therefore confirm that the SMC has a large extension along the line-of-sight, as was also found by \\citet{crowl01} based on its star clusters.  \\begin{figure} \\epsscale{1.2} \\plotone{f43.eps} \\caption{Age vs distance to the sun (projected distance) including different symbols for different metallicity ranges. For five clusters we found reliable distances \\citet{crowl01}, ages \\citep{alsa99,pia01,crowl01}, and metallicities (Kayser et al. 2008, in prep.). All values for BS90 were adopted from \\citet{sabbi07}. The dashed line represents  the SMC distance modulus of $(m-M)_0$ = $18.88 \\pm 0.1$~mag \\citep{Storm04}.} \\label{fig:distribution} \\end{figure}   In Figure~\\ref{fig:distribution} we show the distribution of age vs the distance to the sun of the clusters in our sample. The locations are shown relative to our adopted SMC distance and indicate that the clusters generally are distributed within $\\pm$6-7~kpc of the SMC centroid. Interesting exceptions are the younger clusters Kron\\,28, NGC\\,411, and NGC\\,419 that in projection appear near the center of the SMC. In fact, they could be located considerably closer to us (see also Fig.~\\ref{fig:proj}). Further measurements of the distance of younger clusters thus would be worthwhile. Moreover, we included five clusters for which we found reliable ages, distances, and metallicities in the literature. We divided the cluster metallicities into four groups and use different symbols for each group in the plot. Even though our plot contains only 11 clusters, we can see trends in the distributions of their properties. Age and distance from the sun appear to be correlated. The closest cluster, NGC\\,419, is also the youngest and most metal rich cluster, while the most distant cluster, Lindsay\\,38, is also the most metal poor, in spite of not being the oldest cluster.  \\begin{figure} \\epsscale{1.3} \\plotone{f44.eps} \\caption{Three dimensional distribution is shown for SMC star clusters with ages and distances derived from isochrone fits to CMDs derived from HST observations. Note that the intermediate age clusters are distributed throughout much of the extended body of the SMC. As discussed in the text, the selection of clusters is biased in that our observations generally avoided clusters in locations with high field star densities. However, this incomplete sample suggests that age and radial distance from the center of the SMC are not correlated; e.g. the younger cluster Kron\\,28, NGC\\,411, and NGC\\,419 are at large radial distances and cover a range in metallicity. The yellow star symbolizes the SMC center.} \\label{fig:proj} \\end{figure}   One could speculate that in regions at the outskirts of the double LMC-SMC system the star formation activity has been lower/slower than elsewhere, possibly with more unenriched gas, thus allowing for a more moderate enrichment. The oldest object, NGC\\,121, is not the most metal poor cluster, but the second metal poorest and the second farthest one. Its low metallicity could be the result of both a ''natural'' age-metallicity relation and a ''distance from the system'' effect.  Figure~\\ref{fig:proj} illustrates the distribution of SMC star clusters with high quality distances derived from isochrone fits to CMDs derived from HST observations. This is a highly biased sample; star clusters seen in the direction of the SMC `bar' are not preferred for these projects because of their large levels of field star contamination.  The exception in this case is the cluster BS90 that was accidently included in observations of NGC\\,346 (see Sabbi et al. 2007). The present limited data for clusters in this project show that the SMC is quite extended along the line of sight, consistent with other studies of individual stars and star clusters (see discussion in $\\S$~\\ref{sec:dist}).  This three dimensional distribution of the clusters also demonstrates the lack of trends in cluster age or metallicity with radial distance from the center of the SMC.  \\begin{deluxetable}{ccc} \\tablecolumns{3} \\tablewidth{0pc} \\tablecaption{Distances} \\tablenote{The projected distances were calculated in this paper and adopted from \\citet{crowl01,sabbi07}. Using these values and the cluster coordinates, we determined the cluster distances to the SMC center ($\\alpha = 0^h52^m44.8^s$, $\\delta = -72\\arcdeg49'43''$).} \\tablehead{ \\colhead{Cluster} & \\colhead{Projected Dist.} & \\colhead{Dist. to SMC Center}  \\\\ \\colhead{} & \\colhead{kpc} & \\colhead{kpc} } \\startdata NGC\\,121     & $64.9 \\pm 1.2$ & $8.76 \\pm 1.1$  \\\\ Lindsay\\,1   & $56.9 \\pm 1.0$ & $13.28 \\pm 1.0$ \\\\ Kron\\,3      & $60.6 \\pm 1.1$ & $7.19 \\pm 1.1$ \\\\ NGC\\,339     & $57.6 \\pm 4.1$ & $0.73 \\pm 2.0$  \\\\ NGC\\,416     & $60.4 \\pm 1.9$ & $3.94 \\pm 1.4$  \\\\ Lindsay\\,38  & $66.7 \\pm 1.6$ & $6.27 \\pm 1.3$  \\\\ NGC\\,419     & $50.2 \\pm 2.6$ & $10.83 \\pm 1.6$ \\\\ NGC\\,411     & $50.1 \\pm 1.7$ & $11.1 \\pm 1.3$  \\\\ NGC\\,152     & $59.0 \\pm 1.8$ & $5.58 \\pm 1.3$  \\\\ Kron\\,28     & $45.2 \\pm 1.7$ & $14.78 \\pm 1.3$ \\\\ Kron\\,44     & $57.7 \\pm 1.8$ & $4.37 \\pm 1.3$  \\\\ BS90\t     & $60.3$\t      & $1.23$  \t\\\\ \\enddata \\label{tab:lit_dist} \\end{deluxetable}  \\subsection{Age distribution and cluster formation history} \\label{sec:agedistribution}  By combining the ages of our sample with 9 literature ages for intermediate-age SMC star clusters listed in Table~\\ref{tab:lit_ages}, we obtain a well-observed sample of intermediate-age and old star clusters in the SMC. The cluster NGC\\,361 was excluded from the sample, because the cluster is almost certainly younger than the assumed $\\sim$8~Gyr \\citep{migh98}.  The age distribution is shown in Figure~\\ref{tab:age_hist}. In each panel we show our resulting age distribution using ages of different isochrone models (black histrograms) and the combined sample (white histograms). Since the cluster ages from the literature were derived using different data and methods, their distribution does not change.  In all three plots of Figure~\\ref{tab:age_hist}, the small age gap between $\\sim$8 and 10~Gyr can clearly be seen. In the first panel we used ages derived with the Dartmouth isochrones. \\citet{rich00} based on HST/WFPC2 found two brief cluster formation intervals with the oldest set $8 \\pm 2$~Gyr ago and the second $2 \\pm 0.5$~Gyr ago, and argued that there were gaps in between. During the older burst the clusters NGC\\,339, NGC\\,361, NGC\\,416, and Kron\\,3, and during the younger burst the clusters NGC\\,411, NGC\\,152, and NGC\\,419 have formed according to \\citet{rich00}.  Even though they used the same HST/WFPC2 data as \\citet{rich00}, \\citet{migh98} found no evidence for such cluster formation bursts. We also find no evidence for two significant bursts of star cluster formation in our SMC age distribution, but we do see a slightly enhanced cluster formation activity around 6~Gyr. In the second and the third panel we used our derived Teramo and Padova ages, respectively. The cluster formation at 6~Gyr is even more obvious for both isochrone models than in the upper panel.  Apparently, between $\\sim$5 and 6~Gyr no star cluster with sufficient mass to survive has formed, but if Lindsay\\,113 is older than the assumed 4~Gyr adopted from the literature, the cluster lies within the gap. We suggest that the SMC has formed its clusters during its entire lifetime with some epochs of more intense cluster formation activity. More detailed information about the age distribution requires additional deep observations of all remaining intermediate-age SMC star clusters.  As shown in Figure~\\ref{fig:proj} there appears to be no simple relationship between cluster position and metallicity in any age range. This perhaps is to be expected given that tidal interactions may have perturbed the orbits of star clusters after they formed or provided opportunities for clusters to form at large radii, as in the present-day SMC wing. We have to emphasize that the cluster sample shown in Figure~\\ref{fig:proj} is not complete. Only for 12 clusters reliable distances have been measured this far (see $\\S$~\\ref{sec:dist}), and these are shown in the Figure. The question of the metallicity distribution of the clusters and how this relates to age and position is more complex and beyond the scope of this paper.  \\subsection{Evolutionary history of the SMC as a whole}  Looking at the metallicities of our star clusters (Tab.~\\ref{fig:results}), we see that the SMC did not experience a smooth age-metallicity relation, even though the SMC is believed to be well-mixed at the present day \\citep[but see ][]{Grebel92,gonzalez99}. The oldest SMC star cluster, NGC\\,121, has a metallicity of [Fe/H] = $-1.46 \\pm 0.10$ and an age of 10.5-11.5~Gyr, while Lindsay\\,38 is more metal-poor with  [Fe/H] = $-1.59 \\pm 0.10$ but has an age of $6.5 \\pm 0.5$~Gyr. SMC star clusters of similar age may differ by several tenths of dex in metallicity (see also Da Costa et al. 1998, Kayser et al. 2008, in preparation). The probably most reasonable explanation involves the infall of unenriched, or less enriched gas. The Magellanic Clouds are surrounded by an extensive HI halo \\citep[e.g.,][]{dickey96}, therefore this possibility may be plausible. Another speculative explanation for the existence of those metal-poor clusters is that the SMC acquired these clusters in a past interaction with another dwarf galaxy, similar to the clusters from the Sagittarius dwarf galaxy being acquired by the Milky Way \\citep[e.g., ][]{carraro07}.  The SMC, LMC, and MW form an interacting triple system, which affects each other's star formation history (SFH). However, recent studies have suggested that the Magellanic Clouds only entered the vicinity of the MW fairly recently (e.g., Kallivayalil et al. 2006a/b). It is intriguing that the LMC has a significant age gap between $\\sim$4-9~Gyr, while the SMC formed its clusters continuously during the same time period. Moreover, the SMC appears to have a ''delayed'' globular cluster formation history and formed its first and only globular cluster, NGC\\,121, 2-3~Gyr later then the LMC or the MW.  \\begin{figure} \\epsscale{1.2} \\plotone{f45.eps} \\caption{The age distribution of 15 intermediate-age and old SMC clusters (excluding NGC\\,361) with ages derived in this paper, in Paper~I (plotted as black histograms), and adopted from \\citep[Mighell et al. 1998a/b,][]{mould92,alsa99,rich00,pia01,crowl01} (plotted as white histrograms). Since the cluster ages from the literature were derived using different data and methods, their distribution does not change. In the first panel we used the ages found using the Dartmouth models, in the second we use the Padova ages, and in the last panel the Teramo ages were used. The literature ages of NGC\\,152, NGC\\,411, and Lindsay\\,113, are based on HST/WFPC2 data \\citep[Mighell et al. 1998a/b,][]{alsa99,rich00}, while the adopted ages of Kron\\,28, Kron\\,44, Lindsay\\,11, Lindsay\\,32, and Lindsay\\,116 are derived from ground-based photometry \\citep{mould92,crowl01,pia01}. NGC\\,361 is not considered due to its uncertain age. The age distribution illustrates the continuous cluster formation with the small age gap between $\\sim$8 and 10~Gyr.} \\label{tab:age_hist} \\end{figure}  Possible orbits of the SMC, LMC, and MW have been modelled by several authors \\citep[e.g., Kallivayalil et al. 2006a/b;][]{Bekki05}. Strong tidal perturbations due to interactions could have triggered the cluster formation \\citep[e.g., ][]{Whitmore99} in the SMC. In the LMC, we find that star clusters have formed in evident bursts. The LMC has two main epochs of cluster formation \\citep[e.g., ][]{Bertelli92} and a well-known age-gap of several billion years, in which no star clusters have formed. A few globular clusters are found with coeval ages like the Galactic globular clusters \\citep[e.g., ][]{ols91,olsen98,john99}. We know only of one star cluster, ESO 121-SC03, that lies within the age-gap, which has an age of 8.3-9.8~Gyr \\citep{mackey06}. A correlation between young star clusters in the LMC and putative close encounters with the SMC and MW have been found by e.g. \\citet{gir95}, although the most recent proper motion measurements indicate that the Magellanic Clouds are currently on their first passage around the MW.  For young SMC clusters a relation between close encounters and the cluster formation history is not as obvious as for LMC clusters probably due to a smaller number of clusters \\citep{Chiosi06}. The age distribution in Figure~\\ref{tab:age_hist} shows that a slightly enhanced number of star clusters with ages around 1~Gyr is located in the SMC, which might have been produced through a cloud-cloud collision after a pericenter passage $\\sim$0.5~Gyr ago \\citep{Bekki05}. But evidently massive star clusters older than 1~Gyr have formed continuously until $\\sim$7.5~Gyr ago (Lindsay\\,1). It is not yet understood why populous star clusters older than 4~Gyr have not formed and survived continuously in the LMC, while in the SMC they did. \\citet{Bekki05} explained the different cluster formation histories of the Clouds as a difference in birth locations and initial mass of the host galaxies.  Kallivayalil et al. (2006a/b, see also Piatek et al. 2008) measured proper motions for the SMC and the LMC and used Monte Carlo simulations to model the orbits of the Clouds and the MW. While they found bound orbits for the Clouds, they also found that it was difficult to keep the Clouds bound to each other for more than 1~Gyr in the past. It is possible that the Clouds are not a bound system \\citep[see also e.g., ][]{Bekki05}, and that they are making their first passage close to the Milky Way."
    },
    "0807/0807.4508_arXiv.txt": {
        "abstract": "Measurements of the SNe~Ia Hubble diagram which suggest that the universe is accelerating due to the effect of dark energy may be biased because we are located in a 200-300 Mpc underdense `void' which is expanding 20-30\\% faster than the average rate. With the smaller global Hubble parameter, the WMAP-5 data on cosmic microwave background anisotropies can be fitted without requiring dark energy if there is some excess power in the spectrum of primordial perturbations on 100~Mpc scales. The SDSS data on galaxy clustering can also be fitted if there is a small component of hot dark matter in the form of 0.5 eV mass neutrinos. We show however that if the primordial fluctuations are gaussian, the expected variance of the Hubble parameter and the matter density are far too small to allow such a large local void. Nevertheless many such large voids have been identified in the SDSS LRG survey in a search for the late-ISW effect due to dark energy. The observed CMB temperature decrements imply that they are nearly empty, thus these real voids too are in gross conflict with the concordance $\\Lambda$CDM model. The recently observed high peculiar velocity flow presents another challenge for the model. Therefore whether a large local void actually exists must be tested through observations and cannot be dismissed {\\em a priori}. ",
        "introduction": "The Einstein-de Sitter (E-deS) universe with $\\Omega_\\mathrm{m}=1$ is the simplest model consistent with the spatial flatness expectation of inflationary cosmology. However, Type Ia supernovae (SNe~Ia) at redshift $z \\simeq 0.5$ appear $\\sim25\\%$ fainter than expected in an E-deS universe \\citep{Riess:1998cb,Perlmutter:1998np}. Together with measurements of galaxy clustering in the Two-degree Field survey \\citep{Efstathiou:2001cw} and of cosmic microwave background (CMB) anisotropies by the Wilkinson Microwave Anisotropy Probe (WMAP) \\citep{Spergel:2003cb}, this has established an accelerating universe with a dominant cosmological constant term (or other form of `dark energy') which presumably reflects the present microphysical vacuum state. This `concordance' $\\Lambda$CDM cosmology (with $\\Omega_\\Lambda \\simeq 0.7$, $\\Omega_\\mathrm{m} \\simeq 0.3$, $h \\simeq 0.7$) has passed a number of cosmological tests, including baryonic acoustic oscillations \\citep{Eisenstein:2005su} and measurements of mass fluctuations from clusters and weak lensing \\citep[e.g.][]{Contaldi:2003hi}. Further observations of both SNe~Ia \\citep{Riess:2004nr,Astier:2005qq,WoodVasey:2007jb} and the WMAP 3-year results \\citep{Spergel:2006hy} have continued to firm up the model. However there is no physical basis for this model, in particular there are two fundamental problems with the notion that the universe is dominated by vacuum energy. The first is the notorious fine-tuning problem of vacuum fluctuations in quantum field theory --- the energy scale of the cosmological energy density is $\\sim10^{-12}$ GeV, many orders of magnitude below the energy scale of $\\sim10^2$ GeV of the Standard Model of particle physics, not to mention the Planck scale of $\\sim10^{19}$ GeV \\citep[see][]{Weinberg:1988cp}. The second is the equally acute coincidence problem: since $\\rho_\\Lambda/\\rho_\\mathrm{m}$ evolves as the cube of the cosmic scale factor $a$, there is no reason to expect it to be of ${\\cal O}(1)$ {\\em today}, yet this is apparently the case. In fact what is actually inferred from observations is {\\em not} an energy density, just a value of ${\\cal O}(H_0^2)$ for the otherwise unconstrained $\\Lambda$ term in the Friedmann equation. It has been suggested that this may simply be an artifact of interpreting cosmological data in the (oversimplified) framework of a perfectly homogeneous universe in which $H_0 \\sim 10^{-42} \\mathrm{GeV} \\sim (10^{28} \\mathrm{cm})^{-1}$ is the {\\em only} scale in the problem \\citep{Sarkar:2007cx}.  In fact the WMAP results alone do {\\em not} require dark energy if the assumption of a scale-invariant primordial power spectrum is relaxed. This assumption is worth examining given our present ignorance of the physics behind inflation. We have demonstrated \\citep{Hunt:2007dn} that the temperature angular power spectrum of an E-deS universe with $h \\simeq 0.44$ matches the WMAP data well if the primordial power is enhanced by $\\sim 30\\%$ in the region of the second and third acoustic peaks (corresponding to spatial scales of $k \\sim 0.01-0.1~h~\\mathrm{Mpc}^{-1}$). This alternative model with {\\em no} dark energy actually has a slightly better $\\chi^2$ for the fit to WMAP-3 data than the `concordance power-law $\\Lambda$CDM model' and, inspite of having more parameters, has an {\\em equal} value of the Akaike information criterion used in model selection. Other E-deS models with a broken power-law spectrum \\citep{Blanchard:2003du} have also been shown to fit the WMAP data.  Moreover, an E-deS universe can fit measurements of the galaxy power spectrum if it includes a $\\sim 10\\%$ component of hot dark matter in the form of massive neutrinos of mass $\\sim 0.5$~eV \\citep{Hunt:2007dn,Blanchard:2003du}.  Clearly the main evidence for dark energy comes from the SNe~Ia Hubble diagram.  A mechanism that sets $\\Lambda=0$ is arguably more plausible than one which leads to the tiny energy density $\\rho_\\Lambda\\simeq 10^{-47}$ GeV$^4$ associated with the concordance cosmology.\\footnote{`Quintessence' models, which attempt to address the coincidence problem, also {\\em assume} that every other contribution to the vacuum energy cancels apart from that of the quintessence field.} If $\\Lambda$ is indeed zero then perhaps some effect fools us into wrongly deducing the existence of dark energy by \\emph{mimicking} a nonzero cosmological constant. It is natural to connect this effect with inhomogenities since cosmic acceleration and large scale nonlinear structure formation appear to have commenced simultaneously. This approach offers the possibility of solving the cosmological constant problems within the framework of general relativity and keeps the introduction of new physics to a minimum.\\footnote{In models that seek to explain the observations through modifications of gravity, the relevant scale of $H_0^{-1}$ has to be introduced {\\em by hand}, just as in quintessence models the quintessence field has to be given a mass of order $H_0$ --- these are technically {\\em unnatural} choices since this is an infrared scale for any microphysical theory.}  Several different ways in which inhomogenities could potentially mimic dark energy have been considered in the literature --- for reviews see \\citet{Celerier:2007jc,Buchert:2007ik,Enqvist:2007vb}. In an inhomogeneous universe averaged quantities satisfy modified Friedmann equations which contain extra terms corresponding to `backreaction' since the operations of spatial averaging and time evolution do not commute \\citep{Buchert:1999er}. The backreaction terms depend upon the variance of the local expansion rate and hence increase as inhomogenities develop. Whether backreaction can indeed account for the apparent cosmological acceleration is hotly debated and remains an open question at present \\citep{Wetterich:2001kr,Ishibashi:2005sj,Vanderveld:2007cq,Wiltshire:2007fg,Khosravi:2007bq,Leith:2007ay,Behrend:2007mf,Rasanen:2008it,Paranjape:2008jc}.  Another possibility is that inhomogeneities affect light propagation on large scales and cause the luminosity distance-redshift relation to resemble that expected for an accelerating universe. This has been investigated for a `Swiss-cheese' universe in which voids modelled by patches of Lema\\^{i}tre-Tolman-Bondi (LTB) space-time are distributed throughout a homogenous background. However, the results seem to be model dependent: some authors find the change in light propagation to be negligible because of cancellation effects \\citep{Biswas:2007gi,Brouzakis:2007zi,Brouzakis:2008uw}, whereas \\citet{Marra:2007pm} claim it can partly mimic dark energy if the voids have radius 250 Mpc \\citep{Marra:2007gc}. \\citet{Mattsson:2007tj} has noted that observers may preferentially choose sky regions with underdense foregrounds when studing distant objects such as SNe~Ia, so the expansion rate along the line-of-sight is then greater than average; such a selection effect he argues can allow an inhomogeneous universe to fit the observations without dark energy.  In this paper we are mainly interested in a `local void' (sometimes referred to as ``Hubble bubble'') as an explanation for dark energy; to prevent an excessive CMB dipole moment due to our peculiar velocity we must be located near the centre of the void. An underdense void expands faster than its surroundings, thus younger supernovae inside the void would be observed to be receding more rapidly than older supernovae outside the void. Under the assumption of homogeneity this would lead to the mistaken conclusion that the expansion rate of the Universe is accelerating, although both the void and the global universe are actually decelerating. Henceforth we use the `Hubble contrast' $\\delta_H\\equiv \\left(H_\\mathrm{in}-H_\\mathrm{out}\\right)/H_\\mathrm{out}$ to characterise the void expansion rate, where $H_\\mathrm{in}$ and $H_\\mathrm{out}$ are the Hubble parameters inside and outside the void respectively. (Other authors have used the `jump' ${\\mathcal J}\\equiv\\,H_\\mathrm{in}/H_\\mathrm{out} = 1 + \\delta_H$ to characterise the void.)  The reduced Hubble parameter $h$ is defined as usual by $H_\\mathrm{out} = 100h$ km $\\mathrm{s}^{-1}$ $\\mathrm{Mpc}^{-1}$ throughout.  The local void scenario has been investigated by several authors using a variety of methods \\citep{Celerier:1999hp,Tomita:1999qn,Tomita:2000rf,Tomita:2000jj,Tomita:2001gh,Iguchi:2001sq,Tomita:2002df,Moffat:2005yx,Moffat:2005zx,Moffat:2005ii,Mansouri:2005rf,Vanderveld:2006rb,Garfinkle:2006sb,Chung:2006xh,Alnes:2005rw,Alnes:2006uk,Alexander:2007xx,Biswas:2006ub,Caldwell:2007yu,Clarkson:2007pz,Uzan:2008qp,GarciaBellido:2008nz,Clifton:2008hv,GarciaBellido:2008gd}. In a series of papers, Tomita modelled the void as a open Friedmann-Robertson-Walker (FRW) region joined by a singular mass shell to a FRW background and found that a void with radius 200 Mpc and $\\delta_H=0.25$ fits the supernova Hubble diagram without dark energy \\citep{Tomita:2001gh}. \\citet{Alnes:2005rw} showed that a LTB region which reduces to a E-deS cosmology with $h=0.51$ at a radius of 1.4~Gpc with $\\delta_H=0.27$ can match both the supernova data and the location of the first acoustic peak in the CMB. \\citet{Alexander:2007xx} attempted to find the smallest possible void consistent with the current supernova results --- their LTB-based `minimal void' model has a radius of 350 Mpc and ${\\mathcal J} \\simeq 1.2$ i.e.  $\\delta_H \\simeq 0.2$; a void of similiar size but with $\\delta_H=0.3$ had been discussed earlier \\citep{Biswas:2006ub}. Unfortunately, since this model is equivalent to an E-deS universe with $h=0.44$ {\\em outside} the void where the Sloan Digital Sky Survey (SDSS) luminous red galaxies lie, as it stands it is unable to fit the measurements of the baryonic acoustic oscillation (BAO) peak at $z\\sim0.35$ \\citep{Blanchard:2005ev}.  LTB models of much larger voids were considered by \\citet{GarciaBellido:2008nz} (with radii of 2.3 Gpc and 2.5 Gpc and Hubble contrasts of 0.18 and 0.30 respectively) and it was demonstrated they can fit the supernova data, BAO data and the location of the first CMB peak. \\citet{Clifton:2008hv} found the best fit to the SNe Ia data for a void of radius $1.3 \\pm 0.2$ Gpc and an underdensity of about $70\\%$ at the centre and \\citet{Bolejko:2008cm} confirmed that such a void provides an excellent fit to the latest `Union dataset' \\citep{Kowalski:2008ez}. Moreover \\citet{Inoue:2006rd} have shown that the unexpected alignment of the low multipoles in the CMB anisotropy can be attributed to the existence of a local void of radius $300$~$h^{-1}\\,\\mathrm{Mpc}$. These authors also suggested that the anomalous `cold spot' in the WMAP southern sky is due to a similar void at $z \\sim 1$ and some evidence for this has emerged subsequently \\citep{Rudnick:2007kw}.  Recently, a large number of voids of varying sizes have been identified in the SDSS Luminous Red Galaxy (LRG) catalog in a search for the late integrated Sachs-Wolfe (ISW) effect due to dark energy \\citep{Granett:2008ju}.  How likely is the existence of such huge voids according to standard theories of structure formation? Statistical measures of the void distribution such as the void probability function and underdense probability function have been estimated from the 2dfGRS, SDSS and DEEP2 galaxy redshift surveys \\citep{Hoyle:2003hc,Croton:2004ac,Patiri:2005ys,Conroy:2005cx,Tikhonov:2006it,Tinker:2007zi,Tikhonov:2007di,vonBendaBeckmann:2007wt}. Void probability statistics have also been examined theoretically using analytical methods \\citep{Sheth:2003py,Furlanetto:2005cc,Shandarin:2005ea} and N-body simulations \\citep{Little:1993fe,Schmidt:2000md,ArbabiBidgoli:2001jp,Benson:2002tq,Padilla:2005ea}. However such studies have been restricted to voids with radii of 10-30 Mpc. The scales of the large voids we are considering lie in the linear regime where the variance of the Hubble contrast is directly related to the matter power spectrum $\\mathcal{P}_\\mathrm{m}\\left(k\\right)$. It has been noted (using results from \\citet{Turner:1992yi}) that above 100 Mpc linear theory predictions agree well with N-body simulation results, although on smaller scales the Hubble contrast is underestimated due to non-linear effects \\citep{Shi:1995nq}. Applying linear theory and using the measured CMB dipole velocity, \\citet{Wang:1997tp} obtained the model-independent result $\\langle\\delta_H\\rangle^{1/2}_R < 10.5\\,h^{-1}\\,\\mathrm{Mpc}/R$ in a sphere of radius $R$. (This ought to be an acceptable procedure up to scales of order 800 $h^{-1}$~Mpc --- on larger scales, relativistic corrections become increasingly important.) In this paper we update these results by determining the probability distribution of $\\delta_H$ and the density contrast on various scales using constraints on $\\mathcal{P}_\\mathrm{m}\\left(k\\right)$ from WMAP 5-year data \\citep{Komatsu:2008hk} and the SDSS galaxy power spectrum \\citep{Tegmark:2003uf}.  We find that even the `minimal local void' is {\\em extremely} unlikely if the primordial density perturbation is indeed gaussian as is usually assumed and the other LTB model voids even less so. However by the same token, the ISW effect due to the voids seen in the SDSS LRG survey \\citep{Granett:2008ju} appears to be too strong. Moreover, observed large-scale peculiar velocities appear to be much higher than expected \\citep{Kashlinsky:2008ut,Watkins:2008hf}. It would appear that the standard model of structure formation itself needs reexamination hence the existence of a large local void cannot be dismissed on these grounds. ",
        "conclusions": "A void with $\\delta_H\\simeq 0.2-0.3$ and a radius exceeding $100~h^{-1}\\,\\mathrm{Mpc}$ is required to fit the supernova data without dark energy \\citep{Tomita:2001gh,Biswas:2006ub,Alexander:2007xx}. The probability that we are situated in such a void is less than $10^{-12}$ as can be seen from Fig.\\ref{prob1}. The probability is exponentially smaller for the larger voids of Gpc size that have also been considered \\citep{Alnes:2005rw,GarciaBellido:2008nz,Clifton:2008hv}.\\footnote{ There is a further constraint on Gpc scale voids from the observed absence of a `$y$-distortion' in the spectrum of the CMB \\citep{Caldwell:2007yu} and from the `kinetic Sunyaev-Zeldovich' effect observed for X-ray emitting galaxy clusters \\citep{GarciaBellido:2008gd}. However this has no impact on smaller voids.}  However before we dismiss the possibility of a local void on these grounds we should also evaluate the probability of voids which have actually been claimed to exist elsewhere in the universe. For example it has been argued that a void with radius $200-300~h^{-1}\\,\\mathrm{Mpc}$ and an density contrast of $\\delta=-0.3$ at $z\\sim1$ can account for the WMAP `cold spot' in a $\\Lambda$CDM universe \\citep{Inoue:2006rd}. Even if we conservatively take the radius to be $150~h^{-1}\\,\\mathrm{Mpc}$ (and the same underdensity), the probability that one or more such voids lie within the volume out to $z=1$ is only $1.05_{-0.93}^{+5.24}\\times10^{-10}$.  It has been argued that the WMAP cold spot may not be a localized feature \\citep{Naselsky:2007yd} and there may be no matching void in the NVSS radio source catalogue \\citep{Smith:2008tc}, however an equally striking anomaly arises if we consider the large number of voids which have been identified in the SDSS LRG survey in a search for the late ISW effect \\citep{Granett:2008xb,Granett:2008ju}. These are of angular radius $\\sim 4^0$ corresponding to a (comoving) radius of $\\sim 50~h^{-1}\\,\\mathrm{Mpc}$ and are tabulated as having $1\\sigma$, $2\\sigma$ or $3\\sigma$ underdensities. These numbers relate to the detection significance (the likelihood of detecting the void by chance out of a Poisson distribution) rather than the likelihood of finding such underdensities in a gaussian field which we have computed in this paper (B. Granett, private communication). Moreover the observed LRGs are biased with regard to the dark matter hence the underdensities in dark matter are likely to be smaller than the quoted values.  However if \\citet{Granett:2008xb,Granett:2008ju} have indeed detected the late ISW effect as they assert, we can simply circumvent these uncertainties by requiring that the voids be large enough and/or underdense enough to yield the {\\em observed} CMB temperature decrements. To calculate the late ISW effect we consider the propagation of CMB photons to us from the last scattering surface through an intervening void. The photon temperature change caused by the void is \\begin{equation} \\frac{\\Delta T}{T} = -\\frac{2}{c^2}\\int_{a_\\mathrm{far}}^{a_\\mathrm{near}} \\frac{\\mathrm{d}\\Phi}{\\mathrm{d}a}\\,\\mathrm{d}a, \\end{equation} where $a_\\mathrm{far}$ is the scale factor when the photon crossed the far side of the void and $a_\\mathrm{near}$ is the scale factor when the photon crossed the near side of the void. The gravitational potential of a void with proper radius $r$ is \\begin{equation} \\Phi=\\frac{4\\pi G}{3}r^2\\rho_\\mathrm{b}\\delta\\left(a\\right). \\end{equation} Here the background density is given by $\\rho_\\mathrm{b}=3H^2_0\\Omega_\\mathrm{m}/8\\pi G a^3$ and the density perturbation is given by $\\delta\\left(a\\right)=D\\left(a\\right)\\delta\\left(a_0\\right)$ where $D$ is the linear growth factor. Hence \\begin{equation} \\frac{\\Delta T}{T} = \\Omega_\\mathrm{m}\\left(\\frac{R}{c/H_0}\\right)^2 \\left[\\frac{D(a_\\mathrm{far})}{a_\\mathrm{far}} - \\frac{D(a_\\mathrm{near})}{a_\\mathrm{near}}\\right]\\delta. \\label{isw} \\end{equation} Using this we calculate the expected ISW signal for the 50 highest significance voids in Table 4 of \\citet{Granett:2008xb}, employing the concordance $\\Lambda$CDM cosmology to determine $a_\\mathrm{far}$ and $a_\\mathrm{near}$ for each void from the void redshift measurements. The ISW signal is found to be only $-0.42 \\,\\mu \\mathrm{K}$ on average if the dark matter underdensities are smaller than the observed underdensities in the LRG counts by the bias factor of 2.2 (taking $\\sigma_8=0.8$). This is in contrast to the detected mean signal of $-11.3 \\,\\mu \\mathrm{K}$ which is over 20 times bigger! We must therefore conclude that the void radii and/or underdenities have been significantly underestimated. The void radii can at most be increased by a factor of 1.75 within the quoted uncertainties so the observed signal of $-11.3 \\,\\mu \\mathrm{K}$ can be matched only if the underdensities are increased by a factor of 5 (implying a bias factor of 0.2). The CMB temperature decrements of such model voids calculated using eq.(\\ref{isw}) are shown in Fig.\\ref{hist} and are (by construction) similar to the actual measurements shown in Fig.2 of \\citet{Granett:2008ju}. While such an {\\em underbias} for the observed LRGs may seem implausible, we emphasise that this is the only way in which the temperature decrements observed by \\citet{Granett:2008xb,Granett:2008ju} can be accounted for as being due to the late ISW effect.  Fig.\\ref{hist} displays a histogram of the probabilities for finding such voids in the SDSS LRG survey volume ($5\\,h^{-3}\\,\\mathrm{Gpc}^{3}$ in the redshift range $0.4 < z < 0.75$). The most improbable void is at $z=0.672$ --- in order to yield the {\\em observed} average CMB temperature decrement it must have a density contrast of -0.72 (quoted galaxy underdensity of -0.316 multiplied by 5/2.2) and a radius of $230~h^{-1}\\,\\mathrm{Mpc}$ (radius derived from the quoted volume of $10^7~h^{-3}\\,\\mathrm{Mpc}^3$ and multiplied by 1.75). The probability of such a void is $1.9 \\times 10^{-247}$ according to our calculations. Although linear theory may not be applicable for such a deep void, it is clear that its existence is in gross conflict with the standard theory of structure formation from {\\em gaussian} primordial density perturbations.  \\begin{figure*}% \\begin{minipage}{\\textwidth} \\centering \\begin{tabular}{cc} \\includegraphics*[angle=0,scale=0.5]{fig29.eps} & \\includegraphics*[angle=0,scale=0.5]{fig30.eps} \\end{tabular} \\caption{\\label{hist} The left panel shows the ISW signals of the 50 voids detected by \\citet{Granett:2008xb}, calculated using eq.(\\ref{isw}); in order to match the observed average ISW signal of $-11.3 \\,\\mu \\mathrm{K}$ it has been necessary to increase the void radii by a factor of 1.75 and the underdensities by a factor of 5. The right panel show the probability of such voids occurring in the SDSS LRG survey volume according to the concordance $\\Lambda$CDM model.} \\end{minipage} \\end{figure*}  This conclusion is strengthened by the recent detection of very large peculiar velocities on large scales. As seen in Fig.\\ref{var}, the expected variance of the peculiar velocity as calculated by eq.(\\ref{var3}) is about 200 km~s$^{-1}$ on a scale of 100 $h^{-1}$ Mpc, whereas the measured value is at least 5 times higher, and the discrepancy is even bigger on larger scales up to 300 Mpc \\citep{Kashlinsky:2008ut}.  It is also seen from Fig.\\ref{prob1} that if a determination of the Hubble constant is required with say 1\\% accuracy, then measurements extending out to at least $150~h^{-1}\\,\\mathrm{Mpc}$ must be made to overcome local fluctuations. A similar estimate was made by \\citet{Li:2008yj} who noted that the observed variance in measurements of $h$ is in accord, thus consistent with the assumption of a gaussian density field.  However the voids observed in the SDSS LRG survey \\citep{Granett:2008ju} call this assumption into question. In particular whether there is a large local void is then an issue that must be addressed observationally and not dismissed on the grounds that it is inconsistent with gaussian perturbations. The Hubble flow is presently poorly measured in the redshift range $0.1 \\la z \\la 0.3$ --- just where the effects of such a local void would be most apparent \\citep{Alexander:2007xx}. Given that dark energy may well be an artifact of such a void, this issue needs urgent attention.  The question of how such voids can have been generated without conflicting with the CMB observations is beyond the scope of the present work. Some suggestions have been made in the context of multi-field inflationary models \\citep{Occhionero:1997eh,DiMarco:2005zn,Itzhaki:2008ih}."
    },
    "0807/0807.3328_arXiv.txt": {
        "abstract": "We report HST/NICMOS coronagraphic images of the HD 15115 circumstellar disk at 1.1\\micron.  We find a similar morphology to that seen in the visible and at H band--an edge-on disk that is asymmetric in surface brightness.  Several aspects of the 1.1\\micron\\ data are different, highlighting the need for multi-wavelength images of each circumstellar disk.  We find a flattening to the western surface brightness profile at 1.1\\micron\\ interior to 2\\arcsec\\ (90~AU) and a warp in the western half of the disk.  We measure the surface brightness profiles of the two disk lobes and create a measure of the dust scattering efficiency between 0.55-1.65\\micron\\ at 1\\arcsec, 2\\arcsec, and 3\\arcsec.  At 2\\arcsec\\ the western lobe has a neutral spectrum up to 1.1\\micron\\ and a strong absorption or blue spectrum $>$1.1\\micron, while a blue trend is seen in the eastern lobe.  At 1\\arcsec\\ the disk has a red F110W-H color in both lobes. ",
        "introduction": "Raw material from the interstellar medium is processed through circumstellar disks into many different types of planetary bodies.  Debris disks are a useful environment to study the composition of planetesimals that form around stars different from our own Sun.  Multi-wavelength scattered light observations of disks provide one avenue for determining the compositon of dust caused by collisions of planetesimals.  HD~15115 possesses one of a growing number of spatially resolved debris disks amenable to multiwavelength observations.  HD 15115 is an F2 star at a distance of 45$\\pm$1~~pc \\citep{vanl} and was observed to have an IR excess \\citep{silverstone00}.  Unresolved emission from its disk was detected at 60, 100, and 850 \\micron, implying a ring at $\\sim$35~AU with a temperature of 62~K \\citep{zuckerman04,decin03,williams06}.  Recently, \\citet[][hereafter KFG07]{kalas07} discovered circumstellar emission from a combination of Hubble Space Telescope/Advanced Camera for Surveys (HST/ACS) and Keck near-IR AO coronagraphic imaging.  Surface brightness (SB) profiles were reported for V and H bands.  Furthermore, they observed an extreme size and brightness difference between the two lobes of the disk and shape asymmetries around the disk midplane similar to that seen around $\\beta$ Pictoris and indicative of a second disk at a different inclination \\citep{kalas95,golimowski06}.  We have imaged HD 15115 at 1.1\\micron\\ with HST; we present our observations in \\S\\ref{s:obs}, and analyze the results in \\S\\ref{s:analysis}.  We discuss the implications of our work in \\S\\ref{s:conc} ",
        "conclusions": "\\label{s:conc}  The new features we have observed in the disk at 1.1\\micron\\ add questions to the nature of the disk and the dust.  The increase in PA towards the inner part of the disk could be caused by a genuine warp, or perhaps the presence of a second disk inclined at $>$12$^\\circ$ from the main disk as is seen in $\\beta$-Pictoris.  This is supported by the detection of an asymmetry in the northern part of the disk seen in the ACS and H data of KFG07.  KFG07 noted an overall blue V-H color for the Western lobe of the disk.  The presence of a neutral scattering efficiency out to 1.1\\micron\\ as well as a red F110W-H color interior to 2\\arcsec\\ suggests that the spectrum of the HD~15115 disk is more complicated than the V and H data suggest.  The abrupt change in scattering efficiency beyond 2\\arcsec\\ from F110W to H may be due to the presence of strong absorption at $\\sim$1.65\\micron.  One candidate would be strong absorption due to water ice, something not yet detected in the scattered light of a disk.  If that is the case, even deeper absorption should be present at 2 and 3.5\\micron. Two possibilities exist to explain the red F110W-H color of the inner disk--either there is evidence of strong absorption due to olivine (which has an absorption feature at 1\\micron), or the presence of a very red material in the disk, possibly similar to what is seen in HD~100546 \\citep{ardila07} or HR~4796A \\citep{debes08}.  The age of HD~15115 is not well constrained, which would be helpful in understanding the origins of the debris disk.  Kinematically, it is marginally consistent with membership in the 12~Myr-old $\\beta$-Pictoris Moving Group (BPMG) \\citep{moor06}.  However, this is not borne out by backtracking the position of HD~15115 based on its proper motion and radial velocity using the methods of \\citet{ortega02} and \\citet{song03}, nor is kinematics a reliable selector of a moving group \\citep{song02}.  Most other indicators point to a $>$ 100~Myr age.  \\citet{nordstrom04} determined an age of 900~Myr, but this age is uncertain given their methodology--the lower limit for the age is 0, while the upper limit is 2.2~Gyr.  Ca~II H and K line indicators point to an age of 500~Myr \\citep{silverstone00}.  Although not an accurate age indicator for early F-type stars, the Lithium 6708\\AA\\ feature shows an equivalent width of 40~m\\AA\\ (Song~2008, in prep.), which is consistent with $\\sim$100~Myr Pleiades stars. As a comparison, previously recognized 30~Myr old (Tucana-HorA) stars or younger early F-type stars all show equivalent widths for Li of $\\sim$100m\\AA.  The evidence to date points to an older age for HD~15115.  HD~15115 is a prime example of the potential complexity of nearby circumstellar debris disks.  It possesses a strange brightness asymmetry in its outer reaches that extends into about 2\\arcsec (90~AU).  This asymmetry, however, is a function of wavelength, and at least part of the cause may be compositional or grain population differences between the two sides of the disk.  KFG07 investigated whether the close passage of a star might explain the strange morphology of the disk.  They found a candidate object, HIP~12545, that might have passed very close to the debris disk of HD~15115, but while its motion is consistent with being a co-moving object, backtracking of the the two star's positions back in time based on their Hipparcos/Tycho-2 astrometry and radial velocities shows that they are more widely separated in the past than they are now and could not have interacted.  The fact that we measure a warp only in the western part of the disk around 1\\arcsec\\ suggests that HD~15115 has an even more complex structure.  It would be useful to confirm the presence of this structure in the disk at other wavelengths.  This is not easily done in the visible, unless STIS is refurbished in the next servicing mission--it is the only instrument able to image to within $\\sim$1\\arcsec\\ in the visible.  Follow-up H and K (or F160W or F205W images with HST) will be useful to look for the warp as well as constrain the ultimate nature of the composition of the grains in the disk.  If water ice is indeed present, it will show strong absorption at around 2 and 3.5\\micron.  Therefore, K and L band images of the scattered light disk are cruicial.  What are the other possibilities to explain both the warp, the asymmetry, and the varying nature of the dust grains?  One possibility is that the asymmetry is caused by a relatively recent collision that has injected many particles recently into the disk. Such an event is rare given the collisional times of the parent bodies and the timescale for such an event to be evenly spread in azimuth.  Similar asymmetries have been seen in the mid-IR with $\\beta$-Pictoris \\citep{weinberger03}, and it is striking that HD~15115 shares many of traits with this well known disk."
    },
    "0807/0807.0810_arXiv.txt": {
        "abstract": "Measuring the statistics of galaxy peculiar velocities using redshift-space distortions is an excellent way of probing the history of structure formation. Because galaxies are expected to act as test particles within the flow of matter, this method avoids uncertainties due to an unknown galaxy density bias. We show that the parameter combination measured by redshift-space distortions, $f\\sigma_8^{\\rm mass}$ provides a good test of dark energy models, even without the knowledge of bias or $\\sigma_8^{\\rm mass}$ required to extract $f$ from this measurement (here $f$ is the logarithmic derivative of the linear growth rate, and $\\sigma_8^{\\rm mass}$ is the root-mean-square mass fluctuation in spheres with radius $8h^{-1}$Mpc). We argue that redshift-space distortion measurements will help to determine the physics behind the cosmic acceleration, testing whether it is related to dark energy or modified gravity, and will provide an opportunity to test possible dark energy clumping or coupling between dark energy and dark matter. If we can measure galaxy bias in addition, simultaneous measurement of both the overdensity and velocity fields can be used to test the validity of equivalence principle, through the continuity equation. ",
        "introduction": "Observations of the accelerating nature of the Universe show that there is fundamental physics at work that we do not understand \\cite{riess98,perlmutter99}. Many possibilities have been postulated including a new form of the vacuum which is not present in the contemporary high energy physics, or a modification of gravity which would revolutionize our understanding of space and time. Discovering which mechanism is correct is one of the key challenges for 21st century science.  The detection of acceleration was obtained by using supernovae as standard candles, and therefore relies on measuring the cosmological geometry. The physical process causing the acceleration could also affect structure formation, which provides a complementary way of distinguishing between models. In particular, models in which general relativity is unmodified have different Large-Scale Structure (hereafter LSS) formation timescales compared with Modified Gravity (hereafter MG) models~\\cite{dvali00,carroll03}. Because of this, many previous studies have considered testing the signature on LSS formation predicted by MG models~\\cite{Lue:2003ky,ishak05,jain07}. Direct observations of LSS growth as traced by galaxies are of limited value because galaxies are not expected to be simple tracers of the underlying matter density field, although we return to this point later.  Maps of galaxies where distances are measured from spectroscopic redshifts show anisotropic deviations from the true galaxy distribution. These differences arise because galaxy recession velocities include components from both the Hubble flow and peculiar velocities from the motions of galaxies in comoving space.  Although these ``redshift-space distortions'' are a nuisance when trying to reconstruct the true distribution of galaxies, they provide a mechanism to measure the build-up of structure, which drives these peculiar velocities on large-scales.  In linear theory, and in the absence of bias, a distant observer should expect a multiplicative enhancement of the overdensity field $\\delta$ that is proportional to $1+f\\mu^2$, where $f$ is the logarithmic derivative of the linear growth rate, and $\\mu$ is the cosine of the angle to the line-of-sight \\cite{kaiser87}. With a local linear bias, the real-space galaxy density field is affected, while the peculiar velocity term is not, so the multiplicative factor is changed to $1+\\beta\\mu^2$, where $\\beta\\equiv f/b$. Because of the $\\mu$ dependence, this information can be extracted from galaxy redshift surveys, and a number of methods and their application have previously been considered. Analyses have been undertaken using the 2degree-Field Galaxy Redshift Survey (2dFGRS) \\cite{colless03}, measuring redshift-space distortions in both the correlation function \\cite{peacock01,hawkins03} and power spectrum after decomposing into an orthonormal basis of spherical harmonics and spherical Bessel functions \\cite{percival04}. Using the Sloan Digital Sky Survey \\cite{York00}, an Eigenmode decomposition has been performed to separate real and redshift-space effects \\cite{tegmark04,tegmark06}. In a recent paper, these low redshift analyses were extended to $z\\simeq1$ \\cite{guzzo08} using the VIMOS-VLT Deep Survey (VVDS) \\cite{lefevre05,garilli08}.  In addition to measuring $\\beta$ at $z=0.8$, this work has pushed explicitly the idea of using large-scale peculiar velocities for constraining models of cosmic acceleration.  These data, and the resulting cosmological constraints, are considered further in Section~\\ref{sec:current}. In particular we argue that it would be better to present results in terms of $b(z)\\sigma_8(z)$ and $f(z)\\sigma_8(z)$, rather than $\\beta$ for local linear bias models. We show in Section~\\ref{sec:pec_vel_meas} that bias-independent constraints on $f(z)\\sigma_8(z)$ are able to discriminate between some models of acceleration as well as $f(z)$, which is commonly extracted from $\\beta$ by applying an independent (and difficult) measurement of bias. Galaxy bias measurements tend to have the same fractional error as the redshift-spce distortion measurements: for example comparing redshift-space distortion results from the 2dFGRS\\cite{hawkins03} with bias measurements from measurements of the 3-pt function\\cite{verde02}.  In addition to a direct measurement of $f(z)\\sigma_8(z)$, redshift-space distortion measurements can be used to test diverse aspects of LSS, as proposed by Song $\\&$ Koyama~\\cite{song08a}: geometrical perturbations can be reconstructed from the evolution history of peculiar velocities.  With the assumption of an additional measurement of galaxy bias, the continuity equation can be tested, and anisotropic stress can be constrained.  Those diverse tests strengthen our power to constrain theoretical models, and are considered further in Section~\\ref{sec:further_tests}. Before we do this, we first review the physics that we hope to test using peculiar velocities (Section~\\ref{sec:models}), and then consider the measurements themselves and what they can directly tell us about structure formation (Section~\\ref{sec:pec_vel_meas}). ",
        "conclusions": "One of the best ways of measuring this structure growth is to use observations of redshift-space distortions. This is an interesting and timely subject given current observations and those planned on a 10--20 year timescale \\cite{Wang08,sapone07,guzzo08}. In this paper we have reviewed the importance of redshift-space distortion measurements given that they provide a measurement of structure growth that is independent of galaxy density bias. We have argued that peculiar velocity measurements are best presented in terms of $\\sigma_{8}^{\\theta}$, or $f\\sigma_8^{\\rm mass}$ for models where the continuity equation holds. The independence from galaxy density bias has not been widely covered in previous literature. Most previous analyses have considered measuring $\\beta$ and $b\\sigma_8^{\\rm mass}$, although there are some exceptions\\cite{percival04}. Although extremely simple, we have focused on the density bias independent constraint resulting from multiplying these together. The physical origin of such a constraint is that linear velocities, which scale with the derivative of the growth factor, depend only on the matter velocity field.  The primary conclusion from our work is that constraints on $\\sigma_{8}^{\\theta}$ or $f\\sigma_8^{\\rm mass}$, are extremely good at helping to distinguish between the dark energy models that we reviewed in Section~\\ref{sec:models}. In fact, as we show in Fig.~\\ref{fig:f_v}, these constraints are equivalent to similar percent measurements of $f$ for some models of cosmic acceleration. They also have the simplicity of not having to model galaxy bias. The simple formula of \\cite{guzzo08} has been adapted to determine constraints on $f\\sigma_8^{\\rm mass}$, and shows future constraints in Fig.~\\ref{fig:sigma8_theta}. Although the \\cite{guzzo08} formula is simplistic, and might not be believed at the percent level, it shows that we can expect a huge step forwards in redshift-space distortion measurements with the next generation of surveys.  Going beyond simply obtaining a single measurement of $\\sigma_{8}^{\\theta}$ or $f\\sigma_8^{\\rm mass}$, we have considered how the underlying perturbation evolution can be tested using peculiar velocity measurements. Peculiar velocity measurements are important because they can be used to reconstruct Newtonian potential $\\Psi$ which sources the dynamics of a galaxy given by Euler equation. Weak-lensing only measures $\\Psi$ in the combination $\\Phi-\\Psi$, so redshift-space distortions offer a complementary test of perturbations.  We have considered how peculiar velocities can be used to test the continuity equation, which is worthwhile since there are many theoretical models which fail to satisfy this relation. If dark energy couples to matter, then current constraints show that it must couple to the dark matter and not to baryonic material. The coupling of dark energy to dark matter modifies the Euler equation for dark matter, and breaks the equivalence principle between dark matter and baryon. This difference in free-fall breaks the continuity equation in which the peculiar velocity of matter is estimated using baryons, while we consider the growth of fluctuations in all matter. In addition, this test can tell if there is dark energy clustering which deepens the curvature potential well because we measure the peculiar velocity of the matter not of the total energy density. Finally we have considered how we can constrain the anisotropic stress by comparing $\\Phi$ and $\\Psi$, reconstructed from the density fields and peculiar velocity respectively."
    },
    "0807/0807.2354_arXiv.txt": {
        "abstract": "The purpose of this work is the computation of the cosmic Type Ia supernova rates, namely the frequency of Type Ia supernovae per unit time in a unitary volume of the Universe. Our main goal in this work is to predict the Type Ia supernova rates at very high redshifts and to check whether it is possible to select the best delay time distribution model, on the basis of the available observations of Type Ia supernovae. We compute the cosmic Type Ia supernova rates in different scenarios for galaxy formation and predict the expected number of explosions at high redshift ($z\\geq 2$). Moreover, we adopt various progenitor models in order to compute the Type Ia supernova rate in typical elliptical galaxies of initial luminous masses of $10^{10}$M$_{\\odot}$, $10^{11}$M$_{\\odot}$ and  $10^{12}$M$_{\\odot}$, and compute the total amount of iron produced by Type Ia supernovae in each case. In this analysis we assume that Type Ia supernovae are caused by thermonuclear explosions of C-O white dwarfs in binary systems and we consider the most popular frameworks: the single degenerate and the double degenerate scenarios. The two competing schemes for the galaxy formation, namely the monolithic collapse and the hierarchical clustering, are also taken into account, by considering the histories of star formation increasing and decreasing with redshift, respectively. We calculate the Type Ia supernova rates through an analytical formulation which rests upon the definition of the SN Ia rate following an instantaneous burst of star formation as a function of the time elapsed from the birth of the progenitor system to its explosion as a Type Ia supernova (i.e. the delay time). What emerges from this work is that: \\emph{i)} we confirm the result of previous papers that it is not easy to select the best delay time distribution scenario from the observational data and this is because the cosmic star formation rate dominates over the distribution function of the delay times; \\emph{ii)} the monolithic collapse scenario for galaxy formation predicts an increasing trend of the SN Ia rate at high redshifts (mainly due to the contribution by massive spheroids), whereas the predicted rate in the framework of a decreasing cosmic star formation rate, more in agreement with the hierarchical scenario, drops dramatically at high redshift; \\emph{iii)} for the elliptical galaxies we note that the predicted maximum of the Type Ia supernova rate depends on the initial galactic mass. The maximum occurs earlier (at about 0.3 Gyr) in the most massive ellipticals, as a consequence of the assumed downsizing in star formation. In addition, we find that the Type Ia supernova rate per unit mass at the present time is higher in bluer ellipticals (i.e. the less massive ones). ",
        "introduction": "The supernovae (SNe) of Type Ia are fundamental for understanding a number of astrophysical problems of primary importance, such as \\emph{i)} the SN progenitors, \\emph{ii)} the determination of cosmological parameters, \\emph{iii)} the chemical enrichment of galaxies and \\emph{iv)} the thermal history of the interstellar (ISM) and intracluster medium (ICM). The evolution of the rate of Type Ia SNe with cosmic time is a fundamental ingredient for the study of all these issues. The observed features of SNe Ia suggest that the majority of these objects may originate from the thermonuclear explosion of a C-O white dwarf (WD) of mass $\\sim 1.4$ M$_{\\odot}$ (Chandrasekhar mass) in binary systems (Chandra exploders). So each SN Ia should be the result of the explosion of the same mass. However, Phillips (1993) pointed out that there is a significant intrinsic dispersion in the absolute magnitudes at maximum light of local Type Ia SNe. This result was interpreted to arise from a possible range of masses of the progenitors or from variations of the explosion mechanism. \\\\ Here we will consider only Chandra exploders for which two scenarios have been proposed: a) The \\emph{Single Degenerate} (SD) scenario, i.e. the accretion of matter via mass transfer from a non-degenerate companion, a red giant or a main sequence star (e.g. Whelan \\& Iben 1973). In this scenario the mass range for the secondary components of the binary system is $0.8-8$M$_{\\odot}$, while the primary masses should be in the range $2-8M_{\\odot}$. The upper limit is given by the fact that stars with masses $M>8M_{\\odot}$ ignite carbon in a non-degenerate core and do not end their lives as C-O WDs. The lower limit is instead obviously due to the fact that we are only interested in systems which can produce a Type Ia SN in a Hubble time. The clock to the explosion is given by the lifetime of the secondary component. b) The \\emph{Double Degenerate} (DD) scenario, i.e. the merging of two C-O WDs which reach the Chandrasekhar mass and explode by C-deflagration (e.g. Iben \\& Tutukov 1984). The merging is due to the loss of orbital angular momentum due to gravitational wave radiation. In the Iben \\& Tutukov paper, the progenitor masses were defined in the range 5-8M$_{\\odot}$ to ensure two WDs of $\\sim 0.7$M$_{\\odot}$ and then reach the Chandrasekhar mass. The clock for the explosion in this model is given by the lifetime of the secondary star plus the time necessary to merge the system due to gravitational wave radiation. This scenario requires the formation of two degenerate C-O WDs at an initial separation less than $\\sim 3 R_{\\odot}$ and this can occur by means of two different precursor systems: a close binary, and a wide binary. The two scenarios differ for the efficiency of the common envelope phase during the first mass transfer, and therefore for the separation attained at the end of the first common envelope phase.  Different arguments can be found in favor or against both SD and DD scenarios and the issue of DD vs. SD is still debated (e.g. Branch et al. 1995, Napiwotzki et al. 2002, Belczynski et al. 2005). The Type Ia SN rate is the convolution of the distribution of the explosion times, usually called the time delay distribution function (DTD), with the star formation history. Several attempts of comparing the Type Ia SN rate and the cosmic star formation rate evolution with redshift have already appeared (e.g. Gal-Yam \\& Maoz, 2004; Dahlen et al. 2004; Cappellaro et al. 2005; Neill et al. 2006; Forster et al. 2006; Barris \\& Tonry, 2006;  Poznansky et al. 2007; Botticella et al. 2008, Blanc \\& Greggio, 2008; among others), but  no clear conclusions arised yet on this point.\\\\ In this paper, besides the above mentioned scenarios for Type Ia SN progenitors we will test the empirical DTD suggested by Mannucci et al. (2006),together with different star formation histories. In particular, we will study the Type Ia SN rate in ellipticals and different cosmic star formation rates, including a strongly increasing cosmic star formation with redshift, based on the monolithic scenario of galaxy formation and never considered in previous studies.\\\\ The paper is organized as follows: in Section 2 the adopted formulation for the SN Ia rate is presented, together with the description of the adopted star formation histories. In section 3 the models for the distribution function of the delay times are presented and in section 4 the predicted SN Ia rates are discussed and compared with the availale data. Finally, in section 5 some conclusions are drawn. ",
        "conclusions": "We calculated the cosmic SN Ia rate densities (i.e. the rate per unit comoving volume) and the rate of the explosion of SNe Ia in typical elliptical galaxies adopting the formalism proposed by G05. This formulation rests upon the definition of the SN Ia rate following an instantaneous burst of star formation as a function of the time elapsed from the birth of the progenitor system to its explosion as a Type Ia SN (i.e. the delay time). This function has been termed as DTD and accounts for the SNe Ia progenitor scenarios and for the initial mass function (IMF). Different DTDs and histories of SF have been considered. In all cases, a Salpeter IMF  was adopted. For all galaxies it was assumed a galaxy  formation epoch at redshift $z_{f}=6$. The reason for chosing $z_f=6$ is given by the fact that the SFR density has been measured up to this redshift (Cole et al. 2001). We have also shown the results obtained assuming $z_{f}=10$ as epoch of galaxy formation. The evolution of the cosmic SN rate with redshift contains, in principle, unique information on the star formation history of the Universe, the IMF of stars, and the Type Ia SN progenitors. These are essential ingredients for understanding galaxy formation, cosmic chemical evolution, and the mechanisms which determined the efficiency of the conversion of gas into stars in galaxies at various epochs (e.g. Madau et al. 1996; Madau, Pozzetti, \\& Dickinson 1998; Renzini 1997). We computed the cosmic Type Ia SN rates for different cosmic SFR histories. Our main result is the prediction of the expected number of explosions of Type Ia SNe at high redshift ($z\\geq 2$).  Here we summarize our  conclusions in detail: \\begin{itemize} \\item we analyzed the effects of the various parameters entering the computation of the SN Ia rate and concluded that the realization probability $A_{Ia}$ (the actual fraction of systems which are able to give rise to a Type Ia SN) should be in the range $10^{-4}-10^{-3}$, where the lower value is appropriate for the DD model. This quantity, in all models, is a free parameter and it was chosen by reproducing the present time Type Ia SN rate in galaxies. Other parameters also play a key role in the computation of the Type Ia SN rate and these are: \\emph{i)} the mass range for the secondaries; \\emph{ii)} the minimum mass for the primaries; \\emph{iii)} the efficiency of the accretion and \\emph{iv)} the distribution of the separations at birth. Actually, all the different channels could contribute to the SN Ia events, each with its own probability ($A_{Ia}$). \\item We computed the Type Ia SN rate for specific elliptical galaxies of different initial luminous masses ($10^{10}$M$_{\\odot}$, $10^{11}$M$_{\\odot}$, $10^{12}$M$_{\\odot}$) for all the studied DTDs and for the SFRs suggested by chemical evolution models which best reproduce the characteristics of local ellipticals (Pipino \\& Matteucci 2004). All the DTDs predict an early maximum in the Type Ia SN rate, between 0.3 and 1 Gyr, according to the galaxy mass. In particular, the maximum predicted at about 0.3 Gyr corresponds to a typical elliptical of initial mass of $10^{12}$M$_{\\odot}$ while, for an elliptical of initial mass of $10^{10}$M$_{\\odot}$ the peak is at about 1 Gyr. The dependence of the maximum on the initial galactic mass is due to the efficiency of star formation, which is assumed to be higher for the most massive ellipticals . We have also considered the predicted Type Ia SN rates per unit mass (SNuM) at the present time vs. $(B-K)$ color relations for the three different ellipticals. We found that the bluer is the color (and hence the lower is the galactic mass), the higher are the predicted SN Ia rates in SNuM at the present time. Unfortunately, a real comparison with data is not possible since the data for ellipticals of different masses are not available. \\item We computed the cosmic Type Ia SN rate in a unitary volume of Universe by adopting different cosmic SFRds, as predicted by both models and observations (Hopkins \\& Beacom, 2006). We found that the SF history largely dominates over the assumed DTD in the calculation of the Type Ia SN rate. Therefore, unless we know the cosmic star formation history, we cannot safely decide which DTD is better on the basis of the observed cosmic SN Ia rate, in agreement with previous works (see also Forster et al. 2006; Blanc \\& Greggio 2008, Botticella et al. 2008). In particular, it is not possible to predict the delay time for the explosion of SNe Ia on the basis of the cosmic star formation rate, since for example, in the hierarchical framework, it is predicted a decrease of the cosmic Type Ia SN rate for $z>1$ irrespective of the chosen DTD, even for that of Mannucci et al. (2005; 2006) DTD,  where the fraction of prompt Type Ia SNe is 50$\\%$!   \\item The cosmic Type Ia SN rates for different SFRds (increasing, decreasing, constant with redshift) differ mostly for redshift $z>1$. We compared our results with the available data, although the highest redshift points are still very uncertain. Therefore, it was not possible to decide which scenario should be preferred, but high redshift data will greatly help to draw conclusions on this point. In particular, we predict that the Type Ia SN rate for $z>2$ should be very high if the monolithic scheme for the formation of ellipticals is assumed and even higher if combined with the empirical DTD of Mannucci et al. (2005), as opposed to models with decreasing or constant SFRds, which all predict a significant drop of the Type Ia SN rate for any chosen DTD. Future observations with JWST of high redshift SNe will help in sheding light on this subject.    \\end{itemize}"
    },
    "0807/0807.2448_arXiv.txt": {
        "abstract": "We present the 3-point function $\\xi_3$ and  $Q_3=\\xi_3/\\xi_2^2$ for a spectroscopic sample of luminous red galaxies (LRG) from the Sloan Digital Sky Survey DR6 \\& DR7. We find a strong (S/N$>$6) detection of $Q_3$ on scales of 55-125 Mpc/h, with a well defined peak around 105 Mpc/h in all $\\xi_2$, $\\xi_3$ and $Q_3$, in excellent agreement with the predicted shape and location of the imprint of the baryon acoustic oscillations (BAO). We use very large simulations (from a cubic box of L=7680 Mpc/h) to asses and test the significance of our measurement. Models without the BAO peak are ruled out by the $Q_3$ data with 99\\% confidence. This detection demonstrates the non-linear growth of structure  by gravitational instability between $z=1000$ and the present. Our measurements  show the expected shape for $Q_3$ as a function of the triangular configuration. This provides a first direct measurement  of  the non-linear mode coupling coefficients of density and velocity fluctuations which, on these large scales, are independent of cosmic time, the amplitude of fluctuations or cosmological parameters. The location of the BAO peak in the data indicates $\\Omega_m =0.28 \\pm 0.05$ and $\\Omega_B=0.079 \\pm 0.025$ (for $h=0.70$) after marginalization over  spectral index ($n_s=0.8-1.2$) linear $b_1$ and quadratic $c_2$ bias, which are found to be in the range: $b_1=1.7-2.2$ and $c_2=0.75-3.55$.  The data  allows a hierarchical contribution from primordial non-Gaussianities in the range $Q_3=0.55-3.35$. These constraints are independent and complementary to the ones that  can be obtained using the 2-point function, which are presented in a separate paper. This is the first detection of the shape of $Q_3$ on BAO scales, but our errors are shot-noise dominated and the SDSS volume is still relatively small, so there is ample room for future improvement in this type of  measurements. ",
        "introduction": "The galaxy  three-point function $\\xi_3$ provides a crucial test and valuable statistical tool  to investigate the origin of structure formation and the relationship between galaxies and dark matter (see  \\citet{bernar} for a review). We will concentrate here on the reduced 3-point function $Q_3 \\simeq \\xi_3/\\xi_2^2$ defined in Eq.\\ref{fiftheq} as the scaling expected from non-linear couplings (with $Q_3 \\simeq 1$). Measurements of the three-point function and other higher-order statistics in galaxy catalogs have a rich history (eg \\citet{peeblesgroth}, \\citet{frypeebles}, \\citet{baumgartfry}, \\citet{gazta1992}, \\citet{bouchet1993}, \\citet{frygazta1994}). In the past decade, three-point statistics have confirmed the basic picture of gravitational instability from Gaussian initial conditions (\\citet{friemangazta1994}, \\citet{jingborner1998}, \\citet{friemangazta1999}, \\citet{feldman2001}). The connection between these observables and theoretical predictions is best  done on large scales, where the physics (gravity) is best understood, but the surveys  previously available were not large enough to have good statistics on sufficiently large scales.  With the completion of large redshift surveys such as 2dFGRS (\\citet{colless2001}) and SDSS (\\citet{York00}) we expect measurement of higher-order  statistics to provide tighter constraints on cosmology (\\citet{colombietal1998}, \\citet{szapudietal1999}, \\citet{matarreseetal1997}, \\citet{scoccimarro2004}, \\citet{sefusatti2005}). First measurements of the redshift space $\\xi_3$ in the 2dFGRS (\\citet{gazta2005}) and SDSS (eg \\citet{nichol2006}) show good agreement with expectations (see also \\citet{kayo2004}, \\citet{nishimichi2007}, \\citet{kulkarni2007} and references therein).   We will assume here that the initial conditions are Gaussian. Current models of structure formation predict a small level of initial non-Gaussianities that we can neglect here. A popular parametrization of this effect is to assume that initial curvature perturbations, given by the  gravitational potential, $\\Phi$,  are given by $\\Phi = \\Phi_L + f_{NL} ~(\\Phi_L^2 - <\\Phi_L^2>)$, where $\\Phi_L$ is Gaussian and $f_{NL}$ is a non-linear coupling parameter of order unity. This produces non-Gaussianities in the matter density perturbations at wavenumber $k$ which, using the Poisson equation, are suppressed by the square of the horizon scale $k_H \\equiv H_0/c$, so that:  $Q_3 \\simeq 3 f_{NL} (k_H/k)^2 T(k)/D(a) $, where $T(k) \\simeq 1$ is the so called CDM transfer function and $D(a)$ is the growth factor. In our analysis $(k_H/k)^2 \\simeq 10^{-3}$ so that these type of primordial non-Gaussianities produce negligible contribution to $Q_3$ for models with $f_{NL} \\simeq 1$. A more detailed analysis of this will be presented elsewhere (see \\citet{sefusattikomatsu} for a detailed forecast for this model). If the non-Gaussianities come from a non-linear coupling in the matter density field rather than in the gravitational potential, the resulting 3-point function will have a non-Gaussian contribution similar to that produced by non-linear bias $c_2$ (ie see Eq.\\ref{eq:Q3G} below  and Conclusions).  The shape and amplitude of $Q_3$ depends on galaxy bias, ie how galaxy light traces the dark matter (DM) distribution. This is both a problem and an opportunity. A problem because biasing can confuse our interpretation of the observations. An opportunity because one can try to measure the biasing parameters out of $Q_3$. This idea was first proposed by \\citet{frygazta1993} and has been applied to the 3-point function and bispectrum of  real data (\\cite{friemangazta1994,fry1994,feldman2001,scoccimarro2001,verdeetal2002,gazta2005,nichol2006}).   This paper is the third on a series of papers on clustering of LRG. In the first two papers \\cite{paper1,paper2} we studied redshift space distortions on the 2-point correlation function. The reader is referred to these papers for more details on the LRG samples, simulations and the systematic effects. Similar LRG samples from SDSS have already been used by different groups to study the 2-point function (eg \\cite{detection,hutsia,percival,padmanabhan2007,blake}) and found good agreement with predictions in the BAO scales, where density fluctuations are at a level of few percent. This is encouraging and indicates that this sample is large and accurate enough to investigate clustering on such large scales.  In this paper we follow closely the methodology presented in 3 previous analysis. \\citet{barrigagazta2002}  presented a comparison of the predictions for the two and three-point correlation functions of density fluctuations, $\\xi_2$ and $\\xi_3$, in gravitational perturbation theory (PT) against large Cold Dark Matter (CDM) simulations. Here we use the same method and codes to estimate the clustering in simulations. \\citet{gaztascocci2005} extend these results into the non-linear regime and focus on the effects of redshift distortions and the extraction of galaxy bias parameters in galaxy surveys. \\citet{gazta2005} apply this methodology to the 2dFGRS. Here we apply the very same techniques to the LRG data, so the reader is referred to these papers for more details.  \\citet{kulkarni2007} have also estimated $\\xi_3$ using LRG galaxies from SDSS DR3, but focusing on smaller scales. We use DR6  which has 3 times the area (and volume) of DR3. We also use a volume limited sample and a different estimator for the correlation functions and errors, focusing on the largest scales. ",
        "conclusions": "We have studied the large scale 3-point correlation function for luminous red galaxies from SDSS, and particularly the reduced $Q_3=\\xi_3/\\xi_2^2$, which measures the scaling expected from non-linear couplings. We find a well-detected peak at 105Mpc/h separation that is in agreement with the predicted position of BAO peak. This detection is significant since it is also imprinted in $\\xi_2$ and $\\xi_3$ separately. We focus our interpretation in $Q_3$ because it is a measure independent of time, $\\sigma_8$ or growth factor. It only depends on the shape of the initial 2-point function and the non-linear coupling of the gravitational interaction. Our result for $Q_3$ is in excellent agreement with predictions from Gaussian initial conditions. When we use the $Q_3$ data alone (with no fit to the 2-point function) we are able to break the strong degeneracy between $\\Omega_m$ and $\\Omega_B$ (see Fig.\\ref{fig:OmOb}). Our detection shows a clear preference for a high value of $\\Omega_B=0.079 \\pm 0.025$. This value is larger, but still consistent at 2-$\\sigma$ with recent results of WMAP ($\\Omega_B=0.045$). At 3-sigma level, the $\\Omega_m-\\Omega_B$ becomes degenerate. Models with no BAO peak are ruled out at 99\\% confidence level.  We have used very large realistic mock simulations to study the errors. These simulations show  that $Q_3$ is not significantly modified in redshift space so we can use real-space perturbation theory (see Fig.\\ref{fig:q3sim}). This agreement also indicates that loop corrections are small on BAO scales \\cite{Bernardeau08}. This analysis is independent from 2-point statistics, which tests the linear growth of gravity, since 3-point statistics test the non-linear growth. A high value for $\\Omega_B$ is also consistent with the analysis of the peak in the 2-point correlation function shown in \\citet{detection} and in Paper IV (\\citet{paper4}) of this series, which detect a slightly higher peak than expected. We have done all the analysis with just one set of triangle configurations, with fixed sides of $r_{12}=33 \\pm 5.5$ Mpc/h and $r_{13}=88 \\pm 5.5$ Mpc/h, to center our attention to BAO scale. This is about optimal, but we notice that there is much more to learn from $Q_3$, which will be presented in future analysis.  Results on smaller scales are consistent with what we find here.   Data is in excellent agreement with Gaussian initial conditions, for which $Q_3=0$. But note that our quadratic bias detection  $c_2=0.75-3.55$ is degenerate with a primordial non-Gaussian (hierarchical) contribution. Indeed the mean value of $c_2 \\simeq 2$ seems larger in observations than in halo simulations, for which we find $c_2^h \\simeq 0.2$. We believe that this indicates that halos are sometimes occupied by more than one galaxy, which increases the effective value of $c_2$ \\cite{bhalo}. But if we are conservative we can not rule out a primordial non-Gaussian contribution in the range $Q_3(Primordial)= c_2-c_2^h=[0.55,3.35]$.  Note that we have pixelized our data in cubical cells of side $dr=11$ Mpc/h. This results in some lost of small scale information but allows for a very fast method to estimate 3-point function \\cite{barrigagazta2002}. This is important for data, but more for simulations. In the MICE7680 simulation there are close to $N=10^{11}$ particles. A brute force method to estimate 3-point correlation would require $N^3=10^{33}$ operations, while our method based on pixels just needed $5 \\times 10^{12}$ operations.  \\begin{figure}[t] \\includegraphics[width=7cm]{figures/fitOmObPAU.eps} \\caption{Same as Fig.\\ref{fig:OmObsim} for a future photometric Survey with photo-z error of $\\Delta z<0.003 (1+z)$, volume of $V = 10 Gpc^3/h^3$ and number density of $\\bar{n}=10^{-3} h^3/Mpc^3$ LRG galaxies.} \\label{fig:OmObPAU} \\end{figure}  Future surveys will be able to improve much upon our measurement here. A photometric survey with $\\Delta z<0.003 (1+z)$ precision (corresponding to $dr<9$ Mpc/h at z=0), such as in the PAU Survey \\cite{PAU} should have enough spatial resolution to measure $Q_3(\\alpha)$ as presented in this paper (recall that we are binning our radial distances in $dr=11$ Mpc/h). Such survey could sample over 10 times the SDSS DR6 volume (ie to z=0.9) with 20 times better LRG number density (ie for $L>L_*$). Fig.\\ref{fig:OmObPAU} shows the forecast for such a survey, which we have simulated with the MICE7680 mocks in redshift space with a photo-z of $\\Delta z<0.003 (1+z)$ and for the same triangles as shown in our SDSS analysis. This is just illustrative, as we have not marginalized over biasing and other cosmological uncertainties. But note  that the improvement is substantial and shows the potentiality of this method to constrain cosmological parameters and models of structure formation.     EG wish to thank Bob Nichol for suggesting the test with the EH no-wiggle model. We acknowledge the use of simulations from the MICE consortium (www.ice.cat/mice) developed at the MareNostrum supercomputer (www.bsc.es) and with support form PIC (www.pic.es), the Spanish Ministerio de Ciencia y Tecnologia (MEC), project AYA2006-06341 with EC-FEDER funding, Consolider-Ingenio PAU project CSD2007-00060 and research project 2005SGR00728 from Generalitat de Catalunya. AC acknowledge support from the DURSI department of the Generalitat de Catalunya and the European Social Fund."
    },
    "0807/0807.2481_arXiv.txt": {
        "abstract": "In traditional seeing-limited observations the spectrograph aperture scales with telescope aperture, driving sizes and costs to enormous proportions. We propose a new solution to the seeing-limited spectrograph problem. A massively fiber-sliced configuration feeds a set of small diffraction-limited spectrographs. We present a prototype, tunable, J-band, diffraction grating, designed specifically for Astronomical applications: The grating sits at the heart of a spectrograph, no bigger than a few inches on a side. Throughput requirements dictate using tens-of-thousands of spectrographs on a single 10 to 30 meter telescope. A full system would cost significantly less than typical instruments on 10m or 30m telescopes. ",
        "introduction": "\\label{sec:intro} The next generation of large telescopes is designed, among other scientific requirements, to measure fundamental properties of galaxies from high-redshift ($z$) to today.  However, the instrumentation cost for the next generation of telescopes will be enormous.  Here, we present an approach that combines several modern technologies to significantly reduce the cost of astronomical spectrographs.  For any two functionally equivalent spectrographs, their volume increases, roughly, with the telescope diameter cubed\\cite{Schroeder87}.   There is a power-law relationship between telescope diameter (volume) and instrument cost, where the exponent sits between two and three. For example, DEIMOS, on a 10m telescope cost US\\$10m dollars\\cite{faber03} (though it was originally budgeted for less.)  The preliminary estimate for an instrument with less multiplexing capability, on a 30m telescope, is US\\$60m\\cite{paz06}.  This growth of telescope diameter, and instrument cost and complexity, is schematically illustrated in Figure \\ref{fig:cost}.  \\begin{figure}[htp] \\includegraphics{Figure1.pdf} \\caption{\\label{fig:cost}A schematic illustration of the evolution in size and cost of optical spectrgraphs.  The top two rows show pictures of the Sloan digital Sky Survey  (SDSS; 2.5m primary) multi-object  spectrograph, and the Keck (10m primary) spectrograph, DEep Imaging Multi-Object Spectrograph.  Notice the relationship between telescope size (going down the page) and instrument volume.  At 10m, the spectrograph dwarfs the PI and engineers.  For the proposed 30m telescope and wide field optical spectrograph (WFOS), a CAD model is rendered next to DEIMOS.  The WFOS is the size of a small house. The device in the bottom left is a Fabry-Perot spectrograph built by Axsun Technologies, it illustrates the principle behind this design: build thousands of miniature spectrographs.} \\end{figure}  In this paper we describe the design of a spectrograph, built with tens of thousands of mini-spectrographs.  By clever use of fiber multiplexing and slicing, enormous optics disappear, simplifying alignment and test procedures, which ultimately drive down the cost of the instrument package.  A single, MEMS, spectrograph, including detector, costs manufacturers roughly US\\$20 per part.  It is well understood that industrial replication reduces the cost of an individual component.  However, a surprising recent result, is that replicated spectrographs reduce the {\\em total} cost of an instrument.  For example, the recently designed Visible Integral-Field Replicable Unit Spectrographs (VIRUS), which is replicated by $\\sim100\\times$, saves approximately a factor of two or three compared to a monolithic design\\cite{hill06}.  Micro-electrical-mechanical-systems (MEMS) technology, allows the production of spectrographs even smaller than the VIRUS.  The drive to reduce spectrograph size is the drive reduce total instrument cost.  Though these MEMS spectrographs may sound exotic, we emphasize that this technology is not new\\cite{onat05, Musca05}, commercial, off-the shelf MEMS spectrographs exist today.  However, these commercial spectrographs, are typically tuned for specific applications (e.g. laser communication, television displays and molecular identification.)   This work leverages the research of MEMS spectroscopy, to design an instrument tuned for a specific astronomical application. ",
        "conclusions": ""
    },
    "0807/0807.1161_arXiv.txt": {
        "abstract": "Recent observations have discovered a giant HI bridge that appears to connect between the outer halo regions of M31 and M33. We propose that  this HI bridge can be formed as a result of the past interaction between M31 and M33 based on test particle simulations with different orbits of M31 and M33 for the last $\\sim 9$  Gyr. We show that strong tidal interaction between M31 and M33 about $4-8$ Gyr ago can strip HI gas from  M33 to form HI streams around M31 which can be observed as a HI bridge if they are projected onto the sky. We show  that the number fraction of models reproducing well the observed HI  distribution of the bridge is only $\\sim 0.01$\\% (i.e., $\\sim 10$ among $\\sim 10^5$ models) and thus suggest that the observed structure and kinematics of the HI bridge can give some  constraints on the past orbits of M31 and M33. We  suggest that the observed outer  HI warp in M33 could be fossil evidence for the past M31-M33 interaction. We also  suggest  that some of high velocity clouds (HVCs) recently found in M31 could be HI gas originating from M33. We briefly discuss other possible scenarios for the formation of the HI bridge. ",
        "introduction": "Recent observational studies on HI gas around M31 based on the Green Bank Telescope (GBT) have reported the presence of at least 20 discrete  HI clouds within  50 kpc of the  M31 disk (e.g., Thilker et al. 2004). Furthermore possible candidates of HVCs around M31 with heliocentric velocities of $-520$ to $-160$ km s$^{-1}$ have been discovered in a blind HI survey for the disk and halo regions of M31 (e.g., Westmeier et al. 2007). Braun \\& Thilker (2004, hereafter BT04) have discovered a faint bridge-like HI structure that appears to connect between  M31 and M33. Although these recent results have provided  new clues to formation and evolution of M31 (BT04), no theoretical works have yet clarified the origin of the HI properties surrounding  M31, in particular, the intriguing bridge-like structure between  M31 and M33.  Corbelli et al. (1989) revealed that the HI gas disk of M33 extends out to twice the optical radius and shows complicated kinematical properties indicative of the presence of a warped disk. By using a tilted ring model, Corbelli  \\& Schneider (1997) demonstrated that the observed distribution and rotation curve profile of HI in the outer part of M33 are  consistent with the presence of a warped gas disk. Although they suggested the formation of the gaseous warp in M33 due to tidal force from M31, the observed unique HI properties of  M33 have not been discussed by theoretical and numerical studies in terms of the past M31-M33 interaction.   The purpose of this Letter is to first show that the observed HI bridge between M31 and M33 can be formed from the past tidal interaction between M31 and M33 based on a larger number ($>10^6$) of test particle simulations for the past interaction. We here try to choose  orbital models of M31 and M33 (among a large number of those) which can reproduce well the observed distribution of the HI bridge (BT04). We therefore do not intend to use hydrodynamical and chemodynamical simulations that are numerically costly and were used in our previous studies for the formation of HI streams and bridges  around  galaxies (e.g., Bekki et al. 2005a, b; 2008). ",
        "conclusions": "Although the present study has  presented a scenario that the observed apparent HI bridge between M31 and M33 is a tidal stream of HI gas around  M31 formed from the past M31-M33 interaction, this scenario would be only one of several possible scenarios. For example, BT04 suggested that the observed bridge can be a   ``cosmic web'' which extends between massive galaxies as predicted in previous numerical simulations on structure formation. It would be also possible that the bridge is a tidal stream originating not from M33 but from other gas-rich dwarfs that may have already been destroyed by the strong tidal field of M31. If fully self-consistent hydrodynamical simulations for the bridge formation in our future studies can also reproduce the observed sharp HI edge and outer HI warps in M33 (e.g., Corbelli \\& Salpeter 1993; Corbelli et al. 1997), then the present interaction scenario can be regarded as more realistic and reasonable.   Ibata et al. (2007) have recently found an extended metal-poor stellar halo around M33 in the deep photometric survey of M31, though the projected distribution of the halo is not so clear owing to the very limited spatial extent of the survey. If the stellar halo of M33 has a relatively homogeneous spatial distribution with no signs of disturbance and extends more than 35 kpc from the center of M33,   then M33 is highly unlikely to have lost HI gas initially within 35 kpc: the present tidal interaction scenario needs to be dramatically modified. It would be possible that the observed outer metal-poor halo of M33 can be the very outer part of the M31 stellar halo. Future kinematical studies of metal-poor stars around M33, which can confirm that the stars belong to M33 rather than to  M31, will enable us to discuss how far the M33 stellar halo extends and  thereby assess the viability of the present tidal interaction scenario.   The present models have shown  that HI gas clouds originating from M33 can be located in  tidal streams within $\\sim 100$  kpc of M31 and show systematic  rotation with respect to M31. One of implications from this result is that at least some of the observed HVCs around M31 (e.g., Westmeier et al. 2007) can originate from M33. Metallicities of the M31 HVCs originating from M33 may well be  as small as those of HI gas in the outer part of the gas disk in M33. Also the HVCs from M33 can have systematic rotation with the amplitude significantly smaller than that of the M31 disk ($V_{\\rm c} \\sim 250$ km s$^{-1}$). Thus future observational studies on  chemical abundances and kinematical propertied  of HVCs will help us to reveal the M31 HVCs originating from M33.  Although the present  gas-rich late-type disk galaxies are observed to have extended HI disks (e.g., Broeils \\& van Woerden 1994), it remains observationally unclear  how and when the extended HI disks were formed. The present model M1 (with the first M31-M33 interaction about 8 Gyr ago) would not be reasonable, if the extended HI disk in M33 was formed relatively recently (4-5 Gyr ago). It should be thus stressed that the viability  of the present scenario for the origin of the M31-M33 HI bridge would  depend on whether {\\it M33 had already developed its extended HI disk before it experienced the first tidal interaction with M31}.    Loeb et al. (2005) proposed a proper motion amplitude of $100\\pm 20$ km s$^{-1}$ form M31 based on orbital models of M31 and M33 and on the observed proper motion of M33 by B05. Recently van der Marel and Guhathakurta (2007) have investigated line-of-sight velocities of the seventeen M31 satellite galaxies in order to derive the M31 transverse velocity. They have found that the Galactocentric tangential velocity of M31 is highly likely to be less than 56 km s$^{-1}$ and suggested that M31 and M33 is in a tightly bounded system. These results appear to be consistent with the present tidal interaction scenario for the HI bridge. As shown in the present study, the observed location of the HI bridge alone can not allow us to provide very precise predictions on the present 3D velocities of M31 and M33. We plan to use the observed kinematical data sets for the HI bridge in order to give much stronger constraints on the 3D orbits of M31 and M33 by comparing the observations with more sophisticated numerical simulations."
    },
    "0807/0807.3955_arXiv.txt": {
        "abstract": "Traditionally, the distance to \\n4038/39 has been derived from the systemic recession velocity, yielding about 20 Mpc for $H_0 = 72$ \\kmsmpc. Recently, this widely adopted distance has been challenged based on photometry of the presumed tip of the red giant branch (TRGB), which seems to yield a shorter distance of $13.3\\pm 1.0$ Mpc and, with it, nearly 1 mag lower luminosities and smaller radii for objects in this prototypical merger. Here we present a new distance estimate based on observations of the Type Ia supernova (SN) 2007sr in the southern tail, made at Las Campanas Observatory as part of the Carnegie Supernova Project. The resulting distance of $\\dsnia = \\DIa$ Mpc [$(m-M)_0 = 31.74\\pm 0.27$ mag] is in good agreement with a refined distance estimate based on the recession velocity and the large-scale flow model developed by Tonry and collaborators, $\\dflow = 22.5\\pm 2.8$ Mpc. We point out three serious problems that a short distance of 13.3 Mpc would entail, and trace the claimed short distance to a likely misidentification of the TRGB. Reanalyzing {\\em Hubble Space Telescope} (\\hst\\,) data in the Archive with an improved method, we find a TRGB fainter by 0.9 mag and derive from it a preliminary new TRGB distance of $\\dtrgb = 20.0\\pm 1.6$ Mpc. Finally, assessing our three distance estimates we recommend using a conservative, rounded value of $D = 22\\pm 3$ Mpc as the best currently available distance to The Antennae. ",
        "introduction": "\\label{sec1}  The Antennae (\\n4038/39) are the nearest example of a major merger involving two gas-rich disk galaxies of comparable mass. They allow us to study processes of dissipational galaxy assembly from close up, thus providing a valuable glimpse of what must have been more frequent events in the early universe. The Antennae have been observed extensively at all wavelengths (e.g., X-rays: \\citealt{fabb04}; UV: \\citealt{hibb05}; optical: \\citealt{whit99}; IR: \\citealt{gilb07}; and 21-cm line: \\citealt{hibb01}) and have also been repeatedly modeled via N-body and hydrodynamical simulations \\citep[e.g.,][]{tt72,barn88,miho93,hibb03}. An accurate distance to this prototypical merger is of obvious importance in assessing its structure and dynamics, as well as for determining the physical properties of its myriad of stars, star clusters, peculiar objects such as ULXs (ultraluminous X-ray sources), and gas clouds.  Traditionally, the distance to \\n4038/39 has most often been derived from the systemic recession velocity and an adopted Hubble constant $H_0$, with or without corrections for deviations of the Hubble flow from linearity due to various attractors. A frequently used modern value for the distance is 19.2 Mpc \\citep{whit99}. This value is based on a systemic recession velocity relative to the Local Group of $\\czlg = 1439$ \\kms \\citep{ws95,rc3}, a linear Hubble flow, and $H_0 = 75$ \\kmsmpc. Updated to the \\hst\\ Key Project's derived $H_0 = 72\\pm 8$ \\kmsmpc\\ \\citep{freed01}, this often-used distance to The Antennae becomes $20.0_{-2.0}^{+2.5}$ Mpc.  A significantly shorter distance of $13.3\\pm 1.0$ Mpc has recently been determined by \\cite{savi08} via photometry of the TRGB in a region near the tip of the southern tidal tail. This photometry---performed on new, deep images obtained with \\hst\\,'s Advanced Camera for Surveys (ACS)---seems to support the short distance derived earlier from \\hst/WFPC2 images of the same region by the same method \\citep{savi04}. If the short distance is correct, The Antennae ``diminish'' in physical size, mass, and luminosity, as do their stars, clusters, and gas clouds. However, their recession velocity then deviates from the best large-scale-flow model \\citep[hereafter TBAD00]{tonr00} by about 500 \\kms\\ or 2.7\\,$\\sigma$ \\citep{savi08}.  Clearly, new measurements of the distance to The Antennae by independent methods are highly desirable. In the present paper we derive a new distance from observations of the Type Ia SN 2007sr, which appeared in the midsection of the southern tail in 2007 December. Section~\\ref{sec2} presents our observations, data analysis, and new distance. Section~\\ref{sec3} discusses problems with the short distance and points out a likely error in the identification of the true TRGB by \\citet{savi08}. Finally, \\S~\\ref{sec4} presents our conclusions and recommendation.  \\begin{figure*} \\centering \\includegraphics[scale=0.545]{n4038newdist_fig01_aph.eps} \\caption{ Supernova 2007sr in the southern tidal tail of \\n4038. {\\em Left:} Image of The Antennae reproduced from the Digital Sky Survey; the box marks the field of view corresponding to the right panel. {\\em Right:} Image of SN 2007sr in the $V$ passband, produced by stacking 29 exposures of 50~s duration each obtained over many nights with the CCD camera of the Swope 1.0~m telescope at Las Campanas Observatory. \\label{fig01}} \\end{figure*} ",
        "conclusions": "\\label{sec4}  We have presented $u'Bg'Vr'i'JH$ light curves of SN 2007sr beginning 6 days after $B$ maximum and obtained during four months by the Carnegie Supernova Project. Analysis of these light curves in conjunction with \\sneia\\ data from the first-year campaign by the Project yields a formal true distance modulus of $\\mu_0=31.74\\pm 0.07$ mag. With all known systematic errors and uncertainties included, this distance modulus becomes $\\dmodsnia=31.74\\pm 0.27$ mag, corresponding to a distance of $\\dsnia=\\DIa$ Mpc to \\n4038/39. (Differential depth effects between SN 2007sr and the center of mass of the galaxies are negligible).  This new \\snia\\ distance agrees well with the distance $D_{\\rm flow}= 22.5\\pm 2.8$ Mpc estimated from the systemic recession velocity of The Antennae and the large-scale flow model by TBAD00, where the quoted error reflects the cosmic random radial velocity ($\\sigvelth = 187$ \\kms) of the model. On the other hand, \\dsnia\\ disagrees strongly with the short distance of $D = 13.3\\pm 1.0$ Mpc estimated by \\citet{savi08} from the TRGB. We have discussed three serious problems with such a short distance, and have pointed out the likely misidentification of the TRGB by these authors as the cause of the short distance.  Reprocessing their \\hst/ACS frames and using an improved method of TRGB detection, we have located what we believe to be the true TRGB at $I_0 \\approx \\trgbmag = 27.46\\pm 0.12$ mag, fully 0.9~mag below (i.e., fainter than) the TRGB claimed by \\citeauthor{savi08}\\ \\ Analyzing photometry of this new, fainter TRGB we have derived a preliminary distance of $\\dtrgb = 20.0\\pm 1.6$ Mpc.  Clearly, additional distance estimates to The Antennae (e.g., from Cepheids, planetary nebulae, etc.) will be very valuable. However, for the moment we see no reason for adopting any distance shorter than 20 Mpc. Given the concordant new $\\dsnia = \\DIa$ Mpc and recession-velocity based $\\dflow = 22.5\\pm 2.8$ Mpc, both supported by our new, preliminary $\\dtrgb = 20.0\\pm 1.6$ Mpc, we recommend using a conservative, rounded value of $D = 22\\pm 3$ Mpc as the best currently available distance to The Antennae."
    },
    "0807/0807.3352_arXiv.txt": {
        "abstract": "A number of recent numerical investigations concluded that the remnants of rare structures formed at very high redshift, such as the very first stars and bright redshift $z\\approx 6$ QSOs, are preferentially located at the center of the most massive galaxy clusters at redshift $z=0$. In this paper we readdress this question using a combination of cosmological simulations of structure formation and extended Press-Schechter formalism and we show that the typical remnants of Population III stars are instead more likely to be found in a group environment, that is in dark matter halos of mass $\\lesssim 2 \\cdot 10^{13}~h^{-1} \\mathrm{M_{\\sun}}$. Similarly, the descendants of the brightest $z \\approx 6$ QSOs are expected to be in medium-sized clusters (mass of a few $10^{14}~h^{-1} \\mathrm{M_{\\sun}}$), rather than in the most massive superclusters ($M>10^{15}~h^{-1} \\mathrm{M_{\\sun}}$) found within the typical 1~Gpc$^3$ cosmic volume where a bright $z\\approx 6$ QSO lives.  The origin of past claims that the most massive clusters preferentially host these remnants is rooted in the numerical method used to initialize their numerical simulations: Only a small region of the cosmological volume of interest was simulated with sufficient resolution to identify low-mass halos at early times, and this region was chosen to host the most massive halo in the cosmological volume at late times. The conclusion that the earliest structures formed in the entire cosmological volume evolve into the most massive halo at late times was thus arrived at by construction. We demonstrate that, to the contrary, the first structures to form in a cosmological region evolve into relatively typical objects at later times. We propose alternative numerical methods for simulating the earliest structures in cosmological volumes. ",
        "introduction": "Rare dark matter halos at high redshifts host interesting astrophysical objects, especially before or at the end of the reionization epoch. One example is given by the very first Population~III (PopIII) stars formed in the universe at $z \\gtrsim 40$, which started the metal enrichment of the interstellar medium and the reionization process \\citep{abel02,san02,bromm04,naoz06}, and possibly also produced intermediate mass black hole seeds that grow to become super-massive black holes ($M_{BH}> 10^9 M_{\\sun}$) within the first billion year after the Big Bang \\citep{vol03}. Another example are bright $z \\approx 6$ QSOs \\citep{fan04}, considered to be hosted in the most massive dark matter halos at that time \\citep{MR05}. The luminosity of such an object is powered through accretion onto a supermassive black hole \\citep[e.g. see][]{hop06}, which may be the descendant of one of the first generation of PopIII stars formed in the universe, at $z \\gtrsim 40$ (\\citealt{mad01}; see also \\citealt{ts07a}).  Many numerical and theoretical investigations have aimed at characterizing the properties of both PopIII stars \\citep[e.g. see][]{abel02,bromm04,ciar,ree05,gao05,osh07} and of high redshift QSOs \\citep[e.g. see][]{hop06,MR05,dimatteo05,li07}. However, the rarity of these structures makes a fully self-consistent treatment even of the formation of the underlying dark matter halos typically outside the current capabilities of a standard cosmological simulation, as the dynamic range that needs to be resolved is too large. The main limitation at high mass resolution is the box size, $l_{box}$: The enforcement of periodic boundary conditions limits the power spectrum of density fluctuations to modes with wavelengths shorter than $l_{box}$. This results in a severe bias on the measured number density of rare, massive halos, especially before the epoch of the reionization of the universe when the abundance of galaxies derived from numerical simulations can be underestimated by up to an order of magnitude \\citep{bar04}. While the number density of rare objects can be estimated in the context of Press-Schechter-type modeling \\citep{lacey93,mo_white96,she99,bar04}, studying the details of the formation histories of the first galaxies and QSOs, possibly including hydrodynamics and radiative feedback processes, requires high-resolution numerical simulation.  Thus, given a simulation box that is large enough to contain one or more of the high-redshift halos of interest, the problem is to identify a sub-region that contains one of them for high-resolution simulation. In passing we note that this challenge is different from that of developing a simulation code that is able to adaptively increase the resolution when following the gas collapse that leads to the formation of the first stars. This has been successfully implemented with adaptive mesh refinement codes (for example ENZO - \\citealt{bry98}). \\citet{gao05} proposed a method to identify rare structures formed at very high redshift by recursively resimulating successively smaller, nested, sub-regions of a simulation box at progressively higher resolutions. Specifically, their method resimulates the region of the box centered on the most massive dark matter halo, reidentified at higher redshift and lower mass within the sub-region at each resimulation step, until the desired resolution is achieved in a small fraction of the entire simulation box. This procedure is well-motivated by the fact that number density of massive halos is increased in regions of large-scale over-density \\citep{bar04}, and thus structure formation in the sub-region on small scales at early times is accelerated by sitting at the top of an over-density, an over-density that is known to exist because it collapsed into the massive halo identified in the previous resimulation step. The \\citet{gao05} method extends high-resolution resimulation techniques used previously \\citep[e.g. see ][]{nav94,whi00}, and has been used recently by several authors to study the first stars and QSOs \\citep{gao05,ree05,gao07,li07}.  The recursive resimulation method for volume selection is very effective at identifying a region with a very high-redshift halo when compared to random selection of a region of equal volume within the parent simulation box, but still it does not lead to the rarest high-redshift halos in the volume. This is due to the stochastic nature of the formation and growth of dark matter halos \\citep[e.g. see][]{PS,bond,she99}. The most massive dark matter halo in a given volume at some early time does not in general evolve into the most massive halo at a later time. The probability of the most massive halo to evolve into the most massive halo, quantified following \\citet{lacey93}, decreases as time passes and the characteristic mass scale increases. As we demonstrate, the probability falls far below unity for the evolution of the first stars and QSOs to the present time. The stochastic nature of growth histories of dark matter halos is addressed and discussed in \\citet{gao05}, but unfortunately some of the subsequent studies using their method do not take into account this stochasticity and assume a very strong correlation between the location of the richest clusters at $z=0$ and that of rare halos at very high redshift.  In this paper we clarify this issue by presenting some basic, but often overlooked, results of Gaussian random fields in extended Press-Schechter theory to quantify the locations at later times of the first dark matter halos to form in a simulation box. We compare these analytical and numerical results to those published using recursive resimulation and highlight the benefits and limitations of that method. In addition, we propose an alternative method, based on analysis of the density field in the initial conditions, to select the sub-region of a simulation box that contains some of the earliest structure in the box without being biased toward the regions hosting the largest halos at $z=0$. The paper is organized as follows. In Sec.~\\ref{sec:QSOs} we investigate where the descendants of $z=6$ QSOs are today, while in Sec.~\\ref{sec:PopIII} we extend the result to the first Population III stars in the Universe. In Sec.~\\ref{sec:Ic} we discuss the implications of our results for the recursive mesh refinement method and propose a viable alternative. Sec.~\\ref{sec:conc} summarizes and concludes. ",
        "conclusions": "\\label{sec:conc}  By means of both analysis of numerical simulations and of extended Press-Schechter modeling we investigated the relation between the most massive dark matter halos at different redshifts. The main conclusion of this work is that---contrary to expectations from many recent works (e.g. see \\citealt{MR05,ree05,gao05,gao07,li07})---the most massive halo at a redshift $z_1>z_2$ does not necessarily evolve into the most massive at $z=z_2$. This is a robust conclusion that can be naturally understood in the context of growth of dark matter density perturbations, for example by constructing merger trees through the \\citet{lacey93} model. Rare high-redshift objects, such as the remnants of the first PopIII stars and QSOs, are not hosted at $z=0$ in the most massive halos, but rather live in a variety of environments. For example, the typical $z>40$ PopIII star remnant lives in a dark matter halo that at $z=0$ has a mass of $\\approx 2 \\cdot 10^{13} h^{-1} \\mathrm{M_{\\sun}}$, typical for a galaxy group, and not within rich clusters as claimed by \\citet{whi00}. Similarly the descendant of a typical $z\\approx 6$ QSO is not located within the most massive clusters at $z=0$ as assumed by \\citet{li07}.  These conclusions have important consequences on the application of the recursive simulation method introduced by \\citet{gao05} to identify high-redshift rare halos based on progressive refinement of regions centered around the most massive $z=0$ cluster. In fact, while the recursive method is indeed effective at identifying a sub-region of the simulation with earlier-than-average structure formation, it finds neither the earliest structures in the box, which are dominated by small-scale density fluctuations, nor typical early structure, as it preferentially identifies objects located in the regions with the highest bias.  These limitations may have only a minor effect when the goal is to investigate the formation of one rare Population III halo in the simulation box, as it is done for example in \\citet{gao05} and in \\citet{ree05}. However in different physical scenarios it is important to correctly estimate the rarity of the halo simulated and to assess how typical their growth histories are. This is critical if additional physics beyond gravitational interactions is included, such as star formation and radiative feedback. One such example is the formation of rare high-redshift QSOs: \\citet{li07} use the \\citet{gao05} method to identify at $z\\approx 6.5$ ``the most massive halo in a $\\approx$ 3 Gpc$^3$ volume'' and then conclude that the QSO formed in this halo reproduces the properties of observed QSOs with the same number density. From our analysis in Sec.~\\ref{sec:QSOs} it is clear that the halo identified by \\citet{li07} as progenitor of the largest $z=0$ cluster is not likely to be the most massive at $z \\approx 6.5$. Thus other similar or more massive halos are expected to be present at $z \\approx 6.5$ in their $\\approx$ 3 Gpc$^3$ simulation volume: in principle any of these halos could host a bright QSOs, with important consequences for the comparison between observed and simulated QSO number densities. In addition, when the goal is to study the environment in which QSOs live, selecting host halos with the resimulation method introduces systematic effects in the results difficult to quantify and correct for, because these halos would have above-average growth (and merging) histories.  To avoid selecting only the halos with atypical accretion histories and in an attempt to improve over the identification of some of the rarest high-redshift halos in a box, we propose instead to select the initial conditions for high-resolution resimulation based on the analysis of the linear density field at uniform resolution. The method, described in Sec.~\\ref{sec:Ic}, identifies subregions of the simulation box with high-redshift halos as those with the highest peaks in the density field, requiring a mass resolution in the field comparable to that of the mass of the halos one wishes to select.  The applicability of the method is thus limited only by the size of the largest density grid that can be accommodated in the available memory, otherwise requiring only a modest amount of computing time compared to the \\citet{gao05} method. This is because our method bypasses the need of a series of N-body simulations to be carried out in addition to the final run.  Unfortunately, our method does not guarantee identification of \\emph{the first} halo on the desired mass scale, as highlighted by some preliminary testing we presented in Sec.~\\ref{sec:valid}. This seems still an elusive goal. When the density field is defined over a fixed grid, there is not a perfect match between the halo catalog constructed from the density field and that obtained by increasing the resolution of the field and then following the full non-linear dynamics with an N-body simulation. An extensive validation of our linear density field initial conditions generation and its application to the formation of the first bright QSOs will be discussed in a subsequent paper."
    },
    "0807/0807.3813_arXiv.txt": {
        "abstract": "In this paper we fully calculate the non-Gaussianity of primordial curvature perturbation of island universe by using the second order perturbation equation. We find that for the spectral index $n_s\\simeq 0.96$, which is favored by current observations, the non-Gaussianity level $f_{NL}$ seen in island will generally lie between 30 $\\sim$ 60, which may be tested by the coming observations. In the landscape, the island universe is one of anthropically acceptable cosmological histories. Thus the results obtained in some sense means the coming observations, especially the measurement of non-Gaussianity, will be significant to make clear how our position in the landscape is populated. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0807/0807.2139_arXiv.txt": {
        "abstract": "The ARGO-YBJ experiment has been designed to study the Extensive Air Showers with an energy threshold lower than that of the existing arrays by exploiting the high altitude location (4300 m a.s.l. in Tibet, P.R. China) and the full ground plane coverage. The lower energy limit of the detector (E $\\sim$ 1 GeV) is reached by the scaler mode technique, i.e. recording the counting rate at fixed time intervals. At these energies, transient signals due to local (e.g. Forbush Decreases) and cosmological (e.g. Gamma Ray Bursts) phenomena are expected as a significant variation of the counting rate compared to the background. In this paper the performance of the ARGO-YBJ detector operating in scaler mode is described and discussed. ",
        "introduction": "The study of gamma ray sources in the 1 - 100 GeV energy range has so far been only partially covered by observations by the satellite-based EGRET, with a maximum detectable energy of 30 GeV. On the other hand, present ground-based detectors such as Cherenkov Telescopes are still trying to reduce their threshold to energies well lower than 100 GeV. \\\\ In this energy range, Gamma Ray Bursts (GRBs) are certainly the most intriguing transient phenomena. GRBs have been deeply studied in the keV-MeV energy range, but only EGRET, in the last decade, provided  some information in the GeV range. Three bursts have been detected at energies $>$1 GeV, with one photon of E=18 GeV \\cite{catelli,hurley}. The study of the high energy emission of GRBs could provide extremely useful constraints to the emission models and the parameters of the surrounding medium. \\\\ The interest for transient emission in the GeV region is not only restricted to extragalactic phenomena such as GRBs but also includes events closer to us - inside our Solar System - such as gamma rays from Solar Flares or high energy solar protons producing short time increases of the detector counting rate at the ground (Ground Level Enhancements). \\\\ The study of these transient phenomena can be successfully performed by ground-based experiments such as Extensive Air Shower (EAS) detectors working in ``single particle mode''\\cite{vernetto}. EAS arrays usually detect air showers generated by cosmic rays of energy E $\\geq$ (1 $-$ 100) TeV; the arrival direction of the primary particle is measured by the time delay of the shower front in different counters, and the primary energy is evaluated by the number of secondary particles detected. \\\\ EAS arrays can work in the energy region E $<$1 TeV operating in single particle mode, i.e. counting all the particles hitting the individual detectors, independently of whether they belong to a large shower or they are the lone survivors of small showers. Because of the cosmic ray spectrum steepness, most of the events detected with this technique are in fact due to solitary particles from small showers generated by 1 $-$ 100 GeV cosmic rays. \\\\ Working in single particle mode, an air shower array could in principle detect a transient emission in the interval 1 $-$ 100 GeV (both by gamma rays or charged primaries) if the secondary particles generated by the primaries give a statistically significant excess of events over the background due to cosmic rays. \\\\ The power of this technique is in its extreme simplicity: it is sufficient to count all the particles hitting the detectors during fixed time intervals (i.e. every second or more frequently, depending on the desired time resolution) and to study the counting rate behaviour versus time, searching for significant increases. The observation of an excess in coincidence with a GRB detected by a satellite would be an unambiguous signature of the nature of the signal. \\\\ The single particle technique does not allow one to measure energy and direction of the primary gamma rays, because the number of detected particles (often only one per shower) is too small to reconstruct the shower parameters. However, it can allow to study the temporal behaviour of the high energy emission, and, with some assumptions on the spectral slope (possibly supported by satellite measurements at lower energies), can give an evaluation of the total power emitted. \\\\ The background rate depends on the altitude and, to a lesser extent, on the geomagnetic latitude; it is not strictly constant and several effects both environmental and instrumental are responsible for variations on different time scales with amplitudes up to a few per cent. For this technique, an accurate knowledge of the detector and its sensitivity to both environmental and instrumental parameters is of crucial importance in order to correctly evaluate the statistical significance of the detected signals. \\\\ Obviously the sensitivity of this technique increases with the detection area and the energy threshold decreases at higher observation levels, so the most suitable detectors have a large area and work at very high altitude. \\\\ The ARGO-YBJ experiment, located at the YangBaJing High Altitude Cosmic Ray Laboratory (Tibet, P.R. China, 4300 m a.s.l.) with a detection area of $\\sim$6700 $m^2$, successfully exploits the full coverage approach at very high altitude with the aim of studying the cosmic radiation with a low energy threshold. \\\\ In order to cover many physics items in different energy ranges, two main operation modes have been designed: (1) shower mode, with an energy threshold of a few hundreds of GeV, and (2) scaler mode, with a threshold of a few GeV, this latter differing from the single particle mode for different particle multiplicities (i.e. not only the single particles) are put in coincidence and counted on different scaler channels. \\\\ The large field of view ($\\gtrsim$2 $sr$, limited only by atmospheric absorption) and the high duty cycle (close to 100$\\%$) typical of EAS detectors, make ARGO-YBJ suitable to search for unknown sources and unpredictable events such as GRBs or other transient phenomena. The site location (latitude 30$^{\\circ}$ 06${'}$ 38${''}$ N, longitude 90$^{\\circ}$ 31${'}$ 50${''}$ E) allows sky monitoring in the declination band $-10^{\\circ}<\\delta<70^{\\circ}$. \\\\ In this paper the performance of ARGO-YBJ working in scaler mode will be described, with some examples of application of this technique in solar physics and in the search for GeV emission from GRBs. ",
        "conclusions": "The ARGO-YBJ detector operated in scaler mode has a high duty cycle and optimal stability, fulfilling the fundamental requirements of the scaler mode technique. \\\\ The main goal is the measurement of the high energy tail of GRB spectra, and the detection of few events during the experiment life time is expected if the spectrum of the most energetic GRBs extends at least up to 100 GeV. In this case, with four measurement channels sensitive to different energies, valuable information on the high energy spectrum slope and possible cutoff may be obtained. \\\\ In any case, even if no signal is detected, the large field of view and the high duty cycle, allowing for  continuous monitoring overhead, make the ARGO-YBJ experiment one of the most sensitive detectors for the study of the high energy spectrum of GRBs, with typical fluence upper limits down to $10^{-5} erg/cm^2$ in the 1 - 100 GeV region, well below the energy range explored by the present generation of Cherenkov telescopes. \\\\ On the other hand, due to the high rigidity cut-off, the large area of the detector and its high mountain location allow the study of the solar modulation of galactic cosmic rays and energetic relativistic particles coming from the Sun with high sensitivity and fine time resolution. This has been shown by the observation of the January 2005 FD with a preliminary discussion on the recovery phase. \\\\ Data collected up to May 2008 have been analyzed according to the procedures outlined in this work and the results of these analyses will be presented in later papers."
    },
    "0807/0807.3765_arXiv.txt": {
        "abstract": "We present results from $N$-body + magnetohydrodynamical simulations of merging clusters of galaxies. We find that cluster mergers cause various characteristic magnetic field structures because of the strong bulk flows in the intracluster medium. The moving substructures result in cool regions surrounded by the magnetic field. These will be recognized as magnetized cold fronts in the observational point of view. A relatively ordered magnetic field structure is generated just behind the moving substructure. Eddy-like field configurations are also formed by Kelvin-Helmholtz instabilities. These features are similarly seen even in off-center mergers though the detailed structures change slightly. The above-mentioned characteristic magnetic field structures are partly recognized in Faraday rotation measure maps. The higher absolute values of the rotation measure are expected when observed along the collision axis, because of the elongated density distribution and relatively ordered field structure along the axis. The rotation measure maps on the cosmic microwave background radiation, which covers clusters entirely, could be useful probes of not only the magnetic field structures but also the internal dynamics of the intracluster medium. ",
        "introduction": "Clusters of galaxies have plenty of hot plasma as well as galaxies and dark matter (DM), which is called intracluster medium (ICM). Some observational evidence indicates that ICM is magnetized. For example, some clusters have diffuse non-thermal synchrotron radio emission that is called radio halos or relics, which shows that there are magnetic fields as well as relativistic electrons in the intracluster space \\citep{Giov99, Kemp01}.  In addition, Faraday rotation measure observations of polarized radio sources such as radio lobes and AGNs behind and/or in clusters indicate that the magnetic field structures are quite random \\citep{Clar01, Vogt03, Govo06}.  Comparing the synchrotron radio flux with the hard X-ray one (or its upper limit) due to the inverse Compton scatterings of cosmic microwave background (CMB) photons, we are able to estimate the volume averaged magnetic field strength (or its lower limit). Typically, the strength of $\\sim 0.1 \\mu$G is obtained with this method \\citep{Fusc99, Fusc05} though those detections of non-thermal hard X-ray components are still controversial \\citep{Ross04, Fusc07}. On the other hand, somewhat higher values ($\\sim$ a several $\\mu$G) tend to be obtained with the Faraday rotation measure method \\citep{Clar01, Vogt03, Govo06} though these results depend on the detailed magnetic field structures that are not fully understood \\citep{Enss03, Murg04}.  Although the magnetic energy density is typically $\\sim $ a few percents of the thermal one in the intracluster space, it is believed that the magnetic fields play a crucial role in various aspects of ICM. It is likely that non-thermal particles are accelerated via shocks \\citep{Sara99, TakiNait00, ToKi00, Mini01, Ryu03, Taki03, Inou05} and/or turbulence \\citep{Rola81, Schl87, Blas00, Ohno02, Fuji03, Brun04}, where the magnetic fields play a definitive role in most theoretical models of the particle acceleration. Heat conduction is probably depends on the magnetic field configurations because charged particles cannot freely move in the direction perpendicular to the field lines. As we wrote before, magnetic field seems to be a minor component in global ICM dynamics. However, it is possibly important in relatively small scales, where fluid instabilities such as Rayleigh-Taylor and Kelvin-Helmholtz ones might be suppressed by magnetic tension.  Cluster mergers have a significant impact on ICM magnetic field evolution. Turbulent motion excited by mergers would amplify the field strength via dynamo mechanism. Moving substructures are expected to sweep the field lines and form the cold subclumps surrounded by the field lines \\citep{Vikh01}. \\citet{Asai04} and \\citet{Asai07} performed two and three dimensional magnetohydrodynamical(MHD) simulations of moving cold subclumps in hot ICM in rather idealized situations, respectively, which confirmed that expectation. This field structure might be responsible for the suppression of heat conduction and fluid instabilities, which is essential for the maintenance of cold fronts.  A lot of numerical simulations about merging clusters have been done, most of which are $N$-body + hydrodynamical simulations \\citep{Roet96, Taki99, Taki00, Rick01, Ritc02, Asca06, Taki06, McCa07, Spri07}. Although these simulations give us a great deal of understanding of the structures, evolution, and observational implications for merging clusters of galaxies, they make only a limited contribution to investigate the magnetic field structures. MHD simulations are essential in this regard. However, $N$-body + MHD simulations are rather rare though \\citet{Roet99} did pioneering work. It is true that Lagrangian particle methods based on smoothed particle hydrodynamics \\citep[see][]{Mona92} are extensively used in cosmological MHD simulations \\citep{Dola99, Dola02}. Considering that Eulerian mesh codes are essentially better at following evolution of magnetic fields, simulations based on such codes are highly desirable. In this paper, we present the results from $N$-body + MHD simulations of merging clusters of galaxies and investigate characteristic magnetic field structures during mergers and their implications.  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 simulations. In \\S \\ref{s:resu} we present the results. In \\S \\ref{s:summ} we summarize the results and discuss their implications. ",
        "conclusions": "\\label{s:summ} We perform $N$-body and MHD simulations of merging clusters of galaxies. We find that cluster mergers cause various characteristic magnetic field structures because of strong bulk flow motion. The magnetic field component perpendicular to the collision axis is amplified especially between a bow shock and contact discontinuity. As a result, a cool region wrapped by the field lines appears. A relatively ordered field structure along the collision axis appears just behind the moving substructure. Eddy-like field structures are also generated by Kelvin-Helmholtz instabilities. Although the detailed structures change slightly, similar features are seen in off-center mergers. RM maps have some information about the magnetic field structure. The above-mentioned characteristic structures are partly recognized in the RM maps when we observe the merging systems in the direction perpendicular to the collision axis. On the other hand, RM maps observed nearly along the collision axis are less informative in this respect. Typical absolute values of the RM become higher when the system is observed along the collision axis because of both the density distribution elongated toward the axis and ordered magnetic field along the same direction. In minor mergers, amplification of the magnetic field component perpendicular to the collision axis becomes less prominent. The results about magnetic field along the collision axis does not change very much. Although numerical resolution effects have the significant impact on the small scale structures such as widths of a bow shock and contact discontinuity and detailed small scale fluctuations of the magnetic field, our results about overall global structures are reasonably reliable.    Some observational and theoretical studies suggest that heat conduction in the ICM is suppressed from the Spitzer value. For example, a sharp temperature jump across a cold front in A2142 requires that the heat conductivity is reduced by a factor of between 250 and 2500 at least in the direction across the front \\citep{Etto00}. The spatial temperature variations in the central region of A754 also suggest that the conductivity is at least an order of magnitude lower than the Spitzer value \\citep{Mark03}. It is well-known that fine-tuning of heat conductivity is necessary in order to reproduce observational natures of cool cores assuming that radiative cooling in the cores balances with the conductive heating from the outer part of the clusters \\citep{Zama03}. Although the detailed process involved there is still unclear, magnetic fields likely play an important role. Heat conduction in the direction across the magnetic field lines is likely suppressed because electrons cannot travel freely in that direction. Using two dimensional MHD simulations with anisotropic heat conduction, \\citet{Asai04} shows that a moving subclump naturally form the magnetic field structure along the contact discontinuity and that the temperature jump there is maintained if the heat conduction perpendicular to the field line is sufficiently suppressed. Similar results are obtained in three dimensional cases \\citep{Asai07}. Basically, our simulations confirmed their results about magnetic field configurations in more realistic situations, though the heat conduction is not included.  We show that Faraday rotation measure maps are useful to obtain information about characteristic magnetic field structures formed by cluster mergers. This could be another probe of internal dynamics of clusters. However, there is a serious problem in this method at present. We need polarized radio sources such as radio lobes and/or AGNs in and/or behind the clusters of galaxies to do that. In other words, we are able to obtain RM maps in the limited regions where we have suitable polarized sources by chance. Naturally, these do not always correspond to the regions that we are interested in. However, this difficulty could be overcame if we have suitable polarized sources that covers a cluster entirely. One possible solution is to use the CMB as the polarized source \\citep{Ohno03}, though this is still challenging in the present status of CMB observations. There are a lot of observations that are ongoing and planned for measuring the CMB polarization. We hope that future observations of the CMB polarizations will enable us to make RM maps that cover clusters entirely, which would give us an important clues to understand the internal dynamics as well as the magnetic field structures in clusters of galaxies.  In actual clusters, the magnetic field have random components in small scales. It is very difficult to treat such scales by numerical simulations that concern the global structures. Clearly, observed RM maps are influenced by the small scale magnetic field fluctuations, which are not considered in RM maps calculated from our results. As a result, coherent length of the magnetic fields tends to be overestimated effectively, which means absolute values of RM from the simulations are also overestimated. However, it is probable that global spatial patterns of RM maps are nearly independent of such small scale fluctuations. In addition, it is also probably robust that higher RM are expected when the merging systems are observed along the collision axis."
    },
    "0807/0807.4373_arXiv.txt": {
        "abstract": "We discuss the possible dynamical role of extended cosmic defects on galactic scales, specifically focusing on the possibility that they may provide the dark matter suggested by the classical problem of galactic rotation curves. We emphasize that the more standard defects (such as Goto-Nambu strings) are unsuitable for this task, but show that more general models (such as transonic wiggly strings) could in principle have a better chance. In any case, we show that observational data severely restricts any such scenarios. ",
        "introduction": "Introduction}  Symmetry breaking phase transitions in the early universe are expected to have produced networks of topological defects \\cite{Vilenkin}. The possible roles of these defect networks in key cosmological scenarios will depend both on the type of defect considered and on the corresponding dynamics. For example, if defects are to significantly contribute to the dark energy, they should have a negative equation of state. If so the best possible situation is that of a frustrated network, in which case \\begin{equation} w\\equiv\\frac{p}{\\rho}=-\\frac{N}{3} \\end{equation} where $N$ is the defect's spatial dimension ($N=1,2$ respectively for cosmic strings and domain walls, while $N=3$ corresponds to a cosmological constant).  It is clear that only the cases with $N=2$ and $N=3$ can lead to the recent acceleration of the Universe. A crucial property shared by these two cases is that the ratio between the dark energy density and the background density grows rapidly with time, with dark energy being dynamically dominant only around today. The defect network should therefore have a present density close to critical ($\\Omega^0_{de}\\sim1$), but compatibility with other cosmological observables requires it to have a characteristic scale several orders of magnitude below the horizon $\\xi \\ll H^{-1}$.  Defects can also act as seeds for structure formation \\cite{Silk}. This may be the case if the ratio between the average defect energy density and the background density is approximately a constant and the characteristic scale of the network is roughly proportional to the Hubble radius ($\\xi \\propto H^{-1}$). Moreover, the defect fluctuations on the Hubble scale must be small ($\\delta\\lsim 10^{-5}$), and if the characteristic scale of the defects is of the order of the Hubble scale then it follows that their average energy density must also be very small ($\\Omega_{def}\\lsim 10^{-5}$) for consistency with cosmic microwave background anisotropies.  It is well known that a scaling cosmic string network has all these properties \\cite{Vilenkin}. Notice that the required network properties are mutually incompatible: although there are unified dark energy scenarios in which dark matter and dark energy are described by a single entity, it is not possible for a given defect network to simultaneously be the dark energy and act as a seed for structure formation.  The possible role of domain walls as dark energy candidates has been investigated in detail in \\cite{IDEAL1,IDEAL3} where it was shown that the dynamics of realistic domain wall networks appears to be incompatible with a dark energy role. Here we take a closer look at the role of defects on smaller scales. Specifically we will be mostly interested in kiloparsec (that is, galactic) scales. In particular, in doing this we will also explore another possibility that has been recently put forward \\cite{Redington,Alex}, \\textit{viz.} that cosmic strings could be the dark matter and explain the observed flat rotation curves of spiral galaxies. ",
        "conclusions": ""
    },
    "0807/0807.3553_arXiv.txt": {
        "abstract": "Various radio observations have shown that the hot atmospheres of galaxy clusters are magnetized. However, our understanding of the origin of these magnetic fields, their implications on structure formation and their interplay with the dynamics of the cluster atmosphere, especially in the centers of galaxy clusters, is still very limited. In preparation for the upcoming new generation of radio telescopes (like EVLA, LWA, LOFAR and SKA), a huge effort is being made to learn more about cosmological magnetic fields from the observational perspective. Here we present the implementation of magneto-hydrodynamics in the cosmological SPH code GADGET \\citep{springel2001,springel2005}. We discuss the details of the implementation and various schemes to suppress numerical instabilities as well as regularization schemes, in the context of cosmological simulations. The performance of the SPH-MHD code is demonstrated in various one and two dimensional test problems, which we performed with a fully, three dimensional setup to test the code under realistic circumstances. Comparing solutions obtained using ATHENA \\citep{2008arXiv0804.0402S}, we find excellent agreement with our SPH-MHD implementation. Finally we apply our SPH-MHD implementation to galaxy cluster formation within a large, cosmological box. Performing a resolution study we demonstrate the robustness of the predicted shape of the magnetic field profiles in galaxy clusters,  which is in good agreement with previous studies. ",
        "introduction": "\\label{sec:intro}  Magnetic fields have been detected in galaxy clusters by radio observations, via the Faraday Rotation Signal of the magnetized cluster atmosphere towards polarized radio sources in or behind clusters \\citep[see][for recent reviews]{2002ARA&A..40..319C,2004IJMPD..13.1549G} and from diffuse synchrotron emission of the cluster atmosphere \\citep[see][ for recent reviews]{2004IJMPD..13.1549G,2008SSRv..134...93F}. Our understanding of the origin of cosmological magnetic fields is particularly limited. But their evolution and possible implications for structure formation are also not yet fully understood. In addition their interplay with the large-scale structure formation processes, as well as their link to additional dynamics within the cluster atmosphere is unclear, especially their role in the cool core regions and the influence of these regions on the evolution of the magnetic fields.  The upcoming, new generation of radio telescopes (like EVLA, LWA, LOFAR and SKA) will dramatically increase the volume of observational data relevant for our understanding of cosmological seed magnetic fields in the near future. To investigate the general characteristics of the magnetic fields in and beyond galaxy clusters at the level required for a meaningful comparison to current and forthcoming observations, numerical simulations are mandatory. Non-radiative simulations of galaxy clusters within a cosmological environment which follow the evolution of a primordial magnetic seed field have been performed using Smooth-Particle-Hydrodynamics (SPH) codes \\citep{1999A&A...348..351D,2002A&A...387..383D,dolag2005b}  as well as Adaptive Mesh Refinement (AMR) codes \\citep{2005ApJ...631L..21B,2008A&A...482L..13D}. Although these simulations are based on quite different numerical techniques they show good agreement in the predicted properties of the magnetic fields in galaxy clusters. When radiative cooling is included, strong amplification of the magnetic fields inside the cool-core region of clusters is found \\citep{2008A&A...482L..13D}, in good agreement with previous work \\citep{2000cucg.confE..75D}. Cosmological, magneto-hydrodynamical simulations were also performed using finite-volume and finite-difference methods. Such simulations are used to either follow a primordial magnetic field \\citep{2008ApJS..174....1L} or the creation of magnetic fields in shocks through the so-called Biermann battery effect \\citep{1997ApJ...480..481K,Ryu..1998}, on which a subsequent turbulent dynamo may operate. The latter predict magnetic field strength in filaments with somewhat higher values \\citep[e.g. see ][]{2004PhRvD..70d3007S} than predicted by simulations which start from a primordial magnetic seed field, but are in line with predictions of magnetic field values from turbulence \\citep{2008Sci...320..909R}. Therefore further investigations are needed to clarify the structure, evolution and origin of magnetic fields in the largest structures of the Universe, their observational signatures, as well as their interplay with other processes acting in galaxy clusters and the large scale structure.  The complexity of galaxy clusters comes principly from their hierarchical build up within the large-scale structure of the Universe. In order to study their formation it is necessary, to follow a large volume of the Universe. However, one must also describe cosmic structures down to relatively small scales, thus spanning 5 to 6 orders of magnitudes in size. The complexity of the cluster atmosphere reflects the infall of thousands of smaller objects and their subsequent destruction or survival within the cluster potential. Being the source of shocks and turbulence, these processes directly act on the magnetic field causing re-distribution and amplification. Therefore realistic modelling of these processes critically depends on the ability of the simulation to resolve and follow correctly this dynamics in galaxy clusters.  Starting from a well-established cosmological n-body smoothed particle hydrodynamic (SPH) code GADGET \\citep{springel2001,springel2005} we present here the implementation of magneto-hydrodynamics, which allows us to explore the full size and dynamical range of state of the art cosmological simulations. GADGET also allows us to turn on the treatment of many additional physical processes which are of interest for structure formation and make interesting links with the treatment of magnetic fields for future studies. This includes thermal conduction \\citep{2004MNRAS.351..423J,2004ApJ...606L..97D}, physical viscosity \\citep{2006MNRAS.371.1025S}, cooling and star-formation \\citep{springel2003}, detailed modelling of the stellar population and chemical enrichment \\citep{2004MNRAS.349L..19T,2007MNRAS.382.1050T} and a self consistent treatment of cosmic rays \\citep{2007A&A...473...41E,2007MNRAS.378..385P}. The MHD implementation presented here is fully compatible with all these extensions, but here we want to focus on non-radiative simulations. All such processes are expected to increased the complexity and lead to interplay with the evolution of the magnetic field. This would make it impossible to critically check the numerical effects caused by the different SPH-MHD implementations and therefore we will ignore such additional processes in this work.  The paper is structured as follows: In section 2 we present the details of the numerical implementation, whereas in section 3 we present various code validation tests, all performed in fully three dimensional setups. In section 4 we present the formation of a galaxy cluster as an example for a cosmological application before we present our conclusions in section 5. In addition we present a convergence test for the code  in the appendix. ",
        "conclusions": "\\label{sec:conc}  We presented the implementation of MHD in the cosmological, SPH code GADGET. We performed various test problems and discussed several instability correction and regularization schemes. We also demonstrated the application to cosmological simulations, the role of resolution and the role the regularization schemes play in cosmological simulations.  Our main findings are:  \\begin{itemize}  \\item The combination of many improvements in the SPH implementation, like the correction terms for the variable smoothing length \\citep{springel02} as well as the usage of the signal velocity in the artificial viscosity \\citep{1997JCoPh..136....298S} together with its generalization to the MHD case \\citep{2004MNRAS.348..123P} improve the handling of magnetic fields in SPH significantly.  \\item Correcting the instability by explicitly subtracting the contribution of a numerical non-zero divergence of the magnetic field to the Lorenz force from the Maxwell tensor as suggested by \\citet{2001ApJ...561...82B}  seems to perform well. Specifically in three dimensional setups where it seems to work much better than other suggestions in the literature.  \\item The SPH-MHD implementation performs very well on simple shock tube tests as well as on planar test problems. We performed all tests in a fully three-dimensional setup and find excellent agreement of the results obtained with the SPH-MHD implementation compared to the results obtained with ATHENA in one or two dimensions.  \\item With a convergence study we demonstrate that the SPH-MHD results when increasing the resolution are converging to the true solution, especially in the sharp features. However, in some regions it seems that small but systematic differences converge only very slowly to the correct solution.  \\item Regularization schemes help to further suppresses noise and $\\mathrm{div}(\\vec{B})$ errors in the test simulations, however one has to carefully select the numerical parameters to avoid too strong smoothing of sharp features. Performing a full set of individual shock tube tests allows one to tune the numerical schemes and to determine optimal values. However they reflect an optimal choice for problems where the local timescales are mostly similar to the global timescale of the problem. For cosmological simulations it turns out that regularization by artificial dissipation leads to questionable results, whereas the regularization by smoothing the magnetic field (which is applied on global timescales) produces reasonable results.  \\item The SPH-MHD implementation allows us to perform challenging cosmological simulations, covering a large dynamical range in length-scales. For galaxy clusters, only the shape of the predicted magnetic profiles is, (with the exception of the central part of clusters) converged in resolution and in good agreement with previous studies. Also the structures obtained in synthetic Faraday Rotation maps are in good agreement with previous findings and compare well with observations.  \\end{itemize}  The results obtained with artificial dissipation in cosmological simulations indicate that physical dissipation could play a crucial role in determining the exact shape of the predicted, magnetic field profiles in galaxy clusters. Future work, especially when including more physical processes at work in galaxy cluster -- as can be done easily with our SPH-MHD implementation -- will reveal an interesting interplay between dynamics of the cluster atmosphere and amplification of magnetic fields. Thus having the potential to shed light on many, currently unknown aspects of cluster magnetic fields, their structure and their evolution."
    },
    "0807/0807.0146_arXiv.txt": {
        "abstract": "We report the discovery of a circumbinary disk around the Herbig Ae/Be system v892~Tau.  Our detailed mid-infrared images were made using segment-tilting interferometry on the Keck-1 Telescope and reveal an asymmetric disk inclined at $\\sim$60$\\arcdeg$ with an inner hole diameter of 250~mas (35~AU), approximately 5$\\times$ larger than the apparent separation of the binary components.  In addition, we report a new measurement along the binary orbit using near-infrared Keck aperture masking, allowing a crude estimate of orbital parameters and the system mass for the first time.  The size of the inner hole appears to be consistent with the minimum size prediction from tidal truncation theory, bearing a resemblance to the recently unmasked binary CoKu Tau/4.  Our results have motivated a re-analysis of the system spectral energy distribution, concluding the luminosity of this system has been severely underestimated. With further study and monitoring, v892~Tau should prove a powerful testing ground for both predictions of dynamical models for disk-star interactions in young systems with gas-rich disks and for calibrations of pre-main-sequence tracks for intermediate-mass stars. ",
        "introduction": "v892~Tau (Elias-1, Elias 3-1) is a young stellar object in the Taurus-Aurigae star forming region \\citep[$d=140~pc$,][]{kdh1994}. Its visible spectrum is classified as spectral type B8V \\citep{hernandez2004}, making v892~Tau one of the few Herbig Ae/Be stars in Taurus (in addition to AB~Aur, MWC 480).  The spectral energy distribution (SED) of v892~Tau is dominated by bright thermal emission in the mid-infrared \\citep{hillenbrand1992}, with relatively large line-of-sight extinction in the visible \\citep[one estimate is A$_{\\rm V}\\sim5.9$,][]{kh1995}.  Despite being one of the closest Herbig stars and after decades of multi-wavelength observations, the nature of v892~Tau is still uncertain.  The SED suggests v892~Tau is either a young embedded Class I source \\citep{lada1987} still surrounded by its nascent envelope or a more evolved, Class II object (i.e, Herbig Ae/Be star) seen through its edge-on disk. Spatially-resolved imaging can easily distinguish between these two scenarios, and early speckle interferometry \\citep{kataza1991,haas1997,leinert2001} suggested the presence of an extended and elongated nebula in the near-infrared that could be due to scattering in bipolar lobes.  The high-resolution speckle imaging of \\citet{smith2005} clearly resolved the K band ($\\lambda=2.2\\mu$m) emission to be coming from two unresolved stars in v892~Tau with little or no sign of a nebula or extended emission.  The two stars of roughly equal brightness had an apparent separation of 55~milliarcseconds (7.7~AU) and some evidence of orbital motion was seen between epochs separated by 7 years.  Although this binary should have a dramatic effect on the infrared emission, carving out a large hole in the circumbinary disk, recent analysis of the SED of v892~Tau including ISO data \\citep{acke2004} uncovered no distinct signature of the underlying binary\\footnote{\\citet{acke2004} did note strong 11.2$\\mu$m PAH emission and anomalous infrared colors.}.  Other workers \\citep{liu2005,liu2007} made use of the technique of nulling interferometry in the mid-IR ($\\lambda=10.3\\mu$m) to marginally resolve v892~Tau (FWHM $\\sim$20~AU) along PA~164$\\arcdeg$ consistent with normal (single-star) disk emission.  Here, we report new mid-iIR imaging of the v892~Tau system which resolves these mysteries, discovering very extended and resolved emission that we interpret to be coming from a circumbinary disk. We also confirm the presence of the binary at K band and our new data allow first crude estimates of the orbital elements. Lastly, we report a new SED analysis which provides an improved determination of the system luminosity and our viewing geometry. ",
        "conclusions": "We have discovered an extensive circumbinary disk around v892~Tau in the mid-infrared.  We also independently confirm the binary nature of the underlying stellar system and our new measurement allows us to fit an astrometric orbit, finding a period $\\sim$14~years for system mass of $\\sim$6$\\msun$. Our limited orbital phase coverage and some ambiguity in position angles allow only crude estimates and we strongly urge continued monitoring of this system using near-IR speckle, aperture masking, or adaptive optics.  We have proposed a new SED decomposition with a line-of-sight extinction ($A_V\\sim11$) higher than previously thought, implying a system luminosity $\\sim$400$\\lsun$.  This result highlights the fact that circumbinary disks (and transitional disks) have much larger opening angles since the puffed-up inner wall causes enhanced extinction of the central stars even at inclinations of 60$\\arcdeg$.  v892~Tau is another case where high-resolution imaging has motivated a fundamental shift in our understanding of an individual object.  As spectral energy distributions are used to discover ``transitional'' disks implicating planet formation, we must be cognizant of the important role of binarity and the associated circumbinary disks that can mimic signs of planet formation \\citep[e.g.,][]{ireland2008}. Far from merely being spoilers to planet finders, new circumbinary disks offer fresh laboratories for studying dynamical interactions between gas-rich disks and the massive orbiting bodies embedded within them as well as critical opportunities to calibrate pre-main sequence tracks."
    },
    "0807/0807.2469_arXiv.txt": {
        "abstract": "The Arcminute Microkelvin Imager is a pair of interferometer arrays operating with six frequency channels spanning $13.9$--$18.2$~GHz, with very high sensitivity to angular scales $30\\arcsec$--$10\\arcmin$. The telescope is aimed principally at Sunyaev-Zel'dovich imaging of clusters of galaxies. We discuss the design of the telescope and describe and explain its electronic and mechanical systems. ",
        "introduction": "\\label{sec:intro}  Clusters of galaxies are the best samplers of matter on the largest scales and are sensitive probes of structure formation, both in the linear regime of growth and when merging, shocking and gradual virialization occur. Determining the structures and physics of clusters, their mass function and their evolution is therefore of basic importance.  Observationally, however, relatively little is known about the properties of the cluster population at redshifts $z\\gtrsim 1$, due to the fundamental problem of the dimming of surface brightness with redshift. Optical observations suffer from confusion from foreground galaxies, and both optical and X-ray observations suffer from biases towards strong concentrations of mass.  Ionized matter at any redshift out to recombination will inverse-Compton scatter the cosmic microwave background (CMB) radiation and so imprint spatial fluctuations upon it. Hot ionized gas in the deep potential wells of galaxy clusters (statistically) up-scatters CMB photons passing through it, giving rise to the Sunyaev-Zel'dovich (SZ) effect (\\citealt{sz70}, \\citealt{sz72}; see \\eg \\citealt{birkinshaw99} and \\citealt*{carlstrom02} for reviews); at frequencies below 217~GHz, one sees a decrease in the CMB temperature towards the cluster. This dip in temperature $\\Delta T$ is given in the Rayleigh-Jeans region by $\\Delta T = -2yT_{\\rm{CMB}}$, where the Comptonization parameter $y=\\int n_{\\rm{e}} \\sigma_{\\rm{T}} \\frac{k_{\\rm{B}}T_{\\rm{e}}}{m_{\\rm{e}} c^2} {\\rm d}l$, $\\sigma_{\\rm{T}}$ is the Thomson cross-section, $l$ is the line-of-sight path, and $m_{\\rm{e}}$, $n_{\\rm{e}}$ and $T_{\\rm{e}}$ are respectively the electron mass, density and temperature. The dip $\\Delta T$ is thus proportional to the line integral of pressure but does \\textit{not} depend on redshift $z$. One way of seeing this is that in such a inverse-Compton scattering process, the power lost by the electrons and given to the CMB photons is proportional to the energy density of the CMB radiation which is proportional to $(1+z)^{4}$; this exactly cancels out the cosmological drop in bolometric surface brightness proportional to $(1+z)^{-4}$. This is an extremely important result --- one can get \\textit{directly} to high redshift using SZ without any intermediate steps.  The total SZ decrement in flux density, $ \\Delta S_{\\rm SZ}$, is proportional to the integral of the brightness temperature $ \\int \\Delta T_{\\rm SZ}\\,{\\rm d}\\Omega$ over the solid angle $\\Omega$ subtended by the cluster. Thus $\\Delta S_{\\rm{SZ}} \\propto D_{\\rm{A}}^{-2} \\int n_{\\rm{e}}T_{\\rm{e}} \\, \\rm{d}V, $ where $D_{\\rm{A}}$ is the angular-size distance and $\\rm{d}V$ is an element of volume. For concordance cosmologies, for redshifts between, say, 0.3 and 3, $D_{\\rm{A}}$ is only weakly dependent on redshift. Since $D_{\\rm{A}}$ is always a weaker function of redshift than luminosity distance $D_{\\rm{L}}$, SZ clusters will have a selection function that is much less dependent on redshift then most self-luminous objects. $\\Delta S_{\\rm{SZ}}$ measures the total thermal energy of the cluster, and since temperature is strongly correlated to mass (\\eg $T \\propto M^{2/3}$ for a population of clusters, modelled as virialized gravitationally-collapsed systems), the integrated SZ flux density measures gas mass directly, and without bias to concentrated structure. Detailed studies of clusters in SZ and other wavebands will help to establish the cosmological relationship between $M$ and $T$ and the thermal history of the cluster, and will help to calibrate the cluster scaling relations.  These points immediately emphasize the importance of SZ surveys for clusters. These have long been advocated (see \\eg \\citealt*{ksy86}, \\citealt{bm92}, \\citealt{bs94}, et seq.) and it is now widely recognized that such surveys are key measurements for cosmology. For example, it is critical to measure and understand the comoving number density of clusters, $N(M,z)$, as a function of mass $M$ and redshift. $N(M,z)$ is a very strong function of the key cosmological parameter $\\sigma_8$ --- the density contrast on scales of $8\\,h^{-1}$Mpc now --- and of the physics of cluster assembly.  The X-ray flux from a cluster is proportional to $n_{\\rm{e}}^{2}f(T_{\\rm{e}})$, where $f(T_{\\rm{e}})$ is not a strong function of $T_{\\rm{e}}$. The SZ flux is proportional to $n_{\\rm{e}} T_{\\rm{e}}$, so the combination of good X-ray and SZ data on a cluster gives robust determinations of density and temperature distributions and the related physics. Since the SZ flux density is proportional to the integrated gas pressure rather than to the density squared, it is relatively more sensitive to the outer parts of clusters than are X-rays, and not biased towards dense regions.  All these points also emphasize the importance of pointed, high-resolution SZ observations of clusters known from X-ray, optical/IR or SZ surveying, as well as X-ray/optical/IR observations of SZ-selected targets. For example, cluster observations from XMM-Newton and Chandra show clear entropy structures in a cluster's gas, with contact discontinuities and sloshing motions (see \\eg \\citealt{rossetti07}), and bubbling and reheating of cooling gas by FRI radio jets of low luminosity but considerable bulk kinetic power \\citep[\\eg ][]{sanders07}. Observations in different wavebands tease out different selection effects and allow for cross-calibration.  The first unequivocal SZ detections were made using the OVRO 40-m telescope and the NRAO 140-foot telescope \\citep*{birkinshaw81,birkinshaw-plus84,uson86}, and the first SZ image was obtained with the Ryle Telescope (RT; \\citealt{jones93}). The RT and OVRO/BIMA interferometers have since then made many SZ observations of known clusters (see \\eg \\citealt{grainge96}, \\citeauthor{grainge02a} \\citeyear{grainge02b}a,b, \\citeauthor{cotter02a} \\citeyear{cotter02b}a,b, \\citealt{grainger02}, \\citealt{saunders03}, \\citealt{jones05}; \\citealt*{carlstrom96}, \\citealt{grego00}, \\citealt{patel00}, \\citealt{reese00}, \\citealt{joy01}, \\citealt{reese02}, \\citealt{bonamente06}, \\citealt{laroque06}). Each of these observations has however been long (1--60 days), because even the shortest baselines resolved out most of the SZ signal: some 90 per cent of the SZ flux from a massive cluster at $z=0.2$ is missed by a $600$-$\\lambda$ baseline, the smallest possible given the existing dish sizes of the RT and OVRO/BIMA.  Several other instruments have successfully observed the SZ effect, including the OVRO 5.5-m Telescope \\citep{herbig95}, SuZIE \\citep{holzapfel97}, ACBAR \\citep{cantalupo02}, CBI \\citep{udomprasert04}, VSA \\citep{lancaster05} and OCRA \\citep{lancaster07}. It has, however, long been clear that, for detailed imaging and blind surveying, a new generation of fast SZ imagers is required, with angular-scale and brightness-temperature sensitivities far better matched to the features of interest in clusters. Of course, such instruments should also be able to measure the CMB power spectrum at high angular scales ($\\ell > 2000$) and thus investigate the origins of, for example, the excess CMB power detected by CBI \\citep{cbi-excess}. They also are pertinent to the search for other non-Gaussian features such as topological defects.  Accordingly, we have designed, built and now operate the Arcminute Microkelvin Imager (AMI, see \\citealt{rk01} and AMI Collaboration \\citeyear{ami-a1914}), a pair of interferometer arrays operating around 15~GHz near Cambridge, UK. Before describing AMI in detail, we next review other new SZ instruments, as well as the advantages and disadvantages of interferometers.   \\subsection{Fast SZ Instruments} \\label{sec:intro:instruments}  Several powerful SZ instruments are now coming online. The majority are direct-imaging, focal-plane arrays of bolometers, operating at frequencies above 90~GHz: ACT (\\eg \\citealt{ACT06}) in Chile uses a 6-m antenna with a resolution of $1\\farcm44$ at \\eg 150~GHz; APEX (\\eg \\citealt{APEX06}) in Chile has a 12-m antenna with $1\\farcm0$ resolution at \\eg 150~GHz; SPT (\\eg \\citealt{SPT04}) at the South Pole has a 10-m antenna with resolution of $1\\farcm06$ at \\eg 150~GHz; the Planck satellite, due for launch in 2009 January, will detect clusters in SZ over the whole sky, but detecting distant clusters will be challenging due to its limited resolution (see \\eg \\citealt*{geisbuesch05}).  There are also three interferometers, operating at lower frequencies: AMI, see below; AMiBA (see \\eg \\citealt{amiba06}) in Hawaii, with the array comounted on a hexapod and currently at 90~GHz; and the SZA (see \\eg \\citealt{sza07}) in California with 3.5-m diameter antennas --- the SZA has been operating at 30~GHz with a resolution of $1\\arcmin$ and will be operated at 90~GHz as part of CARMA (\\texttt{www.mmarray.org}).   \\subsection{Advantages and Disadvantages of Interferometers} \\label{sec:intro:ints}  Interferometry has some significant advantages for observations where high sensitivity and low systematics are required (see \\eg \\citealt{church95}, \\citealt{lay00}, \\citealt{katy-thesis} and \\citealt{zwart2007}):  \\begin{enumerate}  \\item Stability of receivers. Short-term fluctuations in the gains of the front-end amplifiers are uncorrelated and therefore contribute only to the random noise.  \\item Emission from the atmosphere is largely uncorrelated (provided the antenna beams do not overlap within the troposphere), and associated receiver power fluctuations are completely uncorrelated and thus contribute only to the system noise.  \\item Fringe-rate filtering. Astronomical signals are modulated at the celestial fringe rate by the Earth's rotation. Unwanted interference from sources at terrestrial fringe rates, such as ground spill and geostationary satellite broadcasts, can be attenuated. Residual correlated atmospheric emission also has very little power in components that are synchronous with the astronomical fringe rate.  \\item  Interferometers have zero response to power which is uniformly spatially distributed (\\eg most of the atmospheric emission and the CMB average temperature), and hence do not respond to temporal fluctuations in that power.  \\item Each interferometer baseline responds to only a narrow range of spatial frequencies; as well as filtering out the large-scale atmospheric signals this can be used to \\eg minimize the response to the primary CMB anisotropies while retaining sensitivity to clusters.  \\item Foreground removal. Radio point sources can be detected by baselines chosen to be sensitive to smaller angular scales. This allows them to be measured simultaneously and subtracted.   \\end{enumerate}   \\noindent The main disadvantages of interferometers are complication (with the number of correlators going as the square of the number of detector elements) and a restriction on observing frequency. The SZ decrement has a maximum amplitude in intensity at approximately 130~GHz; contaminating synchroton emission from radio sources \\textit{generally} decreases with increasing frequency (falling in temperature as $\\nu^{-(\\alpha + 2)}$ where $\\alpha$ is the flux-density spectral index), although the effect of dust in galaxies on SZ observation at, say, 150~GHz is presently uncertain. Radio interferometers employing heterodyne receivers have costs and amplifier noise that rise with frequency, which in the past has limited interferometers to observing frequencies of around 30~GHz.   \\subsection{The Rest of this Paper} \\label{sec:intro:rest}  We next describe the design of the instrument (section \\ref{sec:design}), including the choice of frequency and the need for two interferometric arrays, the Large Array (LA) and the Small Array (SA). We describe the antennas, optics and array configurations of the SA and LA in section \\ref{sec:ant}. The amplifiers and cryostats are outlined in section \\ref{sec:cryo}, with the signal chain expounded in section \\ref{sec:if}. Section \\ref{sec:corr} covers correlator and readout systems, and section \\ref{sec:control}, the telescope control systems. Some lessons learned are outlined in section \\ref{sec:comm}.    \\begin{table} \\centering \\caption{AMI technical summary.}\\label{table:summary} \\begin{tabular}{{l}{c}{c}} \\hline% & SA                 & LA                  \\\\ \\hline% Antenna diameter           & 3.7~m              & 12.8~m              \\\\ Antenna efficiency         & 0.75               & 0.67                \\\\ Number of antennas         & 10                 & 8                   \\\\ Number of baselines        & 45                 & 28                  \\\\ Baseline lengths (current) & 5--20~m            & 18--110~m           \\\\ Primary beam (15.7~GHz)    & $20\\farcm1$        & $5\\farcm5$          \\\\ Synthesized beam           & $\\approx3\\arcmin $ & $\\approx30\\arcsec$  \\\\ Flux sensitivity           & 30~mJy~s$^{-1/2}$  & 3~mJy~s$^{-1/2}$    \\\\ Observing frequency        & \\multicolumn{2}{c}{13.9--18.2{~GHz}}     \\\\ Bandwidth                  & \\multicolumn{2}{c}{4.3{~GHz}}            \\\\ Number of channels         & \\multicolumn{2}{c}{6}                    \\\\ Channel bandwidth          & \\multicolumn{2}{c}{0.72{~GHz}}           \\\\ System temperature         & \\multicolumn{2}{c}{25~K}                 \\\\ Declination range          & $>-15\\degr$        & $>-20\\degr$         \\\\ Elevation limit            & $+20\\degr$         & $+5\\degr$           \\\\ Polarization measured      & \\multicolumn{2}{c}{I+Q}                  \\\\ \\hline% \\end{tabular} \\end{table} ",
        "conclusions": "\\begin{tabular}{{l}{c}{c}} \\hline% & SA                 & LA                  \\\\ \\hline% Antenna diameter           & 3.7~m              & 12.8~m              \\\\ Antenna efficiency         & 0.75               & 0.67                \\\\ Number of antennas         & 10                 & 8                   \\\\ Number of baselines        & 45                 & 28                  \\\\ Baseline lengths (current) & 5--20~m            & 18--110~m           \\\\ Primary beam (15.7~GHz)    & $20\\farcm1$        & $5\\farcm5$          \\\\ Synthesized beam           & $\\approx3\\arcmin $ & $\\approx30\\arcsec$  \\\\ Flux sensitivity           & 30~mJy~s$^{-1/2}$  & 3~mJy~s$^{-1/2}$    \\\\ Observing frequency        & \\multicolumn{2}{c}{13.9--18.2{~GHz}}     \\\\ Bandwidth                  & \\multicolumn{2}{c}{4.3{~GHz}}            \\\\ Number of channels         & \\multicolumn{2}{c}{6}                    \\\\ Channel bandwidth          & \\multicolumn{2}{c}{0.72{~GHz}}           \\\\ System temperature         & \\multicolumn{2}{c}{25~K}                 \\\\ Declination range          & $>-15\\degr$        & $>-20\\degr$         \\\\ Elevation limit            & $+20\\degr$         & $+5\\degr$           \\\\ Polarization measured      & \\multicolumn{2}{c}{I+Q}                  \\\\ \\hline% \\end{tabular} \\end{table}"
    },
    "0807/0807.3145_arXiv.txt": {
        "abstract": "A special type of Hamamatsu MPPC, with a sensitive area of 1.3$\\times$1.3~mm$^2$ containing 667 pixels with 50$\\times$50~$\\mu$m$^2$ each, has been developed for the near neutrino detector in the T2K long baseline neutrino experiment. About 60 000 MPPCs will be used in total to read out the plastic scintillator detectors with wavelength shifting fibers. We report on the basic performance of MPPCs produced for T2K. ",
        "introduction": "T2K (Tokai-to-Kamioka)~\\cite{T2K} is a long baseline neutrino oscillation experiment in Japan, using an intense beam from J-PARC accelerator at Tokai and the massive Super-Kamiokande detector 295~km away. The main goals of T2K are a sensitive search for the $\\nu_e$ appearance from $\\nu_\\mu$, which is related to the mixing angle $\\theta_{13}$, and precise measurements of `atmospheric' oscillation parameters. In order to achieve the aimed precision, good understanding of the beam properties and $\\nu$-nucleus interaction are indispensable. The `near detector' (ND) complex will be placed in Tokai to provide this information.  The T2K-ND~\\cite{T2K-ND280} consists of several sub-detectors with specific and complimentary functions. As the basic elements for particle detection, most of detectors will use the plastic scintillator read out by wavelength shifting (WLS) fibers. This is a widely used technique, especially in recent accelerator neutrino experiments~\\cite{sci-wls}. In those experiments, multi-anode PMTs (MAPMTs) have been used as the photosensor. For T2K, MAPMT is not a good candidate because some of detectors have to operate inside a magnetic field of 0.2~T and cope with a limited space available.  The following are major requirements for photosensors in T2K: \\begin{itemize} \\item More or equal photon detection efficiency than that of a MAPMT. \\item Compact to fit the limited space inside the magnet. \\item Operational in a magnetic field. \\item Good stability and low cost for a large number of readout channels (60 000). \\end{itemize}% We decided to use the Multi-Pixel Photon Counter (MPPC)~\\cite{MPPC, MPPC2} in August 2005. Since then, we continued the development collaborating with Hamamatsu and KEK Detector Technology Project. In this paper, we report on the performance of the MPPC developed for T2K. ",
        "conclusions": "New type of MPPCs have been developed for the T2K experiment. The T2K-MPPC is designed to have 1.3$\\times$1.3~mm$^2$ sensitive area in order to minimize the light loss when coupled to a WLS fiber of 1.0~mm diameter. The size of pixel is 50$\\times$50~$\\mu$m$^2$ and the number of pixels is 667.  In T2K, about 60 000 MPPCs will be used in total. The mass production has started in February 2008. The gain, breakdown voltage, noise rate, photon detection efficiency, and cross-talk and afterpulse rate of T2K-MPPCs are measured for each device. The device uniformity is found to be excellent based on the measurement of 5820 MPPCs. All of MPPCs satisfy our requirements. We have established techniques necessary for a large scale application of MPPC to the WLS readout.  The T2K experiment is planed to start in 2009. It will be the first experiment to use MPPCs in a large scale."
    },
    "0807/0807.3832_arXiv.txt": {
        "abstract": "We study the role of the unstable equilibrium points in the transfer of matter in a galaxy using the potential of a rotating triaxial system. In particular, we study the neighbourhood of these points for energy levels and for main model parameters where the zero velocity curves just open and form a bottleneck in the region. For these energies, the transfer of matter from the inner to the outer parts and vice versa starts being possible. We study how the dynamics around the unstable equilibrium points is driven, by performing a partial normal form scheme and by computing the invariant manifolds of periodic orbits and quasi-periodic orbits using the reduced Hamiltonian. In particular, we compute some homoclinic and heteroclinic orbits playing a crucial role. Our results also show that in slow rotating and/or axisymmetric systems the hyperbolic character of the equilibrium points is cancelled, so that no transfer of matter is possible through the bottleneck. ",
        "introduction": "\\label{sec:intro} This paper focuses on the study of the dynamics around the hyperbolic equilibrium points, $L_1$ and $L_2$, of a non-axisymmetric galactic potential, i.e. a potential whose principal axes on the $(x,y)$ plane are different. In particular, we analyse the role of the invariant manifolds associated with the unstable periodic orbits and with invariant tori around $L_1$ and $L_2$ in the large scale transfer of matter within the system. For this purpose, we use both semi-analytical and numerical techniques to compute the invariant manifolds and to study the role they play in the global morphology. This is a modern approach using dynamical systems which has successfully been applied in celestial mechanics and astrodynamics (e.g. \\cite{gom91,jor99,gom01,koo00,gom04}). A similar technique has already been used to explain the spiral arm and ring morphology in barred galaxies \\cite{rom06,rom07,ath08}. Previous theories believe that spiral arms are density waves in a disc galaxy \\cite{lin63}. The density waves propagate from the centre towards the principal resonances of the galaxy, where they damp down \\cite{too69}. Other replenishment theories have been proposed, therefore, to obtain long-lived spirals (see \\cite{ath84} for a review). Using this innovative approach, we obtain outer rings and long-lived spirals. This paper is intended to explain the geometrical behaviour of such structures in a general manner.  Such studies have many applications in galactic dynamics. Elliptical galaxies can be triaxial, (i.e. their principal axes can be all different from each other) while disc galaxies can contain several triaxial components, rotating or non-rotating, and with widely varying degree of non-axisymmetry, such as haloes, bulges, bars, or oval discs. In this paper, we consider a galactic model that describes a triaxial system. By studying its dynamics within a wide range of parameters, we will explain the dynamics of all the components and make links with their morphology. The logarithmic potential is a very suitable model for our purposes, because it has a simple expression, it is analytic, its series expansion can be calculated up to a high order, and it admits both semi-analytical and numerical treatments.  We choose a reference frame such that the origin of coordinates coincides with the centre of the galaxy and the triaxial system is fixed, i.e. a reference frame rotating with the system. The idea of the present analysis is based on the fact that if we consider a range of energies for which the particles are confined in the inner region defined by the zero velocity curve, in our case, for energies lower or equal to that of the unstable equilibrium points ($E_J(L_1)$), then there is no possible transfer of matter from the inner region of the galaxy to the outer region, or vice versa (see Fig.~\\ref{fig:bottleneck}(a) and (b)). In this paper, we focus on energies slightly larger than that of the unstable equilibrium point, for which an opening in form of a bottleneck appears. We will hereafter refer to this aperture as bottleneck. Transfer of matter can be possible through this bottleneck (see Fig.~\\ref{fig:bottleneck}(c)), and we study the objects that drive the motion in this region.  \\begin{figure} \\begin{center} \\includegraphics[scale=0.43,angle=-90.0]{plotcv0.ps} \\caption{Location of the equilibrium points and the zero velocity curves. {\\bf (a)} Zero velocity curves at an energy level smaller than the energy of $L_1$ and $L_2$. Three regions are thus defined, namely an inner (solid dark grey), an outer (solid white) and the forbidden region (hatched light grey), and no transfer of matter between the inner and outer regions is possible. {\\bf (b)} Zero velocity curves at the energy level of the unstable equilibrium points $L_1$ and $L_2$. {\\bf (c)} Zero velocity curves at an energy level larger than the energy of $L_1$ and $L_2$. A bottleneck appears and matter can transit from the inner to the outer regions (solid white) and vice versa.} \\label{fig:bottleneck} \\end{center} \\end{figure}  Therefore, our first goal is to study the neighbourhood of the unstable equilibrium points in the energy range where the bottleneck appears. This energy range is suitable to apply the normal form technique, which is known as the reduction to the centre manifold. As previously mentioned, it has been successfully applied to celestial mechanics problems, in particular to the \\rtbp, e.g. \\cite{gom91,gom01,jor99}. The reduced Hamiltonian obtained from the reduction to the centre manifold process gives a good qualitative description of the phase space near the equilibrium points and uncouples the centre manifold from its hyperbolic behaviour. The procedure is similar to the Birkhoff normal form, usually used to study the stability properties around the central equilibrium point of the galaxy and to compute the families of periodic orbits around it, e.g. \\cite{gus66,mir89,bel07}.  Our second goal is to study the behaviour of the hyperbolic invariant manifolds associated with the orbits contained in the centre manifold; these are essentially invariant manifolds of periodic orbits and invariant tori. As is well-known, stable and unstable invariant manifolds are dynamical features, that are responsible for the global dynamics in a dynamical system. In order to obtain the global picture of the transfer of mass, it is necessary to compute the invariant manifolds and to look for intersections between different parts, that is, obtaining the possible heteroclinic and homoclinic connections. These type of connections have been recently studied in the \\rtbpe and applied to obtaining transit and non-transit orbits in celestial mechanics \\cite{koo00,gom04}.  In Section~\\ref{sec:mod}, we describe the characteristics of the galactic model. In Section~\\ref{sec:mot}, we present the equations of motion and we study the effect of the main parameters on the linear behaviour of the equilibrium points. In Section~\\ref{sec:red}, we explain in detail the reduction to the centre manifold in the particular case of the logarithmic potential and perform a study of the practical convergence of the reduced Hamiltonian. In Section~\\ref{sec:res}, we compute the invariant manifolds associated with periodic orbits and quasi-periodic orbits and we study the role they play in the transfer of matter. We also perform an analysis on the variation of the free parameters of the potential. Finally, in Section~\\ref{sec:conc} we conclude. ",
        "conclusions": "\\label{sec:conc} In this paper we use a suitable galactic potential (the logarithmic potential) to perform a semi-analytic study of the neighbourhood of the unstable equilibrium points $L_1$ and $L_2$. For a detailed study of their neighbourhood, we use a partial normal form scheme and we find that the main objects are the planar and vertical families of Lyapunov orbits and invariant tori. We compute the invariant manifolds associated with both periodic orbits and quasi-periodic orbits using the reduced Hamiltonian and we find that the motion around $L_1$ and $L_2$ is mainly driven by the invariant manifolds of the planar Lyapunov orbits. However, we are also interested in the global structure of the galaxy. Thus, we study the possible homoclinic and heteroclinic connections between the planar periodic orbits. For such a purpose, we use suitable Poincar\\'e surfaces of section. We note that this approach has successfully been used in celestial mechanics and in this paper we apply it to a galactic dynamics problem, namely the formation of spiral arms and rings in barred galaxies, in comparison to other theories given so far.  Here we are interested in the hyperbolic behaviour of $L_1$ and $L_2$ and, particularly, in determining the role the invariant manifolds play in the transfer of matter for energy levels where the zero velocity curves are open (i.e. a range of energies somewhat larger than the energy of $L_1$ and $L_2$) and a bottleneck appears around $L_1$ and $L_2$. For such purpose, we also study the influence of the main model parameters, namely the pattern speed, $\\Omega$, and the shape parameters, $\\p$ and $\\q$. The logarithmic models are suitable for describing triaxial systems such as haloes, bars, bulges in disc galaxies or elliptical galaxies. The structures we have constructed using the invariant manifolds, however, are not globally dependent on the model characteristics \\cite{rom06,rom07}. This implies, for example, that if a system had some degree of rotation, these kind of structures should be present. Elliptical galaxies are triaxial systems that do not present any external feature, i.e. they are an ellipsoidal distribution of matter with different degrees of ellipticity and they present neither spiral arms nor rings. Observations show that elliptical galaxies barely rotate or do not rotate at all as a figure \\cite{bin87,sch82}. Our results are in agreement with this statement. We have shown that models that rotate slowly or do not rotate cancel the hyperbolic behaviour of the equilibrium points and, thus, no transfer or escape of matter is possible. On the other hand, bars in disc galaxies are non-axisymmetric components usually characterised in the literature by elliptical distributions of density \\cite{pfe84} although both observations and simulations show they might have more rectangular endings \\cite{ath90,ath02}. It is well-known that bars rotate and observations show that spiral arms or rings emanate from the ends of the bar. This characteristic is also consistent with our results which show that if the system rotates at a given angular velocity, the hyperbolic equilibrium points are present, so that the invariant manifolds drive the motion, and therefore, set the global morphology to the galaxy. We have seen that, depending on the rotation velocity and the shape of the bar, the morphology will be that of a barred spiral galaxy or that of a barred ringed galaxy."
    },
    "0807/0807.1538_arXiv.txt": {
        "abstract": "We analyze the photometric redshift catalog of the Sloan Digital Sky Survey Data Release 5 (SDSS DR5) to estimate the Fisher information in the galaxy angular power spectrum  with the help of the Rimes-Hamilton technique. It is found that the amount of Fisher information contained in the galaxy angular power spectrum is saturated at lensing multipole scale $300\\le l\\le 2000$ in the redshift range $0.1\\le$photo-z$<0.5$. At $l=2000$, the observed information is two orders of magnitude lower than the case of Gaussian fluctuations. This supports observationally that the translinear regime of the density power spectrum contains little independent information about the initial cosmological conditions, which is consistent with the numerical trend shown by Rimes-Hamilton. Our results also suggest that the Gaussian-noise description may not be valid in weak lensing measurements. ",
        "introduction": "Cosmic shear, which refers to the weak gravitational lensing by large-scale structure in the universe, is an invaluable tool in cosmology since it directly probes the gravitational clustering of dark matter in the universe without concerns about light-to-matter bias.  Very recently, the three dimensional map of dark matter distribution has been indeed reconstructed on cosmological scales from the measurements of cosmic shear \\citep{Massey07}, which marks an advent of {\\it lensing cosmology}.  Among many lensing observables, the angular power spectrum is regarded particularly important, since it is in principle capable of constraining the key cosmological parameters with high precision, including the enigmatic nature of dark energy \\citep[see][for a recent review]{HJ08}.  The success of the angular power spectrum as a cosmological probe, however, is subject to one critical issue: how much information does it preserve about the initial conditions of the universe?  A natural expectation is that some but not all information might have been destroyed in the subsequent nonlinear evolution. The amount of information contained in the power spectrum about the initial cosmological conditions is directly related to the precision of the cosmological parameters constrained by using the matter power spectrum. In other words, the non-Gaussian errors caused by the loss of information in the power spectrum would propagate into large uncertainties in the determination of the cosmological parameters. The error propagation and precision on cosmological parameters can be quantified in terms of the Fisher information contained in the matter power spectrum \\citep{TTH97}.  \\citet[][hereafter, RH05]{RH05} have for the first time estimated the Fisher information in the matter power spectrum. Basically, RH05 have calculated the information content about the amplitude of the linear power spectrum averaged over an ensemble of many N-body realizations and found very little independent information at the translinear scale. Somewhat surprisingly, however, RH05 also found that there is a sharp rise in the amount of information in the nonlinear scale \\citep[see also][]{RH06,RHS06}. This phenomenon has been also noted by the analytic work of \\citet{NSR06}.  Yet, according to the recent results derived by \\citet{NS07} based on the halo model, the Fisher information in the density power spectrum about all key cosmological parameters including the initial amplitude is highly degenerate both in the translinear and the nonlinear regime. Their work has indicated that it might not be possible to extract the initial cosmological conditions from the nonlinear dark matter power spectrum to a high statistical accuracy.  Thus, the previous numerical and analytic results forecast that the accuracy in the determination of the cosmological parameters in lensing cosmology may be lower due to the non-Gaussian errors. As the next generation of large surveys optimized for weak lensing will soon be on the pipeline, it is imperative to test observationally the information content in the lensing power spectrum. In this Letter we attempt to do this test by applying the RH05 technique to the photometric galaxy catalogs from the Sloan Digital Sky Survey Data Release 5 \\citep[SDSS DR5,][]{AM07} at typical lens redshifts $z\\sim 0.3$. ",
        "conclusions": "It is worth mentioning here that what you have calculated using the SDSS data is not rigorously the information about the amplitude of the power spectrum. But, it only approximates the information in the linear power spectrum to the extent that $d\\ln P/d\\ln A$ is near unity.  Our result provides an observational evidence for the previous theoretical clues that at translinear scale there is very little independent information in the matter power spectrum, consistent with light tracing matter till this scale \\citep{RH05,RH06,RHS06, NSR06,NS07}. Since the multipole scales where the information saturation is detected correspond to the weak lensing regime, this result has a direct impact on the weak lensing analyses. The loss of information in the angular power spectrum would lead to increasing non-Gaussianity in sampling errors. Unlike the usual assumption adopted in most weak lensing analyses that the non-Gaussianity contribution to the sampling errors for the angular power spectrum is marginal \\citep[][and references therein]{WH00,CH01}, our result suggests that it should be quite substantial and thus the Gaussian-noise description for the lensing power spectrum should not be valid."
    },
    "0807/0807.3904.txt": {
        "abstract": "In this series of four lectures, I discuss four important aspects of AGN host galaxies. In Lecture \\#1, I address the starburst-AGN connection. First, I briefly review the primary diagnostic tools that are used to quantify and distinguish star formation and nuclear activity.  Next I describe the best evidence for a connection between these two processes, first at low luminosity and then at high luminosity. In the last section, I summarize the main results and offer possible explanations. In Lecture \\#2, I discuss our current understanding of ultraluminous infrared galaxies (log[L$_{IR}$/L$_\\odot$] $\\ge$ 12; ULIRGs). First, I describe the general properties of ULIRGs, comparing the local sample with their distant counterparts. Then I discuss the role of ULIRGs in the formation and evolution of spheroids and their massive black holes. The discussion of their possible role in the metal enrichment of the IGM through superwinds is postponed until Lecture \\#3. In this third lecture, I discuss the importance of feedback processes in the local and distant universe.  The emphasis is on mechanical feedback. I describe the basic physics of winds, a few classic examples of winds in the local universe, the statistical properties of winds, near and far, and their impact on galaxy formation and evolution. A list of potential thesis projects is given at the end.  The fourth and final lecture is on elemental abundances as tracers of star formation. First, I explain the basic principles behind chemical evolution, and describe three simple models whose predictions are compared with observations in the Milky Way. Next I discuss and give an interpretation of the results of abundance determinations in local quiescent and starburst galaxies before discussing elemental abundances in the more distant universe. ",
        "introduction": "\\subsection{Introduction} \\label{}  The apparent connection between black hole driven nuclear activity and starburst activity on large scale has been the topic of debates for many years (e.g., see references in review by Veilleux 2001). More than ever, this topic is relevant to help us understand galaxy formation and evolution, the global star formation and metal enrichment history of the universe, and the origin of nuclear activity and associated black hole growth.  The existence of an apparently tight relation (Fig. 1) between black hole masses and spheroid masses (or velocity dispersions; Gebhardt et al.  2000; Ferrarese \\& Merritt 2000) points to a causal connection between spheroid formation (via a starburst) and black hole growth (via nuclear activity). A flurry of theoretical papers have tried to make sense of these results. In many scenarios, gas or radiation pressure from a starburst- and/or AGN-driven wind helps shut off the fuel supply to the black hole and terminate star formation in the surrounding galaxy (e.g., Murray, Quataert, \\& Thompson 2005). Regardless of the exact process involved in regulating the black hole and spheroid growths (this topic of negative feedback is covered in Lecture \\#3 of this series, \\S 3), the correlation indicates that the starburst-AGN connection is alive and well and has had a cosmologically important impact on galaxy formation and evolution.  \\begin{figure} %\\centerline{\\epsfig{figure=fig1_1.eps,width=1.0\\textwidth,angle=0}} %\\plotone{fig1.eps} \\caption{Black hole mass versus bulge luminosity (left) and the luminosity-weighted aperture dispersion within the effective radius (right). Green squares denote galaxies with maser detections, red triangles are from gas kinematics, and blue circles are from stellar kinematics. Solid and dotted lines are the best-fit correlations and their 68\\% confidence bands. (From Gebhardt et al. 2000)} \\end{figure}  To better understand this connection, one first needs to discuss the diagnostic tools that are used to detect and distinguish star formation and nuclear activity. This is done in \\S 1.2. In \\S 1.3, I describe key results from recent studies of low- and high-luminosity AGNs. In \\S 1.4, I summarize the results and offer a few possible explanations.  \\subsection{Star Formation Diagnostics} \\label{}  A very useful paper here is Kennicutt (1998). The material in \\S\\S 1.2.1 -- 1.2.4 is taken directly from that review and is therefore not described in detail.  \\subsubsection{Ultraviolet}  Hot, young stars emit copious amounts of UV radiation (e.g., Leitherer et al. 1999). The strength of the UV (1500 -- 2800 \\AA) continuum scales linearly with the luminosity of young stars and therefore with the star formation rate. For solar abundances and a Salpeter Initial Mass Function ($\\equiv$ IMF, 0.1 -- 100 M$_\\odot$):  \\begin{eqnarray*} SFR (M_\\odot~{\\rm yr}^{-1}) = 1.4 \\times 10^{-28} L_\\nu ({\\rm ergs}~{\\rm s}^{-1} {\\rm Hz}^{-1}) \\end{eqnarray*}  \\subsubsection{Recombination Lines}  Hot, young stars emit radiation that ionizes the surrounding ISM. The Str\\\"omgren sphere is the spherical volume of this HII region where the rate of ionizations balances the rate of recombinations.  The nebular lines produced in the HII region effectively re-emit the integrated stellar luminosity shortward of the Lyman limit ({\\em i.e.} $\\ge$ 13.6 eV). The intensity of these lines scales linearly with the number of hot, young stars and therefore the star formation rate. For solar abundances and a Salpeter IMF (0.1 -- 100 M$_\\odot$):  \\begin{eqnarray*} SFR (M_\\odot~{\\rm yr}^{-1}) = 7.9 \\times 10^{-42} L_{H\\alpha} ({\\rm ergs}~{\\rm s}^{-1}) \\end{eqnarray*}   \\subsubsection{Forbidden Lines}  H$\\alpha$ is redshifted out of the visible window beyond $z \\sim 0.5$. In principle, H$\\beta$ and the higher order Balmer emission lines could be used to estimate star formation rates, but these lines are weak and stellar absorption more strongly influences their emission-line fluxes than that of H$\\alpha$.  Neutral oxygen has the same ionization potential as hydrogen (13.6 eV). This means that ionized oxygen coexists with ionized hydrogen and therefore lines produced by ionized oxygen scales with the number of hot, young stars in HII regions. The strongest emission feature in the blue is the [O II] $\\lambda\\lambda$3726, 3729 forbidden-line doublet. The strength of these collisionally excited lines is sensitive to abundance and ionization state (electron temperature) of the gas, more so than the recombination lines. A rough calibration is:  \\begin{eqnarray*} SFR (M_\\odot~{\\rm yr}^{-1}) = (1.4 \\pm 0.4) \\times 10^{-41} L_{[O II]} ({\\rm ergs}~{\\rm s}^{-1}) \\end{eqnarray*}  Given its blue wavelength, this doublet is also more sensitive to dust extinction than H$\\alpha$.   \\subsubsection{Far-Infrared Continuum}  A significant fraction of the bolometric luminosity of a galaxy may be absorbed by interstellar dust and re-emitted in the thermal infrared (10 -- 300 $\\mu$m). The absorption cross-section of the dust is strongly peaked in the UV, so to first order the far-infrared emission scales with the star formation rate. In the limiting case of a dust cocoon surrounding a star-forming galaxy:  \\begin{eqnarray*} SFR (M_\\odot~{\\rm yr}^{-1}) = 4.5 \\times 10^{-44} L_{FIR} ({\\rm ergs}~{\\rm s}^{-1}) \\end{eqnarray*}  \\subsubsection{AGN Contamination}  In many cases (particularly at higher redshifts), it is difficult to spatially resolve the emission produced by star formation from that produced by the AGN. In this situation, one must use spectroscopic methods to disentangle the two processes. In some instance, one may use the strengths of stellar atmospheric features to quantify the starburst. These include the Balmer series, Ca I triplet $\\lambda\\lambda\\lambda$8498, 8542, and 8662 in the visible/deep-red and Si IV $\\lambda$1400, C IV $\\lambda$1550, and He II $\\lambda$1640 in the UV (Robert et al. 1993).  The emission lines produced in the ionized gas also bear the signature of the source of energy. The ionizing spectra of all but the hottest O stars cut off near the He II edge (54.4 eV), while AGNs are generally strong X-ray emitters. These high-energy photons have two effects on the emission line spectrum: (1) the material near the AGN is more highly ionized and emit strong high-ionization lines; (2) due to the strong energy dependence of the absorption cross-section ($\\sigma_\\nu \\propto \\nu^{-3}$), these high-energy photons are absorbed deeper into the gas clouds and produce extended partially ionized zones. These zones are strong emitters of collisionally excited low-ionization lines. One therefore expects an enhancement of {\\em both} high- and low-ionization lines in AGN relative to those in HII regions. Several diagnostic diagrams at optical and near-infrared wavelengths have been designed to specifically take advantage of these differences (e.g., Veilleux \\& Osterbrock 1987; Osterbrock, Tran, \\& Veilleux 1992; Kewley et al. 2001).  In deeply obscured galaxies, the UV, optical, and near-infrared diagnostics cannot be used to distinguish between starbursts and AGN. One must rely on the relative intensities of the mid-infrared fine structure lines and/or the strengths (equivalent widths) of the polycyclic aromatic hydrocarbon (PAH) features. The principles behind the use of the fine structure lines are roughly the same as for the optical/UV lines {\\em i.e.} use line ratio diagrams that take advantage of the fact that AGN are copious emitters of low- and high-ionization lines. The use of the PAH features nicely complements that of the fine structure lines since they are generally easier to detect in the fainter, more distant galaxies. The PAH features are less visible in AGN due to the much stronger continuum in these objects and possible PAH destruction.  \\subsection{Evidence for a Starburst-AGN Connection} \\label{}  Since the triggering mechanism for AGN activity probably depends on the luminosity of the AGN, I make a distinction in the following discussion between the nearby, low-luminosity Seyferts and Fanaroff-Riley type I (FR I) radio galaxies and the more distant and powerful quasars, Fanaroff-Riley type II (FR II) radio galaxies, and ultraluminous infrared galaxies (ULIRGs; log [L$_{IR}$/L$_\\odot$] $\\ge$ 12 by definition - this is the subject of Lecture \\#2; \\S 2).   \\subsubsection{Low-Luminosity AGN}  Direct evidence for recent nuclear star formation exists in a number of Seyfert 2 galaxies ({\\em i.e.} Seyferts without broad recombination lines).  Optical and ultraviolet spectroscopy of the nuclear regions of these galaxies often reveals the signatures of young and intermediate-age stars.  The stellar Ca II triplet feature at $\\lambda\\lambda\\lambda$8498, 8542, 8662 in Seyfert 2s has an equivalent width similar to that in normal galaxies while the stellar Mg~Ib $\\lambda$5175 is often weaker (Terlevich, Diaz, \\& Terlevich 1990; Cid Fernandes et al. 2004).  This result is difficult to explain with a combination of an old stellar population and a featureless power-law continuum from an AGN.  The most natural explanation is that young red supergiants contribute significantly to the continuum from the central regions.  Evidence for intermediate-age (a few 100 Myrs) stars in Seyfert galaxies is also apparent in the blue part of the spectrum, where the high-order Balmer series and He I absorption lines appear to be present in more than half of the brightest Seyfert 2 galaxies (e.g., Cid Fernandes \\& Terlevich 1995; Joguet et al. 2001; Gonz\\'alez Delgado, Heckman, \\& Leitherer 2001; Fig. 2). A few of these objects may even harbor a broad emission feature near 4680 \\AA, possibly the signature of a population of young (a few Myrs) Wolf-Rayet stars (Gonz\\'alez Delgado et al.  2001).  The ultraviolet continuum from some of the brightest UV Seyfert 2s also appears to be dominated by young stars based on the strength of absorption features typically formed in the photospheres and in the stellar winds of massive stars (e.g., Heckman et al. 1997; Gonz\\'alez Delgado et al.  1998; Fig. 2). The bolometric luminosities of these nuclear starbursts ($\\sim$ 10$^{10}$ L$_\\odot$) are similar to the estimated bolometric luminosities of their obscured Seyfert 1 nuclei.  This explains why UV-bright stellar clusters are more frequently detected in Seyfert 2s than in Seyfert 1s (Mu\\~noz Mar\\'in et al.  2007). The recent detection of near-infrared CN bands in Seyferts brings support to the idea that star formation is indeed connected to the AGN in these objects (Riffel et al. 2007).  \\begin{figure} %\\plottwo{fig2a}{fig2b} \\caption{($left$) UV spectrum of NGC 7130. It is displayed in log ($F_\\lambda$) to show the emission and absorption lines. The most important stellar wind and photospheric absorption lines are labeled. ($right$) Normalized optical spectrum of NGC 7130 (dashed line) plotted with the normalized spectrum of a B0 V star combined with a G0 V star (thick line). 60\\% of the light is from a B0 V and 40\\% from a G0 V star. The comparison shows that most of the stellar features in NGC 7130 are well reproduced by a combination of young (B0 V) and old (G0 V) stars. (From Gonzalez Delgado et al. 1998)} \\end{figure}  In recent years, SDSS has contributed significantly to our knowledge of local AGNs. For instance, Kauffmann et al. (2003b) have studied a sample of 22,623 narrow-line AGN with 0.02 $<$ z $<$ 0.3. They find that the hosts of low-luminosity AGN have a stellar population similar to that of normal early-type galaxies, while the hosts of high-luminosity AGN have much younger mean stellar ages. Indeed, young ($<$ 1 Gyr) stellar population appears to be a general property of AGN with high [O~III] luminosity (Figs. 3 and 4). This is true regardless of the presence of broad recombination lines ({\\em i.e.} types 1 and 2). The young stars are spread out over scales of at least a few kpc.  \\begin{figure} %\\plotone{fig3.eps} \\caption{ The strengths of the 4000 \\AA\\ break and H$\\delta$ absorption feature are plotted as a function of log $L[O~III]$. The solid line shows the median, while the dashed lines indicate the 16-84 percentiles of the 1/V$_{\\rm max}$ weighted distribution. (From Kauffmann et al. 2003b)} \\end{figure}  \\begin{figure} \\epsscale{0.5} %\\plotone{fig4.eps} \\caption{The fraction $F$ of AGN with H$\\delta_A$ values that are displaced by more than 3 $\\sigma$ above the local of star-forming galaxies is plotted as a function of log $L[O~III]$. The dashed line indicates the fraction of such systems in the subsample of normal massive galaxies. The dotted line indicates the fraction of such systems in the subsample of normal massive galaxies with $D_n(4000) < 1.6$. (From Kauffmann et al. 2003b)} \\end{figure}  \\subsubsection{High-Luminosity AGN}  Abundant molecular gas has been detected in radio galaxies and quasars (e.g., Evans et al. 2001), but is this gas forming stars? The extensive multiwavelength data set on these objects seems to indicate that starbursts are indeed present in local and distant quasars. Approximately 20 -- 30\\% of all PG~QSOs show an infrared excess L$_{IR}$/L$_{blue}$ $>$ 0.4.  IR-excess QSOs tend to have large dust and H$_2$ masses, suggesting that the infrared -- submm ``bump'' in the spectral energy distribution of PG~QSOs is due to star formation.  More direct evidence for a starburst-AGN connection has recently been found in local quasars from Spitzer mid-infrared spectroscopy (Schweitzer et al. 2006; Shi et al. 2007). PAH emission is detected in 11 of 26 PG QSOs and in the average spectrum of the other 15 PG QSOs (Schweitzer et al. 2006). The strength of the PAHs in these quasars is consistent with the far-infrared luminosity being produced primarily by U/LIRG-like starbursts with star formation rates of order 2 -- 300 M$_\\odot$ yr$^{-1}$ (Fig. 5). The strength of the starburst (measured by the FIR or PAH luminosity) correlates with that of the QSO (based on the 5100 \\AA\\ luminosity, a direct indicator of the mass accretion rate onto the black hole; Netzer et al. 2007; Fig. 6). This suggests a strong starburst -- AGN connection in these objects.  \\begin{figure} \\epsscale{1.0} %\\plotone{fig5.eps} \\caption{PAH fluxes $F$(PAH 7.7 $\\mu$) vs. $F$(60 $\\mu$m) for local QSOs and starburst-dominated ULIRGs. (From Schweitzer et al. 2006)} \\end{figure}  \\begin{figure} %\\plotone{fig6.eps} \\caption{$Top:$ Correlation of the optical (5100 \\AA) and FIR (60 $\\mu$m) continuum luminosities. $Bottom:$ $L(5100)$ vs. $L$(PAH 7.7 $\\mu$m) showing detections (filled squares) and upper limits (open squares). (From Netzer et al. (2007)} \\end{figure}  A recent Spitzer study of high-z quasars by Maiolino et al. (2007a) fail to detect PAH emission in these objects.  This would indicate that the correlation between star formation rate and AGN power ``saturates'' at high luminosities (Fig. 7). The ``flattening'' of the relation also seems to be present for CO emission (Maiolino et al. 2007b). It may therefore be that not enough fuel is available at high $z$ to fuel the starburst and match the strength of the AGN in these quasars.  \\begin{figure} %\\plotone{fig7.eps} \\caption{ ($a$) PAH(7.7~$\\mu$m) luminosity as a function of the QSO optical luminosity. Blue diamonds are data from Schweitzer et al. (2006). The red square is the upper limit obtained by the average spectrum of luminous, high-z QSOs.  ($b$) Distribution of the PAH(7.7~$\\mu$m) to optical luminosity ratio in the local QSOs sample of Schweitzer et al. (2006). The hatched region indicates upper limits.  The red vertical line indicate the upper limit inferred from the average spectrum of luminous QSOs at high-z. (From Maiolino et al.  2007a) } \\end{figure}  \\subsection{Discussion} \\label{}   As explained in \\S 1.3, starbursts often coexist with actively accreting supermassive black holes, but the starburst-AGN relation appears to be tighter at high luminosity than at low luminosity. This suggests that black hole fueling in low-luminosity AGN is more stochastic and does not necessarily scale with the surrounding starburst.  The fueling of AGN requires mass accretion rates $\\dot{M}$ $\\approx$ 1.7 (0.1 / $\\epsilon$)(L/10$^{46}$ ergs s$^{-1}$) M$_\\odot$ yr$^{-1}$, where $\\epsilon$ is the mass-to-energy conversion efficiency. A modest accretion rate of order $\\sim$ 0.01 M$_\\odot$ yr$^{-1}$ is therefore sufficient to power a Seyfert galaxy. Only a small fraction of the total gas content of a typical host galaxy is therefore necessary for the fueling of these low-luminosity AGNs. A broad range of mechanisms including intrinsic processes (e.g., stellar winds and collisions, dynamical friction of giant molecular clouds against stars; nuclear bars or spirals produced by gravitational instabilities in the disk) and external processes (e.g., minor galaxy interaction or mergers) may be at work in these objects.  So it may not be surprising after all that the power of these AGN does not necessarily scale with the surrounding starburst.  The stringent requirements on the mass accretion rates for luminous AGNs almost certainly require external processes such as major galaxy interactions or mergers to be involved in triggering and sustaining this high level of activity over $\\sim$ 10$^8$ years.  Starbursts are a natural consequence of major mergers. This topic will be discussed in more detail in Lecture \\#2 (\\S 2).  A lack of ``fuel'' (= molecular gas) may explain the apparent break in the starburst-AGN correlation at very high AGN luminosities. ",
        "conclusions": ""
    },
    "0807/0807.3017_arXiv.txt": {
        "abstract": "{One of the most dramatic events in the life of a low-mass star is the He flash, which takes place at the tip of the red giant branch and is followed by a series of secondary flashes before the star settles on the zero-age horizontal branch (ZAHB). Yet, no stars have ever been positively identified in this key phase in the life of a low-mass star (hereafter the ``pre-ZAHB'' phase).} {In this paper, we investigate the possibility that at least some pre-ZAHB stars may cross the instability strip, thus becoming variable stars whose properties might eventually lead to positive identifications. In particular, it has been suggested that some of the RR Lyrae stars with high period change rates ($\\dot{P}$) may in fact be pre-ZAHB stars. Here we present the first theoretical effort devoted to interpreting at least some of the high-$\\dot{P}$ stars as pre-ZAHB pulsators.} {We constructed an extensive grid of evolutionary tracks using the Garching Stellar Evolution Code (GARSTEC) for a chemical composition appropriate to the case of the globular cluster M3 (NGC~5272), where a number of stars with high $\\dot{P}$ values are found. We follow each star's pre-ZAHB evolution in detail, and compute the periods and period change rates for the stars lying inside the instability strip, also producing pre-ZAHB Monte Carlo simulations that are appropriate for the case of M3.} {Our results indicate that one should expect of order 1 pre-ZAHB star for every 60 or so bona-fide HB stars in M3. Among the pre-ZAHB stars, approximately 22\\% are expected to fall within the boundaries of the instability strip, presenting RR Lyrae-like pulsations. On average, these pre-ZAHB pulsators are expected to have longer periods than the bona-fide HB pulsators, and 76\\% of them are predicted to show negative $\\dot{P}$ values. While the most likely $\\dot{P}$ value for the pre-ZAHB variables is $\\approx -0.3$~d/Myr, more extreme $\\dot{P}$ values are also possible: 38\\% of the variables are predicted to have $\\dot{P} < -0.8$~d/Myr.} {It appears likely, therefore, that some~-- but certainly not all~-- of the RR Lyrae stars in M3 with high (absolute) $\\dot{P}$ values are in fact pre-ZAHB pulsators in disguise.} ",
        "introduction": "Low-mass stars, after exhausting their central hydrogen supply, develop increasingly more massive and more highly degenerate helium cores as they continue their evolution accross the Hertzsprung-Russell diagram (HRD). As the mass of the core increases due to the continued supply of He-rich material from the H-burning shell that surrounds the core, it also contracts and becomes increasingly denser and hotter. Eventually, the conditions in the core become ripe for the ignition of the He-burning triple-$\\alpha$ reactions. Since the core at this point is electron-degenerate, the process does not take place quiescently, but rather in the form of a thermonuclear runaway~-- the so-called {\\rm He flash}. The primary He flash, which does not take place at the center proper (due to the very efficient neutrino cooling there), is followed by a series of secondary flashes, each occurring increasingly closer to the star's center, until degeneracy is lifted throughout the core and quiescent He burning finally commences~-- and a so-called {\\em horizontal branch} (HB) star is born \\citep[see][for a recent review]{mc07}.  In the process just described, a low-mass star will typically change its luminosity by several orders of magnitude; its temperature may also increase very significantly. The inner structure of the star undergoes dramatic changes in the process as well. And yet, in spite of the importance of this evolutionary phase~-- hereafter the pre-zero-age HB (or pre-ZAHB) phase, to the best of our knowledge not a single such star has ever been positively identified. This dramatically limits the possibility of directly testing current models of pre-ZAHB evolution.  The reasons why it has proven so challenging to pinpoint pre-ZAHB stars are twofold: first, the pre-ZAHB phase is very short, compared to other major evolutionary phases in the life of a low-mass star. Second, pre-ZAHB stars are largely expected to overlap the loci occupied by asymptotic giant branch (AGB), HB, and red giant branch (RGB) stars in observed color-magnitude diagrams (CMD).  Bearing these difficulties in mind, \\citet{mc05,mc07} suggested an alternative approach for the detection of pre-ZAHB stars, namely, through {\\em stellar pulsations}. Indeed, some pre-ZAHB stars are expected to rapidly cross the Cepheid/RR Lyrae instability strip on their route from the RGB tip to the ZAHB, thus becoming pulsating stars along the way. At least some of these variables may present anomalously large period change rates ($\\dot{P}$), as a consequence of their very high evolutionary speeds, and so at least some pre-ZAHB stars could be singled out in view of their high $\\dot{P}$ values. As well known, high-$\\dot{P}$ values (e.g., $|\\dot{P}| \\gtrsim 0.1-0.15$~d/Myr) have indeed often been observed among field and cluster RR Lyrae stars \\citep[see, e.g.,][for a review]{hs95}, whereas canonical stellar evolution theory does not predict such high $\\dot{P}$ values for these stars, except towards the end of the HB phase when the star is already approaching the AGB \\citep[e.g.,][]{ywl91}. On the other hand, small, random mixing events associated with the composition redistribution in the cores of HB stars have also been proposed as possible causes for such high period changes in RR Lyrae stars \\citep{sr79}.  The main goal of the present paper is, accordingly, to perform the first systematic study of the expected properties of pre-ZAHB stars, with a view towards their detection using a combination of pulsation and CMD properties. In this paper we shall, in particular, attempt to describe the expected pre-ZAHB stellar evolution in the case of the Galactic globular cluster M3 (NGC~5272), which is well known for being a particularly RR Lyrae-rich cluster, and also for containing a number of high-$\\dot{P}$ RR Lyrae stars \\citep[see, e.g., the recent work by][and references therein]{cc01}.  In \\S\\ref{metodos} we describe the method employed to compute the evolutionary tracks, along with our Monte Carlo technique to produce synthetic pre-ZAHB distributions. In \\S\\ref{empirico} we present the empirical data that we use to compare our theoretical predictions with the observations. In \\S\\ref{resultados} we show the results from this comparison. Conclusions and final remarks are provided in \\S\\ref{conclusiones}. ",
        "conclusions": "In this paper, we have performed a detailed theoretical study of the expected evolutionary and pulsation properties of pre-ZAHB stars for a metallicity appropriate to the case of the Galactic globular cluster M3. Our results indicate that there should be around one pre-ZAHB star for every 60 or so bona-fide HB stars in the cluster. Based on extensive Monte Carlo simulations we find that, among the pre-ZAHB stars, of order 20\\% should exhibit RR Lyrae-like pulsations~-- to be compared with 21\\% and 57\\% blue HB- and red HB-like pre-ZAHB stars, respectively, and but a very small (of order $1-2\\%$) population of AGB/RGB-like pre-ZAHB stars.  The pre-ZAHB, RR Lyrae-like pulsators are characterized by relatively long periods (typically longer than the average for the cluster), and (especially) by period change rates falling mostly in the $[-1,0]$~d/Myr range, with $-0.3$~d/Myr being the most likely value. However, more extreme $\\dot{P}$ values are also possible; in fact, 38\\% of the variables in our simulation have $\\dot{P} < -0.8$~d/Myr. Relatively few (i.e., around 24\\%) of the pre-ZAHB pulsators are found with positive period change rates.  Our pre-ZAHB Monte Carlo simulations show, in addition, that for an M3-like HB mass distribution the pre-ZAHB CMD positions should be basically indistinguishable from those of bona-fide HB stars in the cluster. We predict that, in the specific case of M3, $\\sim 9$ pre-ZAHB stars should be present (assuming a total population of 530 HB stars), $\\sim 2$ of which most likely disguised as RR Lyrae stars with relatively long (and negative) period change rates. Since the number of stars in M3 with high $\\dot{P}$ is significantly larger \\citep[e.g.,][]{cc01}, we conclude that pre-ZAHB evolution may explain the high-$\\dot{P}$ phenomenon for some (but not all) of the RR Lyrae stars with high period change rates in this cluster.  In the future, we will extend our calculations to other clusters (and dwarf galaxies) with different metallicities and HB morphologies. These calculations should also prove important for the identification of candidate pre-ZAHB pulsators among field stars in the Galaxy and neighboring galaxies. With the era of the survey telescopes quickly approaching, the number of such stars available for study should dramatically increase, thus making empirical studies of pre-ZAHB evolution increasingly more feasible~-- at least in the long term."
    },
    "0807/0807.3684_arXiv.txt": {
        "abstract": "We have selected a small sample of post-AGB stars in transition towards the planetary nebula and present new Very Large Array multi-frequency high-angular resolution radio observations of them. The multi-frequency data are used to create and model the targets' radio continuum spectra, proving that these stars started their evolution as very young planetary nebulae. In the optically thin range, the slopes are compatible with the expected spectral index (-0.1). Two targets (IRAS 18062+2410 and 17423-1755) seem to be optically thick even at high frequency, as observed in a handful of other post-AGB stars in the literature, while a third one (IRAS 20462+3416) shows a possible contribution from cold dust. In IRAS 18062+2410, where we have three observations spanning a period of four years, we detect an increase in its flux density, similar to that observed in CRL 618.  High-angular resolution imaging shows bipolar structures that may be due to circumstellar tori, although a different hypothesis (i.e., jets) could also explain the observations. Further observations and monitoring of these sources will enable us to test the current evolutionary models of planetary nebulae. ",
        "introduction": "The evolution of planetary nebulae remains a challenging topic in astrophysics. Several studies have focused on the link between the post-Asymptotic Giant Branch (AGB) and planetary nebula phases to investigate the changes that occur in the evolutionary stage between these object classes. Planetary nebulae (PN) have been observed over a wide wavelength range, from X-ray to radio frequencies. Their complex morphologies and the shaping mechanisms that produced them are still a matter of debate.  Companion stars, jets from central stars, magnetic fields, dust tori, and interacting winds are some of the possible shaping agents suggested as being responsible for various PN morphologies, and an overlap of their actions cannot be ruled out.  Hubble Space Telescope (HST) imaging has demonstrated that pre-PN show non--spherical symmetry even at late spectral types, which implies that the shaping process in these sources starts very early after the AGB \\citep{balick&frank}. Non-spherically symmetric structures have also been observed in AGB stars that are part of a binary system \\citep{karovska}, which could indicate that some shaping agents are already active in this evolutionary phase. In general, the current theory of PN evolution is based on the Interacting Stellar Wind (ISW) model \\citep{kwok1978} and its generalised version \\citep{kahn}. However, this model does not account for shaping because it assumes that an asymmetric distribution of matter is already present when the wind interaction occurs. Other models take into account the possible role of jets, and have been successfully applied to some nebulae \\citep{sahai98}. It has been suggested that the interaction with a companion object, even a massive planet, may provide the necessary asymmetry \\citep{soker06}. Also, large scale magnetic fields, influencing or determining the shapes, might be sustained by a dynamo process \\citep{blackman}.  The birth of a PN is defined by the observation of an ionisation front, which is itself a shaping agent and could heavily influence the morphology established in earlier evolutionary phases, disrupting the molecular and dust circumstellar shells. Near- and mid-IR images of PN, which trace the molecular and warm dust emission, have been compared to optical line images, which trace the ionised elements, and have shown the presence of similar structures \\citep{latter}. Therefore it seems that molecular and ionised gas and dust grains can spatially coexist in these sources. This makes the investigation of the spatial distribution and physical properties of the ionised component in these envelopes even more compelling.  In this context, important information can be provided by observations of very young PN, or pre-PN, where the physical processes associated with PN formation are still occurring. To investigate the properties of these rare objects in transition from the post-AGB to PN, we have selected a sample of pre-PN and searched for radio emission from ionised shells. The targets were selected from stars classified in the literature as hot post-AGB candidates, showing strong far-IR excess and B spectral type features. We detected radio emission in 10 sources in the selected sample \\citep{umana}. The detection of radio continuum emission is proof of the presence of free electrons and therefore of an ionised shell.  For a better understanding of these sources, we have performed a multi-frequency follow-up to build up radio continuum spectra and collect information about the physical conditions in the nebulae. Several nebular parameters can be determined by radio flux density measurements (for example, electron density and ionised mass) but most calculations rely on the assumption of an optically thin emission. Although this is presumably the case for observations at frequencies higher than 5 GHz, only the construction of the whole centimetre continuum spectrum can enable us to distinguish between optically thick and thin regimes.  VLA A array observations have also given us insight in the morphologies of the envelopes. ",
        "conclusions": "We have presented radio continuum observations of a small sample of pre-PN. The observations were carried out to inspect both the morphology and the physical conditions of the ionised component in these nebulae. Compact planetary nebulae are, in principle, only one possible explanation for radio-frequency continuum emission from evolved stars, the others being stellar photosphere, chromosphere, and circumstellar dust \\citep{knapp}. Our observations rule out the possibility of a stellar wind or dust origin for the emission, since the obtained spectra are not compatible with such hypotheses. The dust would exhibit an approximate black body spectrum ($S_\\nu \\propto \\nu^2$ in the radio range), and a stellar wind would have a 0.6 spectral index over a wide range of frequencies \\citep{panagia}.  %  We have modelled the observed spectra assuming that the emission arises in spherical shells around the central stars, with the density decreasing as $r^{-2}$. The values of electron densities and ionised masses found for our targets confirm their nature as young PN, since they fall within the ranges expected for such objects. The data seem to point to a somewhat steeper than expected slope beyond 5~GHz but, by spectral fitting, we calculated spectral indexes that match, to within errors, the theoretical value of -0.1.  IRAS~18062+2410 appears to be optically thick up to 8.4~GHz and also shows flux density variations on a time scale of a few years. Such variations are rarely observed in the post-AGB/young-PN transition phase and indicate that this star belongs to a group of objects having such properties as those observed in CRL 618. IRAS 17423-1755 also seems to be optically thick at high frequency, but the complete spectrum is needed to confirm this result. Quite interestingly, IRAS 20462+3416, whose overall centimetre spectrum is very flat, has a larger flux density at 22.4 GHz than at lower frequencies. We speculate that this may be due to a contribution from cold dust.  The high-angular resolution observations that we have presented show that the ionisation in these targets has already involved the walls of the cavity around the central star and is not limited to the tenuous gas within the cavity. All of the targets have been at least partly resolved, except IRAS 18062+2410. Almost all of the envelopes show two peaks of emission in a somewhat extended nebulosity that can be explained as an opacity effect, possibly due to a circumstellar torus, or the onset of jets. Unlike the other targets, the envelope in IRAS 22023+5249 shows more of an elongation along an axis than two bright peaks. Such a morphology would be better explained by the action of jets rather than the presence of a torus."
    },
    "0807/0807.0152_arXiv.txt": {
        "abstract": "\\noindent \\citet{Venemans2005} found evidence for an overdensity of Ly$\\alpha$ emission line galaxies associated with the radio galaxy MRC 0316--257 at $z=3.13$ indicating the presence of a massive protocluster. Here, we present the results of a search for additional star-forming galaxies and AGN within the protocluster. Narrow-band infrared imaging was used to select candidate [O{\\sc iii}] emitters in a 1.1$\\times$1.1 Mpc$^2$ region around the radio galaxy. Thirteen candidates have been detected. Four of these are among the previously confirmed sample of Ly$\\alpha$ galaxies, and an additional three have been confirmed through follow-up infrared spectroscopy. The three newly confirmed objects lie within a few hundred km s$^{-1}$ of each other, but are blueshifted with respect to the radio galaxy and Ly$\\alpha$ emitters by $\\sim2100$ km s$^{-1}$. Although the sample is currently small, our results indicate that the radio--selected protocluster is forming at the centre of a larger, $\\sim60$ co-moving Mpc super-structure. On the basis of an HST/ACS imaging study we calculate dust-corrected star-formation rates and investigate morphologies and sizes of the [O{\\sc iii}] candidate emitters. From a comparison of the star formation rate derived from UV-continuum and [O{\\sc iii}] emission, we conclude that at least two of the [O{\\sc iii}] galaxies harbour an AGN which ionized the O$^{+}$ gas. ",
        "introduction": "To understand the formation and evolution of galaxy clusters, it is desirable to find and study their high redshift progenitors. Although galaxy clusters have been found out to a redshift of $z=1.5$ \\citep{Mullis2005,Stanford2005}, their higher redshift progenitors are sparse and difficult to find. A successful technique for finding more distant structures that by--passes the need for surveying very large areas of the sky is to search for emission-line galaxies in the neighbourhood of luminous high--redshift radio galaxies (HzRGs) using narrow-band imaging.  Multiwavelength studies of HzRGs have resulted in strong evidence that they are massive forming galaxies  \\citep[e.g.][]{Seymour2007,Villar-Martin2006}  and are frequently associated with overdensities of emission-line galaxies \\citep{Venemans2007}. These overdense regions in the early universe are the likely progenitors of local galaxy clusters or groups and are termed ''protoclusters\". PROCESS (PROtoCluster Evolution Systematic Study) is a project designed to use a few key radio-selected protoclusters with $2 \\le z \\le 5$ to investigate the formation of and evolution of various populations of galaxies in dense environments \\citep{Overzier2006,Overzier2007,Overzier2008,Venemans2005,Venemans2007}.  This article presents observations of the protocluster surrounding the PROCESS radio galaxy MRC\\,0316-257. The associated $1.5 \\ {\\rm Jy}$ radio source was listed in the 408 MHz Molonglo Reference Catalogue \\citep{Large1981} and was optically identified with a galaxy at $z=3.13$ by \\citet{McCarthy1990}. \\citet{LeFevre1996} spectroscopically confirmed two Ly$\\alpha$ emitting companions to the HzRG, indicating that the radio galaxy is located in a dense environment.  Recently, \\citet{Venemans2005,Venemans2007} confirmed 31 Ly$\\alpha$ emitters at a similar redshift of MRC\\,0316-257. The corresponding overdensity is approximately 3.3 times the galaxy field density at this redshift. The protocluster redshift of $z \\sim 3.13$ corresponds to an epoch when both the cosmic star formation rate and the quasar luminosity function were at their peak, indicating that this is a key epoch for studying the evolution of different populations of galaxies. We have identified and studied additional galaxies in the MRC\\,0316-257 protocluster on the basis of their redshifted [O{\\sc iii}] emission. The observations that we shall discuss here consist of infrared imaging and spectroscopy with ESO's Very Large Telescope (VLT) and deep optical imaging with the Advanced Camera for Surveys (ACS) on the Hubble Space Telescope (HST).  Section 2 of this article is an outline of the observations and the data reduction. In Section 3 we present results from the VLT search programme and the deep ACS images. Corrected star formation rates derived from the UV fluxes are used to discriminate between star-forming galaxies and obscured AGNs and the morphologies and sizes of the candidate emitters are discussed. The implications of our results for the space density of [O{\\sc iii}] emitting galaxies and the origin of the [O{\\sc iii}] emission are discussed in Section 4 and the conclusions of the article are presented in Section 5. We assume  a flat cosmology  with H$_0$=71 [km sec$^{-1}$ Mpc$^{-1}$] and $\\Omega_{\\rm m}$=0.27 \\citep{Spergel2003}. At the distance of MRC\\,0316-257 an angular scale of 1\\,arcsec corresponds to a projected linear scale of 7.73\\,kpc. All magnitudes are given in the AB system \\citep{Oke1974}. ",
        "conclusions": "Searching for [O{\\sc iii}]--emitting galaxies is a new feasible method for detecting protocluster members.  We have detected a new population of [O{\\sc iii}] emitting galaxies in the neighbourhood of the radio galaxy MRC\\,0316-257 at $z$=3.13. About half of the [O{\\sc iii}] candidates emitters are LBGs and a third were also previously detected by the Ly$\\alpha$ selection technique. The [O{\\sc iii}] technique complements narrow-band searches using Ly$\\alpha$ and H$\\alpha$ emission, and observations of the Lyman and Balmer breaks, for finding members of protoclusters. All of these different galaxy selection techniques are needed to study the different galaxy populations and to obtain a complete understanding of protocluster evolution. 13 candidate [O{\\sc iii}] emitters were detected, including the radio galaxy, and 8 of these were spectroscopically confirmed. Three [O{\\sc iii}] emitting galaxies lie in a small redshift interval at  $3.095<z<3.105$, which is blushifted by 2100\\kmps\\ with respect to the overdensity of Ly$\\alpha$ emitting galaxies associated with the protocluster structure surrounding the radio galaxy.  These three [O{\\sc iii}] emitters may trace a substructure that lies in front of the previously known 0316-257 protocluster. Further narrow-band imaging and spectroscopy of this field are necessary to establish the existence of such a structure. A possible explanation for the [O{\\sc iii}]--Ly$\\alpha$ redshift separation are the difference in the filter band pass together with the existence of a possible big supersructure around the radio galaxy. The observed space density of [O{\\sc iii}] galaxies is consistent with the overdensity of 3.3 derived for Ly$\\alpha$ galaxies in the MRC\\,0316-257 protocluster by \\citet{Venemans2005} The detected [O{\\sc iii}] candidate emitters exhibit a wide range of morphologies. About half have compact structures with an unresolved core on the HST ACS image. All three spectroscopically confirmed [O{\\sc iii}] emitters are compact. The median size of these objects is typically $\\le 0\\farcs24$, which at the average redshift $z=3.11$ correspond to $\\sim 2$ kpc. The remaining sources are clumpy, consistent with merging or interacting systems. Using the next generation of wide-field infrared imagers and spectrographs, it will be possible to increase the number of high-z detected [O{\\sc iii}] emitters in the MRC\\,0316-257 by an order of magnitude and carry out comprehensive infrared studies of several protoclusters at high redshift. Not only will such observations result in greatly improved statistics, but they will also facilitate detailed mapping of the cosmic web at high redshifts."
    },
    "0807/0807.4221_arXiv.txt": {
        "abstract": "{The fast spinning B-star Regulus has recently been found to be orbited by a fainter companion in a close circular path with orbital period  $P_{\\rm b} = 40.11(2)$ d. Being its equatorial radius $R_{\\rm e}$ $32\\%$ larger than the polar one $R_{\\rm p}$, Regulus possesses a remarkable quadrupole mass moment $Q$.} {We investigate the effects of $Q$ on the orbital period $P_{\\rm b}$ of its companion in order to see if they are measurable, given the present-day level of accuracy in measuring $P_{\\rm b}$. Conversely, we will look for deviations from the third Kepler law, attributed to the quadrupole mass moment $Q$ of Regulus, to constrain the ratio $\\gamma=m/M$ of the system's masses.} {The impact of $Q$ on the orbital period is analytically worked out with a straightforward perturbative approach. The resulting correction $P^Q$ is compared to other competing dynamical effects. $P^Q$ and the Keplerian period $P^{\\rm Kep}$ are expressed in terms of the phenomenologically determined system's parameters; $\\gamma$ is treated as an unknown. $P^Q$ is compared to the observational accuracy in measuring the orbital period $\\delta P_{\\rm b}=0.02$ d and to the systematic uncertainty  $\\delta(P^{\\rm Kep})$ due to the errors in the system's parameters entering it. The discrepancy $\\Delta P=|P_{\\rm b}-P^{\\rm Kep}|$ is examined in order to see for which values of $\\gamma$ it becomes statistically significant. The physical meaning of the obtained range of values for $\\gamma$ is discussed in terms of $Q$.} {$P^Q$ is larger than $\\delta P_{\\rm b}$ but still smaller than the systematic uncertainty in $P^{\\rm Kep}$ by two orders of magnitude. The major sources of bias are the velocity semiamplitude $K$ of the motion of the primary and its mass $M$. Assuming edge-on configuration, i.e. $i=90$ deg, if $\\gamma\\gtrsim 0.096$ $Q$ would be positive, i.e. Regulus would be prolate, contrary to the observations. If $\\gamma\\lesssim 0.078$ $Q$ would be negative, but its magnitude would be one-two orders of magnitude larger than the approximate estimate $Q\\approx M(R^2_{\\rm p}-R^2_{\\rm e})=-2.4\\pm 0.5\\times 10^{49}\\ {\\rm kg\\ m}^2$. } {Regulus is the first extrasolar binary system in which the orbital effects of the asphericity of the primary are larger than the observational sensitivity; moreover, no other competing aliasing orbital effects are present. Thus, it is desirable that it will become the object of future intensive observational campaigns in order to reduce the systematic uncertainty due to the system's parameters below the measurability threshold.} ",
        "introduction": "Regulus ($\\alpha$ Leo, HR 3982, HD 87901) is a nearby ($d=24.3\\pm 0.2$ pc \\citep{vanL07}) intermediate mass star of spectral class B \\citep{Joh53,Gr03} and has the peculiarity of being animated by a very fast rotation. Indeed, combined interferometric and spectroscopic studies \\citep{McAL05} have shown that its equatorial radius $R_{\\rm e}$ is $32\\%$ larger than the polar one $R_{\\rm p}$, and its rotation period is 15.9 hr, which corresponds to an equatorial rotation speed that is $86\\%$ of the critical breakup velocity. Other fast spinning stars whose oblateness has been interferometrically measured  are Altair ($\\alpha$ Aquil{\\ae}, HD 187642 \\citep{Alt}) and Achernar ($\\alpha$ Eri, HD 10144 \\citep{Acher}).  In principle, a companion orbiting not too far from such  highly deformed stars would experience relevant orbital effects induced by the primary's oblateness which could be used  to dynamically put constraints either on such an important stellar physical parameter or on the orbital/physical parameters of the system. In fact, the Be star Achernar recently turned out to be orbited by a faint companion \\citep{Acherdoppia} which should be a A1V-A3V star \\citep{AcherB}, but since its orbital period should amount to about 15 yr it has not yet been possible to obtain its orbital parameters. No companions are known for Altair. The situation is more favorable for Regulus since spectroscopic observations over the last few years have demonstrated that it has a close\\footnote{In fact, Regulus has a known wide companion $\\alpha$ Leo B at a separation of $\\approx 175''$ which is itself a binary \\citep{McAL05}, but it has a far too great separation to have ever interacted directly with Regulus. } companion in a circular orbit described in 40.11 d \\citep{Gies08}.  In this paper we will investigate the possibility of detecting some dynamical orbital effects induced by the oblateness of Regulus on its companion and will use them to put some constraints on the mass of the secondary. ",
        "conclusions": "In this paper we investigated the impact of the huge quadrupole mass moment $Q$ of Regulus on the orbital period of its faint companion.  It turns out that, for a reasonable estimate of its value, $Q$ induces a correction to the Keplerian period which could, in principle, be measured because it is larger than the error in phenomenologically measuring the orbital period $P_{\\rm b}$. However, according to the present-day level of knowledge of the system's parameters, the uncertainty in the Keplerian period is larger than the correction due to $Q$ by about two orders of magnitude; the most important sources of errors are the velocity semiamplitude $K$ and the mass $M$ of Regulus.  An analysis of deviations from the third Kepler law as due to $Q$ for the edge-on orbital configuration and $M=3.4(2)$M$_{\\odot}$ has shown that values of the ratio $\\gamma$ of the secondary's mass to the primary's one larger than about 0.096 are to be ruled out because they would yield a positive quadrupole mass moment for Regulus, i.e. it would be prolate, contrary to the observations. On the other hand, for $\\gamma\\lesssim 0.078$ $Q$ would be negative, but with an unlikely large value.  Further observational campaigns with different techniques would be of great significance in the attempt of reducing the overall uncertainty in the system's parameters down to the level required for detecting the effect of the shape of Regulus."
    },
    "0807/0807.1890_arXiv.txt": {
        "abstract": "We present the first results of a search for \\lya\\ emitters (LAEs) in the DEEP2 spectroscopic database that uses a search technique that is different from but complementary to traditional narrowband imaging surveys.  We have visually inspected $\\sim$20\\% of the available DEEP2 spectroscopic data and have found nine high-quality LAEs with clearly asymmetric line profiles and an additional ten objects of lower quality, some of which may also be LAEs.  Our survey is most sensitive to LAEs at $z$=4.4--4.9 and that is indeed where all but one of our high-quality objects are found. We find the number density of our spectroscopically-discovered LAEs to be consistent with those found in narrowband imaging searches. The combined, averaged spectrum of our nine high-quality objects is well fit by a two-component model, with a second, lower-amplitude component redshifted by $\\sim$420~\\kmpersec\\ with respect to the primary \\lya\\ line, consistent with large-scale outflows from these objects.  We conclude by discussing the advantages and future prospects of blank-sky spectroscopic surveys for high-$z$ LAEs. ",
        "introduction": "Over the last decade, studies of how galaxies form and evolve have been pushed to unprecedented look-back times, allowing us to observe galaxies at epochs when the universe was only a fraction of its present age.  A number of surveys for high-redshift ($z$$>$2) galaxies has been undertaken, including many in the rest-frame UV/optical (e.g., Steidel et al.\\ 1996, 1999, 2003; Sawicki et al.\\ 1997; Franx et al.\\ 2003; Daddi et al.\\ 2004; Sawicki \\& Thompson 2005; Iwata et al.\\ 2007) and far-IR (e.g., Barger et al.\\ 1998; Hughes et al.\\ 1998; Blain et al.\\ 1999; Smail et al.\\  2002; Chapman, et al.\\ 2003; Webb et al.\\ 2003; Sawicki \\& Webb, 2005; Coppin et al.\\ 2006). The study of galaxies selected by their strong \\lya\\ emission allows an investigation of a population that's potentially very different from these other high-$z$ galaxy populations.   \\lya\\ selection not only gives us the prospect of probing galaxies with very faint spectral continuum levels, but also may select systems in very early stages of a starburst (e.g., Malhotra \\& Rhoads 2002).  It is thus not surprising that over the last several years much observational effort has gone into \\lya\\ surveys (e.g., Ajiki et al.\\ 2003; Hu et al.\\ 1998, 2004; Kudritzki et al.\\ 2000; Ouchi et al.\\ 2003;  Rhoads et al.\\ 2000, 2003; Rhoads \\& Malhotra 2001;  Santos et al.\\ 2004). Indeed, the short rest-frame wavelength of the \\lya1216 line allows galaxies to be detected and spectroscopically confirmed to very high redshifts, $z$$\\lesssim$7, without the need for infrared observations.  The ability to push to such high redshifts raises the intriguing possibility of probing the tail end of the epoch of reionization, regarded by many as marking the onset of galaxy formation in the Universe.  The expected evolution of the IGM neutral fraction across the reionization epoch has led to the prediction that the observed \\lya\\ luminosity function (LF) will undergo rapid evolution as IGM opacity to \\lya\\ photons drops with decreasing redshift (e.g., Haiman \\& Spaans 1999; Malhotra \\& Rhoads 2004; Kashihawa et al.\\ 2006) and attempts have been made to detect the signature of this effect albeit with contradictory results (Malhotra \\& Rhoads 2004; Kashikawa 2006).  At slightly lower redshifts, the LF of high-$z$ galaxies may be evolving due to evolution in the intrinsic properties of galaxies (e.g., Sawicki \\& Thompson 2006; Iwata et al.\\ 2007; see also Bouwens et al.\\ 2007). Consequently even if the observed evolution of the \\lya\\ LF proves to be real, it will remain an open question as to whether it reflects an evolution in the opacity of the IGM, or evolution in the \\emph{intrinsic} properties of galaxies at a time when they were young.  To believe the former we must understand the latter, and thus it is important to study the properties of Lyman Alpha Emitters (LAEs) in detail not only at very high redshift, $z$$>$6, but also as a function of time.  To date, narrowband imaging searches for high-$z$ LAEs have been the most fruitful approach to finding LAEs and have yielded significant samples of galaxies at \\zs3--6.5, with several dozen objects now confirmed spectroscopically by various teams (e.g.,  Rhoads et al.\\ 2003; Hu et al.\\ 2004; Dawson et al.\\ 2004, 2007; Shimasaku et al.\\ 2006; Kashikawa et al.\\ 2006).   An alternative to narrowband imaging surveys lies in direct, blank-sky spectroscopic searches whose strategy is to bypass the photometry stage and find galaxies directly in blank-sky spectra. One advantage of this approach is the potential saving of the time associated with deep narrowband imaging; another is the gain in sensitivity that results from working at the natural linewidth of the line ($\\sim$10--20\\,\\AA) rather than fighting against the sky background admitted by the full width of the narrowband filter ($\\sim$100--200\\,\\AA) (see Martin \\& Sawicki 2004).  Until recently, the chief disadvantage of spectroscopic searches has been the small sky area accessible to most spectrographs. Serendipitous spectroscopic discoveries of single emission-line objects (e.g., Franx et al.\\ 1997; Dawson et al.\\ 2002), while useful, do not yield {\\it samples} of galaxies suitable for statistical analysis.  Spectroscopic searches of regions enhanced by gravitational lensing have met with some success (Santos et al.\\ 2004; Stark et al.\\ 2007), probe very deep into the luminosity function but over very small areas.  Until recently, multislit blank-sky searches (e.g., Crampton \\& Lilly 1999; Martin \\& Sawicki 2004; Tran et al.\\ 2004) have found many line emitters but no confirmed LAEs;  only now are the sky areas accessible to this technique becoming large enough to report successful detections (Martin, Sawicki, Dressler, \\& McCarthy 2008).  In the present work we take an approach that is hybrid to the dedicated blank-sky spectroscopic surveys and the serendipitous discoveries.  Specifically, we carry out a dedicated search for serendipitous LAEs in the extremely large, existing spectroscopic database of the DEEP2 survey.  The present paper reports our first discovery of a sample of \\lya\\ emitters using these data and illustrates some of the constraints that this approach can place on the properties of these objects.  The paper is organized as follows.  In \\S\\ref{data.sec} we describe our DEEP2 data and their suitability to the task at hand.  In \\S\\ref{deep2serendips.sec} we describe our search for LAEs and report on the objects that we have found. In \\S\\ref{ConfidenceTests.sec} we present a number of tests with which we build confidence in the high-redshift nature of our objects. In \\S\\ref{numberdensity.sec} we constrain the number density and luminosity function of high-$z$ LAEs and in \\S\\ref{SpectralProperties.sec}  we study the spectral properties of the \\lya\\ emission line at high redshift. In \\S\\ref{discussion.sec} we discuss future prospects for finding LAEs using deep blank-field spectroscopy and summarize our main findings.  Throughout this paper we adopt the $\\Omega_M$=0.3, $\\Omega_\\Lambda$=0.7, $H_0$=70~km~s$^{-1}$~Mpc$^{-1}$ cosmology. ",
        "conclusions": "\\label{discussion.sec}  Deep LAE surveys are an important tool for the study of formation and evolution of galaxies. Despite their many very important successes, narrowband imaging surveys for LAEs have their limitations:  the numbers of spectroscopically-confirmed objects are still modest and the lack of detailed spectroscopy limits the ability to study in detail such properties as galactic outflows. Moreover,  the lack of spectroscopic redshifts for all the LAEs introduces biases in luminosity function and correlation function studies.  Thus, while we must continue to press with narrowband imaging surveys, we must also develop other, complementary  techniques.  Martin et al.\\ (2008) compare in detail the merits of spectroscopic and narrowband imaging imaging LAE surveys and discuss the future prospects for these techniques in the coming era of extremely large telescopes; here, we focus more narrowly on the near-term prospects of exploiting the remaining $\\sim$80\\% of the DEEP2 spectroscopic database.  In this paper we have demonstrated that a large blank-sky spectroscopic survey can find significant numbers of LAEs.  Blank-sky spectroscopic searches have been attempted in the past (e.g., Crampton \\& Lilly 1999; Martin \\& Sawicki 2004; Tran et al.\\ 2004; Martin et al.\\ 2008), although the number of LAEs fund by them to date is very small.   Here we have piggy-backed on a large intermediate-$z$ survey that is DEEP2 and in just $\\sim$20\\% of these data we have found 9 objects that are bona fide LAEs (and have a further 10 of which some may also be such).  So far, ours is a modest sample, but we have been able to use it to show that spectroscopic LAE searches can do well at constraining properties of the population.  Specifically, we found that the number density of objects in our survey is consistent with those found by narrowband imaging searches (\\S~\\ref{numberdensity.sec}), and --- using only our stacked discovery spectra with no additional follow-up data, we confirmed the interpretation that LAEs are associated with large-scale outflows of material (\\S~\\ref{SpectralProperties.sec}). Much more could be done with a bigger sample and with deep follow-up spectroscopy.  A simple extension of the by-eye search described in this paper to the remaining $\\sim$80\\% of the DEEP2 data can be expected to yield five times more LAEs in total, i.e. 42 class 3 and 2 objects (or 88 if class 1 objects are also included).  However, further improvements in efficiency over that achieved here are possible.  Specifically, in the present by-eye search we are sure to have missed many objects since the human eye is not an optimal tool when confronted with vast amounts of data.  An automated search of the available dataset should be much more efficient.  Our first experiments with implementing such an automated search are very encouraging and show that we can recover the objects found by eye as well as find additional, missed line emitters.  We speculate that with an automated search we should be able to double the number of LAEs we find per unit area --- as is also suggested by a comparison of our number counts with those of narrowband imaging surveys (Fig.~\\ref{LF.fig}).  An automated search of DEEP2 would thus give us a sample of $\\sim$90 class 3 and 2 objects, all with spectroscopic redshifts and with high-resolution spectra suitable for the study of the details of gas kinematics.  Finally, the spectral coverage of DEEP2 gives us the potential to probe LAEs all the way up to $z$$\\sim$6.6 (Fig.~\\ref{volume_vs_z.fig}), thus allowing us to study the evolution of the population from near the epoch of reionization to $z$$\\sim$4.5 using a single, uniform dataset. Another clear advantage of an automated search is that it would allow us to accurately calibrate the detection efficiency through simulations and thus correct for it in our LF analysis.  Moreover, analyzing the full four-field DEEP2 dataset would also allow us to constrain the importance of field-to-field variance.  Clustering effects  appear important even in large fields such as those from Subaru/Suprimecam and constraining them using the four DEEP2 fields should be a worthwhile endeavor.  In summary, we have shown that the systematic examination of a large spectroscopic database can yield significant numbers of faint high-$z$ line emitters which are immediately suitable for further studies such as the determination of their number density (\\S~\\ref{numberdensity.sec}) or kinematics (\\S~\\ref{SpectralProperties.sec}). Automating such a search and applying it to an even larger spectroscopic database should be possible and should yield a large sample of high-quality high-$z$ LAEs suitable for a variety of studies."
    },
    "0807/0807.4151_arXiv.txt": {
        "abstract": "name{\\color{red}\\bf Abstract} \\begin{abstract} {\\footnotesize We propose a conceptually and computationally simple method to evaluate the neutrinos emitted by supernova remnants using the observed $\\gamma$-ray spectrum. The proposed method does not require any preliminary parametrization of the gamma ray flux; the gamma ray data can be used as an input. In this way, we are able to propagate easily the observational errors and to understand how well the neutrino flux and the signal in neutrino telescopes can be constrained by $\\gamma$-ray data. We discuss the various possible sources of theoretical and systematical uncertainties ({\\em e.g.}, neutrino oscillation parameters, hadronic modeling, {\\em etc.}), obtaining an estimate of the accuracy of our calculation. Furthermore, we apply our approach to the supernova remnant RX J1713.7-3946, showing that neutrino emission is very-well constrained by the H.E.S.S. $\\gamma$-ray data: indeed, the accuracy of our prediction is limited by theoretical uncertainties. Neutrinos from RX J1713.7-3946 can be detected with an exposure of the order ${\\rm km}^2\\times\\,{\\rm year}$, provided that the detection threshold in future neutrino telescopes will be equal to about 1 TeV.} ",
        "introduction": "Introduction}   Under-water and under-ice neutrino telescopes are instruments aiming to  discoveries. They could reveal an effective acceleration of cosmic rays (CR) in galactic sources (such a supernova remnants \\cita{ginz} and micro-quasars \\cita{carla}) and/or in extragalactic sources (such as AGN and gamma ray bursts \\cita{halzen}). However, at present it is possible to obtain reliable expectations only for few of these sources, such as the supernova remnants (SNR) discussed in this paper.  The idea that CR could originate in supernovae has been put forward already in 1934 but the first quantitative formulation of the conjecture that the young SNRs refurnish the Milky Way of cosmic rays, compensating the energy losses, is due to Ginzburg \\& Syrovatskii \\cita{ginz}. In fact, the turbulent gas of SNRs is a large reservoir of kinetic energy and this environment can support diffusive shock waves acceleration~\\cita{fermi}. The theory of CR acceleration in SNRs is still in evolution, but the generic expectations are stable. The CR flux in SNRs is expected to have a power law spectrum with spectral index $\\Gamma=2.0-2.4$ at low energies with a cutoff at an energy $E_{\\rm c}$ which depends on the details of the acceleration mechanism and on the age of the system. In specific implementations the cutoff energy $E_{\\rm c}$ can be as large as several PeV and, thus, is consistent with the ``knee\" in the CR spectrum at $E\\sim3\\times 10^{15}\\, {\\rm eV}$, which is believed to mark the transition from galactic to extra-galactic origin of CR~\\cita{berez}.   In recent times, great progress has been made in the observation of SNRs. In particular, the High Energy Stereoscopic System (H.E.S.S.) \\cita{hessWEB} has determined quite precisely the gamma ray spectra of few SNRs showing that they extend above 10 TeV. In the context of Ginzburg \\& Syrovatskii hypothesis, it is natural to postulate that the observed gammas are produced by the decay of $\\pi^0$ (and $\\eta$) resulting from the collision of accelerated hadrons with the ambient medium. New and crucial observations are being collected and the hadronic origin seems to be favored for certain SNRs, such as Vela Jr \\cita{vela} and RX J1713.7-3946 \\cita{rxj,rxjhadr}.\\footnote{ See also \\cita{waxman} for a recent analysis leading to a different conclusion.} It is not yet possible, however, to exclude that (part of) the observed radiation is produced by electromagnetic processes. The definitive proof that SNRs effectively accelerate CR could be obtained by the observation of high energy neutrinos in neutrino telescopes presently in operation, in construction or in project \\cita{nutel}.  As well known, there is a strict connection between photon and neutrino fluxes produced by hadronic processes in transparent sources (see, {\\em e.g.},~\\cita{gaisserbook}), which results from the fact that the same amount of energy is roughly given to $\\pi^0$, $\\pi^+$ and $\\pi^-$ in hadronic collisions. On this basis, one can estimate the neutrino fluxes expected from a sources with known $\\gamma$-ray spectra, trying to identify detectable sources and/or to optimize the detection strategies. The gamma-neutrino connection has been described in various recent papers \\cita{x1,Costantini:2004ap,x2,Vissani:2006tf,kappes,x4,x5,Vissani:2008zz} at a different level of accuracy, relying on different assumptions on the primary cosmic ray spectrum and/or hadronic interaction model.  In this particular moment, when high energy gamma ray astronomy is flourishing and the neutrino telescopes are finally becoming a reality, it is clearly important to have solid and transparent predictions. We continue, thus, the work started in \\cita{Vissani:2006tf,Vissani:2008zz,Villante:2007mh}, proposing a method to calculate neutrinos fluxes which is, at the same time, simple, accurate and model-independent. Our results are in essence a straightforward applications of standard techniques \\cita{lipari}, but we believe that that they will be useful since they improve the existing calculations in various respects. In particular, we provide simple analytic expressions for the neutrino fluxes which have a general validity and can be applied directly to gamma ray data since they do not require any parametrization of the photon spectrum. This allows us to propagate easily the observational errors in the gamma ray flux and, thus, to understand how well the neutrino flux and the signal in neutrino telescopes can be constrained by $\\gamma$-ray data. We also discuss the various possible sources of theoretical and systematical uncertainties, obtaining an estimate of the accuracy of our method. This kind of analysis is relevant in the present situation, since the number of events expected in neutrino telescopes from SNRs is quite low (see, {\\em e.g.}, \\cita{Vissani:2006tf}) and even small (downward) revisions of the expected signals may be important and/or require different detection strategies. It is thus important to understand the relevance of different assumptions in the calculations and the origin of (apparently) contrasting results appeared in the literature.  The plan of the paper is as follow. In Sect.~\\ref{method} we review the method presented in our previous paper~\\cita{Vissani:2006tf}. The relations obtained in this work -- Eqs.~(\\ref{allkernels}) -- are physically equivalent to those presented in \\cita{Vissani:2006tf}. However, they are more compact and more convenient for numerical computations since we have been able to recast them in such a way that they require only one numerical integration. In Sect.~\\ref{oscillations}, we discuss the effect of neutrinos oscillations (see \\cita{Strumia:2006db} for a review) and we discuss the relevance of uncertainty in neutrino mixing parameters for the predicted neutrino flux. In Sect.~\\ref{finalrelations} we present our main results, {\\em i.e.}, Eqs.~(\\ref{numuflux}) and (\\ref{antinumuflux}) which relate the (oscillated) muon neutrino and antineutrino fluxes to the $\\gamma$-ray flux. As it is explained, the only necessary information to predict neutrino fluxes are the relative production rates of the various mesons in hadronic processes, which are robust predictions of hadronic interaction models. In this paper, we adopt the results from Pythia \\cita{pythia}, estimated by using the parametrization of hadronic cross sections presented in \\cita{Koers:2006dd}. A comparison with SYBILL~\\cita{sybill} and DPMJET-III \\cita{dpmjet} is performed (when possible) by using the parametrization and tabulations of hadronic cross sections presented by \\cita{Kelner:2006tc} and \\cita{Huang:2006bp}. In Sect.~\\ref{fluxes} and \\ref{rates} we propose a procedure to predict neutrino fluxes and event rate in $\\nu$-telescopes directly from $\\gamma$-ray observational data, and we apply it to RX J1713.7-3946 which is presently the best studied SNR. In Sect.~\\ref{conclusions} we summarize our main results. ",
        "conclusions": ""
    },
    "0807/0807.3547_arXiv.txt": {
        "abstract": "Hyper-accreting discs occur in compact-object mergers and in collapsed cores of massive stars. They power the central engine of $\\gamma$-ray bursts in most scenarios. We calculate the microphysical dissipation (the viscosity and resistivity) of plasma in these discs, and discuss the implications for their global structure and evolution. At the temperatures ($k_{\\rm b}T > \\mec2$) and densities ($\\rho\\sim 10^{9}-10^{12}~ {\\rm gr~cm^{-3}}$) characteristic of the neutrino-cooled innermost regions, the viscosity is provided mainly by mildly degenerate electrons, while the resistivity is modified from the Spitzer value due to the effects of both relativity and degeneracy. Under these conditions the magnetic Reynolds number is very large ($R_{\\rm eM} \\sim 10^{19}$) and the plasma behaves as an almost ideal magneto-hydrodynamic (MHD) fluid. Among the possible non-ideal MHD effects the Hall term is relatively the most important, while the magnetic Prandtl number, $\\Pr$ (the ratio of viscosity to resistivity), is typically larger than unity: $10 \\lsim \\Pr  \\lsim 6 \\times 10^3$. Inspection of the outer radiatively inefficient regions indicates similar properties, with magnetic Prandtl numbers as high as $\\sim 10^{4}$. Numerical simulations of the magneto-rotational instability (MRI) indicate that the saturation level and angular momentum transport efficiency may be greatly enhanced at high Prandtl numbers. If this behaviour persists in the presence of a strong Hall effect we would expect that hyper-accreting discs should be strongly magnetised and highly variable. The expulsion of magnetic field that cannot be dissipated at small scales may also favour a magnetic outflow. We note that there are limited similaries between hyper-accreting discs and X-ray binary discs -- which also have a high magnetic Prandtl number close to the black hole -- which suggests that a comparison between late-time activity in $\\gamma$-ray bursts and X-ray binary accretion states may be fruitful. More generally, our results imply that the possibly different character of high Prandtl number MHD flows needs to be considered in studies and numerical simulations of hyper-accreting discs. ",
        "introduction": "The fluid dynamics of dense accretion discs is described by the Navier-Stokes and magneto-hydrodynamic equations.  Within such a description, the microscopic properties of the fluid are parameterized (in part) via the kinematic viscosity $\\nu$, which describes the rate at which inter-particle collisions damp fluid motions, and the resistivity $\\eta$, which quantifies the dissipation of currents via Ohmic losses. Timescales can be associated with each of these quantities. The viscous timescale of a disc $t_{\\rm acc} = r^2 / \\nu$, for example, is the timescale on which the microscopic viscosity would lead to angular momentum transport over a scale $r$. As is well known \\citep{pringle81}, for astrophysical discs the timescales for evolution driven by the microscopic viscosity are many orders of magnitude larger than the observed evolution timescales. It is therefore generally assumed that evolution is driven by an ``anomalous\" or ``turbulent\" viscosity $\\nu_{\\rm t} \\gg \\nu$ \\citep{shakura73}, possibly augmented in some cases by angular momentum loss in a wind \\citep{blandford82}. In most astrophysical discs $\\nu_t$ is likely generated from turbulence driven by the magneto-rotational instability (MRI; see \\citet{balbus98} and references therein). Since $\\nu_{\\rm t} \\gg \\nu$ most studies of discs make the implicit assumption that the actual values of $\\nu$ and $\\eta$ are unimportant, and consider instead evolution under the action of the effective transport coefficients $\\nu_t$ and $\\eta_t$.  Provided that the Reynolds ($R_{\\rm e}$) and magnetic Reynolds ($R_{\\rm eM}$) numbers are large enough, the absolute value of either one may indeed be of little import for accretion discs.  Their ratio, however, which we define as the magnetic Prandtl number, \\begin{equation} {\\rm Pm} \\equiv \\frac{\\nu}{\\eta} \\end{equation} is almost certainly {\\em not} ignorable \\citep{balbus98}. An accretion disc serves to convert the large-scale energy present in the shear flow into other energy forms -- turbulent kinetic energy, magnetic field, and heat -- and the dissipation required to generate heat may vary strongly depending upon the microphysical transport coefficients. In particular, if ${\\rm Pm} \\gg 1$ the plasma is highly viscous on the scale at which resistivity operates. In this regime one may expect that dissipation of the magnetic field will be suppressed, with a resulting modification in the saturation level of dynamo generated fields and an inverse cascade of magnetic energy to large scales \\citep{brandenburg01}. The importance of the magnetic Prandtl number for the resulting magnetic field structure has been illustarted analytically \\citep{umurhan07,umurhan07b} and numerically, both in idealised simulations of dynamo action within forced turbulence \\citep{schekochihin04} and in local shearing-box simulations of the MRI within accretion discs \\citep{lesur07,fromang07}\\footnote{Since dissipation in discs typically occurs on small scales for which the fluid is unaware of the large scale shear, one would anticipate that the basic physics of large ${\\rm Pm}$ dynamos is similar. Disc simulations are required, however, to assess the impact of the small-scale physics on MRI saturation and the large scale flow properties. It is unclear whether the shearing-box approach is adequate for study of these effects.}.  For discs the current limited set of simulations indicates that the efficiency of angular momentum transport, parameterized via the Shakura-Sunyaev $\\alpha$ parameter, increases with ${\\rm Pm}$, and can approach $\\alpha \\sim 1$ given the combination of modest ${\\rm Pm} = 8$ and a weak net vertical magnetic field. Motivated by these results, \\citet{balbus08} calculated the expected radial variation of ${\\rm Pm}$ within standard Shakura-Sunyaev models of geometrically thin, radiatively efficient accretion discs. They found that ${\\rm Pm}$ exceeds unity within $\\sim 50$ Schwarzschild radii of the central object, and suggested that the qualitatively different behaviour of the MRI in the high ${\\rm Pm}$ limit might be responsible for the time-dependent outbursts and state changes that are observed in X-ray binaries.  In this paper, we compute the magnitude and radial dependence of the magnetic Prandtl number in  hyper-accreting discs. Such discs, which can support accretion rates of the order of $1 \\ M_\\odot \\ {\\rm s}^{-1}$, form from compact object mergers \\citep[e.g.][]{ruffert99} and when the cores of rapidly rotating massive stars collapse \\citep{woosley93}. The disc structure is shaped by neutrino emission in the regions close to the accreting object, while in the outer parts neither neutrinos (the temperature is too low) nor photons (which are trapped within the flow) can provide cooling. Hyper-accreting discs are invoked to provide the initial energy injection in many models for $\\gamma$-ray bursts (GRBs). Our ultimate motivation for studying the magnetic Prandtl number in these discs is to understand the complex phenomenology (jets, highly variable prompt emission,  X-ray flares, plateau phases etc) that is observed in GRBs \\citep[e.g.][]{burrows05} and which may well derive from the evolution of the disc.  We anticipate different results  from those found by \\citet{balbus08} for photon-cooled discs. In the neutrino-cooled region the differences arise from two reasons. First, since neutrino opacities are vastly smaller than photon opacities, hyper-accreting discs are much cooler and denser than an extrapolation of Shakura-Sunyaev discs would imply, and this on its own would alter the magnetic Prandtl number.  Second, the temperatures and densities in hyper-accreting discs are such that the electrons are mildly relativistic and mildly degenerate, while the elastic collision cross-section for the nuclei exceeds the Coulomb cross-section. As a consequence, the classical \\citet{spitzer62} expressions for $\\nu$ and $\\eta$, which are adequate for ordinary accretion flows, no longer apply. They do apply, instead, in the outer regions of hyper-accreting discs, where a different $\\Pr$ behaviour occurs since the flow is radiatively inefficient and photon-pressure dominated.  The organization of this paper is as follows. In \\S2 through \\S5 we describe the formalism for calculating the structure of hyper-accreting discs. This is required here in order to determine consistently the plasma conditions within these discs -- including the temperature, density, nuclear composition and degree of degeneracy. Readers familiar with vertically averaged hyper-accreting disc models will find that our approach generally follows standard practice. The structure of the resulting disc models is presented in \\S6. Given these conditions, we then calculate in \\S7 and \\S8 the resisitivity and viscosity in the neutrino cooled inner region of the disc. These Sections contain the principle new results of this paper. Combining the viscosity and resistivity, we show in \\S10 that ${\\rm Pm} > 1 ~{\\rm or} \\gg 1$ across the entire radial range relevant for GRB models. In \\S11 we estimate the strength of other non-ideal MHD effects, and in \\S12 we study the structure of the outer non-radiative region of the disc. \\S13 and \\S14 discuss and summarize our results, and what they may imply for disc evolution and GRB observables. ",
        "conclusions": "In this paper, we calculated the magnetic Prandtl number for plasma conditions in hyper-accreting discs. This dimensionless number $\\Pr = \\nu/\\eta = \\lambda_{\\nu}^{2}/\\lambda_{\\eta}^{2}$ expresses the relative size of two critical turbulent flow scales: the scale on which velocity fluctuations are viscously damped $\\lambda_{\\nu}$, and the scale $\\lambda_{\\eta}$ on which ohmic losses result in magnetic field dissipation. We find that in the inner neutrino-cooled regime: \\begin{itemize} \\item[(1)] Electric resistivity involves relativistic, mildly degenerate electrons. As a consequence, the resistivity is very low and weakly dependent on density and temperature. \\item[(2)] The main source of viscosity is Coulomb collisions between mildly {\\it degenerate electrons} and  protons. It decreases as the flow gets denser or cooler, since the electron fraction decreases. \\item[(3)] The magnetic Prandtl number is always {\\em greater} than unity, unless the angular momentum transport efficiency, parameterized via the Shakura-Sunyaev $\\alpha$, is very small ($\\alpha < 0.01$). $\\Pr$ ranges between a few tens to a few $10^3$, and it typically increases with $\\alpha$. In the optically {\\em thin} regions for neutrinos, the main dependence of $\\Pr$ is on density (inversely, via viscosity) since the disc temperature varies very slowly as a function of $\\dot{M}$ and radius. This causes $\\Pr$ to {\\em decrease} with increasing accretion rates. In the optically {\\em thick} regions, the main dependence is on temperature, that varies more rapidly than density. As a consequence, $\\Pr$ {\\em increases} with $\\dot{M}$ \\end{itemize}  \\no In the outer, radiatively inefficient region:  \\begin{itemize} \\item[(4)] Resistivity and viscosity have the classical ``Spitzer'' dependencies on temperature and density, since electrons are non-relativistic and non-degenerate and viscosity is mainly given by Coulomb interactions between ionised helium particles. The magnetic Prandtl number is thus a much stronger function of temperature $\\Pr \\propto T^{4}/\\rho$. In contrast, it does not depend on the overall magnitude of accretion (via $\\dot{M}$ and $\\alpha$). The Prandtl number is always much larger than unity ($\\Pr \\simeq 10^4 (100~\\rs)/r$) and it is unlikely to fall below unity for the entire extent of the disc. \\end{itemize}  \\no We conclude by noting a number of implications for the global structure of hyper-accreting discs:  \\begin{itemize} \\item[(5)] For all plausible values of the magnetic field strength in the disc, the Hall effect provides the largest non-ideal term in the MHD equations. These discs are perhaps a unique example of a high magnetic Reynolds number flow in which the Hall effect is the dominant non-ideal term. Since there are no simulations where these conditions are accounted for, their effect on the magnetic field evolution and MRI is unclear. \\item[(6)] The large values of the magnetic Prandtl number mean that the evolution of magnetic field within hyper-accreting discs is likely to be qualitatively different from discs with lower magnetic Prandtl numbers. The results of forced dynamo and MRI simulations suggest that, in the high $\\Pr$ regime, small scale field dissipation is suppressed and the saturation level of the magnetic field is enhanced, perhaps dramatically \\citep{lesur07,fromang07}. Numerical simulations of hyper-accreting discs that do not account for microphysical dissipation may therefore have underestimated the magnetic field strength. \\item[(7)] The {\\em immediate} consequence of a high value of $\\Pr$ for models of the central engines of GRBs is that disc formation is likely to be accompanied by vigorous expulsion of magnetic flux. This favours models in which energy is liberated in the form of a magnetic jet or outflow. \\end{itemize}"
    },
    "0807/0807.3292_arXiv.txt": {
        "abstract": "Generic inspirals and mergers of binary black holes produce beamed emission of gravitational radiation that can lead to a gravitational recoil or {\\em kick} of the final black hole.  The kick velocity depends on the mass ratio and spins of the binary as well as on the dynamics of the binary configuration. Studies have focused so far on the most astrophysically relevant configuration of quasi-circular inspirals, for which kicks as large as $\\sim 3,300\\, \\KMS$ have been found. We present the first study of gravitational recoil in {\\em hyperbolic} encounters. Contrary to quasi-circular configurations, in which the beamed radiation tends to average during the inspiral, radiation from hyperbolic encounters is plunge dominated, resulting in an enhancement of preferential beaming. As a consequence, it is possible to achieve kick velocities as large as \\kickmax. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0807/0807.3771_arXiv.txt": {
        "abstract": "We show that the proper motion of the Becklin-Neugebauer (BN) object is consistent with its dynamical ejection from the $\\theta^1$~Ori~C binary, contrary to recent claims by G\\'omez et al. Continued radio observations of BN and future precise astrometric observations of $\\theta^1$~Ori~C with SIM and the Orion Nebula Cluster with GAIA can constrain the properties of this ejection event, with implications for theories of how the nearest example of massive star formation is proceeding. ",
        "introduction": "\\label{S:intro}  Understanding massive star formation remains one of the most challenging and important problems of contemporary astrophysics (Beuther et al. 2007; Zinnecker \\& Yorke 2007). The complexity of the process means that massive star formation theories, such as the turbulent core model (McKee \\& Tan 2003), the competitive accretion model (Bonnell \\& Bate 2006) and stellar coalescence model (Bonnell et al. 1998; Clarke \\& Bonnell 2008) require close testing against observed systems. The closest forming (i.e. accreting) massive star is thought to be radio source I (Menten \\& Reid 1995) within the Orion Nebula Cluster (ONC), at a distance of $414\\pm7$~pc (Menten et al. 2007, adopted throughout), in the Kleinmann-Low (KL) region. As reviewed by Tan (2008), this source has been used as observational evidence in support of all three of the above theories. Part of this confusion is due to the Becklin-Neugebauer (BN) object, 9.9\\arcsec to the NW (Fig.~1), which is a fast moving (radio-ONC-frame proper motion of $\\mu_{\\rm BN}=13.2\\pm 1.1\\:{\\rm mas\\:yr^{-1}}$, i.e. $v_{\\rm 2D,BN}=25.9\\pm2.2\\:{\\rm km\\:s^{-1}}$ towards P.A.$_{\\rm BN}=-27^\\circ.5\\pm4^\\circ$, Plambeck et al. 1995; G\\'omez et al. 2008) embedded B star ($L_{\\rm BN}=(2.1 - 8.5)\\times 10^3L_\\odot$, Gezari, Backman \\& Werner 1998, equivalent to a zero age main sequence mass $m_{\\rm BN,zams} = 9.3\\pm2.0\\sm$). This proper motion implies that BN has been moving through the KL region and made a close, possibly coincident, passage with source {\\it I} about 500 years ago. Thus to understand the nearest example of massive star formation, we need to understand the origin of BN's motion.  Including the $(+21) - (+8) = +13\\:{\\rm km\\:s^{-1}}$ radial velocity of BN with respect to the ONC mean (Scoville et al. 1983; Walker 1983), the 3D ONC-frame velocity of BN is $v_{\\rm 3D,BN}=29\\pm3\\:{\\rm km\\:s^{-1}}$, and its kinetic energy is $E_{\\rm BN} = (8.3\\pm2.3)\\times 10^{46} (m_{\\rm BN}/10\\sm)\\:{\\rm ergs}$. BN is very likely to have formed somewhere in the ONC and then attained its high speed by a close interaction with a massive multiple stellar system followed by dynamical ejection (Poveda, Ruiz \\& Allen 1967).  Tan (2004) proposed BN was launched from the $\\theta^1$~Ori~C binary (also shown in Fig.~\\ref{fig:bn}), since this is the only stellar system in the ONC known to have all the physical properties required by this scenario: (1) a location along BN's past trajectory (\\S\\ref{S:ast}); (2) an (optical)-ONC-frame proper motion ($\\mu_{\\theta^1C}=2.3\\pm0.2\\:{\\rm mas\\:yr^{-1}}$, van Altena et al. 1988, i.e. $v_{\\rm 2D,\\theta^1C} = 4.5\\pm0.4\\:{\\rm km\\:s^{-1}}$, towards $\\rm P.A._{\\theta^1C}=142^\\circ.4\\pm4^\\circ$) that is in the opposite direction to BN (the direction to BN from $\\theta^1$~Ori~C is a P.A.$=-30^\\circ.949$) and is of the appropriate magnitude (the dynamical mass of BN implied by this motion agrees with the estimate of $m_{\\rm BN,zams}$ and is $m_{\\rm BN,dyn}=8.6\\pm1.0\\sm$ assuming negligible error in $m_{\\theta^1C}=49.5\\sm$ and negligible motion of the pre-ejection triple system in this direction; a pre-ejection motion of 0.35~mas/yr along this axis (\\S\\ref{S:high}) would contribute an additional $1.5\\sm$ uncertainty); (3) primary ($m_{\\theta^1C-1}=34\\sm$) and secondary ($m_{\\theta^1C-2}=15.5\\sm$) masses greater than $m_{\\rm BN}$ (Kraus et al. 2007); (4) a semi-major axis of $a=17.0\\pm5.8$~AU (Patience et al. 2008) and thus a total orbital energy ($E_{\\rm tot}=Gm_{\\theta^1C-1}m_{\\theta^1C-2}/(2a)= (2.7\\pm0.9)\\times 10^{47}\\:{\\rm ergs}$) greater than the sum of BN's kinetic energy and $\\theta^1$~Ori~C's kinetic energy ($1.00\\times 10^{46}\\:{\\rm ergs}$) (see Tan 2008 for a review). Note, $\\theta^1$~Ori~C's recoil in this scenario is large enough to remove it from the Trapezium region (see Pflamm-Altenburg \\& Kroupa 2006 for theoretical studies of the dynamical decay of Trapezium-like systems) and may be enough to eject it from the ONC completely, with implications for the effectiveness of its ionizing feedback on disrupting the star cluster formation process.  Rodr\\'iguez et al. (2005) and Bally \\& Zinnecker (2005) proposed BN was launched from an interaction with radio source {\\it I}, which would require this system to be a massive binary, recoiling away from any large scale ($\\gtrsim 100$~AU) gas that it was originally accreting. G\\'omez et al. (2008) used the relative motion to BN with respect to source {\\it I} to claim that BN could not have made a close passage with $\\theta^1$~Ori~C, excluding this possibility at the 5-10~$\\sigma$ level.  We show in \\S\\ref{S:ast} that if BN's motion is considered in the reference frame of the ONC, then a close (coincident) passage with $\\theta^1$~Ori~C is allowed by the data, which permits the scenario of dynamical ejection of BN from $\\theta^1$~Ori~C. In \\S\\ref{S:high} we discuss the potential of future high precision astrometric measurements to constrain the properties of BN's dynamical ejection, which then constrain BN's interaction distance with source {\\it I}, the mass of source {\\it I}, and thus the strength of tidal perturbations on the massive protostar during this encounter.  \\vspace{0.2in} ",
        "conclusions": "We have reviewed the latest evidence that BN was dynamically ejected from the $\\theta^1$~Ori~C binary, finding that $\\theta^1$~Ori~C has all the physical properties expected in this scenario. We showed that the trajectory of BN is also consistent with this scenario, in contrast to recent claims by G\\'omez et al. (2008). We discussed how high precision astrometry of $\\theta^1$~Ori~C with SIM can yield information on the pre-ejection velocity of the system and the size of any subsequent deflections, in particular that of BN caused by close passage with source {\\it I}, the nearest massive protostar."
    },
    "0807/0807.0291_arXiv.txt": {
        "abstract": "The amdlib AMBER data reduction software is meant to produce AMBER data products from the raw data files that are sent to the PIs of different proposals or that can be found in the ESO data archive. The way defined by ESO to calibrate the data is to calibrate one science data file with a calibration one, observed as close in time as possible. Therefore, this scheme does not take into account instrumental drifts, atmospheric variations or visibility-loss corrections, in the current AMBER data processing software, amdlib.  In this article, we present our approach to complement this default calibration scheme, to perform the final steps of data reduction, and to produce fully calibrated AMBER data products. These additional steps include: an overnight view of the data structure and data quality, the production of night transfer functions from the calibration stars observed during the night, the correction of additional effects not taken into account in the standard AMBER data reduction software such as the so-called \"jitter\" effect and the visibility spectral coherence loss, and finally, the production of fully calibrated data products.  All these new features are beeing implemented in the modular pipeline script \\texttt{amdlibPipeline}, written to complement the amdlib software. ",
        "introduction": "\\label{sect:intro}%  The AMBER data reduction software: \\texttt{amdlib}\\cite{2007A&A...464...29T}, features a new type of algorithm, the so-called P2VM algorithm\\cite{2004SPIE.5491.1222M}, which performs a direct fit of the observed fringes using a template fringe pattern recorded during the instrument calibration and produces a series of observables; namely, the squared visibility, the closure phase, and the differential phase\\cite{2006-ITHD_FMillour}. The standard calibration plan handles the AMBER signal from the raw data to raw averaged visibilities OI DATA files. Though a simple calibration script (\\texttt{amdlibDivideOiData}) is included in the \\texttt{amdlib} software, a complete calibration procedure is not yet provided.  In this document, we present our approach to produce calibrated OI data files. It consists of several recipes, just like for the standard \\texttt{amdlib} software, to organize, evaluate, plot, and finally, calibrate the data. This philosophy allows one to switch easily from one recipe to another if necessary. This leads to the modular structure of the example pipeline script \\texttt{amdlibPipeline}\\footnote{currently available at \\url{http://www.mpifr-bonn.mpg.de/staff/fmillour/}}, written in the scientific language \\texttt{yorick}, to perform the final steps of data reduction after an \\texttt{amdlib} data reduction session.  This paper is organized into three different sections: \\begin{itemize} \\item A section about file management and coping with different types of calibrations (different DITs, different modes, use of BCD), \\item The rise of advanced calibration methods: ``coherence length'' and ``jitter'' effects correction, \\item How to cope with calibration stars and transfer function, as well as a template for data calibration. \\end{itemize} ",
        "conclusions": "The standard default calibration scheme proposed for AMBER (calibrating one file with another) is obviously insufficient to get realistic error bars, and an overnight calibration quality estimate.  We presented a proposition for a user-friendly pipeline framework for the amdlib software. We are able, with the proposed scripts and an Internet connection, to produce science-grade data with realistic error bars in a semi-automated way, as it is already done for the first data reduction steps in \\texttt{amdlib}.  This article is thought to be a starting point for further reflexions about AMBER data calibration. Further work could allow one to perform the AMBER data reduction in a fully automated way, allowing in the future the production of archived data products, in the same way as it is done with the ESO imaging and spectroscopic instruments (WFI, FEROS, etc.)."
    },
    "0807/0807.0634_arXiv.txt": {
        "abstract": "The recent advent of laser guide star adaptive optics (LGS AO) on the largest ground-based telescopes has enabled a wide range of high angular resolution science, previously infeasible from ground-based and/or space-based observatories.  As a result, scientific productivity with LGS has seen enormous growth in the last few years, with a factor of $\\approx$10 leap in publication rate compared to the first decade of operation.  Of the 54~refereed science papers to date from LGS AO, half have been published in the last $\\approx$2 years, and these LGS results have already made a significant impact in a number of areas.  At the same time, science with LGS AO can be considered in its infancy, as astronomers and instrumentalists are only beginning to understand its efficacy for measurements such as photometry, astrometry, companion detection, and quantitative morphology.  We examine the science impact of LGS AO in the last few years of operations, largely due to the new system on the Keck~II 10-meter telescope.  We review currently achieved data quality, including results from our own ongoing brown dwarf survey with Keck LGS.  We assess current and near-future performance with a critical eye to LGS AO's capabilities and deficiencies.  From both qualitative and quantitative considerations, it is clear that the era of regular and important science from LGS AO has arrived. ",
        "introduction": "\\label{sec:intro}  %  Astronomers have envisioned using laser guide star adaptive optics (LGS AO) to achieve diffraction-limited observations from ground-based telescopes for over two decades\\cite{1985A&A...152L..29F, 1987Natur.328..229T, 1994OSAJ...11..263H}.  (See Refs.~\\citenum{2004aoa..book.....R} and~\\citenum{1998aoat.conf.....H} for a historical review.)  The realization of this vision has required the dedication, talents, and resources of myriad individuals and organizations. LGS AO is now being used regularly to conduct a broad range of science.  The promise of near diffraction-limited imaging and spectroscopy from the ground over most of the sky is coming to fruition.  This heralds an important new capability for observational astronomy.  The purpose of this paper is to examine the science that has been done using LGS AO, with a critical eye to what has been achieved and what promises and challenges remain.  Much has been accomplished since the last review of science from LGS (Ref.~\\citenum{2006SPIE.6272E..14L}), which focused on publication output.  Here, we go beyond mere productivity and examine the science {\\em impact} of LGS as demonstrated in the record of published research. ",
        "conclusions": "\\subsection{A High Angular Resolution Survey of Field Brown Dwarfs}  To illustrate what is currently possibly with LGS, we examine the on-sky performance of the Keck LGS system.  Since the inception of the Keck system for open science use in 2005, my collaborators and I have been conducting a high angular resolution near-IR study of nearby brown dwarfs.  Our goals are (1) to assess the binary frequency of ultracool dwarfs; (2) to test atmospheric models with these coeval systems, \\eg, as associated with the abrupt spectral transition from the L~dwarfs to the T~dwarfs;\\cite{2006astro.ph..5037L} (3) to search for exceptionally low-temperature companions;\\cite{2008A&A...482..961D} and (4) to find and monitor substellar binaries suitable for dynamical mass determinations.\\cite{liu08-2m1534orbit}  LGS AO represents a major advance for this science area.  Most known brown dwarf binaries have separations of $\\lesssim$0.3\\arcsec,\\cite{2006astro.ph..2122B} hence the need for high angular resolution imaging to find and characterize them via resolved photometry and spectroscopy, but they are too optically faint for natural guide star AO.  Our observations have achieved 3--4$\\times$ the angular resolution at $K$-band (2.2~\\micron) compared to \\HST\\ and thus are more sensitive to close companions.  In addition, the ability of Keck LGS AO to find tighter binaries means that systems with much shorter orbital periods than the current sample can be found and expeditiously monitored.  Finally, many of the key spectral diagnostics of brown dwarfs are in the infrared and thus probing the atmospheric properties of these objects with resolved colors and spectra of binaries is well-suited to current LGS capabilities.   \\begin{figure}[h] \\vskip -0.5in \\begin{center} \\begin{tabular}{c} \\hskip -0.7in \\includegraphics[width=3.5in,angle=90]{plot-obs-summary-fwhm-k} \\hskip -1.6in \\includegraphics[width=3.5in,angle=90]{plot-obs-summary-fwhm-ferror-k.ps}\\\\ \\hskip -0.6in \\includegraphics[width=3.5in,angle=90]{plot-obs-summary-strehl-k} \\hskip -1.6in \\includegraphics[width=3.5in,angle=90]{plot-obs-summary-strehl-ferror-k.ps}\\\\ \\end{tabular} \\end{center} \\vskip -4ex \\caption[fig:kecklgs] { \\label{fig:kecklgs} Summary of Keck LGS AO $K$-band (2.2~\\micron) performance, based on our near-IR imaging survey of brown dwarfs (Liu \\etal, in prep).  No bad data have been censored, so the compilation comprises a mix of seeing conditions, target airmasses, and technical performance.  {\\bf Top panels:} Histogram of $K$-band FWHM and its fractional RMS variation within a given dataset.  Typically each dataset is composed of 6--12 individual images taken over an elapsed time of 10--20~minutes.  The median FWHM is 0.076\\arcsec\\ with a median RMS variation of 5\\%.  {\\bf Bottom panels:} Same histograms for $K$-band Strehl.  The median Strehl is 0.17 with an RMS variation of a factor of 1.19 (\\eg, Strehl = $0.17\\pm0.03$).} \\end{figure}   Our brown dwarf imaging survey provides an excellent dataset for assessing typical Keck LGS performance in the case of off-axis observations, namely the situation where the LGS is pointed to the science target but tiptilt sensing and correction are derived from an adjacent field star.  Brown dwarfs are far too optically faint to serve as their own tiptilt references and hence the need for a nearby tiptilt star -- this is the same observing situation as for many extragalactic LGS applications and thus provides a good reference point.  For Keck, the tiptilt star must be within 60\\arcsec\\ of the science target -- in practice, we find that this results in a sky coverage fraction of about 2/3 for an estimated $K$-band Strehl ratio of $\\gtrsim$0.2.  Since targets for our brown dwarf program span most of the sky (except for avoidance of the galactic plane), this 2/3 sky coverage estimate is a fair representation of the fraction of any set of generic targets that can be imaged with LGS.  Figure~\\ref{fig:kecklgs} summarizes the image quality of a subset of our Keck LGS observations of nearby brown dwarfs, spanning multiple observing runs from 2005--2007.  No bad data have been censored, so a mix of seeing conditions, off-axis tiptilt star properties, and technical performance (e.g., LGS projected power and sodium light return flux) are represented.  (See also Ref.\\ \\citenum{2006SPIE.6272E..14L} for more performance descriptions.)  The median $K$-band image FWHM for our survey is 0.076\\arcsec\\ with a best value of 0.051\\arcsec.  The median $K$-band Strehl is 0.17 with a best value of 0.34.  LGS images are naturally time-variable, and the shape and detailed structure of the PSF changes in every image.  Figure~\\ref{fig:kecklgs} provides one representation of this PSF variability, plotting histograms of the fractional RMS deviations in $K$-band FWHM and Strehl ratios.  Typically, each of our brown dwarf datasets constitutes a series of 6--12~images, each with a integration of about 1~min and total elapsed time of about 15--30~min.  Over these time scales, the plotted histograms show significant FWHM and Strehl variations.  This is obviously a challenge for science programs requiring a stable PSF.   \\begin{figure}[t] \\vskip -0.3in \\begin{center} \\begin{tabular}{c} \\hskip -0.75in \\includegraphics[width=3.75in,angle=90]{make-figure-orbit-skyplot.labels.fancy_color.1orbit.ps} \\hskip -1.8in \\includegraphics[width=3.75in,angle=90]{make-figure-orbit-sep.ps} \\end{tabular} \\end{center} \\vskip -6ex \\caption[fig:2m1534-orbit] { \\label{fig:2m1534-orbit} First dynamical mass determination for a binary T~dwarf, the T5.0+T5.5 binary 2MASS~J1534-2952AB (Ref.\\ \\citenum{liu08-2m1534orbit}).  Six epochs of data from Keck LGS are combined with \\HST\\ archival data to form a dataset from 2000-2008 that spans about 50\\% of the 15-year orbital period. Through careful analysis of the binary images with PSF fitting, the Keck LGS data achieve sub-milliarcsecond relative astrometry of the two components.  The total mass is $59\\pm3~\\Mjup$.  This is the first field binary for which both components are directly confirmed to be below the stellar/substellar limit ($\\lesssim$75~\\Mjup).  This is also the coolest and lowest mass binary with a dynamical mass determination to date. The left plot shows the relative astrometry and best-fitting orbit.  The right plot shows the separation measurements compared to the best fitting orbit (solid line) and two alternative orbital periods of similar total mass (dashed and dotted lines).  The right bottom sub-panel shows the residuals between the observations and the predictions from the orbits.} \\end{figure}   Despite the notable PSF fluctuations, high precision science can be with Keck LGS data.  For example, we have been using LGS to carefully monitor the orbits of binary brown dwarfs in order to determine their dynamical masses and thereby test theoretical models.  Despite the hundreds of brown dwarfs that have been identified and had their spectrophotometric properties characterized, direct measurements of their most fundamental property, namely their mass, are sorely lacking.  Typical estimated orbit periods are a decade a more, but many of the binaries were discovered $\\gtrsim$5~years ago with \\HST\\ and thus multi-epoch LGS followup is well-suited to dynamical mass determinations.  Figure~\\ref{fig:2m1534-orbit} shows an example of what is possible, presenting the orbit of the T5+T5.5 binary 2MASS~J1534-2952AB by Liu, Dupuy and Ireland (Ref.\\ \\citenum{liu08-2m1534orbit}). This is the first dynamical mass determination for a binary T~dwarf, the coolest and least luminous class of brown dwarfs.  There are two stars within 6\\arcsec\\ of \\twomassbin\\ which we use as PSFs to deblend the light of the two binary components.  Through extensive Monte Carlo testing of different deblending methods applied to simulated images of binary stars constructed from the single PSFs, we are able to measure the relative position and orientation of the components to 0.3--0.7~mas and 0.2--0.7\\degs\\ RMS at $K$-band.  The precision of the best measurements is such that atypical sources of uncertainty need to be considered, including: \\begin{enumerate} \\item The calibration of the instrumental plate scale and orientation can introduce additional uncertainty, since these values for the Keck facility near-IR camera are known to about 1~part in $10^{-3}$ and 0.1\\degs, respectively. \\item The differing spectral types of the two components means that the light from each is subjected to a slightly different amount of differential chromatic refraction (DCR).  Given the sky orientation of the binary (just about North-South), the declination of the target ($\\delta=-29\\degs$ meaning the smallest possible airmass for Keck observing is 1.55), and the fact that we observe it near transit (for best AO performance), the DCR causes the separation of the binary to appear slightly smaller at $H$ and $K$-bands and slightly larger at $J$-band compared to the true position as would be observed at zenith.  The amplitude of the DCR effect is about 0.3~mas, much smaller than the measurement errors at most (but not all) of the Keck epochs.  However, the effect is a systematic one so we correct the relative astrometry of the two components based on synthetic photometry from T~dwarf spectra.  (Since T~dwarfs have suppressed flux at the longward portion of the $K$-band, the DCR is smaller than it would be for L~dwarfs observed at the same airmass.) \\end{enumerate} \\noindent The quality of the Keck LGS astrometry equals or exceeds that available for the same target from \\HST, though at the very smallest separations ($\\approx$FWHM) \\HST\\ data still have an advantage since the very stable \\HST\\ PSF can be simulated to high fidelity\\cite{1993ASPC...52..536K} and then fitted to tight binary images (\\eg, Ref.\\ \\citenum{liu08-2m1534orbit}).  However, LGS AO provides the necessary long-term platform for synoptic monitoring of visual binaries, especially since the required amount of observing time at each epoch is relatively modest but many epochs are needed. This in contrast to \\HST\\ where target acquisition is slow and monitoring a populous sample over many epochs is telescope time-intensive.     \\subsection{Review of Published Keck Performance}  To further summarize the current capability, {\\bf Table~3} provides a summary of achieved Keck LGS performance as presented in the published science papers.  Obviously this is a heterogeneous assemblage from both the standpoint of the observations and the subsequent analysis, but it does provide ``ground truth'' of the quality of on-sky data that has proven suitable for publication.  We consider several types of measurements that readily lend themselves to quantitative uncertainties (as opposed to more difficult measurements to quantify such as, \\eg, morphological studies of resolved objects): \\begin{itemize} \\item {\\em Relative photometry:} the flux of a science target relative to other sources in the same image of known brightness (\\eg, from 2MASS); \\item {\\em Absolute photometry:} the absolute flux of a science target as directly measured from the images of the target and then compared to a photometric calibrator star observed contemporaneously (but not simultaneously); \\item {\\em Relative astrometry:} the location of sources in an LGS dataset relative to each other. \\item {\\em Absolute astrometry:} the location of sources in a LGS dataset as referenced to absolute astrometry tied to an external dataset (\\eg, \\HST). \\item {\\em Binary (relative) photometry:} the relative flux ratio of a binary; \\item {\\em Binary (relative) astrometry:} the separation and position angle of a binary; \\item {\\em Crowded field astrometry:} the positions and/or proper motions for a field with many sources (\\eg, star clusters and resolved nearby galaxies); \\item {\\em Crowded field (relative) photometry:} the same as relative photometry above, except for a field with many sources; \\item {\\em Crowded field (absolute) photometry:} the same as absolute photometry above, except for a field with many sources. \\end{itemize}  \\noindent {\\bf Table~3} shows that high quality measurements have been achieved with LGS, especially for relative measurements.  In the published papers, the most measurements are available for relative photometry and astrometry of low-mass binaries, since this is a class of science that readily benefits from LGS."
    },
    "0807/0807.0128_arXiv.txt": {
        "abstract": "We present detailed analysis of three globular cluster X-ray sources in the XMM-Newton extended survey of M31. The X-ray counterpart to the M31 globular cluster Bo\\thinspace 45 (XBo\\thinspace 45) was observed with XMM-Newton on 2006 December 26. Its combined pn+MOS 0.3--10 keV lightcurve was seen to vary by $\\sim$10\\%, and its 0.3--7.0 keV emission spectrum was well described by an absorbed power law with photon index 1.44$\\pm$0.12. Its variability and emission is characteristic of low mass X-ray binaries (LMXBs) in the low-hard state, whether the accretor is a neutron star or black hole. Such behaviour is typically observed at luminosities $\\la$10\\% Eddington. However, XBo\\thinspace 45 exhibited this behaviour at an unabsorbed, 0.3--10 keV luminosity of { 2.5$\\pm$0.2$\\times 10^{38}$} erg s$^{-1}$, or{ $\\sim$140\\%} Eddington for a 1.4 $M_{\\odot}$ neutron star accreting hydrogen. Hence, we identify  XBo\\thinspace 45 as a  new candidate black hole LMXB. { XBo\\thinspace 45 appears to have been consistently bright for $\\sim$30 years, consistent with theoretical prediction for a globular cluster black hole binary formed via tidal capture}. Bo\\thinspace 375 was observed in the 2007, January 2 XMM-Newton observation, and has a two-component spectrum that is typical for a bright neutron star LMXB. Bo\\thinspace 135 was observed in the same field as Bo\\thinspace 45, and could contain either a black hole or neutron star. ",
        "introduction": "Of the 14 bright low mass X-ray binaries (LMXBs) associated with Galactic globular clusters (GCs), 13 are known to contain neutron stars, and the other one probably does as well \\citep[see e.g.][ and references within]{wa01,int04}. However, black hole binaries have been seen in extragalactic GCs, proving that they form in such environments.   For example, \\citet{mac07} discovered the first  GC black hole binary in the giant elliptical galaxy NGC 4472.  { The lack of known GC black hole binaries has been a subject of great interest for theoretical modellers. \\citet{kal04} show that black hole binaries formed through exchange interactions should have duty cycles of $\\sim$0.001, consistent with the absence of GC black hole binaries at the  time. They also predict that black hole binaries formed  through tidal capture of a main sequence star would be bright, persistent X-ray sources, and infer  from their absence that tidal capture probably disrupts the main sequence star. However, neither of these results preclude the discovery of new GC black hole binaries.  }  The bright X-ray sources of M31 have been studied for over 25 years, with Einstein \\citep{tf91}, ROSAT \\citep{S97,S01}, Chandra \\citep[see e.g.][]{dis02, kaa02,will04} and XMM-Newton \\citep[see e.g.][]{shi01,trud05,pfh05,lsg08}.  \\citet{dis02} conducted a Chandra survey of selected regions of  M31, and found that most of their bright X-ray sources were associated with GCs, with $\\sim$10\\% of GC sources exhibiting 0.5--7 keV luminosities $\\ga$10$^{38}$ erg s$^{-1}$. They also showed the M31 GC  population to be significantly different from that of the Milky Way, as $\\sim$30\\% of the M31 GC X-ray sources exhibited luminosities $>$10$^{37}$ erg s$^{-1}$, whereas  only one out of 11  Galactic GC X-ray sources exceeded 10$^{37}$ erg s$^{-1}$ \\citep{verb95}. These results contradict those of \\citet{S01}, who concluded that the GC X-ray populations of the Milky Way and M31 were similar.  \\citet{dis02} also compared the optical properties of M31 GCs with and without X-ray sources, as well as M31 GCs and Milky Way GCs with X-ray sources. They compared optical colours and apparent magnitudes, radial velocities, metallicities and colour excesses, as well as core radii. They found the only possibly significant difference to be in their luminosities: the median luminosity of M31 GCs with X-ray sources was $\\sim$0.55 magnitudes higher than that of M31  GCs without. Similarly, M31 X-ray GCs were slightly brighter in V than Milky Way X-ray GCs. \\citet{dis02} { speculated} that this might associate X-ray sources with  higher GC masses.  We report on three bright X-ray sources in M31, all associated with populous, old globular clusters: Bo\\thinspace 45, Bo\\thinspace 135 and Bo\\thinspace 375. We will argue that Bo\\thinspace 45 contains a good black hole (BH) candidate, while Bo\\thinspace 375 is likely to contain a neutron star (NS).  Bo\\thinspace 135 is more likely to contain a NS, but we cannot rule out a BH accretor.   We first review the emission and variability of LMXBs in their various states in Section~\\ref{bhs}, and then discuss  known properties of our  targets and their host clusters in Section~\\ref{gc}. We describe the observation and analysis in Section~\\ref{obs}, and our results in Section~\\ref{res}. We finally present our discussion and conclusions in Section~\\ref{dis}. ",
        "conclusions": "\\label{dis} We have examined the emission spectra and time variability of three X-ray sources associated with GCs in M31, using the 2006 December 26 and 2007 January 2 XMM-Newton observations. The emission  of XBo\\thinspace 45 is well described by a pure power law with photon index $\\sim$1.4, and is highly variable. This is consistent with a NS or BH LMXB in the low state \\citep{vdk94}, but is not consistent with a  BH LMXB in the high state or steep power law state \\citep[and references within]{mr03}, or  a NS LMXB emitting at $>$10$^{38}$ erg s$^{-1}$ \\citep{wsp88,cbc01,bcb03}.  { \\citet{sg98} calculated a distance to M31 of 784 kpc, $\\pm$13 kpc of statistical error, $\\pm$17 kpc of systematic error. Combining these distance uncertainties adds further uncertainties in the source luminosities of $\\pm$5\\%. Therefore, the 0.3--10 keV luminosity range for XBo\\thinspace  45 is 2.5$\\pm$0.2$\\times$10$^{38}$ erg s$^{-1}$, or 140$\\pm$10\\% of the Eddington limit for a 1.4 M$_{\\odot}$ neutron star primary. However, several LMXBs have been found to contain neutron stars with mass as high as $\\sim$2.1 M$_{\\odot}$ \\citep[see e.g.][]{nice05,ozel06}; XBo\\thinspace 45 has a 0.3--10 keV luminosity of $\\sim$80\\% Eddington for such systems. Since \\citet{glad07} showed that transitions in neutron star LMXBs occur at $\\la$10\\% Eddington, we consider XBo\\thinspace 45 to exhibit low state behaviour at an apparent luminosity too high for a neutron star. We therefore identify the accretor in XBo\\thinspace 45 as a BH candidate. We note that XBo\\thinspace  45  has been observed several times by the Einstein and ROSAT observatories over the last $\\sim30$ years, varying in luminosity only by a factor of $\\sim$2. This behaviour is consistent with that predicted for a GC BH binary formed by tidal capture \\citep{kal04}.  In contrast to XBo\\thinspace 45,} the observed two component emission  of XBo\\thinspace 375 is consistent with a bright NS LMXB, but not a low state NS or BH LMXB, nor a BH in the high or steep power law states. Hence, we classify XBo\\thinspace 375 as a NS LMXB.  Finally, the emission spectrum of XBo\\thinspace 135 is consistent with a pure power law with photon index $\\sim$1.6, but the fit is significantly improved by adding a blackbody component. Hence, XBo\\thinspace 135 is consistent with a NS or BH LMXB, and deeper observation is required for further classification.  { Of the thirteen bright X-ray sources in Galactic GCs, twelve have exhibited X-ray bursts, confiming their natures as NS LMXBs \\citep[see ][ for a review]{vl06}. \\citet{ph05} conducted a survey of  XMM-Newton observations of X-ray sources associated with M31 GCs, looking for X-ray bursts. They found simultaneous pn and MOS detections of bursts in two sources, and several burst candidates that were detected in the pn only. Such bursts would be identifiable in the lightcurves of our target sources. However, X-ray bursts are thought to be forbidden at luminosities $>$ 50\\% Eddington \\citep{lpt93}, hence we do not expect bursts from XBo\\thinspace 135 or XBo\\thinspace 375. }     XBo\\thinspace 45 exhibits low-state behaviour at luminosities $\\sim$10 times higher than expected for 1.4 M$_{\\odot}$ NS LMXBs; XBo\\thinspace 135 may also. A composite of two or three bright X-ray sources would likely result in a different spectrum to that observed in XBo\\thinspace 45; hence, if XBo\\thinspace 45 were a composite, it would be more likely to be made from { $\\sim$10} NS LMXBs in the low state. \\citet{dis02} calculated the probability for a GC to contain multiple bright X-ray sources: they argue that if $p$ is the probability of finding one  bright X-ray source in a GC, then Poisson statistics dictate that probability for two bright X-ray sources should be  $p$/2, while the probability for 3 bright X-ray sources should be  $p^{2}$/6. For a $p$ $\\sim$0.1--0.2, \\citet{dis02}  predict 3--5 GCs with 2 X-ray sources and less than 1 with 3 X-ray sources in M31. It is therefore very unlikely that XBo\\thinspace 45 combines the emission of $\\sim$10  low-state neutron star LMXBs.  It is however possible that the emission from XBo\\thinspace 45  is anisotropic. It exhibits low-state spectra at luminosities $\\sim$10--20 times higher than expected for neutron star LMXBs; hence it could simply be beamed by a factor of 10--20. In this case, the emission would be restricted to a small solid angle; one might expect to observe such beaming in $\\sim$5--10\\% of randomly oriented systems. \\citet{tp04} modeled the spectra of 43 M31 GCs. If we loosely class a low-state spectrum as a power law with $\\Gamma$ $<$1.7, then \\citet{tp04} found 20 GCs consistent with low-state spectra. Five of those,  XBo\\thinspace 5, XBo\\thinspace 82, XBo\\thinspace 135,  XBo\\thinspace 153 and XBo\\thinspace 386 exhibited apparent luminosities $>$10$^{38}$ erg s$^{-1}$. If we include XBo\\thinspace 45, then 6 out of   21 GC systems consistent with low-state spectra (including XBo\\thinspace 45), have measured luminosities $>$10$^{38}$, i.e. $\\sim$30\\%, a significantly larger fraction than expected from beaming.   We note that the host cluster Bo\\thinspace 45, which contains a BH LMXB candidate, is significantly larger than the cluster Bo\\thinspace 375 { (see Section 3.2)}, which we think contains a NS LMXB. \\citet{dis02} describe Bo\\thinspace 375 as not at all unusual, with parameters close to the median of M31 GCs. This suggests that Bo\\thinspace 45 (and also Bo\\thinspace 135) are particularly  massive, and may therefore be more prone to forming BH LMXBs. Therefore we conclude that, unlike the Milky Way, at least one GC in M31 is likely to contain a black hole binary."
    },
    "0807/0807.4558_arXiv.txt": {
        "abstract": "We present abundances of C, N, O, F, Na, and Fe in six giant stars of the tidally disrupted globular cluster NGC 6712. The abundances were derived by comparing synthetic spectra with high resolution infrared spectra obtained with the Phoenix spectrograph on the Gemini South telescope. We find large star-to-star abundance variations of the elements C, N, O, F, and Na. NGC 6712 and M4 are the only globular clusters in which F has been measured in more than two stars, and both clusters reveal F abundance variations whose amplitude is comparable to, or exceeds, that of O, a pattern which may be produced in $M$ $\\gtrsim$ 5$M_\\odot$ AGB stars. Within the limited samples, the F abundance in globular clusters is lower than in field and bulge stars at the same metallicity. NGC 6712 and Pal 5 are tidally disrupted globular clusters whose red giant members exhibit O and Na abundance variations not seen in comparable metallicity field stars. Therefore, globular clusters like NGC 6712 and Pal 5 cannot contribute many field stars and/or field stars do not form in environments with chemical enrichment histories like that of NGC 6712 and Pal 5. Although our sample size is small, from the amplitude of the O and Na abundance variations, we infer a large initial cluster mass and tentatively confirm that NGC 6712 was once one of the most massive globular clusters in our Galaxy. ",
        "introduction": "\\label{sec:intro}  The formation and evolution of our Galaxy remains one of the great unanswered questions in modern astronomy. \\citet{els62} suggested formation via the monolithic collapse of a gaseous protocloud on a timescale of 10$^8$ years. \\citet{sz78} challenged this notion by proposing that the halo formed through the accretion of independent fragments over a longer period, 10$^9$ years. These seminal works studied Galactic archaeology using the kinematics and metallicities of stars and globular clusters in the disk and halo. Today, Galaxy formation is discussed within the context of $\\Lambda$CDM cosmology and hierarchical structure formation \\citep{white78,freeman02} with the ongoing accretion of the Sagittarius dwarf galaxy being the most prominent example \\citep{ibata94}.  Another mechanism for populating the disk and halo is through the destruction of globular clusters via tidal shocks, two body relaxation etc \\citep{gnedin97}. Although the current mass in globular clusters is small, the initial globular cluster population may have been considerably larger than the present population \\citep{gnedin97}. The globular cluster Palomar 5 exhibits large tidal tails that extend over 10 degrees and contain more mass than the remaining cluster \\citep{odenkirchen01,odenkirchen03}. Therefore, Pal 5 is in the process of being tidally disrupted and is currently contributing stars to the disk and halo.  Chemical abundances place strong constraints upon the fraction of halo and disk stars that may come from disrupted globular clusters and/or the types of globular clusters that may populate the disk and halo. Specifically, every well studied Galactic globular cluster exhibits large star-to-star abundance variations for the light elements from C to Al \\citep{smith87,kraft94,gratton04}. Although the amplitude may vary from cluster to cluster, the abundances of C and O are low when N is high, O and Na are anticorrelated as are Mg and Al. Indeed, Str{\\\"o}mgren photometry reveals that every globular cluster has large star-to-star variations in the $c_1$ = $(u-v)-(v-b)$ index at all evolutionary stages \\citep{grundahl00}. The Str{\\\"o}mgren $u$ filter includes the 3360\\AA\\ NH molecular lines, and \\citet{6752nh} recently showed that the N abundances are directly correlated with the $c_1$ index. Therefore, it is likely that all globular clusters possess large N abundance variations at all evolutionary stages. Although hydrogen burning at high temperatures may explain the observed abundance patterns \\citep{langer93,langer95,denissenkov98,karakas03}, the source of the nucleosynthesis and the nature of the pollution mechanism remain unknown. Intermediate-mass ($\\sim$3 to 8$M_{\\odot}$) asymptotic giant branch (AGB) stars were the assumed polluters owing to the mono-metallic nature of most GCs, even though detailed AGB models have so far mostly failed to match the observations \\citep{fenner04,karakas06b}. Nevertheless, these abundance patterns seen in every cluster have rarely, if ever, been observed in field stars to date \\citep{pilachowski96,gratton00b}.  \\citet{smith02} conducted a detailed abundance analysis of four bright giant stars in Pal 5 and found variations of O, Na, and Al. (No abundance measurements have been performed upon stars in the tidal tails of Pal 5.) While most stars lost from a tidally disrupted cluster would be main sequence stars, abundance variations of O, Na, and Al have now been identified on the main sequences of globular clusters \\citep{gratton01,cohen05}. Since no radial gradients are associated with the O to Al abundance variations (with the exception of 47 Tucanae [\\citealt{nf79,briley97}]), observations of red giants in the cluster should be equivalent to observing red giants in the tidal tails. That abundance variations of O, Na, and Al are found in Pal 5 suggests that clusters like Pal 5 cannot provide many field stars and/or field stars do not form in environments with chemical enrichment histories similar to Pal 5. Of great interest for our understanding of Galactic and globular cluster formation would be the identification of clusters undergoing tidal disruption in which no light element abundance variations are detected.  Of the large sample of globular clusters studied by \\citet{paresce00} using the Hubble Space Telescope, all have mass functions (as inferred from their luminosity functions) which peak at 0.25M$_\\odot$. Not surprisingly, the mass function of Pal 5 is flatter than other clusters revealing significant depletions of low mass stars presumably stripped by the Galactic tidal field \\citep{koch04}. The globular cluster NGC 6712 is a small and sparse globular cluster whose mass function peaks at 0.75M$_\\odot$ instead of 0.25M$_\\odot$ \\citep{demarchi99,andreuzzi01}. That is, NGC 6712 is the only cluster whose mass function decreases with decreasing mass. With an orbit penetrating deep into the bulge, $R_{\\rm pericentric}$ = 0.9 kpc \\citep{dinescu99}, tidal forces have stripped away a substantial fraction of NGC 6712's lower mass stellar population. Calculations suggest that NGC 6712 may have lost up to 99\\% of its original mass \\citep{takahashi00}. All that remains of NGC 6712 is a remnant core of a cluster that was probably once one of the most massive in the Galaxy. The presence of a high luminosity x-ray source and a surprisingly large blue straggler population reinforce the idea that NGC 6712 was once much more massive and concentrated \\citep{paltrinieri01}. Therefore, NGC 6712 has almost certainly contributed stars to the disk and/or halo. Previous abundance analyses of NGC 6712 only considered one post-AGB star \\citep{jasniewicz04,mooney04} whose composition may not reflect the composition of the cluster due to the rich nucleosynthesis occurring in the late phases of stellar evolution. In this paper, we present the first detailed chemical abundance analysis of bright red giant stars in this tidally disrupted globular cluster. ",
        "conclusions": "\\label{sec:summary}  Based on high resolution infrared spectra, we derive abundances of C, N, O, F, Na, and Fe in six giant stars of the tidally disrupted globular cluster NGC 6712. For the elements C, N, O, F, and Na, we find large star-to-star abundance variations and correlations between these elements, a characteristic that NGC 6712 shares with every well studied Galactic globular cluster. This is only the second cluster in which F abundances have been measured in useful numbers of stars and both clusters show F variations whose amplitude is comparable to, or exceeds, that of O. Within the limited data, globular clusters appear to have lower F abundances than field and bulge stars at the same metallicity. Of great interest would be measurements of F in additional stars in $\\omega$ Cen and other globular clusters as well as in larger samples of field stars, with both samples overlapping in metallicity. From the amplitude of the O and Na abundance variations, we tentatively confirm that NGC 6712 was once one of the most massive clusters in our Galaxy.  NGC 6712 is a tidally disrupted cluster as revealed through its highly unusual luminosity function. Pal 5 is another tidally disrupted globular cluster. Both NGC 6712 and Pal 5 have almost certainly contributed stars to the disk and halo. Both clusters exhibit large star-to-star abundance variations for light elements, a characteristic which has yet to be identified in field halo stars. Therefore, the light element abundance variations detected in NGC 6712 indicate that clusters like NGC 6712 and Pal 5 have not provided many field stars and/or field stars did not form in environments with chemical enrichment histories like NGC 6712 and Pal 5. As pointed out by \\citet{smith02}, disrupted globular clusters like Pal 5 have lost CN-strong, O-poor, Na-rich, Al-rich stars to the halo field. But where are these stars? Of great interest would be an abundance analysis of stars within the tidal tails of Pal 5 as well as a large-scale dedicated search for O, Na, and Al abundance anomalies in field halo stars."
    },
    "0807/0807.3652_arXiv.txt": {
        "abstract": "{We continue investigation of the hidden plane-mirror symmetry in the distribution of excursion sets in cosmic microwave background (CMB) temperature anisotropy maps, previously noticed in the three-year data of the Wilkinson microwave anisotropy probe (WMAP), using the WMAP 5 years maps. The symmetry is shown to have higher significance, $\\chi^2 < 1.7$, for low multipoles $\\ell < 5$, while disappearing at larger multipoles, $\\chi^2 > 3.5 $ for $\\ell > 10$. The study of the sum and difference maps of temperature inhomogeneity regions, along with simulated maps, confirms its existence.The properties of these mirroring symmetries are compatible with those produced by the Sachs-Wolfe effect in the presence of an anomalously large component of horizon-size density perturbations, independent of one of the spatial coordinates, and/or a slab-like spatial topology of the Universe.} ",
        "introduction": "The properties of the CMB temperature anisotropy, as well as its polarization, are among the basic sources of information on cosmological parameters \\cite{DB1,Sp,Kom}. Their tiny features, such as the local spikes in the multipoles power spectrum, deviation from the statistical isotropy, and non-Gaussianity signatures, may all be the result of various fundamental processes having occurred in the early Universe. Among the reported anomalies are the alignment of the principal (Maxwellian) vectors of low multipoles, the north-south power asymmetry, the southern anomalous cold spot, etc (see de Oliveira-Costa et al. 2004; Copi et al. 2004,2007; Schwarz et al. 2004; Eriksen et al. 2004,2007; Cruz et al. 2005; Morales \\& Saez 2008). In the present paper we continue the study of another deviation from statistical isotropy: the hidden partial plane-mirror symmetry in the distribution of CMB temperature fluctuations excursion sets, i.e. of one-connected pixel sets equal to and higher than the given temperature threshold (lower for negative thresholds) previously found in the WMAP 3-years temperature maps \\cite{mirr}. We use the WMAP 5-years data \\cite{WMAP5}, not only to confirm the mirroring effect found in WMAP3 maps while studying the inhomogeneities in the distribution of the excursion sets, but also to reveal its other properties. By inquiring into the dependence of the mirror symmetry on the angular scale, we show that the effect has the highest significance at low multipoles $\\ell <5$ and that it quickly disappears at higher multipoles.  Similarly, the study of the sum and difference maps of temperature inhomogeneity regions provides additional insight into the mirroring. Namely, when the sum and difference maps are created via reflection of one of the maps, as it should be for mirrored images, anisotropic properties of excursion sets do survive, while they disappear if the sum map is created without reflection. Difference maps from independent radiometers (A-B) have been used to test the role of scan inhomogeneities and noise \\cite{mirr}. The signal-to-noise ratio for the studied excursion sets is about 4:1 and the excursion sets in (A-B) map do not show any specific property observed in the sum (A+B) map. The negligible role of the noise was also checked using the foreground reduced maps available in http://lambda.gsfc.nasa.gov/product/map/current/. Although contamination of a Galactic or interplanetary origin at these multipoles certainly cannot be excluded, following Gurzadyan et al (2007b), in the last section we discuss which properties of the Universe might be responsible for this effect, if it had a cosmological origin. ",
        "conclusions": "We have analyzed the scale dependence of partial mirror symmetry in the distribution of excursion sets, as previously found in the WMAP3 maps, with a nearly antipodal location of symmetry centers. Studies used the WMAP5 W-band maps, and the results obtained using WMAP5 and WMAP3 data agree. The centers lie close to one of the $\\ell=3$ multipole Maxwellian vectors, but not close to the sum of multipoles vectors up to $\\ell=8$ \\cite% {mirr}. Also they are close to the ecliptic pole and are nearly orthogonal to the CMB dipole apexes. They are moving towards the Galactic equator with the increase in the temperature threshold interval.  This symmetry appears to be a large-angle effect, i.e. it is stronger at low multipoles and it weakens rapidly  for larger $\\ell$: its statistical significance is quantified by $\\chi^2 < 1.7$ at $\\ell < 5$ and $\\chi^2 > 3.5$ at $\\ell > 10$. The symmetry was also tested using the following procedure: the sum map of the symmetry regions was obtained first, via rotation over $\\pi$, as is usually the case for mirrored images, and then without such a rotation. The clear mirroring in the first case and its complete absence in the second case make the case for a partial mirror symmetry stronger.  Turning to the origin of this symmetry, an unknown interplay of interplanetary and Galactic foregrounds or another unspecified non-cosmological contribution to the low multipoles certainly remains a possibility. If, however, it has a cosmological origin, then a signature of the simplest non trivial, $T^1$, spatial topology of the Universe is among the options, as discussed in Gurzadyan et al (2007b). For this topology, the points with coordinates $z$ and $z+L$ are identified where $z$ is one of the spatial coordinates. Such a model may be also considered as a limiting case of the Universe with compact flat spatial sections having the $T^3$ topology if the identification scales $L_1,L_2$ along two other spatial coordinates are much more than $L$ -- the slab topology (for early papers on a non trivial spatial topology of the Universe, see Zeldovich (1973), Sokolov \\& Schwartsman (1975), Sokolov \\& Starobinsky (1975), Fang \\& Sato (1983)). Note that, one should not expect any mirror symmetry, even a partial one, for comparable topological scales $L\\sim L_2,L_3$.  As shown in Starobinsky (1993), for this $T^1$ topology, a large-angle pattern of a CMB temperature anisotropy has just the form (\\ref{split}). The first term on its right hand side has the exact mirror symmetry with respect to the $(x,y)$-plane. It originates from the Sachs-Wolfe effect at the last scattering surface from density perturbations that do not depend on $z$. The second term represents a remaining part of anisotropy and does not have any symmetry at all. However, for $a_0L$ on the order of $R_{\\mathrm{hor}}$ or slightly more, where $a_0=a(t_0)$ is the present scale factor of a Friedmann-Robertson-Walker cosmological model, the latter term should somehow be suppressed since the Sachs-Wolfe contribution to it from the last scattering surface comes from perturbations having wave vectors with $|% \\mathbf{k}|\\ge 2\\pi/L$. That is why one expects the total large-angle pattern of $\\Delta T/T$ to have an \\emph{approximate} mirror symmetry in this case. \\footnote{% Note the other effect worsening the symmetry at large angles (low multipoles): a contribution from the integrated Sachs-Wolfe effect at small redshifts due to a cosmological constant (first calculated in \\cite{KS85}) or dynamical dark energy.}  More generically and not connected with a non trivial spatial topology of the Universe, such an approximate mirror symmetry at large angles arises when the large-scale $z$-independent part of density perturbations inside the last scattering surface is anomalously large. In all these cases, the rms amplitude of the first term on the right hand side of Eq. (\\ref{split}) quickly becomes negligible compared to the second one with the growth of $% \\ell$. Weakening of the mirroring symmetry for higher values of $\\ell$ found in Sect. 3 is in a good agreement with this theoretical prediction and may be considered as an additional argument for the reality of the mirroring effect.  Until now, searches for the mirroring effect of the form (\\ref{split}) or directly for the $T^1$ topology gave negative results (see e.g. de Oliveira-Costa et al. 1996; de Oliveira-Costa et al. 2004; Cornish et al. 2004), rasing the lower limit on the physical topological scale $a_0L$ to $\\sim R_{\\mathrm{hor}}=14$ Gpc. The numerical value is given for the standard $\\Lambda$CDM cosmological model with $\\Omega_m=0.3,~% \\Omega_{\\Lambda}=0.7$. However, for higher values of $a_0L$ this topology is not excluded. This follows already from the fact that the much more restrictive cubic $T^3$ topology ($L=L_1=L_2$) with $a_0L > R_{hor}$ is still considered as a viable possibility; see the recent papers \\cite{AJLS07,A08}, where an inconclusive evidence for the latter case with $a_0L \\approx 1.15R_{hor}$ is presented. This shows that topological explanation of the partial mirror symmetry investigated in this paper is possible. From our analysis, it is still too early to speak about the value of $L$, since the secure separation of CMB temperature fluctuations into a mirrored and non-mirrored parts needs higher resolution maps. Also, a non-topological (though still cosmological) explanation of such an effect is possible, as pointed out above. In this respect, see the recent papers \\cite{GK2008} where it was shown that voids can act as hyperbolic lenses in a spatially flat Universe, producing specific signatures in CMB temperature fluctuations. Future observational data will help solve these problems.  Shortly before submission of this manuscript, a paper appeared \\cite{GE} where CMB statistical anisotropy of an axial type was studied with the preferred axis very close to the one defined by our $CE_N-CE_S$ direction."
    },
    "0807/0807.0717_arXiv.txt": {
        "abstract": "A line list of vibration-rotation transitions for $^{13}$C substituted HCN is presented. The line list is constructed using known experimental levels where available, calculated levels and {\\it ab initio} line intensities originally calculated for the major isotopologue. Synthetic spectra are generated and compared with observations for cool carbon star WZ Cas. It is suggested that high resolution HCN spectra recorded near 14 $\\mu$m should be particularly sensitive to the $^{13}$C to $^{12}$C ratio. ",
        "introduction": "\\label{sec:Intro}  Carbon giant stars are though to arise from the third dredge up of asymptotic giant branch (AGB) stars, which pollutes the envelope and atmosphere with nuclear processed material from the interior. This increases the abundance of carbon and the $^{13}$C/$^{12}$C ratio to well above the terrestrial level (0.011). In fact the $^{13}$C/$^{12}$C ratio in carbon giants has been measured to be as high as one-third \\citep{ab97}. There are a number line lists or opacity functions available for \\hcn\\ \\citep{Jorgensen1,Jorgensen2,ao98,ha02b,ha06}. However, when considering the contribution to opacity by molecular species such as HCN it may not be sufficient to account only for \\hcn, but also for the isotopologue \\htcn.  The calculation of a complete triatomic linelist is computationally expensive, requiring several tens of thousands of CPU hours \\citep{te07}. However, within the Born-Oppenheimer approximation the electronic structure of isotopologues is identical. Thus both the potential energy and electric dipole moment functions are identical for all isotopologues. This implies that the vibration-rotation frequencies and transition intensities are likely to be very similar.  For a hetronuclear molecule, such as HCN, the main differences between the spectra of isotopologues are caused by the change in reduced mass. In the harmonic approximation the vibrational contribution to the line frequency is proportional to $\\rho^{-\\frac{1}{2}}$, where $\\rho=m_1m_2/(m_1+m_2)$ is the reduced mass. So the frequency of the CN stretch vibrational mode of \\hcn\\ and \\htcn\\ will differ by the order of 1\\% and the bend and H-C stretch mode by less.  In terms of both the observation of \\htcn\\ and the line blanketing of a model atmosphere this shift in line frequency is more important than the small differences in line intensity between comparable lines of \\hcn\\ and \\htcn.  It is the aim of this work to compute a set of energy levels and line frequencies for \\htcn\\ and \\htnc.  These energy levels will be used in conjunction with the Einstein A coefficients for \\hcn\\ and \\hnc\\ computed by \\citet{ha02b}, to generate a \\htcn/\\htnc\\ linelist. The line frequencies in this \\htcn/\\htnc\\ linelist have been corrected using available laboratory data in order to increase its accuracy. The new linelist has been used to compute synthetic spectra for C-stars, with different $^{13}$C/$^{12}$C ratios using star WZ Cas as a prototype. ",
        "conclusions": "\\label{sec:conc}   We present a new set of {\\em ab initio} rotation-vibration energy levels for \\htcn\\ and \\htnc\\  which were calculated for angular momenta of $J$ = 0, 1, 2, 3, 5, 10, 20, 30, 40, 60 and for both even ($e$) and odd ($f$) parity. The states were computed up to an energy of at least $10000 + B(J(J+1))$~\\cm\\  above the \\htcn\\ ground state, where $B\\sim1.5$~\\cm\\ is the HCN rotational constant. The quantum number assignments were made using a method similar to that described in \\citet{ha06}.   The new linelist has been incorporated into our computations of C-rich synthetic spectra. The detailed analysis of the infrared spectra of C-giant star with high $^{13}$C/$^{12}$C ratios ought to take account of \\htcn\\ and \\htnc\\ species. Moreover in many cases HCN spectra probably provides the best chance of determining the $^{13}$C/$^{12}$C ratios in atmospheres of the coolest stars as the CO bands at 2.3 $\\mu$m are usually saturated and other molecular bands are severely blended. Our upgraded opacity sources can be used for the determination of carbon isotopic ratios in atmospheres of carbon stars. The most promising  regions are those around 3.6 and 14 $\\mu$m where the the  $v_2+v_3$ (bend and CN stretch) combination bands and $v_2$ (bend) fundamental and hot bands of \\hcn\\  and \\htcn\\ molecules are located, respectively. However, to do this requires to use the spectral data of the proper quality.  Finally, the use of \\hcn\\ and \\htcn\\ lines for numerical analysis of infrared spectra of evolved stars is restricted by the incompleteness of presently available opacity sets, see \\citet{te07}. Special attention needs be paid to the computation of other lines for polyatomic molecules such as C$_3$, NH$_3$, CH$_4$ and C$_2$H$_2$, and their isotopologues."
    },
    "0807/0807.2462_arXiv.txt": {
        "abstract": "We carry out a suite of cosmological simulations of modified action $f(R)$ models where cosmic acceleration arises from an alteration of gravity instead of dark energy.  These models introduce an extra scalar degree of freedom which enhances the force of gravity below the inverse mass or Compton scale of the scalar.  The simulations exhibit the so-called chameleon mechanism, necessary for satisfying local constraints on gravity, where this scale depends on environment, in particular the depth of the local gravitational potential.  We find that the chameleon mechanism can substantially suppress the enhancement of power spectrum in the non-linear regime if the background field value is comparable to or smaller than the depth of the gravitational potentials of typical structures.   Nonetheless power spectrum enhancements at intermediate scales remain at a measurable level for models even when the expansion history is indistinguishable from a cosmological constant, cold dark matter model.  Simple scaling relations that take the linear power spectrum into a non-linear spectrum fail to capture the modifications of $f(R)$ due to the change in collapsed structures, the chameleon mechanism, and the time evolution of the modifications. ",
        "introduction": "\\label{sec:intro}   Cosmic acceleration can arise from either an exotic form of energy with negative pressure or a modification to gravity in the infrared. Self-consistent models for the latter are highly constrained by the stringent tests of gravity in the solar system. Additional propagating degrees of freedom must be suppressed by non-linearities in their equations of motion \\cite{Vainshtein:1972sx,Deffayet:2001uk,Dvali:2006su} in a stable manner \\cite{Dolgov:2003px,Sawicki:2007tf,Seifert:2007fr}.  These suppression mechanisms manifest themselves  with the formation of non-linear structure in the Universe \\cite{Lue:2004rj,Koyama:2007ih,hu07b}. Understanding the physical content, phenomenology and even the basic viability of such models thus requires cosmological simulations.  One possibility that has received much recent attention is the so-called $f(R)$ class of models (see \\cite{Sotiriou:2008rp} and references therein). These models generate acceleration through a replacement of the Einstein-Hilbert action by a function of the Ricci or curvature scalar $R$  \\cite{Caretal03,NojOdi03,Capozziello:2003tk}. They also introduce an extra propagating scalar degree of freedom that acts as an effective fifth force on all forms of matter \\cite{Chi03,Chiba:2006jp}. The range of the force depends non-linearly on the local curvature and can be made to become infinitesimal at high curvature.  With an appropriate choice of the function $f(R)$, deep potential regions can trap the field at high curvature leading to a nonlinear ``chameleon mechanism'' \\cite{khoury04a} that suppresses local deviations from ordinary gravity \\cite{Navarro:2006mw,Faulkner:2006ub,hu07a}.  Whether or not solar system tests of gravity are satisfied in an $f(R)$ model then depends on the depth of the gravitational potential including the astrophysical and cosmological structure surrounding it \\cite{hu07a}.  Likewise observable deviations from ordinary gravity for upcoming dark energy probes such as weak lensing, galaxy clustering and clusters of galaxies depend on the whole history of non-linear structure formation.  Previous cosmological simulations ({\\it e.g.}~\\cite{WhiKoc01,Maccio:2003yk,Nusser:2004qu,StaJai06,Laszlo:2007td}) have focused on modifications of the force law with a fixed and density independent range. Such modifications alone are incapable of satisfying local tests of gravity.  In a companion paper \\cite{oyaizu08b}, the numerical methodology for solving the non-linear field equation of $f(R)$ gravity was established. In this second paper of the series, we apply this methodology and carry out a suite of cosmological simulations of $f(R)$ models that are chosen to expose the impact of the chameleon mechanism on the power spectrum of the matter and the lensing potential.  We begin in \\S \\ref{sec:fr} with a review of non-linear gravitational dynamics in $f(R)$ models and proceed to the simulation results in \\S \\ref{sec:results}.  We discuss these results in \\S \\ref{sec:conclusion}. ",
        "conclusions": "\\label{sec:conclusion}  We have carried out the first cosmological simulations of $f(R)$ models for cosmic acceleration that exhibit the chameleon mechanism.  The chameleon mechanism involves a non-linear field equation for the scalar degree of freedom that suppresses the range of the gravitational force modification or Compton scale in the deep gravitational potential wells of cosmological and astrophysical structure.  We have here focused on its impact on the matter power spectrum or equivalently the potential power spectrum relevant for weak lensing surveys.   While we have simulated only one  particular functional form of $f(R)$, we expect that the qualitative behavior of other models that exhibit a chameleon behavior  to yield similar results once scaled to the appropriate Compton scales and field amplitudes.  In the absence of the chameleon mechanism, gravitational interactions would have an enhancement of a factor of 4/3 on all scales smaller than the Compton scale in the cosmological background eventually leading to order unity enhancements in the power at high wavenumber.  The chameleon mechanism turns on when the depth of the local gravitational potential becomes comparable to the field amplitude in the background. We have shown through otherwise identical simulations of structure with the chameleon mechanism artificially turned off that once the chameleon appears, it causes a substantial reduction of the enhanced power.   For example, for a field amplitude of $|f_{R0}|=10^{-6}$ the change in the enhancement of power at $k \\sim 1h^{-1}$Mpc is a factor of four.  Even in cases where current cosmological structures do not possess a chameleon $(|f_{R0}| \\simgt 10^{-5})$, there still is an impact on the power spectrum due to evolutionary effects.  In $f(R)$ models where the field amplitude decreases with curvature, the chameleon can appear again at high redshift when the building blocks of current structure were assembled.   Scaling relations which take the linear power spectrum and map it into the non-linear regime qualitatively misestimate the non-linear power spectrum in several ways. For example, the Smith {\\it et al.}~\\cite{smith03a} prescription assumes that the non-linear power spectrum depends only on the shape of the linear power spectrum near the non-linear scale.  This prescription fails to describe both the chameleon mechanism and the change in the structure and abundance of collapsed objects leading to a severe misestimate at high $k$.  A halo model can in principle do better to model these effects but simple prescriptions that scale the mass function to the linear variance and leave halo profiles unchanged also do not describe the non-linear effects to sufficient accuracy ({\\it cf.}~\\cite{hu07b}).  In the next paper of this series, we intend to study the impact of $f(R)$ modifications on halo properties.     \\smallskip  \\noindent {\\it Acknowledgments}: We thank N. Dalal, B. Jain, J. Khoury, K. Koyama, A. Kravtsov, A. Upadhye, I. Sawicki, F. Schmidt, J. Tinker, and F. Stabenau for useful conversations. This work was supported in part by the Kavli Institute for Cosmological Physics (KICP) at the University of Chicago through grants NSF PHY-0114422 and NSF PHY-0551142 and an endowment from the Kavli Foundation and its founder Fred Kavli. HO was additionally supported by the NSF grants AST-0239759, AST-0507666, and AST-0708154 at the University of Chicago. WH and ML  were additionally supported by U.S.~Dept.\\ of Energy contract DE-FG02-90ER-40560 and WH by the David and Lucile Packard Foundation. Some of the computations used in this work have been performed on the Joint Fermilab - KICP Supercomputing Cluster, supported by grants from Fermilab, Kavli Insititute for Cosmological Physics, and the University of Chicago."
    },
    "0807/0807.2181_arXiv.txt": {
        "abstract": "Single-field models of inflation are analysed in light of the WMAP five-year data. Assuming instantaneous reheating, we find that modular/new inflation models with small powers in the effective inflaton self-interaction are more strongly constrained than previously. The model with a cubic power lies outside the $2\\sigma$ regime when the number of e-folds is $\\caln \\le 60$. We also find that the predictions for the intermediate model of inflation do not overlap the $1\\sigma$ region regardless of the power of the monomial potential. We analyse a number of ultra-violet, DBI braneworld scenarios involving both wrapped and multiple-brane configurations, where the inflaton kinetic energy is close to the maximum allowed by the warped geometry. In all cases, we find that the parameters of the warped throat are strongly constrained by observations. ",
        "introduction": "The inflationary scenario postulates that the universe underwent a phase of very rapid, accelerated expansion in its distant past. Observations -- most notably from the anisotropy power spectrum of the Cosmic Microwave Background (CMB) -- have provided strong support for the paradigm. Despite this success, however, the mechanism which drove the inflationary expansion has yet to be identified. Indeed, there remain many viable versions of the scenario. (For reviews, see, e.g., Refs. \\cite{Lyth:1998xn,Liddle:2000cg,Guth:2005zr,Cline:2006hu,HenryTye:2006uv,McAllister:2007bg,Linde:2007fr,Lyth:2007qh}.) In view of this, it is important to constrain, and ultimately rule out, as many models of inflation as possible.  In this paper, we consider a wide range of inflationary models that are well motivated from particle physics and unified field theory. We focus on models where the cosmic expansion was driven by a single, self-interacting scalar `inflaton' field, $\\phi$. In general, single-field inflationary models are characterised by the action \\begin{equation} \\label{generalaction} S= \\int d^4x \\sqrt{-g} \\left[ \\frac{\\mpl^2}{2} R + P(\\phi , X) \\right], \\end{equation} where $R$ is the Ricci curvature scalar, the `kinetic function' $P(\\phi , X)$ is a function of the inflaton field and its kinetic energy $X \\equiv - \\frac{1}{2} g^{\\mu\\nu}\\nabla_{\\mu} \\phi \\nabla_{\\nu} \\phi$, and $\\mpl = (8\\pi G)^{-1/2}$ is the reduced Planck mass. In the simplest versions of the scenario, the inflaton has a canonically normalised kinetic energy, where $P=X- V(\\phi )$ and $V(\\phi)$ denotes the inflaton potential. In many higher-dimensional models motivated directly from string/M-theory, however, the function $P (\\phi ,X)$ has a more complicated form. This is the case, for example, in the Dirac-Born-Infeld (DBI) inflationary scenario, where inflation is driven by the propagation of one or more ${\\rm D}$-branes \\cite{DBI}.  We compare both canonical and non-canonical models of inflation with the recent observational bounds derived from the combined data of the Wilkinson Microwave Anisotropy Probe (WMAP) \\cite{wmap5}, Baryon Acoustic Oscillations (BAO) \\cite{Percival:2007yw} and Supernovae (SN) surveys \\cite{Riess:2004nr,Astier:2005qq, Riess:2006fw,WoodVasey:2007jb}. In Ref. \\cite{Alabidi:2006qa}, the three-year WMAP data was employed to rule out for the first time a large fraction of viable canonical models at more than $3\\sigma$. We begin by reconsidering this analysis in the light of the WMAP five-year data \\cite{wmap5}. We find that the conclusions of \\cite{Alabidi:2006qa} are confirmed by WMAP5, with the improved bounds on the scalar spectral index strongly constraining models of modular inflation at more than the 2$\\sigma$ level.%  We then proceed to consider non-canonical DBI brane inflationary models driven by a wrapped ${\\rm D}5$- or ${\\rm D}7$-brane \\cite{DBI,BECKER,KOB}. We focus on the `relativistic, ultra-violet' version of the scenario, where the brane is moving towards the tipped region of a warped throat with a kinetic energy close to the upper bound imposed by the warpfactor of the higher-dimensional space. We find that independently of the form of the inflaton potential, the geometry of the warped throat must be strongly constrained if such models are to satisfy the improved observational bounds from WMAP5. We also consider extensions of the scenario to multiple-brane configurations \\cite{thomasward,HLTW} and find that the same challenges arise as for the single-brane case.  The structure of the paper is as follows. In section II, we summarise the current bounds on the observational parameters. In section III, we investigate models of canonical inflation and determine which models are still viable. In Sections IV and V, we derive new observational limits on wrapped and multi-brane UV DBI inflation. We conclude with a discussion in Section VI. ",
        "conclusions": "In this paper, we have considered the status of a wide class of single-field inflationary models in the light of the recent WMAP5 data. For the reasonable range of e-fold values aforementioned, $\\caln = 54 \\pm 7$, we conclude that modular/new inflation models with $2<p\\lesssim3$ are under pressure from the data and that $p<0$ and $|p|\\to\\infty$ satisfy the observations. The required suppression of lower-order terms in the $p>0$ models poses a challenge to model builders, as discussed in Section III. Moreover, it is significant that independent data sets are now converging towards a preferred value of $n_s$, with only tiny differences between WMAP5 only, ${\\rm WMAP5}+{\\rm SDSS}$ \\cite{Percival:2006gt} and ${\\rm WMAP5}+{\\rm BAO} +{\\rm SN}$ data. This convergence implies that we can be more confident as to what represents a viable model from a phenomenological point of view.  We have also shown that the intermediate inflationary scenario can only satisfy the data at $2\\sigma$ for $|\\alpha|\\geq 1.72$, and that the predictions of the model do not overlap the $1\\sigma$ region for any value of $\\alpha$.  We emphasize that the above conclusions follow from assuming that the reheating of the universe immediately after inflation was instantaneous. However, the precise duration of the reheating process is not known. In particular, for a quadratic potential, a matter dominated phase may arise if the inflaton oscillates about its potential minimum for an extended interval before decaying. This would alter the number of e-folds that elapsed from the epoch when observable scales first crossed the Hubble radius during inflation to the time when inflation ended. In general, the number of e-folds depends on the effective equation of state after inflation and this would need to be known if a more precise value of $\\cal{N}$ is to be determined.  Complementary to our work is that of Ref.~\\cite{Peiris:2008be}, where the authors start from the different stand point of aiming to construct the most general inflationary potential using slow-roll reconstruction with a specific e-fold and reheat temperature prior. Reconstructing the inflationary slow-roll parameters and not the spectral parameters from the data has the advantage of directly probing more fundamental parameters. They locate regions in $\\epsilon,\\eta$ space which are consistent with the WMAP5 data, and this allows to reconstruct the Hubble parameters and therefore (via the Hamilton-Jacobi equations) the inflationary potential. The authors also analyse higher-order parameters, but for the purposes of this work, we only need mention that the simplest fit is consistent with the crucial prior of $\\caln>15$, which we interpret as support for the inflationary paradigm.  We have also considered the relativistic, ultra-violet, DBI braneworld scenario, where the inflaton is identified in terms of the radial position of a wrapped ${\\rm D}5$- or ${\\rm D}7$-brane. We  found that when the brane's kinetic energy is maximised, new constraints on the parameters of the bulk geometry can be derived from the WMAP5 data in the $(r , n_s )$ plane. Such constraints are independent of the precise form of the inflaton potential and are stronger than existing limits originating from field-theoretic bounds on the tensor-scalar ratio. In all cases, consistency with the data requires a significant reduction in the volume of the base or the discovery of new classes of Calabi-Yau four-folds which allow for larger Euler numbers.  We then extended our analysis to a recently proposed relativistic, multi-brane DBI scenario. For the case where the inflaton has a quadratic potential, we found that the predicted value of the spectral index is compatible with observations. However, the ratio $\\Vol/N$ is still strongly bounded from above, as in the single brane models, and consequently the same theoretical problems arise in this case also."
    },
    "0807/0807.4522_arXiv.txt": {
        "abstract": "Infrared absorption lines of $\\mathrm{H}_{3}^{+}$, including the metastable $R$(3,3)$^l$ line, have been observed toward eight bright infrared sources associated with hot and massive stars located in and between the Galactic Center Cluster and the Quintuplet Cluster 30~pc to the east. The absorption lines with high velocity dispersion arise in the Galaxy's Central Molecular Zone (CMZ) as well as in foreground spiral arms. The temperature and density of the gas in the CMZ, as determined from the relative strengths of the $\\mathrm{H}_{3}^{+}$ lines, are $T$=200~--~300~K and $n\\leq$50~--~200~cm$^{-3}$. The detection of high column densities of $\\mathrm{H}_{3}^{+}$ toward all eight stars implies that this warm and diffuse gaseous environment is widespread in the CMZ. The products of the ionization rate and path length for these sight lines are 1000 and 10 times higher than in dense and diffuse clouds in the Galactic disk, respectively, indicating that the ionization rate, $\\zeta$, is not less than 10$^{-15}$~s$^{-1}$ and that $L$ is at least on the order of 50~pc. The warm and diffuse gas is an important component of the CMZ, in addition to the three previously known gaseous environments: (1) cold molecular clouds observed by radio emission of CO and other molecules, (2) hot ($T~=$~10$^4$~--~10$^6$~K) and highly ionized diffuse gas ($n_e~=$~10~--~100~cm$^{-3}$) seen in radio recombination lines, far infrared atomic lines, and radio-wave scattering, and (3) ultra-hot ($T~=$~10$^7$~--~10$^8$~K) X-ray emitting plasma. Its prevalence significantly changes the understanding of the environment of the CMZ. The sight line toward GC~IRS~3 is unique in showing an additional $\\mathrm{H}_{3}^{+}$ absorption component, which is interpreted as due to either a cloud associated with circumnuclear disk or the ``50~km~s$^{-1}$ cloud'' known from radio observations. An infrared pumping scheme is examined as a mechanism to populate the (3,3) metastable level in this cloud. ",
        "introduction": "\\subsection{The Central Molecular Zone}  The region of the Galaxy within 200~pc of the center, known as the Central Molecular Zone (CMZ), contains 10\\% of the Galactic interstellar molecular mass \\citep[and references therein]{mor96,gen94,mez96}. The interstellar medium (ISM) in the CMZ is exceptional in many other ways. It is highly turbulent, pervaded by an intense magnetic field \\citep[for a review; ][]{mor06}, contains a multitude of massive stars \\citep[e.g.,][]{mon92,cot98,mun06} and super-massive clusters \\citep[e.g.,][for a short review]{nag90,nag95,fig04}, and also is home to numerous highly energetic objects ranging from supernova remnants to cataclysmic variables \\citep[CVs, e.g.,][]{mun03}. The particle densities in its dense molecular clouds are of order 10$^4$~cm$^{-3}$ and higher. The volume filling factor, \\emph{f$_{\\rm v}$}, of dense molecular clouds has previously been estimated to be $\\sim$ 10\\% \\citep{mor96}, but is more likely $\\sim$ 1\\% \\citep{oka05} in view of the observed visual extinction of $\\mathit{A_V}$~=~25~--~40~mag \\citep{cot00}.  If dense molecular clouds only occupy 1\\% of the CMZ, what interstellar environments make up the remaining 99\\%? There have been some reports of lower density ``diffuse'' molecular gas in the intercloud region. Detailed analysis of the $\\mathit{J}$ = 1 $\\rightarrow$ 0 emission lines of CO and C$^{18}$O by \\citet{dah98}, and the $\\mathit{J}$~=~2~$\\rightarrow$~1 and $\\mathit{J}$ = 1 $\\rightarrow$ 0 emission lines of CO by \\citet{oka98a} have demonstrated the presence of gas with densities of $\\sim$ 10$^{2.5}$ cm$^{-3}$. Indeed \\citet{oka98a} conclude that CO emission from the Galactic center (GC) is dominated by emission from gas of that density.  The volume filling factor of \\emph{f$_{\\rm v}$} $\\sim$ 0.1 claimed by \\citet{mor96} may include such gas. In this paper we present evidence for an even lower density molecular environment in which H$_2$ must still be plentiful but where CO is not abundant. In hindsight, the radio absorption lines with high velocity dispersion of OH \\citep[e.  g.][]{bol64}, H$_2$CO \\citep[e. g.][]{pal69}, HCO$^+$ \\citep{lin81}, and CO \\citep{oka98b} may originate in this newly reported environment.  A second candidate to fill the intercloud medium is the ultra-hot plasma ($T$ $\\sim 10^7$ -- 10$^8$~K) observed by its diffuse X-ray emission at energies of 0.5 -- 10~keV. The radiation from the plasma is observed in all sightlines to the CMZ \\citep[e.g.,][]{koy89,yam93,koy96}. \\citet{laz98} suggested that the highly ionized regions with \\emph{n$_e$} $\\sim 10$ cm$^{-3}$ and $T$ $\\sim 10^6$~K that are required to explain the $\\lambda^2$ dependent radio-wave scattering toward the GC are located at the photon-dominated interfaces between the molecular clouds and this ultra-hot X-ray emitting plasma. They proposed that high density molecular clouds with a volume filling factor of perhaps \\emph{f$_{\\rm v}$} $\\sim$ 0.1 are surrounded by the scattering electron gas and the rest of the space is filled with this plasma (see their Fig. 9).  For an environment such as that described above, which is so dominated by ionized gas, ``Central Plasma Zone\" rather than CMZ would be a more appropriate term for the region. For a variety of reasons the ultra-hot and hot plasma cannot occupy the same regions where molecules abound. How those different categories of gas are distributed and how they coexist in the CMZ is yet to be understood. Observations of a new astrophysical probe of the CMZ, the infrared absorption spectrum of H$_3^+$ \\citep{geb96}, are beginning to shed fresh light on these questions \\citep{bol06}. Earlier measurements of this molecular ion in the Galactic center by \\citet{geb99}, \\citet{got02}, and \\citet{oka05} were made along only a very few sight lines. In this paper we report and analyze observations of seven transitions of H$_3^+$ along eight sight lines, which provide a more comprehensive diagnostic of the extent and nature of the H$_3^+$ - containing gas in the CMZ.   \\begin{table*}[bht] \\begin{center} \\tablewidth{\\textwidth} \\scriptsize \\caption{Target list.}\\label{tb1} \\begin{tabular}{lcccccll} \\hline \\hline Source                &  R.A.      &  Dec.       &   $l$    & $b$      &$L$/$L'$& Other Names               & Reference    \\\\ &     (J2000)&     (J2000) & [\\arcdeg]& [\\arcdeg]& [mag]  & in SIMBAD                 & 1, 2, 3.     \\\\ \\hline  GC~IRS~21        \\dots& 17:45:40.2 & $-$29:00:31 &$-$0.0561 &$-$0.0468 &        &NAME SGR~A~IRS~21, GCIRS~21 & 1, 2, 3.     \\\\ GC~IRS~3        \\dots & 17:45:39.8 & $-$29:00:24 & $-$0.0552&$-$0.0448 &  5.3   &BHA~11, GCIRS~3             & 1, 2, 3.     \\\\ GC~IRS~1W        \\dots& 17:45:40.4 & $-$29:00:27 &$-$0.0549 &$-$0.0471 &  5.5   &NAME SGR~A~IRS~1W, GCIRS~1W & 1, 2, 3.     \\\\ NHS~21           \\dots& 17:42:53.8 & $-$28:51:41 &  +0.0992 &$-$0.0553 &  4.62  &qF~577, [NHS93]~21          & 4, 5, 6.     \\\\ NHS~22           \\dots& 17:46:05.6 & $-$28:51:41 &  +0.1199 &$-$0.0482 &  6.0   & [NHS93]~22                 & 4, 5, 6.     \\\\ NHS~42           \\dots& 17:46:08.3 & $-$28:49:55 &  +0.1481 &$-$0.0426 &  6.05  & qF~578, [NHS93]~42         & 4, 5, 6.     \\\\ NHS~25           \\dots& 17:46:15.3 & $-$28:50:04 &  +0.1592 &$-$0.0657 &  6.2   &NAME PISTOL STAR, [NHS93]~25& 4, 5, 6.     \\\\ GCS~3-2       \\dots   & 17:46:14.7 & $-$28:49:41 &  +0.1636 &$-$0.0606 &  2.5   & GCS~3-2, [NHS93]~24        & 4, 5, 6.     \\\\ \\hline \\end{tabular}  \\tablecomments{Reference: 1. \\citet{kra95}, 2. \\citet{ott99}, 3. \\citet{tan06}, 4. \\citet{nag90}, 5. \\citet{nag93}, 6. \\citet{fig99}. }  \\end{center} \\end{table*}  \\subsection{H$_3^+$ as an Astrophysical Probe}  The crucial information obtained from the observations of H$_3^+$ reported here is the product of the ionization rate of molecular hydrogen \\emph{$\\zeta$}, and the aggregate length of the line of sight through the H$_3^+$ - containing clouds, $L$. The simple chemistry of H$_3^+$ allows us to obtain this product from the observed total H$_3^+$ column densities \\citep{oka05,oka06b}, as described below, under the assumptions that (1) there is a steady-state balance between creation and destruction of H$_3^+$ and (2) the dominant supplier of electrons is carbon, all of which is singly ionized.  Creation of H$_3^+$ begins with ionization of H$_2$.  Because in interstellar regions where H$_2$ is plentiful the ion-neutral reaction H$_2$ + H$_2^+$ $\\rightarrow$ H$_3^+$ + H is rapid, the limit to the formation of H$_3^+$ is the much slower ionization rate of molecular hydrogen,  \\[ {\\rm H_2} + {\\rm R} \\rightarrow {\\rm H_2^+} + e + {\\rm R}^{\\prime}. \\]  \\noindent The ionizing rays, R, are dominated by cosmic rays in the Galactic disk \\citep[][and references therein]{ind07}, but in the CMZ the ionization by X-rays and EUV radiation may be significant. The production rate per unit volume of H$_3^+$ is thus given by \\emph{$\\zeta$}$n$(H$_2$), where \\emph{$\\zeta$} is the effective ionization rate of H$_2$ including all the aforementioned contributors. In steady state, production of H$_3^+$ balances its destruction, which is dominated by dissociative recombination with electrons,  \\[ {\\rm H_3^+} + {\\rm e} \\rightarrow {\\rm 3H~~or~~H_2} + {\\rm H}, \\]  \\noindent and one obtains \\citep{geb99}  \\[ \\zeta n({\\rm H_2}) = k_{\\rm e}n({\\rm H_3^+})n({\\rm e}), \\]  \\noindent where $k_e$ is the electron recombination rate constant, whose temperature dependent value is $\\sim$ 10$^{-7}$ cm$^3$ s$^{-1}$ \\citep{mcc04}. Replacing the electron density with the atomic carbon density after depletion gives  \\[ {n({\\rm e})}/{n({\\rm H_2})} = ({2}/{f}) \\left[{n_{\\rm C}}/{n_{\\rm H}}\\right]_{\\rm SV} R_X, \\]  \\noindent where $R_X$ = 3 -- 10 is the factor increase in relative carbon abundance at the Galactic center over the solar vicinity, $[n_{\\rm C}/n_{\\rm H}]_{\\rm SV}$ \\citep{sod95}, and $f = 2n({\\rm H_2}) /n_{\\rm H}$ is the fractional density of molecular hydrogen relative to the total number of hydrogen atoms, $n_{\\rm H} = 2n({\\rm H_2}) + n({\\rm H})$. Using $N$(H$_3^+$)~=~$L$$n$(H$_3^+$) because of the constancy of $n$(H$_3^+$) \\citep[e.g.,][]{oka06b}, we obtain the equation central to the analysis in this paper relating the product \\emph{$\\zeta$L} to the total H$_3^+$ column density and other measurables,  \\begin{equation} \\zeta L = 2 k_{\\rm e} \\left[{n_{\\rm C}}/{n_{\\rm H}}\\right]_{\\rm SV} R_X N({\\rm H_3^+})_{total}/f. \\end{equation}  As discussed by \\citet{oka05}, in the Galactic center the spectrum of H$_3^+$ provides direct information about the temperature and density of the gas. There the relative population of H$_3^+$ in the ($J$, $K$) = (3,3) metastable level, which is 361 K above the lowest (1,1) level (note that $J$ =~0 does not exist in the ground vibrational state), is measurable and is a useful thermometer for temperatures of $\\sim$ 100 -- 1000~K. Observations of the fractional population in the (2,2) unstable level, from which H$_3^+$ decays to the (1,1) level by spontaneous emission in 27 days \\citep{pan86,nea96}, provides a good measure of cloud densities less than several hundred per cm$^{3}$. The lifetime of H$_3^+$ in diffuse clouds is $1/(k_{\\rm e} n_{\\rm e}) \\approx 10^9$~s, which is two orders of magnitude longer than the collisional timescale with molecular hydrogen. The rotational temperature of H$_3^+$ is therefore maintained via collisions with the ambient H$_2$. The thermalization of H$_3^+$ in interstellar clouds has been modeled by \\citet{oka04}, who provide H$_3^+$ population ratios $n$(3,3)/$n$(1,1) and $n$(3,3)/$n$(2,2) plotted as functions of temperature and density (see their Figs.~2-4). ",
        "conclusions": "%  \\subsection{The Ionization Rate}  An estimate of the ionization rate, $\\zeta$, in the Galactic center is the most important result of these $\\mathrm{H}_{3}^{+}$ observations. To avoid confusion, we use $\\zeta$(H) and $\\zeta$(H$_2$) for the ionization rates of H and H$_2$, respectively, in the Galactic disk where ionization is mainly due to cosmic-rays. We use $\\zeta$ for the rate in the CMZ where the effect of X-rays and FUV radiation may be significant. $\\zeta$(H) and $\\zeta$(H$_2$) are related approximately as $\\zeta$(H$_2$) = 2$\\zeta$(H) \\citep{dal06}.  Early studies of the cosmic ray ionization rate were made in order to understand the physical state of the interstellar medium and its heating mechanisms \\citep{hay61,spi68,fie69}. \\citet{spi68} reported a minimum value of \\emph{$\\zeta_f$}~=~6.8~$\\times$~10$^{-18}$~s$^{-1}$ from an extrapolation of the cosmic ray spectrum observed at Earth, and an upper limit of 1.2 $\\times$ 10$^{-15}$ s$^{-1}$ calculated from an assumed supernova frequency. Their \\emph{$\\zeta_f$} is the total ionization rate for atomic hydrogen including the effect of secondary electrons, written here as \\emph{$\\zeta$}(H). The most recent measurement of \\emph{$\\zeta$}(H), reported by \\citet{web98} using data from the \\emph{Voyager} and \\emph{Pioneer} spacecraft at distances up to 60 AU from the Sun, is \\emph{$\\zeta$}(H)~$\\geq$~(3--4)~$\\times~10^{-17}$~s$^{-1}$, which should replace the lower limit used by \\citet{spi68}.  With the advent of the {\\it Copernicus} satellite and its far-ultraviolet spectrometer, which allowed FUV studies of chemical abundances in the diffuse ISM \\citep[for a review; ][]{spi75}, column densities of molecules and radicals such as HD and OH could be measured.  The reaction sequences leading to the creation of these two molecules start from H$^+$ and are thus useful for estimating \\emph{$\\zeta$}(H). Values of \\emph{$\\zeta$}(H) determined in this way \\citep{odo74,bla77,har78} are mostly a few times 10$^{-17}$ s$^{-1}$, comparable to the lower limit given above. The comprehensive model of the diffuse interstellar medium by \\citet{dis86} suggested a somewhat higher value of \\emph{$\\zeta$}(H)~$\\approx$~7~$\\times$~10$^{-17}$~s$^{-1}$. However they used an H$_3^+$ dissociative recombination rate constant one thousand times lower than the current value which was popular in some circles at the time. They noted that if the thousand times higher constant were to be used, the value of \\emph{$\\zeta$}(H) would need to be increased by a factor of $\\sim$ 10.   The formation and destruction mechanisms for H$_3^+$ are the simplest among all astrophysical probes of the ionization rate, and this makes H$_3^+$, when detectable, a powerful tool for measuring the ionization rate. While the H$_3^+$ column densities observed in dense clouds, \\mbox{(0.4~--~2.3)~$\\times$~10$^{14}$~cm$^{-2}$}, are consistent with \\emph{$\\zeta$}(H$_2$) on the order of 3~$\\times$~10$^{-17}$~s$^{-1}$ \\citep{geb96, mcc99}, the surprisingly high $N$(H$_3^+$) observed in diffuse clouds, comparable to those in dense clouds in spite of their 10 times smaller visual extinction \\citep{mcc98, geb99, mcc02}, require an order of magnitude higher \\emph{$\\zeta$}(H$_2$). Crucial to this conclusion are the laboratory measurements of the dissociative recombination rate constant $k_e$ in Eq.~(1) \\citep{ama88,lar93,lar00,mcc03,mcc04}.  Together, \\citet{mcc02} and \\citet{ind07} have carried out a survey of H$_3^+$ in the diffuse interstellar medium toward 29 nearby stars. H$_3^+$ was detected in 14 of these sight lines with column densities of (0.6~--~6.5)~$\\times$~10$^{14}$~cm$^{-2}$. For these diffuse clouds the product $\\zeta$$L$ is in the range (0.5~--~4.4)~$\\times$~10$^4$~cm~s$^{-1}$. Separating $\\zeta$ and $L$ is not straightforward. Based on various methods of estimating $L$, which gave values of 2.2~-~31 pc for different clouds, the cosmic ray ionization rates were found to be in the range \\emph{$\\zeta$}(H$_2$)~=~(1.2~--~7.4)~$\\times$~10$^{-16}$~s$^{-1}$ with an average of 5 $\\times $10$^{-16}$~s$^{-1}$. Note that the primary ionization rate of H, $\\zeta$$_p$, used by \\citet{ind07} is related to $\\zeta$(H$_2$) by 2.3$\\zeta$$_p$ = $\\zeta$(H$_2$) \\citep{gla74}. This has established that the cosmic ray ionization rate in diffuse clouds is an order of magnitude higher than in dense clouds. This result might be understandable in view of attenuation of soft components of cosmic rays in dense clouds \\citep{cra78}, if a large flux of low energy cosmic rays is ubiquitous in the Galaxy \\citep{mcc03}, or if higher energy cosmic rays are trapped by Alfv\\'{e}n waves in diffuse clouds \\citep{pad05}. The discrepancy between the high \\emph{$\\zeta$}(H$_2$) and previous low \\emph{$\\zeta$}(H) values determined from HD and OH has been explained by \\citet{lis03} as due to neutralization of H$^+$ by charge exchange with small grains which had not been considered in the previous studies, in which the neutralization of H$^+$ was assumed to be due to much slower radiative recombination with electrons.  As seen in Table~\\ref{tb3}, the total column density of H$_3^+$ in the CMZ toward the eight observed stars ranges over (1.8 -- 6.1) $\\times$ 10$^{15}$ cm$^{-2}$, more than an order of magnitude higher than the column densities in diffuse clouds in the Galactic disk, suggesting even higher values of $\\zeta$ and $L$. Since we are interested in setting the minimum value of both $\\zeta$ and $L$, we use the maximum value of $f$~=~1 and the minimum value of $R_X$ = 3 \\citep{sod95, ari96} in Eq.~(1) and obtain the numerical expression,  \\begin{equation} (\\zeta L)_{\\rm min} = 7.4 \\times 10^{-11} {\\rm cm}^3 {\\rm s}^{-1} N({\\rm H}_3^+)_{\\rm total}. \\end{equation}  \\noindent To obtain Eq.~(2), we have used $k_e$~=~7.7~$\\times$~10$^{-8}$~cm$^3$~s$^{-1}$ calculated for $T$ = 250 K from the temperature dependence reported by \\citet{mcc04} assuming the same electron temperature as the gas temperature, which may be justified in view of the fast thermalization of electrons with molecules and atoms. We used the C to H ratio in the solar vicinity of 1.6 $\\times$ 10$^{-4}$ given by \\cite{sof04}, which is close to the value of 1.4 $\\times$ $10^{-4}$, determined directly from infrared spectra of CO and H$_2$ toward an embedded young star NGC~2024~IRS~2 \\citep{lac94}, and $1.3 \\times 10^{-4}$, found by \\citet{usu08} with the same technique but toward NGC~7538~IRS~1 in the outer Galaxy. The observed $N$(H$_3^+$)$_{\\rm total}$ yield lower limits for $\\zeta$$L$, ($\\zeta$$L$)$_{\\rm min}$ = (1.3 -- 4.5) $\\times$ 10$^5$ cm s$^{-1}$. Although lower limits, the values of ($\\zeta$$L$)$_{\\rm min}$ are 1000 times higher than values of $\\zeta$$L$ determined for dense clouds \\citep{mcc99} and 10 times higher than values for diffuse clouds in the Galactic disk \\citep{ind07}. We cannot cleanly separate $\\zeta$ and $L$ based on current understanding of the CMZ, but discuss two extreme cases in the following sections.   \\begin{table*} \\begin{center} \\tablewidth{\\textwidth} \\tiny  \\caption{Column densities of H$_3^+$ in the Galactic center.\\label{tb4}}  \\begin{tabular}{l cccc cc cc c cc} \\hline \\hline &\\multicolumn{4}{c}{$N({\\rm H_3^+})_{\\rm Level}$\\tablenotemark{a}} &\\multicolumn{2}{c}{Relative Population\\tablenotemark{a}} & & & & \\\\  &\\multicolumn{4}{c}{-----------------------------------------------------------------------} &\\multicolumn{2}{c}{--------------------------------------------} & & & & & \\\\ & $N(1,1)$ &$N(1,0)$ &$N(3,3)$ &$N(2,2)$ & $N(3,3)/N(1,1)$    &$N(3,3)/N(2,2)$     % & $n(\\rm{H_2})$  & $T(\\rm{H_2})$ & $N({\\rm H_3^+})_{\\rm Total}$ & $(\\zeta L)_{\\rm min}$     & $L$ \\tablenotemark{b}\\\\ Source &[10$^{14}$cm$^{-2}$]&[10$^{14}$cm$^{-2}$]&[10$^{14}$cm$^{-2}$]&[10$^{14}$cm$^{-2}$] & & & [cm$^{-3}$]                             & [K] & [10$^{14}$cm$^{-2}$] & [$10^3$cm~s$^{-1}$]                            & [pc]\\\\  \\hline   GC~IRS~21\\dots & 28.1$\\pm$12.3 &  9.0$\\pm$ 4.4 & 24.2$\\pm$12.1 & $<$ 8.3 & 0.86$\\pm$0.57 & $>$2.93 & $<$ 125 & 150--450  &   61.3$\\pm$  17.8 &      451$\\pm$     131 &  146$\\pm$  43\\\\ GC~IRS~3 \\dots & 10.8$\\pm$ 2.1 &  3.4$\\pm$ 1.3 &  8.4$\\pm$ 1.9 & $<$ 1.6 & 0.78$\\pm$0.23 & $>$5.36 & $<$ 50  & 225--400  &   22.5$\\pm$   3.2 &      166$\\pm$      23 &   54$\\pm$   8\\\\ GC~IRS~1W\\dots & 18.1$\\pm$ 3.8 &  7.9$\\pm$ 2.4 & 11.7$\\pm$ 3.0 & $<$ 3.3 & 0.65$\\pm$0.21 & $>$3.53 & $<$ 75  & 200--300  &   37.6$\\pm$   5.4 &      277$\\pm$      40 &   90$\\pm$  13\\\\ NHS~21   \\dots &  8.9$\\pm$ 2.2 &  3.7$\\pm$ 1.5 &  5.6$\\pm$ 1.4 & $<$ 2.3 & 0.62$\\pm$0.22 & $>$2.44 & $<$ 125 & 175--275  &   18.2$\\pm$   3.1 &      134$\\pm$      22 &   44$\\pm$   7\\\\ NHS~22   \\dots & 16.9$\\pm$ 5.6 &  7.9$\\pm$ 3.5 &  9.7$\\pm$ 2.7 & $<$ 3.6 & 0.57$\\pm$0.25 & $>$2.73 & $<$ 100 & 150--275  &   34.5$\\pm$   7.2 &      254$\\pm$      53 &   82$\\pm$  17\\\\ NHS~42   \\dots & 17.7$\\pm$ 5.1 &  8.4$\\pm$ 4.1 &  8.6$\\pm$ 4.1 & $<$ 4.0 & 0.49$\\pm$0.27 & $>$2.13 & $<$ 100 & 125--250  &   34.6$\\pm$   7.7 &      255$\\pm$      57 &   83$\\pm$  18\\\\ NHS~25   \\dots & 11.4$\\pm$ 5.9 &  4.0$\\pm$ 3.8 &  7.4$\\pm$ 5.3 & $<$ 3.9 & 0.65$\\pm$0.57 & $>$1.89 & $<$ 175 & 100--300  &   22.8$\\pm$   8.8 &      168$\\pm$      64 &   54$\\pm$  21\\\\ GCS~3-2  \\dots & 17.0$\\pm$ 1.7 &  4.6$\\pm$ 0.8 &  9.8$\\pm$ 1.6 & $<$ 3.0 & 0.57$\\pm$0.11 & $>$3.21 & $<$ 80  & 200--250  &   31.4$\\pm$   2.5 &      231$\\pm$      18 &   75$\\pm$   6\\\\  \\hline  GC~IRS~3~(50~km~s$^{-1}$)\\dots &  3.2$\\pm$ 0.5 &  2.1$\\pm$ 0.4 &  4.9$\\pm$ 0.5 & 1.9$\\pm$ 0.6 & 1.55$\\pm$0.29& 2.61$\\pm$0.50& 150--350    & 350--500  &  10.2$\\pm$    0.8 &       75$\\pm$       6 &   24$\\pm$   2\\\\   \\hline \\end{tabular}  \\tablenotetext{a}{All column densities are for the CMZ only.} \\tablenotetext{b}{Using $\\zeta$ = 1 $\\times$ 10$^{-15}$~s$^{-1}$.}  \\normalsize \\end{center} \\end{table*}   \\subsection{The Case of Large $L$ in the CMZ}  \\indent \\citet{oka05} favored a large $L$ and hence a high volume filling factor for the newly found warm and diffuse molecular gas. This was not only because the $R$(3,3)$^l$ metastable lines with Large velocity widths and similar velocity profiles had been detected along all eight lines of sight toward sources scattered over 30~pc in projected distance from Sgr~A$^\\ast$, but also because the deduced $T$ and $n$ of the gas were similar in all cases. The latter suggests that the absorbing diffuse material is continuous rather than contained in several separate entities. Further support for this interpretation is that the range of H$_3^+$ absorption velocities in each sight line spans the entire range velocities observed at radio and millimeter wavelengths with large apertures and interpreted as coming from the front side of the CMZ. Extending spectroscopy of H$_3^+$ over more widely spaced sight lines in the CMZ will constitute a more definitive test of the filling factor. A long path length may also be favored over high ionization, in that there are limits to the value of $\\zeta$ for efficient production of H$_3^+$ as discussed in the next subsection.  $L$ is unlikely to be significantly larger than the radius of the CMZ, which for the region of interest is $\\sim$ 130 pc, as estimated using Fig.~11b of \\citet{saw04}. Setting $L$~=~100 pc, close to this maximum permissible value, the aforementioned range of values of ($\\zeta$$L$)$_{\\rm min}$ yields the lower limits $\\zeta$$_{min}$ = (0.4 -- 1.5)~$\\times$~10$^{-15}$~s$^{-1}$ for the various sight lines. For lower values of $f$ and higher values of $R_X$, $\\zeta$$_{min}$ will be higher.   The large extent of the warm and diffuse molecular gas postulated above conflicts with the previous concept of the gas in the CMZ, represented, for example, in Fig.~9 of \\citet{laz98} where the ultra-hot X-ray emitting plasma fills the CMZ. The warm and diffuse gas revealed in this paper cannot coexist with the ultra-hot plasma because the plasma electrons will be immediately cooled and molecules will be destroyed. Thus in the above model the volume filling factor of the ultra-hot gas would be reduced by a large factor.  The large extent of the ultra-hot plasma gas in the CMZ was proposed because of observations of extensive X-ray emission near the GC \\citep[e.g.,][]{koy89,yam93,koy96}. \\citet{sun93} and \\citet{mar93} observed similar X-ray emission, but interpreted it differently in view of the difficulty of confining such a hot gas by the gravitational potential of the Galactic center, and the extraordinarily large energy input required to maintain the high temperature. With the advent of the \\emph{Chandra} X-ray satellite, a great many X-ray point sources have been resolved in the GC. These have been interpreted as explaining all \\citep{wan02}, or a significant fraction \\citep{mun03}, of the observed intense X-rays as being due to stellar sources such as cataclysmic variables and young stellar objects. Additional observational reports also argue against the ultra-hot plasma based on the close correlation between the stellar distribution and the observed intensities of the X-rays in the GC \\citep [e.g.,][]{war06} and in the Galactic ridge \\citep [e.g.,][]{rev06}.  However, interpretation of the X-ray emission remains controversial \\citep[e.g.,][]{mun04,moN06,koy07}.  Yet another category of high temperature gas observed in the CMZ is the hot gas observed in recombination lines and other phenomena. \\citet{laz98} studied the $\\lambda^2$-dependent hyper-strong radio-wave scattering and free-free radio emissions and absorptions and interpreted the data as evidence for \\emph{T$_e$} $\\sim 10^6$ K and \\emph{n$_e$} $\\sim 10$ cm$^{-3}$ gas with nearly 100\\% surface filling factor toward the CMZ \\citep[see][for a different interpretation]{gol06}. \\citet{rod05} have observed fine structure and recombination lines toward a sample of 18 sources and reported even higher electron densities of 30 -- 100~cm$^{-3}$. \\citet{laz98} inferred that these regions were at interfaces between molecular clouds and the X-ray plasma and proposed (see their Fig.~9) that the high density molecular clouds with a volume filling factor of perhaps \\emph{f$_{\\rm v}$} $\\sim$ 0.1 are surrounded by the scattering electron gas, and the rest of the space is filled with X-ray emitting plasma. Our observations and analysis \\citep[see also][]{oka05} contradict such a picture \\citep{bol06}.    \\begin{figure*} \\begin{center} \\includegraphics[angle=-90,width=.8\\textwidth]{f4.eps} \\caption{Temperature and density plotted in filled circles as functions of population ratios $n$(3,3)/$n$(1,1) and $n$(3,3)/$n$(2,2) obtained by inverting the diagram of \\citet{oka04}. Temperature is indicated by thick white lines and density by thin lines. The estimated uncertainties in $n$(3,3)/$n$(1,1) are shown in error bars. Velocity resolved data of GCS~3-2 from \\citet{oka05} is shown in filled squares for comparison. \\label{f4}} \\end{center} \\end{figure*}   \\subsection{The Case of High $\\zeta$ in the CMZ}  It is usually assumed that the density of high energy ($E > 100$~MeV) cosmic rays is approximately uniform over the Galaxy because of their long mean free paths. The uniform distribution on the sky of diffuse $\\gamma$ ray emission appears to lend support to this \\citep[e.g.][]{die01}. Recent detailed calculations by \\citet{bue05} typically give a 20~\\% variation depending on location and time. However, the situation is very different for low-energy cosmic rays, which are the most important for ionization of the gas; their energy densities can vary greatly depending on proximity to local cosmic sources \\citep{kul71,spi75,ces75} because of their short mean free paths. For example, a proton of 2~MeV, which effectively ionizes hydrogen, travels only 3~pc before losing its energy in a medium with $n_H$ = 100~cm$^{-3}$ \\citep{cra78}.  The ionization rate by low-energy cosmic rays could be proportional to the number density of the accelerating sources \\citep[SNRs; e.g.][]{koy95} and therefore have strong local variations. If the energy density of cosmic-rays follows the surface density of OB stars, the cosmic ray ionization rate would be an order of magnitude higher in the Galactic center than in the solar neighborhood \\citep{wol03}.  \\citet{yus02} and \\citet{yus07} called for an even higher local enhancement of low-energy cosmic rays in the Galactic center ($\\zeta = 2~\\times~10^{-14}$ s$^{-1}$ to $5~\\times~10^{-13}$s$^{-1}$!) to account for the iron fluorescence line at 6.4~keV. Such values are much larger than those derived in the previous section using H$_3^+$. X-rays and far ultraviolet radiation from stars may also increase $\\zeta$. A factor of 10 or higher increase of $\\zeta$ due to X-ray emission has even been reported in the circumstellar disk of a low mass star \\citep{dot04}.  The enhancements could be large for the CMZ, where there are abundant intense emitters of X-rays and EUV radiation.  For the above high values of $\\zeta$, $L$ would be quite short if Eq.~(2) is applied. However, there are two arguments against adopting such high values of $\\zeta$. First, as argued earlier, the observed high velocity dispersion of H$_3^+$ indicates that the gas is quite extensive rather than confined in many small clumps. This indicates that there are non-local sources of ionization covering an extensive region.  Second, the results of Eqs.~(1) and (2) are based on a simple linear formalism, which is a good approximation for low $\\zeta$. For much higher $\\zeta$ this should be replaced by higher order non-linear formalism in which $N$(H$_3^+$) no longer increases proportionally to $\\zeta$, but tends to saturate. One source of such nonlinearity, revealed by \\citet{lis06,lis07}, is the dissociative recombination of H$_2^+$, H$_2^+$ + e$^-$ $\\rightarrow$ H + H, which competes with the H$_3^+$ production reaction H$_2^+$ + H$_2$ $\\rightarrow$ H$_3^+$ + H. The former reaction, in addition to the dissociative recombination of H$_3^+$, reduces the H$_3^+$ concentration twice. It also increases $n$(H) and decreases $f$, hence the non-linearity. In an example given in the lower figure of Fig.~3 of \\citet{lis06}, a ten times increase of $\\zeta$ from 10$^{-16}$ s$^{-1}$ to 10$^{-15}$ s$^{-1}$ leads to only a fourfold increase of H$_3^+$ (the saturation is more severe for the upper figure). The importance of saturation depends sensitively on the value of $f$, the fractional abundance of H$_2$. Another source of non-linear behavior is the production of H through dissociation of H$_2$ by cosmic rays, X-rays, and EUV radiation. The rates of these processes are comparable to the rate of ionization \\citep{gla73,gla74}. Since the reproduction of H$_2$ on dust grains is slow, H accumulates in the gas and the charge exchange reaction, H$_2^+$ + H $\\rightarrow$ H$_2$ + H$^+$, which has a high rate constant of 6.4~$\\times$~10$^{-10}$~cm$^3$~s$^{-1}$, will compete with the H$_3^+$ producing reaction introducing additional non-linearity. The increase of $n$(H) and resulting reduction of $f$ affect the gas in several other ways, all of which make the increase of $\\zeta$ ineffective in increasing $\\mathrm{H}_{3}^{+}$. While more detailed treatments of these nonlinearities have yet to be worked out, it seems unlikely to us that $\\zeta$ can be larger than $\\sim$ 3 $\\times$ 10$^{-15}$ s$^{-1}$ over extensive portions of the CMZ. Small volumes with higher $\\zeta$ may well exist in the CMZ, but they would not necessarily produce a large fraction of the observed H$_{3}^{+}$.    \\subsection{The $+$50~km~s$^{-1}$ absorption toward GC~IRS~3}  The sight line toward GC~IRS~3 is special in that at $+$50~km~s$^{-1}$ LSR it has a detectable population of H$_3^+$ in the unstable (2,2)~level in addition to the large populations in the (1,1) and (3,3) levels. Also this is the only case for which the population in the (3,3) metastable level is definitely higher than in the (1,1) ground level (see $N$(3,3)/$N$(1,1) in Table 4). The observable population in the (2,2) level indicates a high density and the larger population in the (3,3) level than in the (1,1) level indicates high temperature environment. Our analysis yields $n$~=~175~--~300 cm$^{-3}$ and $T$~=~350~--~500~K. As seen from Fig.~4, this environment is clearly distinct from those elsewhere. We speculate below on two possible locations for this gas: (1) the Circumnuclear Disk (CND) and, (2) the ``50~km~s$^{-1}$ cloud'' behind the GC.  \\subsubsection{Circumnuclear disk}  GC~IRS~3 is only 5\\arcsec~(0.2~pc) from Sgr~A$^\\ast$, which is well within the projected span of the Central Cluster \\citep[the maximum transverse separation of bright stars in the cluster is $\\sim$0.65~pc; ][]{vie05}. The central 3~pc of the GC contains hot, high density gas which shows intricate mini-spiral structures \\citep{eke83,lo83} termed the Western Arc, the Northern Arm, the Eastern Arm, and the Bar \\citep{gue87}. The extension of the Western Arc has been associated with a nearly completely circular, Keplerian circumnuclear disk by analysis of HCN radio emission \\citep{gue87,gen87,jac93,chr05} and by other methods; recent direct observations of the disk via its far infrared dust emission by \\citet{lat99} are particularly noteworthy. There also have been attempts to interpret other three mini-spirals as CNDs which are not necessarily elliptical \\citep[][e. g.]{lis85,qui85,lac91,jac93,lis03}. GC~IRS~3 is located in the cavity between the Northern Arm and the Bar, a region of low gas density. H. S. Liszt (private communication) has suggested that the \\ion{H}{1} absorption at $+$50~km~s$^{-1}$ in Fig.~5 of \\citet{lis85} may be due to the same cloud producing the H$_3^+$ absorption in the sightline toward GC~IRS~3. Not only does the velocity match, but also the high velocity gradient of 100~km~s$^{-1}$ arcmin$^{-1}$ found by \\citet{lis85} is consistent with the width of the H$_3^+$ 50~km~s$^{-1}$ absorption feature.   However, there are two possible problems with this interpretation. First, it calls for an uncomfortably high cosmic ionization rate in order to account for the observed H$_3^+$ column density in the short path length within the CND. Assuming a path length on the order of 0.5 pc, the observed ($\\zeta$$L$)$_{\\rm min}$ of 7.5 $\\times$ 10$^4$~cm~s$^{-1}$ in Table~\\ref{tb4} gives $\\zeta$~$\\sim$~2~$\\times$~10$^{-14}$~s$^{-1}$, an extremely high value.  Such a value is possible \\citep{yus07} in view of the high density of stars and radiation field in the region, but in such conditions the higher order chemistry \\citep{lis06,lis07} discussed in Section 5.3. surely plays a role. Although the detection of H$_3^+$ in the unstable (2,2) level indicates higher density compared with other sight lines, the population of the (3,3) level is still much higher and the distribution is highly non-thermal. The density is definitely much lower than that reported for the CND (10$^4$~--~10$^6$~cm$^{-3}$). This could be due to the fact that GC~IRS~3 is in the northern cavity and is only behind the edge of the CND rather than its densest regions.  The second problem is the absence of the $R$(2,2)$^l$ line in the spectrum of GC~IRS~1W. It is remarkable that the $R$(2,2)$^l$ line is detected toward GC~IRS~3, but not toward GC~IRS~1W which is just 8\\farcs4 apart, or 0.31~pc for a common distance of 7.6~kpc \\citep{eis05,nis06}. If both stars belong to the Central Cluster \\citep{bec78,vie05}, it is not likely that they are much more than 0.5~pc apart. Although the source might well be located in the foreground, GC~IRS~1W overlaps with the high density part of the Northern Arm, while GC~IRS~3 is located in the cavity. It is difficult To understand why only GC~IRS~3 shows $R$(2,2)$^l$ absorption, if the absorption occurs within the CND.    \\subsubsection{The``50~km~s$^{-1}$ cloud''}  The observed velocity of $+$50~km~s$^{-1}$ toward GC~IRS3 immediately suggests that the absorption might arise in the well known ``50~km~s$^{-1}$ cloud'', one of the giant molecular clouds ($> 10^6 M_\\odot$) located near the Galactic center. The ``50~km~s$^{-1}$ cloud'' is sometimes referred to as M$-$0.02$-$0.07, or was called the ``40~km~s$^{-1}$ cloud'' by \\citet{oor77}, and has been discussed in detail by \\citet{bro84}. Many atomic and molecular studies have established that it is located \\emph{behind} the Galactic nucleus, in between the non-thermal radio source Sgr~A East (SNR) and the \\ion{H}{2} region Sgr~A West, which enshrouds the central star cluster \\citep{whi74,gue83,pau96}. \\citet{yus96} argued that the several masers they discovered in the southeast of Sgr~A$^\\ast$ are excited in the shock-heated interface between Sgr~A~East and M$-$0.02$-$0.07, where the expanding supernova remnant hits the giant molecular cloud. \\citet{tsu06} also found clues for an interaction between Sgr~A~East and the ``50~km~s$^{-1}$ cloud'' in the high velocity wings of CS $J=1-0$, although it may not be the case that the ``50~km~s$^{-1}$ cloud'' is the sole cloud that falls in the radial velocity $+$50~km~s$^{-1}$ in the line of sight toward the Galactic center \\citep{zyl92}.  \\citet{geb89} found that GC~IRS~3 and IRS~7 show deep absorption lines of the CO fundamental band at 4.7~$\\mu$m at the radial velocity $+$50~km~s$^{-1}$, while GC~IRS~1 and IRS~2 do not. They attributed the $+$50~km~s$^{-1}$ absorption to the CND but also mentioned the possibility that GC~IRS~3 is not a member of the central cluster, but behind the $+$50~km~s$^{-1}$ cloud and thus not in the Central Cluster. The nature of GC~IRS~3 is controversial \\citep{pot08}; it is among the brightest sources in the Central Cluster region only less lumnous than GC~IRS~1 and GC~IRS~10, and much redder than them in the mid-infrared even though it is (apparently) in the cavity. Because of its location and featureless infrared continuum \\citep[e.g. ][]{kra95}, GC~IRS~3 has widely been believed to be a hot massive star heavily obscured by its own dust (similar to the bright Quintuplet sources), in addition to foreground dust. However, new interferometric observation that resolve the circumstellar dust shell have been interpreted by \\citet{pot08} as evidence that the central object is  a late-type carbon star, the photospheric spectral features of which are completely obscured. This interpretation is consistent with the apparent lack of ionization of the circumstellar shell, which one might expect if the object lies within the Central Cluster, and indeed the lack of ionized gas near GC~IRS~3 \\citep{roc85}. Although somewhat unlikely due to its proximity on the plane of the sky to the Central Cluster, the placement of GC~IRS~3 well behind the cluster and the ``$+$50~km~s$^{-1}$ cloud'' is otherwise a straightforward assignment and avoids the difficulties associated with the CND hypothesis mentioned in the previous section.    \\subsection{Infrared pumping}  During the 2005 Royal Society Discussion Meeting on H$_3^+$ \\citep{oka06a}, J. H. Black proposed an infrared pumping mechanism which can produce a non-thermal rotational level distribution of H$_3^+$ if the infrared flux is sufficiently high. The mechanism is illustrated in Fig.~\\ref{f5}. Infrared radiation pumps H$_3^+$ from the fully populated lowest ortho-level, (1,0), to the (2,0) level in the first excited vibrational state via the $R$(1,0) transition. The molecule decays to the (3,0) level in the ground state by spontaneous emission with a lifetime of 25.3~ms. The molecule further decays to the (3,3) metastable level through the forbidden rotational transition (3,0) $\\rightarrow$ (3,3) emitting a far-infrared photon at 49.62 $\\mu$m in 3.75 hours \\citep{pan86,nea96}.  Thus, H$_3^+$ molecules are pumped from the (1,0) level to the metastable (3,3) level.   \\begin{figure} \\begin{center} \\includegraphics[angle=0,width=.4\\textwidth]{f5.eps} \\caption{Schematic of the infrared pumping mechanism proposed by J. H. Black. H$_3^+$ in the lowest ortho-level (1,0) is excited to the (2,0) level in the first excited vibrational state via the $R$(1,0) transition (red line) by ambient photons. The excited H$_3^+$ spontaneously decays to the (3,0) ground level through the P(3,0) transition (orange line) and then to the metastable (3,3) level (shown in blue). See text for more details. Other transitions used in this paper are also shown. \\label{f5}} \\end{center} \\end{figure}    Since the spontaneous emissions are all very fast, the rate determining process for Black's mechanism is the rate of infrared pumping of the (1,0) level of H$_3^+$. We rewrite the excitation rate given in Eq.~(15) of \\citet{dis82} as \\[ W_{R(1,0)} = \\lambda^2A\\phi_\\nu/8\\pi = 1.51 \\times 10^{-7} \\phi_\\nu {\\rm s}^{-1}, \\]  \\noindent where $\\phi_\\nu$ is the photon flux of the ambient radiation at the wavelength of the transition in the units of photon cm$^{-2}$ s$^{-1}$ Hz$^{-1}$, and the values of $\\lambda$ = 3.6685~$\\mu$m, the Einstein coefficient $A$ = 98.5~s$^{-1}$, and the dilution factor of 0.286 are used. The dilution factor takes into account the spontaneous emission (2,0) $\\rightarrow$ (1,0) back to the ground level (1,1) with $A$ = 35.5 s$^{-1}$, resulting in a reduction of the pumping efficiency. In order for pumping to be significant, the value of $ W_{R(1,0)}$ has to be comparable to the collision rate, \\[ R = n\\sigma v = n k_L \\sim 2 \\times 10^{-7} {\\rm s}^{-1}, \\] \\noindent where we used a number density of $n$ = 100 cm$^{-3}$ and the temperature independent Langevin rate constant, 2~$\\times$~10$^{-9}$~cm$^3$~s$^{-1}$ for the H$_3^+$ + H$_2$ collision.  The energy density of the interstellar radiation field in the solar neighborhood at 4~$\\mu$m has been estimated by \\citet{por06} to be $\\lambda u_\\lambda$ = 0.055 $\\mu$m~eV~cm$^{-3}$~$\\mu$m$^{-1}$ which translates to $\\phi_\\nu$~=~6~$\\times$10$^{-5}$ photons~cm$^{-2}$~s$^{-1}$~Hz$^{-1}$. This is consistent with the value reported by \\citet{mat83} within a factor of 1.5. The corresponding pumping rate is $W_{R(1,0)}$~=~1~$\\times$~10$^{-11}$ s$^{-1}$, about 2$\\times$10$^{4}$ times lower than the collision rate. \\citet{por06} have estimated an average energy density about 20 times higher in the GC. \\citet{lau02} have reported an energy density in the central 120~pc (their Fig.~10) which is another factor of 10 higher. The corresponding pumping rate is still 100 times lower than the collision rate and the infrared pumping is perhaps not very effective for majority of the gas reported in this paper.  In the central parsec of the Galaxy, however, the stellar density is $\\sim$ $10^6$~$M_\\odot$~pc$^{-3}$ \\citep{sch07}, including a large number of luminous red giants and at least 80 hot and massive stars \\citep{pau06}. In this environment infrared pumping likely would overwhelm the collisional relaxation of any H$_3^+$ that is present. The effect, however, must be local to the very central region of the CMZ and probably does not seriously alter the population of the (3,3) level beyond a few tens of parsecs from the center, i.e., in the bulk of the CMZ. However, it may well be the best explanation for the strong $R$(3,3)$^l$ absorption line toward GC~IRS~3 at 50~km~s$^{-1}$ discussed in the previous subsection."
    },
    "0807/0807.3628_arXiv.txt": {
        "abstract": "{Infrared Dark Clouds (IRDCs), condensed regions of the ISM with high column densities, low temperatures and high masses, are suspected sites of star formation. Thousands of IRDCs have already been identified. To date, it has not been resolved whether IRDCs always show star formation activity and, if so, if massive star formation ($\\gtrsim8 M_\\odot$) is the rule or the exception in IRDCs.} {Previous analysis of sub-millimeter cores in the cloud MSXDC G048.65$-$00.29 (G48.65) indicates embedded star formation activity. To characterize this activity in detail, mid-infrared photometry (3--30~$\\mu$m) has been obtained with the \\textit{Spitzer Space Telescope}. This paper analyzes the point sources seen in the 24~$\\mu$m band, combined with counterparts or upper limits at shorter and longer wavelengths.} {Data points in wavelength bands ranging from 1~$\\mu$m up to $850~\\mu$m (\\textit{Spitzer} IRAC, MIPS~24 and 70~$\\mu$m, archival 2MASS data and sub-millimeter counterparts) are used to compare each 24~$\\mu$m source to a set of Spectral Energy Distributions of Young Stellar Object (YSO) models. By assessing the models that fit the data, an attempt is made to identify YSOs as such and determine their evolutionary stages and stellar masses.} {A total of 17 sources are investigated, 13 of which are classified as YSOs, primarily -- but not exclusively -- in an early embedded phase of star formation. The modeled masses of the central stellar objects range from sub-solar to $\\sim$8 $M_\\odot$. Every YSO is at less than 1 pc projected distance from its nearest YSO neighbor.} {IRDC G48.65 is a region of active star formation. We find YSOs in various evolutionary phases, indicating that the star formation in this cloud is not an instantaneous process. The inferred masses of the central objects suggest that this IRDC hosts only low to intermediate mass YSOs and none with masses exceeding $\\sim$8 $M_\\odot$.} ",
        "introduction": "Infrared dark clouds (IRDCs), discovered independently by \\citet{perault1996} and \\citet{egan1998} as dark patches against the Galactic mid-infrared background, are suspected to be sites of clustered star formation \\citep{rathborne2006,simon2006}. Moreover, the suggestion is raised \\citep[see e.g.][]{menten2005,beuther2007} that IRDCs harbor massive young stars. Morphologies of IRDCs range from globular to more filamentary structures. One of the most important properties of IRDCs is their high column density \\citep[$\\gtrsim10^{22}\\ \\mathrm{cm^{-2}}$, e.g.][]{carey2000}. Dust absorbs mid-infrared radiation and makes the clouds stand out as dark features against the mid-infrared background. Temperatures in IRDCs are $< 25 \\ \\mathrm{K}$ and volume densities can exceed $10^5 \\ \\mathrm{cm^{-3}}$ \\citep{egan1998,carey2000}. Typical size and mass scales are $\\sim$5 pc and $10^3$--$10^4~M_\\odot$ \\citep{simon2006}. These low initial temperatures, high molecular densities and total masses typical for IRDCs are precisely what make them suitable candidates for star forming regions. IRDCs are generally not quiescent: they harbor compact cores of sub-millimeter emission \\citep[e.g.][]{carey2000,ormel2005,rathborne2005}. IRDCs are found primarily in the inner Galaxy and close to the Galactic plane. Thousands are known in the first and fourth quadrants of the Galactic plane \\citep{simon_catalog2006}, i.e., the inner Galaxy. \\citet{frieswijk2007} recently identified the first IRDC in the outer Galaxy (in the absence of a bright mid-infrared background), using 2MASS (Two Micron All Sky Survey) color distributions of background stars, followed up by observations of molecular lines and \\textit{Spitzer Space Telescope} photometry.  It has not yet been established whether all IRDCs show active star formation or that some may in fact harbor only starless cores. Moreover, the association of IRDCs to \\emph{massive} ($\\gtrsim8 M_\\odot$) star formation in particular is still a matter of dispute. Studies of star forming regions in general \\citep[e.g.][]{indebetouw2007} and IRDCs in particular \\citep[e.g.][]{beuther2007} prove that star formation activity in a wide stellar mass range and at various phases of evolution can be probed by Spitzer photometry.  The IRDC under investigation in this paper is the cloud at Galactic coordinates $(\\ell,b)=(48\\fdg66, -0\\fdg30)$. It has been previously observed by the Mid-course Space Experiment (MSX), the SCUBA instrument on the James Clerk Maxwell Telescope \\citep[JCMT,][]{ormel2005}, and by the JCMT in CO, $^{13}$CO and HCO$^+$ \\citep{shipman2003}. In addition, it is covered by the Galactic Legacy Infrared Mid-Plane Survey Extraordinaire \\citep[GLIMPSE,][]{benjamin2003}. In the dark cloud catalog of \\citet{simon_catalog2006} this cloud has the designation ``MSXDC G048.65$-$00.29\", hereafter simply ``G48.65\".  Its distance is determined kinematically from molecular line data and an assumed Galactic rotation curve. G48.65 is found to be at a distance of $\\sim$2.5$\\ \\mathrm{kpc}$ \\citep{ormel2005,simon2006}. Its distance to the Galactic Center is $\\sim$7~$\\mathrm{kpc}$ and it is less than 20 pc away from the mid-plane of the Galaxy. Its total mass is estimated at almost 600\\,$M_\\odot$ within a 2 pc area and the molecular (H$_2$) density is $\\sim$$10^3 \\ \\mathrm{cm^{-3}}$ \\citep{simon2006}. \\citet{ormel2005} identified three distinct emission cores at 450 and 850 $\\mu$m; their modeling indicates the presence of central luminosity sources on the order of $10^2$--$10^3 \\,L_\\odot$ in at least two of these cores.  The earliest phases of star formation show emission at $>$30$~\\mu$m as they are embedded in an accreting envelope; the peak wavelength shifts through the mid- and near-infrared to the optical regime as stars evolve toward the main sequence. As mentioned above, sub-millimeter observations indicate ongoing embedded star formation in two or three cores in G48.65 \\citep{ormel2005}. This raises the question whether the star forming activity in G48.65 is limited to the sub-millimeter cores. Additionally, it is realized that observations at shorter wavelengths can improve our understanding of the previously identified star forming cores. This was the motivation to take a closer look in the mid-infrared (roughly 3--30\\,$\\mu$m), where Young Stellar Objects (YSOs) are known to emit the bulk of their energy. This paper describes the interpretation of \\textit{Spitzer Space Telescope} photometry of G48.65, focussing on those sources that show emission at 24\\,$\\mu$m, where the brightest, massive, embedded young stars are expected to be found.  Section \\ref{sec:observations} describes the Spitzer data and their reduction, Section \\ref{sec:results} deals with the classification of twenty sources based on the available data points in the near- and mid-infrared and sub-millimeter regimes. Conclusions and a discussion are presented in Sect.~\\ref{sec:conclusions_and_discussion}. ",
        "conclusions": "\\label{sec:conclusions_and_discussion}  \\subsection{Conclusions} \\label{sec:conclusions}  Of the 20 sources near IRDC G48.65 that are visible at 24\\,$\\mu$m, a total of 13 are classified as YSOs, seven of which are found to be in Stage I, two are in Stage II, two more are in either Stage I or Stage II, and two are uncertain. While all reliably classified YSOs lie along the IRDC extinction filament, the four sources matched to photosphere models all lie away from the filament. Each YSO generally resides within a projected distance of $<$1\\,$\\mathrm{pc}$ from another YSO. We conclude that G48.65 is an example of a dark cloud environment forming a group of stars.  The most important modeled properties of the 24\\,$\\mu$m sources that are classified as YSOs are summarized in Table \\ref{table:allYSOs}. The stellar masses of the modeled objects range from slightly sub-solar to $\\sim$8\\,$M_\\odot$. The inferred evolutionary Stages are predominantly early phases, ranging from early Stage I to late Stage II. This is believed to be caused partly by the selection effect discussed in Sect.~\\ref{sec:MIPSselection}. Therefore, the presented range of evolutionary phases should not be interpreted as a full picture of all star forming activity in this IRDC.  The emission peak P1 identified by \\citet{ormel2005} is resolved by Spitzer into two distinct emission cores: S6 and S5. Both objects are found to be in the earliest Stage of star formation. Stellar masses are not well constrained for these sources but are unlikely to be in excess of $8\\,M_\\odot$. The summed total luminosity of the best-fit models for S5 and S6 is $\\sim$$10^2$--$10^3\\,L_\\odot$, consistent with results from modeling by \\citet{ormel2005} based purely on sub-mm observations. The total luminosity of the best fitting model for objects S12 (EP) and S15 (P2) are consistent with values found by \\citet{ormel2005}.   \\subsection{Robustness of Results} \\label{sec:discussion} It is important to assess a YSO based on data points in a wavelength regime as extended as possible. While a detection of a certain YSO at sub-millimeter wavelengths seems to be a strong indication that it is in Stage I (S5, S6, S12, S15), a detection by 2MASS in the near-infrared does not rule out a Stage I classification, as sources S2 and S7 show. It may therefore lead to erroneous classification of YSOs if one bases it either on only near-infrared or on only far-infrared/sub-millimeter data.  The availability of data points in the 3--30\\,$\\mu$m regime is in some cases enough to constrain the evolutionary stage of a YSO. Source S8 (see Fig.~\\ref{fig:YSOSEDs} for the SED fit) is an example in which the degeneracy of the model fitting is decreased considerably by adding 2MASS data points (see also Sect.~3.3 and Fig.~3 of \\citet{robitaille2007}).  \\subsection{Continuous Low-mass Star Formation}  The Stage I sources in particular show large uncertainties in stellar mass (see Table~\\ref{table:allYSOs}). This is not surprising considering that, for objects that emit mostly through re-radiation from the accreting envelope, it is of course inherently difficult to model parameters of the obscured central heating source. However, the OH maser survey by \\citet{pandian2007} indicates the absence of OH maser emission in the direction of IRDC G48.65. Since OH masers are commonly associated with high-mass star formation ($\\gtrsim$8\\,$M_\\odot$), this agrees with the fact that we find no massive YSOs in G48.65. Considering the above, we are confident that any YSO associated to G48.65 is unlikely to harbor a central star more massive than 8\\,$M_\\odot$.  Based on the range of stellar masses and evolutionary Stages of the YSOs, we conclude that G48.65 has been an active low-mass star forming region for at least $\\sim$$10^6$ years (late Stage II) up to as recent as $\\sim$$10^4$ years ago (early Stage I) and is likely to be in the process of forming still younger stars. Assuming that the IRDC is indeed the birth place of these young stars, the cloud itself must have been stable over a time scale of at least $\\sim$$10^6$ years. This is consistent with modern views on lifetimes of $\\lesssim$$10^6$ years for molecular clouds \\citep[as reviewed in][]{maclow2004} and with lifetimes of $\\lesssim$$10^7$ years for dense cluster forming clouds \\citep[see review by][]{larson2003}.   \\subsection{Mid-infrared Selection Effect} % \\label{sec:MIPSselection} The starting point of the study in this paper -- the set of point sources in the MIPS 24\\,$\\mu$m image -- introduces an observational bias toward the younger evolutionary stages. Early stage YSOs are simply brighter in the mid-infrared, since the total mid-infrared luminosity of a YSO generally decreases with time. Especially in a fairly distant region such as IRDC G48.65, at $\\sim$2.5\\,kpc, one must keep in mind that later stages and intrinsically fainter sources are less likely to be detected. This study in particular is biased toward the earlier stages, since it focusses on the 24\\,$\\mu$m objects.  It is primarily the relatively cold envelope component, dominant in Stage I objects, that gives a YSO SED its large ``bump\" longward of 20~$\\mu$m. Hence, an analysis of objects near the IRDC that are only visible at wavelengths shortward of 8~$\\mu$m is expected to lead to the identification of more evolved (Stages II and III) and less massive YSOs. The IRAC 8~$\\mu$m image, for example, shows at least a factor 3 more point sources than the 24~$\\mu$m image in the same field of view. Such a study will be essential in getting a complete census of the star formation activity in G48.65, both in terms of time scales and in terms of a stellar mass function.  In principle, the set of young stars of various masses presented in this paper could be used to construct an initial mass function (IMF). Extrapolation of such an IMF could provide insight into the absence of massive stars ($>$8\\,$M_\\odot$). However, we refrain from presenting such an exercise here. The first reason for this is the incompleteness of our sample discussed above; a second reason is that the inferred stellar mass of each individual YSO (see Col.~(5) of Table~\\ref{table:allYSOs}) is rather poorly constrained.  \\subsection{Outlook} \\label{sec:outlook}  In the near future, the 2MASS point source catalog (1--3\\,$\\mu$m) and the GLIMPSE catalog (3--9.5\\,$\\mu$m) will be complemented by the MIPSGAL point source list of the Galactic plane at 24 and 70\\,$\\mu$m. Moreover, the UKIRT Infrared Deep Sky Survey \\citep{lawrence2007} is already partly available to the general community and will eventually cover half of the Galactic plane in the $J$, $H$ and $K$ bands, superseding the sensitivity of 2MASS by roughly 3 magnitudes. This will open up the possibility to investigate YSOs in hundreds of other IRDCs in the Galaxy using an approach similar to the one taken in this paper. This approach has shown that source-by-source investigation of emission cores with 1--30\\,$\\mu$m data can lead to the identification of young stars in various phases of evolution. Perhaps positions that show sub-millimeter emission are the most interesting cases to start with if one is attempting to identify very early stages of star formation. However, even with the combination of data from 2MASS, GLIMPSE and MIPSGAL, covering just the 1--70\\,$\\mu$m regime, it is possible to systematically study star formation activity across the Galactic plane. IRDCs may eventually be divided into groups according to whether they show star formation activity, and if they do, whether they form low-mass stars or high-mass stars.  The Stage I sources in Table \\ref{table:allYSOs} are expected to show associated outflows and shocks, traced by e.g. SiO. Detection of such tracers could confirm the star forming activity and reveal kinematic and spatial structure of YSOs. In addition, measurements by the VISIR instrument (8--13\\,$\\mu$m) at VLT will be able to resolve more detailed structure ($<$$10^3\\ \\mathrm{AU}$ at 2.5\\,kpc) of the sources and show envelope and disk components. It may even show that objects that appear as one source in the Spitzer images are in fact multiple YSOs. Source S5, for example, appears slightly elongated in the Spitzer images. In fact, some of the poorer model fits (see Col.~(4) in Table~\\ref{table:allYSOs}) may be explained by the fact that the models do not account for multiple young star-disk-envelope systems inside one telescope beam.  The high spatial and spectral resolution and new wavelength regimes of future observatories, such as ALMA, the Herschel Space Observatory \\citep[specifically the Hi-GAL survey,][]{noriega-crespo2008} and the James Webb Space Telescope, will enable more detailed studies of individual objects associated with G48.65 and other (distant) IRDCs."
    },
    "0807/0807.0301_arXiv.txt": {
        "abstract": "{ Due to its proximity, youth, and solar-like characteristics with a spectral type of K2V, $\\epsilon$ Eri is one of the most extensively studied systems in an extrasolar planet context. Based on radial velocity, astrometry, and studies of the structure of its circumstellar debris disk, at least two planetary companion candidates to $\\epsilon$ Eri have been inferred in the literature ($\\epsilon$ Eri b, $\\epsilon$ Eri c). Some of these methods also hint at additional companions residing in the system. Here we present a new adaptive optics assisted high-contrast imaging approach that takes advantage of the favourable planet spectral energy distribution at 4 $\\mu$m, using narrow-band angular differential imaging to provide an improved contrast at small and intermediate separations from the star. We use this method to search for planets at orbits intermediate between $\\epsilon$ Eri b (3.4 AU) and $\\epsilon$ Eri c (40 AU). The method is described in detail, and important issues related to the detectability of planets such as the age of $\\epsilon$ Eri and constraints from indirect measurements are discussed. The non-detection of companion candidates provides stringent upper limits for the masses of additional planets. Using a combination of the existing dynamic and imaging data, we exclude the presence of any planetary companion more massive than 3 $M_{\\rm jup}$ anywhere in the $\\epsilon$ Eri system. Specifically, with regards to the possible residual linear radial velocity trend, we find that it is unlikely to correspond to a real physical companion if the system is as young as 200 Myr, whereas if it is as old as 800 Myr, there is an allowed semi-major axis range between about 8.5 and 25 AU. ",
        "introduction": "\\label{sec_intro}  With its distance of 3.2 pc and spectral type K2V, $\\epsilon$ Eri is the most nearby solar-like single star. Largely due to this fact, it has been targeted by several attempts to indirectly or directly detect planetary companions. So far, two exoplanet candidates have been reported as companions to $\\epsilon$ Eri -- $\\epsilon$ Eri b (1.55 $M_{\\rm jup}$, $a=3.4$ AU), which was inferred from a periodic radial velocity signal by Hatzes et al. (2000) and subsequently also from astrometric motion of the primary (Benedict et al. 2006), and $\\epsilon$ Eri c (0.1 $M_{\\rm jup}$, 40 AU), which results from an interpretation of observed irregularities in the dust disk surrounding the primary (Quillen \\& Thorndike, 2002). Targeted attempts have been made to directly image both $\\epsilon$ Eri b (Janson et al. 2007) and $\\epsilon$ Eri c (Macintosh et al. 2003, Marengo et al. 2006), but did not yield any detections. $\\epsilon$ Eri has also been observed as a part of large-scale near-infrared adaptive optics surveys, such as Lafreniere et al. (2007), and Biller et al. (2007).  Typically, direct imaging surveys from the ground are performed in narrow-band filters around 1.6 $\\mu$m, since atmospheric opacities cause cool objects (giant planets and brown dwarfs) with temperatures below 800 K to emit a large fraction of their H-band flux just shortward of the CH$_4$ absorption feature at 1.6 $\\mu$m, and since this feature also allows for spectral differential imaging (SDI) to enhance the image contrast. However, the primary-to-secondary brightness contrast is generally smaller at longer wavelengths, such that an improved performance can be gained for some types of high-contrast systems. This was demonstrated with an imaging survey in the L'-band of 22 young, nearby stars by Kasper et al. (2007). The brightness difference is even smaller in M-band (see e.g. Hinz et al. 2006), but here the noise contribution of the thermal background is also significantly higher, making it observationally challenging. Here we examine an intermediate approach, where we use a narrow-band filter towards the red end of the L'-band range, which simultaneously allows for both good brightness contrast and observational simplicity. We test this technique on the $\\epsilon$ Eri system to search for planetary companions at separations intermediate between $\\epsilon$ Eri b and $\\epsilon$ Eri c.  The paper is organized as follows: In Sect. \\ref{sec_obsstrat}, we describe the basic principles of the observing strategy chosen, and in Sect. \\ref{sec_obs} we describe the actual observations performed, in terms of the techniques used and the observing conditions and corresponding instrument performance. This is followed in Sect. \\ref{sec_datared} by a description of the data reduction, and the results as given in Sect. \\ref{sec_contrast}. Our choice of filter is further motivated in Sect. \\ref{sec_comparison}, where it is also explained over which parameter range the method should be preferable to other common methods in high-contrast imaging. In  Sect. \\ref{sec_age}, we discuss the controversial issue of the age of $\\epsilon$ Eri, which has fundamental implications for the detectability of planetary companions in the system. In Sect. \\ref{sec_masslimits} the brightness contrast is translated into a mass contrast. We examine the possible clues for an intermediate companion in Sect. \\ref{sec_dyn}. Finally, we summarize what these issues mean in terms of mass-equivalent detection limits for exoplanet companions to $\\epsilon$ Eri in Sect. \\ref{sec_limits}, and we conclude in Sect. \\ref{sec_conclusion}. ",
        "conclusions": "\\label{sec_conclusion}  We have performed a deep, high-contrast imaging observation of $\\epsilon$ Eri with a new method that allows for detection of lower-mass planets at small separations from bright stars than previously possible. The method is based on ADI in the NB 4.05 $\\mu$m filter, and compares favourably to other methods in the literature. Its advantage is however limited to very bright stars, since the background limit will dominate the contrast limit for faint stars, such that other methods (e.g. SDI+ADI or L'-band ADI) are preferable for such targets. As the final output images show no signs of any convincing companion candidate, we place constraints on any planets that may reside in the system. In particular, we test the hypothesis of the existence of an intermediate orbit planet possibly indicated by radial velocity, Hipparcos $\\Delta \\mu$ astrometry and numerical attempts to reproduce the circumstellar debris disk structure of $\\epsilon$ Eri. In order to emphasize the hypothetical nature of such a companion, it is referred to as $\\epsilon$ Eri x. We find that, due to the uncertain age of the $\\epsilon$ Eri system, which we quote as 200-800 Myr based on a detailed investigation of the existing literature, the presence of $\\epsilon$ Eri x can not be excluded based on our images, but its allowed parameter range can be constrained. In general, a combination of the existing RV data and high-contrast images allows us to conclude that whatever the amount and distribution of planets is in the $\\epsilon$ Eri system, none of them can be more massive than 3 $M_{\\rm jup}$."
    },
    "0807/0807.0137_arXiv.txt": {
        "abstract": "% We present the results of Spectral Energy Distribution (SED) fitting analysis for Lyman Break Galaxies (LBGs) at $z\\sim5$ in the GOODS-N and its flanking fields. From the SED fitting for $\\sim100$ objects, we found that the stellar masses range from $10^{8}$ to $10^{11}M_{\\odot}$ with a median value of $4\\times10^{9}M_{\\odot}$. By using the large sample of galaxies at $z\\sim5$, we construct the stellar mass function (SMF) with incompleteness corrections. By integrating down to $10^{8}M_{\\odot}$, the cosmic stellar mass density at $z\\sim5$ is calculated to be $7\\times10^{6}M_{\\odot}\\textrm{Mpc}^{-3}$. ",
        "introduction": "Recent observations show the gradual increase of the stellar mass density of the universe with time \\citep[e.g.,][]{wilkins08}. However, the studies on the stellar mass of galaxies at $z\\ga5$ are restricted because the lack of sufficiently deep mid-infrared data. With the advent of Spitzer, we can access the rest-frame optical wavelength, and thus, the stellar mass of galaxies at $z\\sim5$. In this work, we explore the stellar mass of Lyman Break Galaxies (LBGs) at $z\\sim5$ using the Subaru/S-Cam imaging data and the Spitzer/IRAC data. ",
        "conclusions": ""
    },
    "0807/0807.2242_arXiv.txt": {
        "abstract": "} \\nc{\\eab}{ We extend the previously described CMB Gibbs sampling framework to allow for exact Bayesian analysis of anisotropic universe models, and apply this method to the 5-year WMAP temperature observations. This involves adding support for non-diagonal signal covariance matrices, and implementing a general spectral parameter MCMC sampler. As a worked example we apply these techniques to the model recently introduced by Ackerman et al., describing for instance violations of rotational invariance during the inflationary epoch. After verifying the code with simulated data, we analyze the foreground-reduced 5-year WMAP temperature sky maps. For $\\ell \\le 400$ and the W-band data, we find tentative evidence for a preferred direction pointing towards $(l,b) = (110^{\\circ}, 10^{\\circ})$ with an anisotropy amplitude of $g_* = 0.15 \\pm 0.039$. Similar results are obtained from the V-band data [$g_* = 0.10\\pm 0.04$; $(l,b) = (130^{\\circ}, 20^{\\circ})$]. Further, the preferred direction is stable with respect to multipole range, seen independently in both $\\ell=[2,100]$ and $[100,400]$, although at lower statistical significance. We have not yet been able to establish a fully satisfactory explanation for the observations in terms of known systematics, such as non-cosmological foregrounds, correlated noise or asymmetric beams, but stress that further study of all these issues is warranted before a cosmological interpretation can be supported. ",
        "introduction": "\\label{secintroduction}  Since the early 1990's, great advances have been made in the field of data analysis techniques for studying the cosmic microwave background (CMB). Observations of the CMB anisotropies, for instance those made by the Wilkinson Microwave Anisotropy Probe (WMAP) experiment \\citep{bennett:2003, hinshaw:2007}, provides the single most powerful probe in contemporary cosmology.  From these, various theoretical universe models may be constrained, and today an effective concordance model based on the inflationary $\\Lambda$CDM framework has been established.  The theory of inflation was initially proposed as a solution to the horizon and flatness problem \\citep{guth:1981}. Additionally, it established a highly successful theory for the formation of primordial density perturbations, thus providing the required seeds for the large-scale structures (LSS), later giving rise to the temperature anisotropies in the cosmic microwave background radiation that we observe today \\citep{starobinsky:1980, guth:1981,linde:1982, muhkanov:1981, starobinsky:1982, linde:1983, linde:1994, smoot:1992, ruhl:2003, runyan:2003, scott:2003}.  A firm prediction of inflation is that the observed universe should be nearly isotropic on large scales. Yet, recent theoretical studies have demonstrated that anisotropic inflationary models are indeed conceivable \\citep{armendariz:2006, contaldi:2007, pullen:2007, kanno:2008, yokoyama:2008}. Two other examples are those presented by \\citet{ackerman:2007} (ACW) and \\citet{erickcek:2008}. The first model considers violation of rotational invariance in the early universe, while the second model describes the effects on the observed perturbation distribution due to a large-scale curvaton field.  The introduction of anisotropic models poses several problems in terms of data analysis. The definition of a proper likelihood function may be non-trivial for a general case, although many models can be described as multivariate Gaussians with non-diagonal covariance matrices. All models mentioned above are examples of this. Yet, even in these relatively simple cases, the numerical evaluation of the likelihood is computationally unfeasible due to the sheer size of the relevant covariance matrix.  In the present paper, we extend the previously described CMB Gibbs sampling framework \\citep{jewell:2004, wandelt:2004, eriksen:2004b} to allow for non-diagonal, but sparse, covariance matrices. As currently described in the literature, this framework allows for exact Bayesian analysis of high-resolution CMB data, but only under the assumption of isotropy, i.e., a diagonal CMB covariance matrix. This method has already been applied several times to the WMAP data \\citep{odwyer:2004, eriksen:2007a, eriksen:2007b, eriksen:2008b}, and has been extended to take into account both polarization \\citep{larson:2007} and internal component separation \\citep{eriksen:2008a}.  The question of isotropy has received considerable attention during recent years, due to unexpected signatures observed in the WMAP sky maps. These data appear to exhibit several significant and distinct signatures of violation of statistical isotropy. First, \\citet{de Oliveira-Costa:2004} found a striking alignment between the two largest harmonic modes in the temperature anisotropy sky, the quadrupole and the octopole. Second, \\citet{vielva:2004} pointed out the presence of a very large cold spot in the southern Galactic sky, apparently incompatible with $\\Lambda$CDM-based simulations. Finally, \\citet{eriksen:2004a} found a significantly anisotropic distribution of power between two hemispheres. The tools developed in the present paper may be able to constrain specific models relevant for these observations. In particular, we use these methods to estimate the anisotropy parameters in the ACW model from the 5-year WMAP temperature data.  The paper is structured as follows: In \\S \\ref{sec:acwmodel}, we review the ACW universe model, and briefly introduce the relevant posterior distribution. Next, we present the method in \\S\\ref{sec:method}, before we apply our tools to simulated data in \\S\\ref{sec:simulations}. In \\S\\ref{sec:wmap} we analyze the five-year WMAP temperature sky maps. Finally, we conclude in \\S\\ref{sec:conclusions}. ",
        "conclusions": "\\label{sec:conclusions}  We have generalized a previously described CMB Gibbs sampler to allow for exact Bayesian analysis of any anisotropic universe models that predicts a sparse signal harmonic space covariance matrix. This generalization involved incorporation of a sparse matrix library into the existing Gibbs sampling code called ``Commander'', and implementation of a new sampling algorithm for the anisotropy parameters given a sky map, $P(\\theta|\\mathbf{s})$.  We then considered a special case of anisotropic universe models, namely the \\citet{ackerman:2007} model which generalizes the primordial power spectrum $P(k)$ to include a dependence on direction, $P(\\mathbf{k})$. Explicit expressions for the resulting covariance matrix is provided in their paper.  We implemented support for this model in our codes, and demonstrated and validated the new tools with appropriate simulations. First, we compared the results from the Gibbs sampler with brute-force likelihood evaluations, and then verified that the input parameters were faithfully reproduced in realistic WMAP simulations.  Finally, we analyzed the five-year WMAP temperature sky maps, and presented for the first time the WMAP posteriors of the ACW model. The results from this analysis are highly intriguing, but we emphasize that the effect of instrumental systematics, particularly in the form of correlated noise, must be better understood before the findings can be given a cosmologically interpretation.  Taken at face value, we find a preferred direction in the W-band WMAP temperature data pointing towards $(l,b) = (110^{\\circ}, 10^{\\circ})$ (Galactic longitude and latitude), with an anisotropy amplitude of $g_* = 0.15\\pm 0.039$, formally corresponding to a $3.8\\sigma$ detection of $g_* \\ne 0$. Similar results for $g_*$ are found for the V-band data, although with a somewhat lower significance ($g_* = 0.10 \\pm 0.04$; $2.5\\sigma$). The preferred direction is very stable with respect to both data set and multipole range. Figure \\ref{figanisotropicmap} illustrates the underlying anisotropic contribution for a simulation with parameters corresponding to the W band posterior.  We have not been able to identify a plausible explanation for this effect in terms of known systematics. First, foregrounds do not appear to have much impact on the results, as consistent results are obtained both from foreground-corrected and raw maps. Second, although correlated noise does lead to a signature similar to the ACW model, its amplitude appears too low in the 5-year data. The least well constrained possibility is that of asymmetric beams, for which we lack proper simulations.  \\begin{figure*} \\includegraphics[width=170mm]{f12.eps} \\caption{A simulated realization drawn from a Gaussian distribution with zero mean and a covariance matrix given by the anisotropic $\\Delta$ term in the ACW model, computed for an asymmetry amplitude of $g_* = 0.14$ and a preferred direction $(l,b) = (110^{\\circ}, 10^{\\circ})$, marked by red dots. Notice the rotational structure about the preferred direction. The amplitude of the anisotropic component is $\\sim\\pm 15\\mu\\textrm{K}$, or $\\sim3$\\% of the isotropic component.} \\label{figanisotropicmap} \\end{figure*}    While this particular signature certainly is highly intriguing, we would like to point out that the main purpose of this paper is the demonstration of a general framework for analyzing anisotropic signal models. This is useful both for studying particular universe models (e.g., the ACW model), but also for understanding systematic effects (e.g., correlated noise) in a given data set. We therefore believe that these methods may be useful in a wide range of applications, only some of which have been demonstrated in this paper."
    },
    "0807/0807.2060_arXiv.txt": {
        "abstract": "We report on a systematic investigation of the \\hb\\ and \\feii\\ emission lines in a sample of 568 quasars within $z < 0.8$ selected from the Sloan Digital Sky Survey.  The conventional broad \\hb\\ emission line can be decomposed into two components---one with intermediate velocity width and another with very broad width. The velocity shift and equivalent width of the intermediate-width component do not correlate with those of the very broad component of \\hb, but its velocity shift and width do resemble \\feii.  Moreover, the width of the very broad component is roughly 2.5 times that of the intermediate-width component. These characteristics strongly suggest the existence of an intermediate-line region, whose kinematics seem to be dominated by infall, located at the outer portion of the broad-line region. ",
        "introduction": "The geometry and kinematics of the broad-line region (BLR) in active galactic nuclei (AGNs) have been studied for about three decades but the details are far from well understood.  It is widely accepted that the BLR is stratified: high-ionization lines originate from small radii and low-ionization lines arise further out \\citep{collin88}. This stratification picture is supported by the results of reverberation mapping, which show that lines of different ionization have different lags (e.g., \\citealt{peterson99}). The dependence of the systemic velocities of the emission lines on ionization \\citep[e.g.,][]{gaskell82,sulentic00a,richards02,shang07} suggests that the BLR may originate from a wind and a disk \\citep[e.g.,][and references therein]{leighly04a,leighly04b}.  However, the profiles of the broad emission lines often contain multiple velocity components, suggesting that the structure of the BLR may be more complex than can be described by a simple stratification or wind $+$ disk model.  It has been known that the broad \\hb\\ line profiles are generally not well described by a single Gaussian. Two Gaussians \\citep[e.g.,][]{netzer07} or a Gauss-Hermite function \\citep[e.g.,][]{salviander07} are often used. Additionally, the profiles show great diversity from object to object.  Sources with narrower \\hb\\ lines tend to have stronger line wings, while those with broader \\hb\\ lines are dominated by the line core \\citep[e.g.,][]{sulentic02}. Some sources have asymmetric and shifted \\hb\\ profiles, suggesting that the BLR has a structure more complex than a single virialized component. The \\hb\\ profile of OQ 208, for example, has an additional redshifted \\hb\\ component of intermediate width that closely resembles the kinematics of \\feii\\ \\citep{marziani93}. Additional evidence that the BLR contains two or more kinematically distinct components comes from differential variability between the line core and wing \\citep[e.g.,][]{ferland90,peterson00}.  The existence of an intermediate-line region and a very broad-line region for \\hb\\ emission was suggested by some previous studies \\citep[e.g.,][]{brotherton96,sulentic00b}.  In a recent spectral decomposition of a large sample of quasars selected from the Sloan Digital Sky Survey (SDSS), Hu et al. (2008, hereinafter Paper I) find that the majority of quasars show \\feii\\ emission that is both redshifted and narrower than \\hb.  Moreover, the magnitude of the \\feii\\ redshift correlates inversely with the Eddington ratio.  These characteristics suggest that \\feii\\ originates from an exterior portion of the BLR, whose dynamics may be dominated by infall.  These findings offer fresh insights into the structure of the BLR.  In light of the trends associated with \\feii\\ emission, we expect that a portion of the \\hb-emitting gas may be related to an inflowing component too. In this Letter, we systematically study the \\hb\\ profiles of SDSS quasars to try to answer two questions: do all quasars have an intermediate-width \\hb\\ component similar to OQ 208, and, if so, is this component also associate with the \\feii\\ emission?  We show that the conventional broad \\hb\\ line actually consists of two kinematically linked components, one of which originates from the same region that emits \\feii. ",
        "conclusions": "\\begin{figure} \\includegraphics[angle=-90,width=0.47\\textwidth]{f3.eps} \\caption{\\footnotesize Correlations between ({\\it a}) \\feii\\ and \\hbi\\ shifts and ({\\it b}) \\feii\\ and \\hbi\\ widths. The solid diagonal lines denote that \\feii\\ and \\hbi\\ have the same shifts and widths. The contours in panel ({\\it a}) show the density of the data points.} \\label{fig-hbife} \\end{figure}  \\begin{figure} \\includegraphics[angle=-90,width=0.47\\textwidth]{f4.eps} \\caption{\\footnotesize Plots of ({\\it a}) \\feii\\ vs. \\hbvb\\ shifts and ({\\it b}) \\feii\\ vs. \\hbvb\\ widths. The two lines have different shifts and widths.} \\label{fig-hbvbfe} \\end{figure}  We find that \\hbi\\ and \\feii\\ emission have similar kinematics. The similarity can be seen not only in individual sources, but also statistically for the whole sample.  From examples {\\it a} to {\\it c} in Figure \\ref{fig-example}, the \\hb\\ profiles change progressively while the \\feii\\ shifts become lower and lower. Figure \\ref{fig-hbife}{\\it a} shows a strong correlation between \\hbi\\ and \\feii\\ shifts. Pearson's correlation coefficient $r_{\\rm P}$ is 0.22, and the probability $P$ of a chance correlation is $< 1\\times10^{-5}$. \\hbi\\ and \\feii\\ have approximately the same shifts. The widths of \\hbi\\ and \\feii\\ are also well correlated and roughly equal (Fig. \\ref{fig-hbife}{\\it b}); $r_{\\rm P}$ = 0.48 and $P < 1\\times10^{-5}$. Except for some sources to the upper left of the solid lines whose errors are large, the majority of sources follow the relation that \\hbi\\ and \\feii\\ have the same shifts and widths. By contrast, \\feii\\ and \\hbvb\\ have different shifts and widths (Fig. \\ref{fig-hbvbfe}). We also find that the kinematic connection between \\hbi\\ and \\hbvb\\ is complicated.  No correlation between \\hbi\\ and \\hbvb\\ shifts is seen (Fig. \\ref{fig-hbihbvb}{\\it a}), but the \\hbi\\ and \\hbvb\\ widths are strongly linked, such that FWHM(\\hbvb) $\\approx$ 2.5 FWHM(\\hbi) (Fig. \\ref{fig-hbihbvb}{\\it b}).  \\begin{figure} \\includegraphics[angle=-90,width=0.47\\textwidth]{f5.eps} \\caption{\\footnotesize Plots of ({\\it a}) \\hbvb\\ vs. \\hbi\\ shifts and ({\\it b}) \\hbvb\\ vs. \\hbi\\ widths. The two components have different shifts and widths. The solid line in panel ({\\it b}) denotes FWHM(\\hbvb) $= 2.5$ FWHM(\\hbi).} \\label{fig-hbihbvb} \\end{figure}  The observational results described above suggest a scenario in which the conventional BLR consists of two components---an intermediate-line region (ILR) and a very broad-line region (VBLR). If both regions are virialized, so that $R \\propto v^{-2}$, then the ILR is about $2.5^2$ or 6.25 times farther from the center than the VBLR.  Because of its redshift, the kinematics of the ILR may be dominated by infall. \\hb\\ emission emerges from both the ILR and VBLR, while most of the \\feii\\ emission comes from the ILR.  The ILR and VBLR defined in this paper are essentially similar to those described in \\citet{corbin95}, \\citet{brotherton96}, \\citet{sulentic00b}, and \\citet{zhu08} but differ from those in \\citet{brotherton94} or \\citet{sulentic99}, which refer to the \\civ-emitting region.  Note that the ILR and VBLR of the \\hb-emitting region are distinct from those of the \\civ-emitting region because the \\civ\\ ILR usually has the systemic redshift while the \\civ\\ VBLR shows a blueshift.  \\begin{figure} \\includegraphics[angle=-90,width=0.47\\textwidth]{f6.eps} \\caption{\\footnotesize Plots of ({\\it a}) \\hbvb\\ vs. \\hbi\\ EWs and ({\\it b}) \\feii\\ vs. \\hbi\\ EWs. No correlations are found.} \\label{fig-ew} \\end{figure}  The lack of correlation between the EWs of \\hbi\\ and \\hbvb\\ (Fig. \\ref{fig-ew}{\\it a}) strongly suggests that the two components are emitted from different regions. If both are photoionized, they must have different covering factors. The relative strength between \\hbi\\ and \\hbvb\\ determines the final \\hb\\ profiles.  In Figure \\ref{fig-hbwfp}, sources broader than $\\sim$5000 \\kms\\ can be fitted well using one Gaussian. This trend can be interpreted in our two-component BLR scenario. It is well known that sources with broad \\hb\\ tend to have weak \\feii/\\hb\\ \\citep[e.g.,][]{bg92,sulentic00a}.  The ILR is weak in these systems because \\feii\\ is weak.  Their profiles show little deviation from a single Gaussian under the typical S/N level of SDSS spectra. The composite spectra of sources with large \\feii\\ redshifts, on the other hand, do show red asymmetry in the \\hb\\ profiles (see Fig. 13 of Paper I).  The variation in the relative strength of the \\hbvb\\ and \\hbi\\ components in different sources reflects the competition between the two components, although they apparently do so in such a manner that their kinematics remained coupled.  It is of interest to note that the strengths of \\hbi\\ and \\feii\\ are not correlated (Fig. \\ref{fig-ew}{\\it b}). The wide range of \\feii/\\hbi\\ ratios reflects either the complexity of the excitation mechanism of \\feii\\ emission \\citep[e.g.,][and references therein]{baldwin04} or large variations in quasar metallicities \\citep[e.g.,][]{netzer07}."
    },
    "0807/0807.0854_arXiv.txt": {
        "abstract": "In this geometrical approach to gravitational lensing theory, we apply the Gauss-Bonnet theorem to the optical metric of a lens, modelled as a static, spherically symmetric, perfect non-relativistic fluid, in the weak deflection limit. We find that the focusing of the light rays emerges here as a topological effect, and we introduce a new method to calculate the deflection angle from the Gaussian curvature of the optical metric. As examples, the Schwarzschild lens, the Plummer sphere and the singular isothermal sphere are discussed within this framework. ",
        "introduction": "The deflection of light by gravitational fields has been studied with great interest in astrophysics as well as in theoretical physics. Fundamental properties such as Fermat's principle for Lorentzian manifolds, conditions on image multiplicity and caustics in spacetime have been discussed in a fully relativistic setting (e.g., see \\cite{perlick} and references therein). In the astrophysical context, however, an impulse approximation with piecewise straight light rays in flat space has proven useful since deflection angles on cosmological scales are very small (for a comprehensive introduction, see e.g. \\cite{schneider} or \\cite{straumann} pp 272--95 and references therein). Despite their different premisses, both treatments have yielded mathematically interesting, general properties which depend on topology. In particular, image counting theorems like the odd number theorem have been established with different versions of Morse theory both in the spacetime lensing and impulse approximation frameworks \\cite{mckenzie}.  In this article, we would like to present another approach to gravitational lensing theory which emphasizes global properties. Specifically, we consider the astrophysically relevant weak deflection limit not in the impulse approximation but treat light rays as spatial geodesics of the optical metric, and use the Gauss-Bonnet theorem. This approach has previously been applied to lensing by cosmic strings \\cite{gibbons}, and we extend it here to static, spherically symmetric bodies of a perfect fluid as simple models for galaxies acting as gravitational lenses. It turns out that the focusing of light rays is, from this point of view, essentially a topological effect. Hence we find, rather surprisingly, that the deflection angle can be calculated by integrating the Gaussian curvature of the optical metric outwards from the light ray, in contrast to the usual description in terms of the mass enclosed within the impact parameter of the light ray. To illustrate, we discuss how this works for three well-known models, namely the Schwarzschild lens, the Plummer sphere and the singular isothermal sphere.  The structure of this article is therefore as follows. In section 2 we give a brief review of the Gauss-Bonnet theorem and introduce two constructions which will be used to investigate the lensing geometry. The optical metric and its Gaussian curvature for static, spherically symmetric systems of a perfect fluid is discussed in section 3, followed by the application to the three lens models mentioned above in section 4.  With regard to conventions, we use metric signature $(-,+,+,+)$, Latin and Greek indices for space and spacetime coordinates, respectively, and set the speed of light $c=1$. $G$ denotes the gravitational constant as usual. ",
        "conclusions": "In this article, we have introduced a geometrical approach to gravitational lensing theory different from the spacetime and impulse approximation treatments. By applying the Gauss-Bonnet theorem to the optical metric, whose geodesics are the spatial light rays, we found that the focusing of light rays can be regarded as a topological effect. We also gave a new expression (\\ref{gb3}) to calculate the deflection angle by integrating the Gaussian curvature of the optical metric outwards from the light ray, in the weak deflection limit. The lens models considered were given by static, spherically symmetric bodies of a non-relativistic, perfect fluid, and we discussed as examples the Schwarzschild lens, the Plummer sphere and the singular isothermal sphere.  It would therefore be interesting to see whether this approach could also be extended and be fruitfully applied to lenses without spherical symmetry, where images are no longer collinear with the lens centre, or to the relativistic strong deflection limit.  \\ack MCW would like to thank Claude Warnick for useful discussions, and the Science and Technology Facilities Council, UK, for funding."
    },
    "0807/0807.0123_arXiv.txt": {
        "abstract": "% We briefly summarize our findings from the unbiased surveys for $z$$\\sim$5 LBGs based on Subaru/Suprime-Cam and follow-up optical spectroscopy. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0807/0807.3316_arXiv.txt": {
        "abstract": "We describe a sample of low-mass Seyfert 2 galaxies selected from the Sloan Digital Sky Survey, having a median absolute magnitude of $M_g = -19.0$ mag.  These galaxies are Type 2 counterparts to the Seyfert 1 galaxies with intermediate-mass black holes identified by Greene \\& Ho (2004).  Spectra obtained with the \\emph{Echellette Spectrograph and Imager} at the Keck Observatory are used to determine the central stellar velocity dispersions and to examine the emission-line properties.  Overall, the stellar velocity dispersions are low ($\\sim40-90$ \\kms), and we find 12 objects having $\\sigmastar < 60$ \\kms, a range where very few Seyfert 2 galaxies were previously known. The sample follows the correlation between stellar velocity dispersion and FWHM([\\ion{O}{3}]) seen in more massive Seyfert galaxies, indicating that the narrow-line FWHM values are largely determined by virial motion of gas in the central regions of the host galaxies, but the [\\ion{O}{3}] emission lines exhibit a higher incidence of redward asymmetries and double-peaked profiles than what is found in typical Seyfert samples.  Using estimates of the black hole masses and AGN bolometric luminosities, we find that these galaxies are typically radiating at a high fraction of their Eddington rate, with a median $\\lbol/\\ledd = 0.4$.  We identify one galaxy, SDSS J110912.40+612346.7, as a Type 2 analog of the nearby dwarf Seyfert 1 galaxy NGC 4395, with a nearly identical narrow-line spectrum and a dwarf spiral host of only $M_g = -16.8$ mag.  The close similarities between these two objects suggest that the obscuring torus of AGN unification models may persist even at the lowest luminosities seen among Seyfert galaxies, below $\\lbol = 10^{41}$ ergs s\\per. Spectropolarimetry observations of four objects do not reveal any evidence for polarized broad-line emission, but SDSS J110912.40+612346.7 has a continuum polarization significantly in excess of the expected Galactic foreground polarization, possibly indicative of scattered light from a hidden nucleus.  Forthcoming observations of this sample, including X-ray and mid-infrared spectroscopy, can provide new tests of the obscuring torus model for active galaxies at low luminosities. ",
        "introduction": "The majority of active galactic nuclei (AGNs) are found in giant galaxies with substantial bulges, and there is a dramatic drop in the AGN fraction for late Hubble types \\citep{hfs97} and for host galaxies with stellar masses below $\\sim10^{10}$ \\msun\\ \\citep{kau03agn}.  This trend is consistent with the generally accepted scenario in which black hole growth and bulge growth are closely coupled, as expected from the correlations between black hole mass and bulge mass \\citep{kr95}, and between black hole mass and stellar velocity dispersion \\citep[the \\msigma\\ relation;][]{fm00, geb00}.  The low AGN fraction in late-type and low-mass galaxies arises from a combination of factors.  Perhaps most important, the black hole occupation fraction in low-mass galaxies is apparently below unity, as demonstrated by the stellar-dynamical non-detections of central black holes in the Local Group galaxies M33 \\citep{geb01, mfj01} and NGC 205 \\citep{val05}.  For those late-type galaxies that do contain a central black hole, shallow central gravitational potential wells may lead to a low efficiency for fueling black hole accretion.  A low-mass black hole ($\\lesssim10^6$ \\msun), even if radiating at its Eddington rate, cannot produce a very luminous AGN, and such objects can only be readily detected as AGNs if they are fairly nearby.  In optical surveys, dust extinction and blending with circumnuclear star-forming regions can further hinder the detection of low-luminosity active nuclei in late-type spirals.  If the \\msigma\\ relation continues toward low masses, then a black hole of $10^6$ \\msun\\ would correspond to a host galaxy velocity dispersion of $\\sigmastar\\approx60$ \\kms\\ \\citep{tre02}.  Thus, to explore the intermediate-mass regime for black holes, it is particularly interesting to search for AGN host galaxies having $\\sigmastar<60$ \\kms.  Such objects are known to exist, but they are rare.  Prior to the Sloan Digital Sky Survey (SDSS), there were only two nearby examples of dwarf galaxies known to contain Seyfert 1 nuclei: POX 52, a dE galaxy with $\\sigmastar=36\\pm5$ \\kms\\ \\citep{ksb87, bar04}, and NGC 4395, an Sd-type spiral with a central velocity dispersion of $\\sigmastar\\lesssim30$ \\kms\\ \\citep{fs89, fh03}.  The availability of SDSS has made it possible for the first time to search systematically for more examples of AGNs with low-mass black holes and small central velocity dispersions.  Greene \\& Ho (2004; hereinafter GH04) carried out the first such survey using Data Release 1 of the SDSS, identifying 19 Seyfert 1 galaxies as likely candidates for having $\\mbh < 10^6$ \\msun.  Although host galaxy properties were not considered in the sample selection, the hosts for this sample turned out to be galaxies of relatively low luminosity, on average about 1 mag fainter than \\lstar.  \\citet{bgh05} measured stellar velocity dispersions for many of the galaxies in this sample and found that they are also low ($\\sim35-80$ \\kms) and fall close to the local \\msigma\\ relation extrapolated toward lower masses.  Most previous searches for low-mass AGNs have concentrated on Type 1 (broad-lined) objects (Greene \\& Ho 2004, 2007; Dong \\etal\\ 2006) because the broad-line widths and AGN continuum luminosity can be used to estimate the black hole mass.  However, at low luminosities, most Seyferts are Type 2 objects \\citep{hfs97, hao05b}, and a full determination of the demographics of AGNs with low-mass black holes must take into account the Type 2 population.  Furthermore, without the glare of a bright AGN point source, the host galaxies of Seyfert 2s are more easily studied.  This can make it possible to examine the central stellar populations and star formation histories of Seyfert 2 host galaxies in a level of detail that would be impossible for Type 1 AGNs \\citep{kau03agn, hec04}.  Since searches for broad-lined AGNs require both the flux and the width of a broad emission line (such as \\hal) to be above some detection threshold, surveys for Type 2 AGNs based on measurements of narrow emission lines can potentially probe lower AGN luminosities, and therefore may be sensitive to AGNs in smaller host galaxies, or containing black holes of lower mass than those that can be found in Seyfert 1 surveys.  Optical spectroscopic surveys such as the Palomar survey \\citep{hfs97} have previously identified some examples of very low-luminosity Seyfert 2 nuclei in late-type galaxies, such as the Sc galaxy NGC 1058, which has a central velocity dispersion of only $31\\pm6$ \\kms\\ \\citep{bhs02}. Other strategies for finding obscured AGNs, such as mid-infrared searches for high-ionization coronal line emission from obscured active nuclei \\citep{sat07}, have contributed further evidence for black holes in late-type spirals.  Since stellar-dynamical searches for low-mass black holes are limited to very nearby galaxies (within a few Mpc at best), AGN surveys remain the best way to examine the demographics of these objects, and to determine the properties of the host galaxies in which low-mass black holes form and grow.  This paper describes new observations of a set of Seyfert 2 galaxies selected from Data Release 2 of the Sloan Digital Sky Survey \\citep[SDSS DR2]{aba04} to have host galaxies with relatively low luminosities and small central velocity dispersions.  We select galaxies spectroscopically identified as having Seyfert 2 nuclei (rather than LINER or AGN/starburst transition or composite types), in order to focus on objects having nuclear emission-line spectra that are clearly and unambiguously dominated by an accretion-powered AGN. These galaxies are in many respects Type 2 counterparts of the low-mass Seyfert 1 sample found by GH04, although the two samples are selected by very different criteria.  Most of the objects in both of these samples are not small or faint enough to be considered dwarf galaxies, but the objects described here are sub-\\lstar\\ galaxies that occupy ranges in mass, luminosity, and central velocity dispersion where very few Type 2 AGNs were known before SDSS.  We present measurements of the properties of these galaxies from high-resolution spectra taken at the Keck Observatory.  The small-aperture, high-resolution Keck spectra make it possible to confirm the AGN classification of these galaxies, to detect weak emission-line features and search for faint broad-line components, to measure stellar velocity dispersions, and to fully resolve the profiles of the narrow emission lines and measure accurate linewidths. We discuss the properties of these galaxies and the Seyfert 1s from the GH04 sample in the context of unified models of AGNs.  We also describe a search for polarized broad-line and continuum emission in four of these low-mass Seyfert 2 galaxies.  In this paper, distance-dependent quantities are calculated assuming $H_0 = 72$ km s\\per\\ Mpc\\per, $\\Omega_m$ = 0.3, and $\\Omega_\\Lambda = 0.7$. ",
        "conclusions": "We have presented an initial, exploratory study of the properties of Seyfert 2 galaxies with sub-\\lstar\\ luminosities selected from the Sloan Digital Sky Survey.  Our measurements reveal very low central stellar velocity dispersions (\\sigmastar $<60$ \\kms) in 12 objects; these are among the smallest velocity dispersions found in any AGN host galaxies, and imply that these galaxies contain some of the least massive black holes in any known AGNs.  We have also identified one of the very few known examples of a high-excitation Seyfert 2 nucleus in a late-type, dwarf spiral host galaxy similar to NGC 4395.  The correlations between [\\ion{O}{3}] linewidth and \\sigmastar\\ established for higher-mass Seyferts continue to hold in this low-mass regime, while the low-mass Seyfert 2 galaxies exhibit an unusually high incidence of peculiarities (such as redward asymmetries and double-peaked profiles) in their forbidden emission lines.  Using rough estimates of black hole mass and bolometric luminosity, we find that the median value of $\\lbol/\\ledd$ for this sample is 0.4, indicating that many of these are objects in which the black hole is undergoing a major growth phase.  Future work on these objects will include \\emph{XMM-Newton} observations and \\emph{Spitzer} mid-infrared spectroscopy, in order to search for evidence of nuclear obscuration and to test whether the obscuring torus of AGN unified models is present in this low-mass, low-luminosity regime."
    },
    "0807/0807.1914_arXiv.txt": {
        "abstract": "Absent any indirect tests on the thermal history of the Universe prior to the formation of light nuclear elements, it is legitimate to investigate situations where, before nucleosyntheis, the sound speed of the plasma was larger than  $c/\\sqrt{3}$, at most equalling the speed of light $c$. In this plausible extension of the current cosmological paradigm, hereby dubbed Tensor-$\\Lambda$CDM (i.e. T$\\Lambda$CDM) scenario, high-frequency gravitons are copiously produced. Without conflicting with the bounds on the tensor to scalar ratio stemming from the combined analysis of the three standard cosmological data sets (i.e. cosmic microwave background anisotropies, large-scale structure and supenovae),  the spectral energy density of the relic gravitons in the T$\\Lambda$CDM scenario can be potentially observable by wide-band interferometers (in their advanced version) operating  in a frequency window which ranges between few Hz and few kHz. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0807/0807.3911_arXiv.txt": {
        "abstract": "This paper explores the mapping between the observable properties of a stellar halo in phase- and abundance-space and the parent galaxy's accretion history in terms of the characteristic epoch of accretion and mass and orbits of progenitor objects. The study utilizes a suite of eleven stellar halo models constructed within the context of a standard $\\Lambda$CDM cosmology. The results demonstrate that coordinate-space studies are sensitive to the recent (0-8 Gyears ago) merger histories of galaxies (this timescale corresponds to the last few to tens of percent of mass accretion for a Milky-Way-type galaxy). Specifically, the {\\it frequency, sky coverage} and {\\it fraction of stars} in substructures in the stellar halo as a function of surface brightness are indicators of the importance of recent merging and of the luminosity function of infalling dwarfs. The {\\it morphology} of features serves as a guide to the orbital distribution of those dwarfs. Constraints on the earlier merger history ($>$ 8  Gyears ago) can be gleaned from the abundance patterns in halo stars: within our models, dramatic differences in the dominant epoch of accretion  or luminosity function of progenitor objects leave clear signatures in the [$\\alpha$/Fe] and [Fe/H] distributions of the stellar halo --- halos dominated by very early accretion have higher average [$\\alpha$/Fe], while those dominated by high luminosity satellites have higher [Fe/H]. This intuition can be applied to reconstruct much about the merger histories of nearby galaxies from current and future data sets. ",
        "introduction": "Phase- and abundance-space substructure in the stellar distribution  around a galaxy is commonly interpreted as a natural outcome of the process of hierarchical structure formation, where large galaxies are built  in part from the accretion of dwarf galaxies. Numerous previous studies have looked at how to understand individual debris features around galaxies in terms of the properties of the progenitor dwarfs  \\citep[e.g.][]{johnston98,helmi99a,johnston01,law05,fardal06,warnick08}. More generally, we might ask: to what extent can the merger history of a galaxy be reconstructed from its surrounding substructure?; and what could  you learn about the nature of merger histories in our Universe by examining the outskirts of a large sample of galaxies to very low surface brightness? These questions have yet to be addressed beyond using simple analytic estimates for the expected scalings in tidal debris \\citep{bullock01,johnston01}.  Motivation for answering these questions in more depth has been amply provided in the last decade by observations which have revealed abundant substructure in the spatial distribution around both the Milky Way and Andromeda galaxies, already seen at a few tens of kpc from their centers, and dominant at 100kpc \\citep[e.g.][and see \\S \\ref{current.sec} for more examples]{ivezic00,newberg02,newberg03,ferguson02,majewski03,belokurov06,ibata07}. The ubiquity of such substructure has become apparent because of dramatic increases in the sky-coverage of halo surveys, the number of tracers (and hence surface brightness to which such surveys are sensitive) and the distances from the parent galaxies which have been probed. In contrast, only slightly more than a decade ago the classical picture of stellar halos around galaxies  was that the stars in them were smoothly distributed in phase-space --- a picture informed by observations of the RR Lyrae and globular cluster distribution around the Milky Way using samples of a few dozen to hundreds  of objects \\citep[e.g.][]{wetterer96}. In our own Galaxy, even halo stars in the Solar neighborhood have been shown to be clumped once their full phase-space coordinates are known and their orbital distribution is considered \\citep{helmi99b,morrison08}.  Satellite accretion is the preferred explanation for apparent phase-space substructure. Debris from the destruction of a satellite will phase-mix along the progenitor's orbit over time to fill the full volume of the original orbit in coordinate space \\citep{johnston98}, while becoming locally colder at each point in velocity-space  \\citep[in order to satisfy Louiville's theorem --- see][for a rigorous description]{helmi99a}. The small range in orbital angular momenta and energy is largely conserved during any gradual evolution of the parent galaxy potential, although the average values can evolve \\citep{penarrubia06}. \\citep[See also][for a discussion in a more violent context.]{knebe05,warnick08} Putting these results for individual accretion events together leads to a spectrum of lumps in phase-space, as has been demonstrated in composite studies of halo distributions \\citep{helmi00,bullock01,johnston01,helmi03,bullock05,delucia08}.  Similarly, substructure is now becoming apparent in metallicity distributions and abundance patterns of halo stars. So far, these studies have relied on using coordinates in both phase- and abundance-space to identify these structures \\citep[e.g.][]{majewski96,navarro04,helmi06,ibata07,gilbert08}. However, abundance space has the advantage over phase-space of preserving a complete memory of a star's initial conditions which cannot be destroyed by subsequent mixing or scattering. In principle,``chemically tagging'' stars could lead to associations that are not apparent in phase-space alone \\citep{blandhawthorn03}. Indeed, it has already been demonstrated that the distinct merging time and mass scales of stellar halo,  satellite and dominant substructure progenitors  means that stars in each of these systems should have distinct patterns in $\\alpha$-element/metallicity space \\citep{robertson05,font06a,font06b,font08}, which implies that distinctions between lower-contrast systems could be possible in higher dimensions. These results also suggest that abundance distributions in a stellar halo should reflect its accretion history \\citep[as has already been shown for metallicity distributions:][]{renda05,font06b}.  This paper takes a first step towards addressing how to broadly  relate substructures apparent today to a galaxy's past by exploring how the observed properties of eleven stellar halo models built entirely from accretion events within the context of a $\\Lambda$CDM universe reflect their accretion histories. Further work will build on the intuition developed here to define statistical measures of substructure that are tuned to be sensitive to the epoch of accretion and mass and orbit type of progenitor satellites. The modeling methods are described in \\S2. In \\S 3 the 1515 simulated accretion events from all eleven model halos are first analyzed individually in order to characterize how the {\\it observed} properties of debris (i.e. spatial and velocity scales, morphology, surface brightness and abundance patterns) are related to the {\\it intrinsic} properties of the progenitor satellite (i.e. accretion time, luminosity and orbit). The intuition developed in \\S 3 is then applied in  \\S 4 to understanding how the {\\it observed} properties of stellar halos (i.e. frequency, morphology, surface brightness and stellar populations of non-uniform features) arise from their unique formation histories. These ideas are illustrated by contrasting our eleven ``standard'' (i.e. built within a $\\Lambda$CDM context) halos with four model halos built from accretion histories that have been artificially constrained to be dominated by ancient/recent and  high/low luminosity events, as well as two more built from events predominantly on radial/circular orbits. The results are reviewed in \\S5, and subsequently applied to ``observations'' of our model halos  in order to demonstrate in principle how such data might be interpreted. In \\S6 we discuss applications of these ideas to both current and future data sets. We summarize our conclusions in \\S7. ",
        "conclusions": "In this paper we have explored the connection between the merger history of a galaxy and the present day structure of its stellar halo in phase- and abundance-space.  We have found that studies of substructure in the surface-density and/or density distribution in stellar halos are sensitive to the recent ($<$ 8 Gyears ago) merging history of a galaxy. These substructures can tell us about the mass, orbit and accretion time of progenitors, while the proportion of mass in the smooth component reflects the importance of early events (whether in the form of direct accretions from minor or major mergers or  {\\it in situ} formation from a heated disk  or monolithic collapse). For a Milky-Way-type galaxy, built within the context of $\\Lambda$CDM, we expect of order few to tens of percent of the halo to be in the form of substructure, for those substructures to be increasingly dominant at large radii and for the inner halo to be relatively smooth. These expectations are broadly consistent with the current data for the Milky Way and Andromeda galaxies .  Our analysis was based on relating the results of a subjective morphological classification of debris morphology from individual accretion events to the properties of the progenitor satellite. As our sample of galaxies and substructures within galaxies grows, the challenging nature of reaching very low surface brightness, and the multiple dependencies of substructure characteristics mean that interpretations of individual features in terms of progenitor properties may not always be unique. Hence, for specific comparisons of data and observations, it makes sense to recast our results in terms of statistical measures of substructure rather than individual interpretations. For example, \\citet{bell08} report on a preliminary analysis of the SDSS halo turnoff-star data which quantifies the level of substructure via the dispersion in counts around a smooth background. The same analysis applied to the model stellar halos presented in this paper suggests a degree of substructure similar to the observations. More generally, these statistics need to be designed to be sensitive to differences in accretion history, armed with the knowledge developed here that: (i) the epoch of galaxy formation sets the percentage of the stellar halo contained in phase-space substructure (and its average [$\\alpha$/Fe]); (ii) the number and mass scale of recently accreted objects set the number, angular scales (and mean metallicities) of substructure; and (iii) the orbit type of progenitors set the morphology of substructure.  In contrast to coordinate-space, signatures in abundance-space are not subject to mixing over time so should last indefinitely. This promises a way to look further back time \\citep{blandhawthorn03}. Our own results demonstrate that variations in [$\\alpha$/Fe] and [Fe/H] among the smooth components of different stellar halos reflect variations is the dominant epoch of accretion and masses of progenitor objects respectively. Moving to higher dimensions in abundance space should allow even more precise interpretations."
    },
    "0807/0807.1255_arXiv.txt": {
        "abstract": "The algorithm of the ensemble pulsar time based on the optimal Wiener filtration method has been constructed. This algorithm allows the separation of the contributions to the post-fit pulsar timing residuals of the atomic clock and pulsar itself. Filters were designed with the use of the cross- and autocovariance functions of the timing residuals. The method has been applied to the timing data of millisecond pulsars PSR B1855+09 and PSR B1937+21 and allowed the filtering out of the atomic scale component from the pulsar data. Direct comparison of the terrestrial time TT(BIPM06) and the ensemble pulsar time PT$_{\\rm ens}$ revealed that fractional instability of TT(BIPM06)--PT$_{\\rm ens}$ is equal to $\\sigma_z=(0.8\\pm 1.9)\\cdot 10^{-15}$. Based on the $\\sigma_z$ statistics of TT(BIPM06)--PT$_{\\rm ens}$ a new limit of the energy density of the gravitational wave background was calculated to be equal to $\\Omega_g h^2 \\sim 3\\cdot 10^{-9}$. ",
        "introduction": "The discovery of pulsars in 1967 \\cite{Hewish68} showed clearly that their rotational stability allowed them to be used as astronomical clocks. This became even more obvious after discovery of the millisecond pulsar PSR B1937+21 in 1982 \\cite{Backer82}. Now a typical accuracy of measuring the time of arrivals (TOA) of millisecond pulsar pulses comprises a few microseconds and even hundreds of nanoseconds for some pulsars. For the observation interval in the order $10^8$ seconds this accuracy produces a fractional instability of $10^{-15}$, which is comparable to the fractional instability of atomic clocks. Such a high stability cannot but used for time metrology and time keeping.  There are several papers considering applicability of stability of pulsar rotation to time scales. The paper \\cite{Guinot88} presents principles of the establishment of TT (Terrestrial Time) with the main conclusion that one cannot rely on the single atomic standard before authorised confirmation and, for pulsar timing, one should use the most accurate realisations of terrestrial time TT(BIPMXX) (Bureau International des Poids et Mesures). The paper of Ilyasov et. al. \\cite{ilyasov89} describes the principles of a pulsar time scale, definition of \"pulsar second\" is presented. Guinot \\& Petit (1991) show that, because of the unknown pulsar period and period derivative, rotation of millisecond pulsars is only useful for investigations of time scale stability a posteriori and with long data spans. The papers \\cite{ilyasov96}, \\cite{kopeikin97}, \\cite{rodin97}, \\cite{ilyasov98} suggest a binary pulsar time (BPT) scale based on the motion of a pulsar in a binary system with theoretical expressions for variations in rotational and binary parameters depending on the observational interval. The main conclusion is that BPT at short intervals is less stable than the conventional pulsar time scale (PT), but at a longer period of observation ($10^2\\div 10^3$ years) the fractional instability of BPT may be as accurate as $10^{-14}$. The paper of Petit \\& Tavella \\cite{petit96} presents an algorithm of a group pulsar time scale and some ideas about the stability of BPT. The paper of Foster \\& Backer \\cite{foster90} presents a polynomial approach for describing clock \\& ephemeris errors and the influence of gravitational waves passing through the Solar system.  In this paper the author presents a method of obtaining corrections of the atomic time scale relative to PT from pulsar timing observations. The basic idea of the method was published earlier in the paper \\cite{rodin06}.  In Sect.~\\ref{sect2}, the main formulae of pulsar timing are described. Sect.~\\ref{sect3} contains theoretical algorithm of Wiener filtering. Sect.~\\ref{sect4} presents the results of numerical simulation, i.e. recovery of harmonic signal from noisy data by Wiener optimal filter and weighted average algorithm. The latter one is used similarly to the paper of Petit \\& Tavella \\cite{petit96}. Sect.~\\ref{sect5} describes an application of the algorithm to real timing data of pulsars PSR B1855+09 and PSR B1937+21 \\cite{kaspi94}. ",
        "conclusions": "An algorithm of forming of ensemble pulsar time scale based on the method of the optimal Wiener filtering is presented.  The basic idea of the algorithm consists in the use of optimal filter for removal of additive noise from the timing data {\\it before} construction of the weighted average ensemble time scale.  Such a filtering approach offers an advantage over the weighted average algorithm since it utilises additional statistical information about common signal in the form of its covariance function or power spectrum. Since timing data are always available relative to a definite time scale, for separation of the pulsar and clock contributions one need to use observations from a few pulsars (minimum two) relative to the same time scale. Such approach allows estimation of the signal covariance function (power spectrum) by averaging all cross-covariance functions or power cross-spectra of the original data.  The availability of two scale differences UTC -- TT and UTC -- PT has resulted in the long awaited possibility of comparing the ultimate terrestrial time scale TT and extraterrestrial ensemble pulsar time scale PT of comparable accuracy. The fractional instability of the terrestrial time scale TT relative to PT and their high correlation have demonstrated that PT scale can be successfully used for monitoring the long-term variations of TT."
    },
    "0807/0807.3066_arXiv.txt": {
        "abstract": "When the universe was about 10 $\\mu $seconds old, a first order cosmological quark - hadron phase transition occurred at a critical temperature of around 200 MeV. In this work, we study the quark-hadron phase transition in the context of brane-world cosmologies, in which our Universe is a three-brane embedded in a five-dimensional bulk, and within an effective model of QCD. We analyze the evolution of the physical quantities, relevant for the physical description of the early universe, namely, the energy density, temperature and scale factor, before, during, and after the phase transition. To study the cosmological dynamics and evolution we use both analytical and numerical methods. In particular, due to the high energy density in the early Universe, we consider in detail the specific brane world model case of neglecting the terms linearly proportional to the energy density with respect to the quadratic terms. A small brane tension and a high value of the dark radiation term tend to decrease the effective temperature of the quark-gluon plasma and of the hadronic fluid, respectively, and to significantly accelerate the transition to a pure hadronic phase. By assuming that the phase transition may be described by an effective nucleation theory, we also consider the case where the Universe evolved through a mixed phase with a small initial supercooling and monotonically growing hadronic bubbles. ",
        "introduction": "The possibility that our $4D$ universe may be viewed as a brane hypersurface embedded in a higher dimensional bulk space has attracted considerable interest lately. A scenario with an infinite fifth dimension in the presence of a brane can generate a theory of gravity which mimics purely four-dimensional gravity, both with respect to the classical gravitational potential and with respect to gravitational radiation \\cite{RS99a}. The gravitational self-couplings are not significantly modified in this model. This result has been obtained from the study of a single 3-brane embedded in five dimensions, with the $5D$ metric given by $ds^{2}=e^{-f(y)}\\eta _{\\mu \\nu }dx^{\\mu }dx^{\\nu }+dy^{2}$, which can produce a large hierarchy between the scale of particle physics and gravity due to the appearance of the warp factor \\cite{RS99a}. Even if the fifth dimension is uncompactified, standard $4D$ gravity is reproduced on the brane. In contrast to the compactified case, this follows because the near-brane geometry traps the massless graviton. Hence this model allows the presence of large or even infinite non-compact extra dimensions. Our brane is identified to a domain wall in a 5-dimensional anti-de Sitter space-time.  The effective gravitational field equations on the brane-world, in which all the matter forces except gravity are confined on the 3-brane in a 5-dimensional space-time with $Z_2$-symmetry have been obtained, by using an elegant geometric approach, in \\cite{SMS00,SSM00}. The correct signature for gravity is provided by the brane with positive tension. If the bulk space-time is exactly anti-de Sitter, then generically the matter on the brane is required to be spatially homogeneous. The electric part of the 5-dimensional Weyl tensor $E_{IJ}$ gives the leading order corrections to the conventional Einstein equations on the brane. This implies a modification of the basic equations describing the cosmological and astrophysical dynamics, which has been extensively considered \\cite{all2}. For reviews of the dynamics and geometry of the brane Universes, as well as for the discussions of the cosmological implications see \\cite{Ma01}.  According to standard cosmology, as it expanded and cooled, the early Universe is expected to have undergone a series of symmetry-breaking phase transitions, at which topological defects may have formed. Phase transitions are labelled first or second order, according to whether the position of the vacuum state in field space changes discontinuously or continuously, as the critical temperature is crossed. A first order phase transition proceeds by bubble nucleation and expansion. When at least $(4-n)$ of these bubble collide, for $n=0,1,2$, an $n$-dimensional topological defect may form in the region between them \\cite{Kajantie:1986hq}. Recent lattice QCD calculations for two quark flavors suggest that QCD makes a transition at a temperature of $T_{c}\\sim 150$ MeV \\cite{TaBo07}. This phase transition, which may have occurred in the early Universe, could lead to the formation of relic quark-gluon plasma objects, which still survive today.  A first order quark-hadron phase transition in the expanding Universe can be described generically as follows \\cite{Kajantie:1986hq}. As the color deconfined quark-gluon plasma cools below the critical temperature $T_{c}\\approx 150$ MeV, it becomes energetically favorable to form color confined hadrons (primarily pions and a tiny amount of neutrons and protons, due to the conserved net baryon number). However, the new phase does not show up immediately. As is characteristic for a first order phase transition, some supercooling is needed to overcome the energy expense of forming the surface of the bubble and the new hadron phase. When a hadron bubble is nucleated, latent heat is released, and a spherical shock wave expands into the surrounding supercooled quark-gluon plasma. This reheats the plasma to the critical temperature, preventing further nucleation in a region passed by one or more shock fronts. Generally, bubble growth is described by deflagrations, with a shock front preceding the actual transition front. The nucleation stops when the whole Universe has reheated to $T_{c}$. This part of the phase transition passes very fast, in about $0.05$ $\\mu $s, during which the cosmic expansion is totally negligible. After that, the hadron bubbles grow at the expense of the quark phase and eventually percolate or coalesce. The transition ends when all quark-gluon plasma has been converted to hadrons, neglecting possible quark nugget production. The physics of the quark-hadron phase transition, as well as the cosmological implications of this process have been extensively discussed in the framework of general relativistic cosmology in \\cite{IgKaKuLa94} - \\cite{IgSc01}  In the context of brane-world scenarios, the Friedmann equation contains deviations to the $4D$ case, which results in an increased expansion in early times. This in general has important effects on the cosmological evolution, and in particular on cosmological phase transitions.  First order phase transitions have also been considered, in the framework of the brane-world model, in \\cite{DaVe01}. Due to the effects coming from the extra-dimensions, phase transitions require a higher nucleation rate to complete, and baryogenesis and particle abundances could be suppressed. The evolution of topological defects is also affected, but the increased expansion cannot solve the monopole and domain wall problems.  In this work, we consider the quark-hadron phase transition in the brane-world scenario. By assuming that the phase transition is of the first order, we study in detail the evolution of the relevant cosmological parameters (energy density, temperature, scale factor, etc) of the quark-gluon and hadron phases, and the phase transition itself. It is important to emphasize that in the early universe the energy density is high, so that one may neglect the terms linearly proportional to the energy density with respect to the quadratic terms. This is carried out in detail in this work. An important parameter to describe the phase transition is the hadron fraction, whose time evolution describes the conversion process. We also consider the effect of the dark radiation term on the phase transition. The brane world effects (the quadratic corrections to the matter energy-momentum tensor, described by the numerical value of the brane tension, and the dark radiation term) lead to an overall decrease of the temperature of the very early universe, and accelerate the transition to the pure hadronic phase.  This paper is organized in the following manner. In Section \\ref{SecII}, we briefly outline, for self-completeness and self-consistency, the field equations in brane-world models and the basics of the quark-hadron phase transition. In the latter, we lay down the equations of state and the relevant physical quantities that are analyzed in the remaining Sections. In Section \\ref{SecIII} we analyze in detail the quark-hadron phase transition. In Section \\ref{SecIV}, we consider bubble nucleation in the brane-world scenario, by assuming that the phase transition may be described by an effective nucleation theory. We discuss and summarize our results in Section \\ref{SecV}. In the present paper we use a system of units so that the speed of light is $c=1$. ",
        "conclusions": "\\label{SecV}  In the context of brane-world scenarios, in the high density cosmological phase the Friedmann equation contains deviations to the $4D$ case, which imposes fundamental phenomenological consequences on the cosmological evolution, and in particular on the cosmological phase transitions. In this work, we have analyze the evolution of the relevant physical quantities, namely, the energy density, temperature and scale factor before, during and after the phase transition. It is important to emphasize that in the early universe the energy density is extremely high, so that one can neglect the terms linearly proportional to the energy density with respect to the quadratic terms. This approximation has been considered in detail throughout this work. In the high density regime the Hubble function is proportional to the energy density of the cosmological matter, which drastically affects the cosmological dynamics of the universe. Moreover, in the early universe the dark radiation term, the projection of the Weyl tensor from the bulk, which appears in form of a radiation-like term in the field equations, may also play an important role. The magnitude of the brane world effects can be characterized by the numerical value of the brane tension  $\\lambda $. In the limit $\\lambda \\to \\infty$ we recover standard general relativity \\cite{Ma01}. On the other hand, the study of the quark-hadron phase transition is also very important from an observational point of view, since the inhomogeneities generated at the QCD phase transition might have a noticeable effect on nucleosynthesis \\cite {IgKaKuLa94a}.  By fully taking into account the brane world effects we have found that the temperature evolution of the universe in the brane world scenario is different from the idealized standard FRW model. The temperature of the early universe in the quark phase is smaller in the brane world scenario, as one can see from Fig.~\\ref{FIG1}, where the long dashed curve corresponds (approximately) to the general relativistic limit. Hence a small value of the brane tension would significantly reduce the temperature of the quark-gluon plasma, and accelerate the phase transition to the hadronic era. An increase of the dark radiation term for fixed bag constant $B$ and brane tension $\\lambda $ gives the same effect, as one can see from Fig.~\\ref{FIG2}. Once the quark-hadron phase transition starts, the hadron fraction $h$ is again strongly dependent on the brane tension. For small values of $\\lambda $, $h(t)$ is much higher than in the standard general relativity, as one can see from Fig.~\\ref{FIG3}. The increase of the dark radiation on the brane also strongly accelerates the formation of the hadronic phase (see Fig.~\\ref{FIG4}), and decreases the time interval necessary for the transition. A small brane tension and a high energy density of the dark radiation also tend to reduce the temperature of the hadronic fluid. Of course, the temperature evolution also depends upon the relativistic degrees of freedom in the equation of state and upon the equation of state. In addition to this, by assuming that the phase transition may be described by an effective nucleation theory, we have also considered the case where the Universe evolved through a mixed phase with a small initial supercooling, and monotonically growing hadronic bubbles. It was shown that at a certain temperature below the critical value, an enormous amount of hadronic bubbles are nucleated, which grow explosively to transform all the plasma into hadrons, indicating that the small supercooling phase transition is very rapid with respect to the simple first order phase transition.  Many details of the QCD phase transition are not yet conclusively understood. Even the order of transition is still a matter of debate. An advance in the understanding of the numerical values of the QCD coupling constants would be very helpful in obtaining accurate cosmological conclusions. Such an advance may also provide a powerful method for testing on a cosmological scale the theoretical predictions of the brane world models and the possible existence of the extra-dimensions."
    },
    "0807/0807.4470_arXiv.txt": {
        "abstract": "{ During the study of a large set of late-type stellar X-ray sources, we discovered a large fraction of multiple systems.} { In this paper we investigate the orbital elements and kinematic properties of three new spectroscopic triple systems as well as spectral types and astrophysical parameters ($T_{\\rm eff}$, $\\log g$, $v\\sin i$, $\\log N$(Li)) of their components.} { We conducted follow-up optical observations, both photometric and spectroscopic at high resolution, of these systems. We used a synthetic approach and the cross-correlation method to derive most of the stellar parameters.} { We estimated reliable radial velocities and deduced the orbital elements of the inner binaries. The comparison of the observed spectra with synthetic composite ones, obtained as the weighted sum of three spectra of non-active reference stars, allowed us to determine the stellar parameters for each component of these systems. We found all are only composed of main sequence stars.} { These three systems are certainly stable hierarchical triples composed of short-period inner binaries plus a tertiary component in a long-period orbit. From their kinematics and/or Lithium content, these systems result to be fairly young.} ",
        "introduction": "\\label{sec:Intro}  Binary and multiple stars are very important astrophysical laboratories. In particular, spectro-photometric and spectro-astrometric binaries offer the unique opportunity to determine, with a high level of accuracy, the basic stellar parameters (mass, radius, and effective temperature) to study stellar structure and evolution. However, the formation and evolution of binary stars are still debated subjects \\citep[e.g.,][]{ZM2001}. Especially, a still unsolved problem is the formation of close binaries with main sequence components separated by few solar radii that in the proto-stellar phase should have been in contact.  In the last years, to answer many of these open questions, relevant observational and theoretical efforts are being done to improve continuously the statistics of binary systems with different periods, mass ratios, etc. \\citep[e.g.,][]{Tok06}.  \\begin{table*} \\caption[RasTyc sources]{Main data of the three {\\it RasTyc} sources from the literature.} \\begin{center} \\begin{tabular}{clccccrrcc} \\hline \\hline \\noalign{\\medskip} \\textsl{RasTyc} Name & Name & $\\alpha$ (2000) & $\\delta$ (2000) & V$_T^{\\,\\,\\rm a}$  & $\\pi^{\\,\\,\\rm b}$ & $\\mu_{\\alpha}^{\\,\\,\\rm a}$~~~ & $\\mu_{\\delta}^{\\,\\,\\rm a}$~~~ & X-ray source & Counts \\\\ & & (h m s) & ($\\degr ~\\arcmin ~\\arcsec$) & (mag) & (mas) & \\multicolumn{2}{c}{(mas\\,yr$^{-1}$)} & 1RXS & (ct\\,s$^{-1}$) \\\\ \\noalign{\\medskip} \\hline \\noalign{\\medskip} RasTyc\\,0524+6739 & \\object{BD+67\\,381}  & 05 24 53.2 & +67 39 39 & 9.065 &  7.8$\\pm$9.3\t& $-0.5$~~  & 26.7~~    & J052454.0+673939 & 4.21$\\times 10^{-1}$ \\\\ RasTyc\\,1828+3506 & \\object{BD+35\\,3261} & 18 28 50.3 & +35 06 34 & 9.049 & 12.2$\\pm$8.2\t& 12.3~~    & $-3.5$~~  & J182849.7+350637 & 6.24$\\times 10^{-2}$ \\\\ RasTyc\\,2034+8253 & \\object{BD+82\\,622}  & 20 34 27.5 & +82 53 35 & 9.730 &  ---\t\t& 61.5~~    & 35.5~~    & J203426.2+825334 & 3.75$\\times 10^{-1}$\\\\ \\hline \\noalign{\\medskip} \\end{tabular} \\begin{tabular}{cc} $^{\\rm a}$ $V$ magnitude and proper motions from the TYCHO-2 catalog \\citep{Hog00}; & $^{\\rm b}$ Parallax from the TYCHO-1 catalog \\citep{Hipp} \\\\ \\end{tabular} \\end{center} \\label{Tab:SourcesRasTyc} \\end{table*}  \\onlfig{1}{ \\begin{figure*}[!th] \\centering \\includegraphics[width=9.0cm]{Evolution_Ha_RXJ0524.eps} \\includegraphics[width=9.0cm]{Evolution_Li_RXJ0524.eps} \\includegraphics[width=9.0cm]{Evolution_Ha_RXJ1828.eps} \\includegraphics[width=9.0cm]{Evolution_Li_RXJ1828.eps} \\includegraphics[width=9.0cm]{Evolution_Ha_RXJ2034.eps} \\includegraphics[width=9.0cm]{Evolution_Li_RXJ2034.eps} \\caption{\\label{fig:EvolutionSpectraAll} High resolution spectra of RasTyc\\,0524+6739 (top panels), RasTyc\\,1828+3506 (middle panels) and RasTyc\\,2034+8253 (lower panels) acquired with the A{\\sc urelie} spectrograph at the  152-cm telescope of the OHP both in the H$\\alpha$ (left panels) and the Lithium spectral regions (right panels). The laboratory wavelengths of the \\ion{Fe}{i}\\,$\\lambda$6546.2 and the H$\\alpha$ lines as well as those of the \\ion{Li}{i}\\,$\\lambda$6707.8, and the \\ion{Ca}{i}\\,$\\lambda$6717.7 are marked with vertical dashed lines in the H$\\alpha$ and Lithium spectral regions, respectively.} \\end{figure*} }  Close binaries containing at least one late-type component, such as RS~CVn and BY~Dra systems, are objects with the strongest magnetic activity (starspots, plages, flares) induced by a dynamo action in the sub-photospheric convection zone. Their strong activity is mainly due to their very fast rotation (spin-orbit synchronization by tidal forces) and to proximity effects.  X-ray sky surveys performed in recent years have allowed to identify thousands of active late-type stars in the field and in open clusters. Follow-up observations of the optical counterparts of X-ray sources have led to discover very young stars far from the typical birth sites, i.e. open clusters and stars forming regions \\citep[e.g.,][]{Wichmann03, Zickgraf05, Torres2006, Guillout2008}, as well as to detect several spectroscopic binaries \\citep[e.g.,][]{Wichmann03b, Frasca2006}. The knowledge of the incidence of binaries and multiple systems in X-ray selected samples of active stars is extremely important to study the recent local star formation history.  One of the largest ($\\sim14\\,000$ active stars) and most comprehensive set of stellar X-ray sources in the field is the so-called \\textsl{RasTyc} sample, which is the result of the cross-correlation of the ROSAT All-Sky Survey (RASS) with the TYCHO catalog \\citep{Guillout1999}. We began to analyze a representative sub-sample of the \\textsl{RasTyc} population in the northern hemisphere \\citep{Guillout2008} to obtain some reliable statistics about the \\textsl{RasTyc} stellar characteristics. For this purpose, we led campaigns of high-resolution spectroscopic observations, with the E{\\sc lodie} \\'echelle spectrograph at the 193-cm telescope and the A{\\sc urelie} spectrograph at the 152-cm telescope of the {\\it Observatoire de Haute Provence} (OHP). For all the sources, we performed a detailed analysis of the cross-correlation function (CCF) and found that single-lined (SB1), double-lined (SB2), and triple-lined (SB3) spectroscopic systems altogether account for more than 35\\,\\% of the sample. In particular, at least 10 sources are clearly identified as triple systems. Our aim is to determine the orbital and physical parameters of these systems. For this reason, we have monitored these new multiple systems, both photometrically and spectroscopically, with the 91-cm telescope of the {\\it Osservatorio Astrofisico di Catania} (OAC).  The majority of the triple systems studied so far are nearby objects and look as visual binaries where one componentis a SB2 system. Here we study three of such newly discovered SB3 late-type systems for which we obtained enough data. A detailed analysis of the entire sample of stellar X-ray sources and of the multiple systems discovered so far is necessary for drawing statistically significant conclusions. Nevertheless, the properties of these three systems can give us some insights into the typical composition of triple systems among stellar X-ray sources.  The paper is organized as follows. We summarize briefly the observations and their reduction in Sect.~\\ref{sec:ObsMet}. The determination of radial velocity and physical parameters ($v\\sin i$'s, Lithium abundances, etc.) are shown in Sect.~\\ref{sec:ObsMet} as well. The photometric modulation, the spectral composition, and the age estimate are discussed in Sect.~\\ref{sec:Discussion}. The conclusions are outlined in Sect.~\\ref{sec:Conclusions}. ",
        "conclusions": "\\label{sec:Conclusions}  This paper is devoted to the analysis of three new triple systems discovered in the \\textsl{RasTyc} sample of stellar X-ray sources. Their spectroscopic and photometric data allow us to conclude that they are almost certainly stable hierarchical triple systems composed of short-period inner binaries plus a tertiary component in a long-period orbit. The orbital periods of the inner binaries range from 3.5 to 7.6 days and the orbits are practically circular. From the high-resolution spectra we also found the spectral composition and the astrophysical parameters of the components that turn out to be all G-K main sequence stars. In all cases, the components of the inner binaries have nearly the same masses, spectral types, and luminosities. From their kinematics and/or Lithium content, these systems result to be fairly young. RasTyc\\,2034+8253 is the only system in which the \\ion{Li}{i}\\,$\\lambda$6707.8 line is strong enough to be clearly visible in the spectra of all the three components and suggests an age in the range $100-300$\\,Myr. It is a possible new member of the IC\\,2391 supercluster. For the remaining systems, the membership to young moving groups is rather uncertain.  Our spectroscopic survey has revealed that multiple systems represent a large fraction of the \\textsl{RasTyc} sources. However, a detailed analysis is absolutely necessary for drawing statistically significant conclusions. Since \\textsl{RasTyc} objects are relatively nearby, the discovery and the study of new triple systems, such as those presented in the present paper, can contribute to a better understanding of the formation and the evolution of close binaries and multiple systems in the solar neighborhood."
    },
    "0807/0807.0473_arXiv.txt": {
        "abstract": "NGC 2770 has been the host of three supernovae of Type Ib during the last 10 years, SN 1999eh, SN 2007uy and SN 2008D. SN 2008D attracted special attention due to the serendipitous discovery of an associated X-ray transient. In this paper, we study the properties of NGC 2770 and specifically the three SN sites to investigate whether this galaxy is in any way peculiar to cause a high frequency of SNe Ib. We model the global SED of the galaxy from broadband data and derive a star-formation and SN rate comparable to the values of the Milky Way. We further study the galaxy using longslit spectroscopy covering the major axis and the three SN sites. From the spectroscopic study we find subsolar metallicities for the SN sites, a high extinction and a moderate star-formation rate. In a high resolution spectrum, we also detect diffuse interstellar bands in the line-of-sight towards SN 2008. A comparison of NGC 2770 to the global properties of a galaxy sample with high SN occurance ($\\geq$ 3 SN in the last 100 years) suggests that NGC 2770 is not particularly destined to produce such an enhancement of observed SNe observed. Its properties are also very different from gamma-ray burst host galaxies. Statistical considerations on SN Ib detection rates give a probability of $\\sim$ 1.5\\% to find a galaxy with three Ib SNe detected in 10 years. The high number of rare Ib SNe in this galaxy is therefore likely to be a coincidence rather than special properties of the galaxy itself. NGC 2770 has a small irregular companion, NGC 2770B, which is highly starforming, has a very low mass and one of the lowest metallicities detected in the nearby universe as derived from longslit spectroscopy. In the most metal poor part, we even detect Wolf-Rayet features, against the current models of WR stars which require high metallicities. ",
        "introduction": "Massive stars end their lives in various ways, as governed by their mass, composition, angular momentum and whether or not they are interacting with a companion star \\citep[e.g.][]{Heger02}. The most massive stars loose their outer layers through winds whose strength strongly depend on the metallicity of the star \\citep[e.g.][and references therein]{Crowther02}. These stripped stars explode as supernovae (SNe) Type Ib (SN Ib) if they have lost their hydrogen envelope or as SNe Type Ic (SN Ic) if they have also lost their He envelope. If the star is also rapidly rotating, it is believed that the star might even produce a Gamma-Ray Burst (GRB) \\citep{MacFadyen99, Woosley06}. SNe Ib/c are much rarer than SNe II, which are produced by stars with masses of 8 to 40 M$_\\odot$, whereas SNe Ib/cs are assumed to require zero age main sequence (ZAMS) masses of $\\gtrsim$ 35 M$_\\odot$ for non-rotating stars \\citep{Woosley02}. GRBs are even less common than SNe Ic which is in line with the general picture that GRB progenitors require some special conditions to produce a GRB.  Only for SNe II has it been possible to identify the progenitor star  - which turned out to be yellow, red or blue giants with masses of around 8 -- 20 M$_\\odot$, e.g. SN 1987A \\citep{Gilmozzi87}, SN 2002ov \\citep{Li07}, SN 2004A \\citep{Hendry06} and SN 2006gl \\citep{Gal-Yam07}. Furthermore, several detections have been made using preexplosion imaging from the {\\it HST} archive \\citep[e.g.][]{Smartt03}. SNe Ib/c have so far evaded an identification with a progenitor star \\citep{Maund05, Crockett07, Crockett08}. This is also the case for GRBs which most often occur at distances that do not even allow us to resolve their host galaxies.  When the progenitor cannot be directly identified, studying the environment can provide important information about what kind of star exploded.  Spatially resolved photometric studies of the sites of different types of SNe show that SNe Type II trace the light of their host galaxies while GRBs are more concentrated towards bright regions \\citep{Fruchter06} and that SNe Type Ib, Ic and GRBs appear to be differently distributed within their hosts \\citep{Kelly07}. Spectroscopic investigations show that the broadlined SNe Ic which are not connected to GRBs are found at sites with higher metallicities than those that are connected to GRBs \\citep{Modjaz08a}. A spatially resolved study of the host of one long-duration GRB which was not connected to a SN also showed a low metallicity in comparison to the rest of the host galaxy \\citep{Thoene08}.  Long duration GRBs have been found to be associated with broadlined SNe Ic. There are four spectroscopically confirmed cases so far, namely GRB\\,980425 \\citep{Galama98}, GRB\\,030329 \\citep{Hjorth03, Matheson03, Stanek03}, GRB\\,031203 \\citep{Cobb04, Thomson04, Malesani04, Gal-Yam04} and GRB\\,060218 \\citep{Pian06,Sollerman06, Modjaz06, Mirabal06}. For all other nearby long GRBs up to 2006 where a SN could have been observed, additional light from the SN component was found in the late time lightcurves \\cite[e.g.,][]{Zeh04}. GRBs therefore offer a unique opportunity to observe a SN from the very onset of the explosion, as indicated by the prompt emission from the GRB. The connection between GRBs and broadlined SNe Ic, however, had to be revised in 2006 when two long-duration GRBs were found not to show any sign of an associated SN \\citep{Fynbo06, Gehrels06, DellaValle06, Gal-Yam06, Ofek07}.   One important supernova that recently caught the attention of the community is SN 2008D which was associated with an X-ray transient (XT) \\citep{Soderberg08}. This supernova was serendipitously detected while the XRT instrument onboard the {\\emph Swift} satellite \\citep{Gehrels04} observed another supernova, SN 2007uy, in the same galaxy. Whether this prompt XT originated from a weak GRB-like event or if it was due to the shock breakout from the star is still under debate \\citep{Soderberg08, Modjaz08b, Malesani08, Xu08}. The X-ray emission, however, is clearly associated with the onset of the explosion, and consequently SN 2008D was one of the earliest observed SNe. The SN spectrum evolved from a smooth spectrum with small undulations characteristic of a high-velocity ejecta into a typical SN Ib \\citep{Soderberg08, Modjaz08b, Malesani08}. Also the host galaxy of this supernova, NGC 2770, has attracted some attention as it produced three SNe within the last 9 years, SNe 1999eh, 2007uy and 2008D. Intriguingly, all three were stripped envelope core-collapse SNe Ib.   In this paper, we present a study on the local properties at the SN sites as well as of other regions in NGC 2770 with longslit spectroscopy and compare the host itself with a sample of other supernova producing galaxies. We want to investigate if the occurrence of three recent SNe Ib can be explained by some physical properties of the host galaxy. In Section 2 we present the observations and data reduction of the spectra, Section 3 studies the global properties of NGC 2770 and properties derived from modeling the spectral energy distribution (SED). In Section 4 we analyze the different regions in the host in terms of metallicity, extinction and star formation rate (SFR). Section 5 compares NGC 2770 to a sample of other nearby galaxies with several recent SNe and discuss the probability to find three SNe Ib in a galaxy within nine years.  Finally, Section 6 investigates the properties of the companion galaxy NGC 2770B. Throughout the paper we use a cosmology with H$_0$=70 km s$^{-1}$ Mpc$^{-1}$, $\\Omega_\\Lambda$=0.7 and $\\Omega_\\mathrm{m}$=0.3. At z = 0.007 this corresponds to 0.13 kpc per arcsecond and we use a distance of 27 Mpc to NGC 2700. ",
        "conclusions": "We have investigated the properties of the three SN sites in NGC 2770, the host of three SNe Ib, and the other regions in the host along 4 longslit positions. Previous observations in all wavelengths from UV to radio allow us to fit the SED of NGC 2770 and derive a range of global properties from it. We then set NGC 2770 in context to a sample of galaxies with frequent SN occurrence (3 or more SNe detected) and also compute the probability to detect three SNe Ib in a galaxy within only 10 years as it was the case for NGC 2770. From this analyses, we then conclude the following: \\begin{itemize} \\item NGC 2770 has global properties similar to the MW, even though it has a higher number of SNe observed. Its SFR and SNR are around the average of other nearby spiral galaxies and our sample of frequent SN galaxies. \\item The only outstanding property of NGC 2770 is its high HI mass which indicates a large reservoir for forming stars. \\item The metallicities at the SN sites in NGC 2770 are around 0.5 solar which is similar to the values observed for nearby GRB sites, but lower than for broadline SN Ic sites. \\item Almost half of the galaxies with SNe II and Ib/c only have one type of SN which is likely connected to the age of the dominant stellar population. Galaxies producing SNe Ib/cs have a higher deVaucouleurs number than those producing SN II. \\item SN and GRB hosts seem to be somewhat different in terms of SFRs and masses. \\item The probability to detect 3 SNe Ib in a galaxy is 0.6 to 1.5 \\% assuming 10,000 monitored galaxies. \\end{itemize} It therefore seems to be likely that observing 3 SNe Ib in NGC 2770 was only a chance coincidence. NGC 2770 is by no means a special galaxy that would be predestined to produce only stripped-envelope SNe. In fact its properties are not typical for galaxies with (frequent) SN Ib/c occurrence. However, it might also imply that the local properties at the SN sites are more important, at least in some galaxies, to produce a certain type of SN than its global properties."
    },
    "0807/0807.0229_arXiv.txt": {
        "abstract": "Measuring the integrated stellar halo light around galaxies is very challenging. The surface brightness of these haloes are expected to be many magnitudes below dark sky and the central brightness of the galaxy. Here I show that in some of the recent literature the effect of very extended Point Spread Function (PSF) tails on the measurements of halo light has been underestimated; especially in the case of edge-on disc galaxies.  The detection of a halo along the minor axis of an edge-on galaxy in the Hubble Ultra Deep Field can largely be explained by scattered galaxy light. Similarly, depending on filter and the shape one assumes for the uncertain extended PSF, 20 to 80 per cent of the halo light found along the minor axis of scaled and stacked Sloan Digital Sky Survey (SDSS) edge-on galaxy images can be explained by scattered galaxy light. Scattered light also significantly contributes to the anomalous halo colours of stacked SDSS images. The scattered light fraction decreases when looking in the quadrants away from the minor axis. The remaining excess light is well modelled with a S\\'ersic profile halo with shape parameters based on star count halo detections of nearby galaxies. Even though the contribution from PSF scattered light does not fully remove the need for extended components around these edge-on galaxies, it will be very challenging to make accurate halo light shape and colour measurements from integrated light without very careful PSF measurements and scattered light modelling. ",
        "introduction": "In recent years we have begun to appreciate the importance of galaxy stellar envelopes as tracers of the hierarchical galaxy formation process. Hierarchical models in a $\\Lambda$CDM context predict that the stellar envelopes around galaxies are created from many disrupted satellites, where size, shape, amount of substructure, and metallicity of the envelope principally depend on the primordial power spectrum, the reionisation epoch, the star formation history of accreted dwarfs, and the total dark matter mass of the host galaxy \\citep[e.g.,][]{BekChi05,BulJoh05,AbaNav06,PurBul07}. The discovery of the Sagittarius Dwarf currently being disrupted by the Milky Way \\citep{Iba94} and the highly structured envelope around M31 \\citep[e.g.,][]{FerIrw02,Iba07} has given much credence to the hierarchical model. While these observations were performed using resolved stars, many measurements of galaxy haloes have been attempted using integrated light \\citep[e.g.,][]{MorBor94,Fry99,WuBur02,ZibWhi04,ZibFer04}. These observations are very difficult, as the halo light is typically at least a factor of 10$^{4}$ below sky level, and therefore careful attention has to be paid to flat fielding and sky subtraction. Here I investigate another effect that has sometimes been underestimated in integrated light studies: the effect of scattered light in extreme PSF tails when examining edge-on galaxies.  The effect of convolving a PSF with a spherical light distribution can be fairly well estimated. For instance, with an elliptical galaxy the light in the central region will be dispersed by the generally broader PSF shape. Further out, the light distribution is nearly unaffected as the PSF shape is steeper than the slowly declining galaxy profile. By over-plotting the PSF on the measured light distribution normalised to the same central brightness, one gets a fairly good impression of which radii are affected by light convolution as modelled by the PSF.  For edge-on galaxies the procedure is not as simple as the luminosity profile perpendicular to the disc can be steeper than the outer tails of the PSF. Close to the midplane of the galaxy the vertical light distribution will still be modified by a point-source-like convolution; however, farther away we are less dominated by the local light. Instead the light scattered from the whole disc --- not just the central core--- significantly contributes to the measured distribution at large scale heights. To estimate this contribution at more than 2 scale lengths above the disc (about 6--15 scale heights) one should not use a PSF profile normalised to the central surface brightness of the edge-on galaxy, but instead normalised to the total brightness of the galaxy. Calculating the actual scattered light contribution at any point is obviously best determined by convolving a 2D model of the intrinsic light distribution with a full 2D PSF.   Another problem arises when studying galaxy haloes from stacked galaxy images, scaled to a common size, in order to reach fainter levels. One can estimate the PSF of the stack by combining appropriate stellar images from each field, scaled by the same factor as the galaxy in the frame. However, for a typical sample selection the distribution of scaling is not symmetric and highly biased toward the scales near the selection limit of the sample. The distribution of scaled PSF images is therefore strongly skewed, meaning that combining the PSFs using median or mean values gives very different results. Assuming that the galaxies are very similar after scaling, no such skewed distribution is present and median and mean combining should give very similar results.  \\begin{figure*} \\includegraphics[width=\\linewidth,trim=1.5cm 0.5cm 1.5cm 0.5cm,clip=true]{udf_profs.eps} \\caption{Minor (left) and major (right) axis surface brightness profiles of the edge-on HUDF galaxy. The data points are derived from \\citet{ZibFer04}. The profiles for the different bands are coded in different colours and symbols and are offset in order to avoid confusion, as indicated by the legend. The input galaxy model is shown by black dashed lines for the different passbands. Thin solid lines indicate crosscuts through the $B$ and $z$-band PSFs (arbitrarily normalised), whilst thick lines show the convolved galaxy model using the same colours for the different passbands as used for the data points. \\label{UDFprofs} } \\end{figure*}   In this research note I investigate a few cases in the literature where the effects of PSF convolution with edge-on galaxies has been underestimated. I will also show that the colour measurements of the envelopes around galaxies can be strongly affected by scattered light. In particular, in Section~\\ref{HUDF} I will show the effect of PSF convolution on the light distribution around an edge-on galaxy in the Hubble Ultra Deep field \\citep[HUDF;][]{BecSti06} as measured by \\citep{ZibFer04}. In Section~\\ref{SDSS} I estimate the scattered light effect on the stacked SDSS image analysis of \\citet{ZibWhi04}. Finally, in Section~\\ref{Concl} I summarise my conclusions and give a few recommendations. ",
        "conclusions": "\\label{Concl}  Measuring the exact halo light distribution around edge-on disc light from the integrated light distribution is very challenging. In addition to careful flat fielding and removal of fore- and background sources, particular attention has to be paid to the effect of scattered light from the central disc. Specifically, for edge-on galaxies it is not sufficient to compare the galaxy minor axis light profile with the PSF profile; at large heights above the disc scattered light is contributed from the whole galaxy, not just the centre, and a full convolution of a 2D model is needed to assess the scattered light contamination. Basic modelling suggests that the minor axis may not be the optimum place to detect a halo, but for a flattened halo one should look in the quadrants away from the major and minor axis where scattered light contribution from the disc is smaller relative to the halo.  I find that the extended component seen along the minor axis of an HUDF edge-on galaxy can almost fully be explained by scattered light. The observations of \\citet{ZibFer04} for this galaxy do show an extended component along the major axis that cannot be explained by scattered light. Similarly, the halo detection around stacked SDSS edge-on galaxies reported by \\citet{ZibWhi04} has a significant scattered light contamination, especially along the minor axis and in the $i$-band. Depending on the assumed best method to create a stacked PSF the contamination along the minor axis in the $i$-band can amount to 80\\% of the light, in the $g$ and $r$-bands 50\\%. The contamination decreases with other PSF assumptions and going to one of the quadrants, but the PSF uncertainties make it hard to derive accurate stellar halo properties. The excess light seen on top of the disk only model is consistent with a S\\'ersic law halo with parameters typical of those found by the GHOSTS survey \\citep[][de Jong et al., in preparation]{deJRad07}.  The anomalous minor axis halo colours reported by \\citet{ZibFer04} and \\citet{ZibWhi04} are consistent with originating from scattered light contamination. In particular, the colour gradients reported by \\citet{ZibWhi04} can be modelled by an extended PSF as calculated from the average of the observed scaled SDSS PSFs. There seems to be no need to invoke extreme stellar population to explain these colours as proposed by \\citet{ZacBer06}. The expected contribution from a typical GHOSTS halo is so small in the $i$-band that it has only a small effect on the predicted colour gradient.  There have been other reports of anomalous galaxy stellar halo colours. \\citet{LeqFor96,LeqCom98} report thick disk or halo {\\em BVI}-band colours of edge-on galaxy NGC\\,5907 that are only consistent with colours of elliptical galaxies, i.e.\\ an old, metal rich stellar population. In this case the errors are unlikely to be due to scattered light; the thick disk is detected at about a factor 10$^3$ below the central brightness at about 80 arcsec above the midplane, while the PSF is already 10$^7$ below peak brightness at 16 arcsec radius. Not even a factor 10 enhancement of the scattered light effect due to the edge-on disc configuration can make up that difference. Similarly, the red haloes reported around Blue Compact Galaxies \\citep[see e.g., ][and references therein]{ZacBer06} are not caused by scattered light. The galaxies are too large compared to the PSF and the haloes show substructure different from the central galaxy that is unlikely to be due to the substructure in the PSF.  Unfortunately, the examples of \\citet{ZibFer04} and \\citet{ZibWhi04} are close to the limit where scattered light is no longer significant and slightly larger objects would have avoided problems. The HUDF does not contain a larger edge-on galaxy that would render scattered light unimportant. However, the SDSS study could now be repeated on a sample of larger galaxies as the area surveyed by SDSS has nearly tripled since the \\citet{ZibWhi04} study. Ideally, such a study would use only a small range of scale sizes to avoid large spatial scaling corrections and uncertainties in the effective PSF. Taking about a two times as large an isophotal selection limit would probably suffice, especially in the $g$ and $r$-bands.  While measuring halo properties from integrated light remains plagued with large uncertainties, we do have a technique that allows us to accurately measure the stellar envelopes around galaxies. By performing star counts of resolved stellar populations we can measure equivalent surface brightnesses to very faint limits as spectacularly demonstrated for M31 and M33 \\citep[e.g.,][]{Iba07}. These observations are not affected by scattered light and only minimally by flat fielding errors. The main limitations of this method lie in fore- and background contaminating objects and low number statistics (too few brights stars at very low surface brightnesses/densities). Very few massive, highly inclined galaxies can be studied with ground-based resolution, but HST allows a much larger sample of galaxies to be studied. Indeed, preliminary results from the GHOSTS survey show conclusively that most large galaxies do have very extended stellar envelopes, with envelope size likely depending on galaxy mass and bulge-to-disc ratio \\citep[][de Jong et al., in preparation]{deJRad07}."
    },
    "0807/0807.1533_arXiv.txt": {
        "abstract": "We describe a new program for determining photometric redshifts, dubbed EAZY. The program is optimized for cases where spectroscopic redshifts are not available, or only available for a biased subset of the galaxies. The code combines features from various existing codes: it can fit linear combinations of templates, it includes optional flux- and redshift-based priors, and its user interface is modeled on the popular HYPERZ code. A novel feature is that the default template set, as well as the default functional forms of the priors, are not based on (usually highly biased) spectroscopic samples, but on semi-analytical models. Furthermore, template mismatch is addressed by a novel rest-frame template error function. This function gives different wavelength regions different weights, and ensures that the formal redshift uncertainties are realistic. We introduce a redshift quality parameter, $Q_z$, that provides a robust estimate of the reliability of the photometric redshift estimate. Despite the fact that EAZY is not \"trained\" on spectroscopic samples, the code (with default parameters) performs very well on existing public datasets. For $K$-selected samples in CDF-South and other deep fields we find a $1\\sigma$ scatter in $\\Delta z/(1+z)$ of 0.034, and we provide updated photometric redshift catalogs for the FIRES, MUSYC, and FIREWORKS surveys. ",
        "introduction": "Accurate redshifts of distant galaxies are crucial for nearly all of observational cosmology. Whereas extensive spectroscopy with multi-object spectrographs on 8-10m class telescopes has yielded redshifts for thousands, and in some cases tens of thousands, of galaxies \\citep[e.g.][]{steidel:03,deep2,vvds}, these galaxies tend to be relatively bright at optical wavelengths. For galaxies fainter than $R\\sim 25$ we rely almost exclusively on photometric redshifts, derived from fitting template spectra to broad- or medium-band photometry \\citep[e.g.][]{lanzetta:96,wolf:03,franx:03,mobasher:04,drory:05}. This situation is not likely to change, even with the advent of efficient spectrographs with very wide fields \\citep[such as WFMOS;][]{wfmos}, multi-object capabilities in the near-infrared \\citep[e.g. MOIRCS;][]{moircs}, or larger telescopes. The signal-to-noise ratio (S/N) per resolution element in the continuum decreases with spectral resolution as ${\\rm S}/{\\rm N} \\propto R^{-0.5}$ for a given exposure time. Therefore, the required integration time to maintain a given S/N per resolution element increases linearly with the spectral resolution, quite independent of the details of the telescope and instruments. As a typical set of broad band filters corresponds to $R\\sim 5$ and typical faint object spectrographs have $R\\sim 1000$, spectroscopy is about two orders of magnitude more time consuming than photometry for a given telescope size. A notable exception is spectroscopy of emission line objects, which can be extremely efficient at faint magnitudes.  The methodology for determining photometric redshifts using the template-fitting approach is essentially straightforward: the photometric data are compared to synthetic photometry for a large range of template spectra and redshifts, and the most likely redshift follows from a statistical analysis of the differences between observed and synthetic data. Several codes exist that perform this task, each employing its own techniques for creating the synthetic photometry and interpreting the residuals in the redshift -- template plane. Popular examples include HYPERZ \\citep{hyperz}, \\textsc{ImpZ} \\citep{impz}, and Le PHARE\\footnote{\\texttt{http://www.oamp.fr/people/arnouts/LE\\_PHARE.html}} (Arnouts \\& Ilbert), which do a straightforward $\\chi^2$ minimization; GREGZ\\footnote{Greg Rudnick did not name his code; the name GREGZ is used for convenience in the present paper.} \\citep{rudnick:01,rudnick:03}, which allows linear combinations of templates and uses Monte Carlo methods to determine the redshift uncertainties; BPZ \\citep{bpz}, which uses Bayesian statistics allowing the use of priors; and ZEBRA \\citep{zebra} and \\texttt{kcorrect} \\citep[][hereafter BR07]{blanton:07}, which include (distinct) iterative template-optimization routines that make use of the extensive spectroscopic databases of the zCOSMOS \\citep{lilly:07} and Sloane Digital Sky Survey \\citep[SDSS; ][]{york:00} projects, respectively.  For obvious reasons photometric redshifts benefit from having high quality photometry in many bandpasses and from sampling strong continuum features in the observed wavelength region (such as a Lyman or Balmer break), irrespective of the methodology. However, given a set of objects with good quality photometry, the aspect that is of paramount importance for obtaining reliable photometric redshifts is the selection of the template set (see Feldmann et al.\\ 2006, \\S\\ref{s:template_set}). \\cite{zebra} obtain very good results by iteratively adapting the templates, minimizing the systematic differences between the best fitting templates and the actual galaxy photometry. This approach not only reduces the random uncertainty in the photometric redshifts but can also eliminate systematic effects in certain redshift ranges \\citep[see][]{zebra}. The disadvantage of this optimization is that its effects can only be assessed when a large sample of galaxies with spectroscopic redshifts is available, and when this sample is a random subset of the entire photometric sample. This assumption may be valid in the case of zCOSMOS, but this is generally not the case in studies of galaxy samples which are significantly fainter than the spectroscopic limit.  In this paper we describe a new photometric redshift code which was written specifically for samples with incomplete and/or biased spectroscopic information (such as, for example, faint $K$-selected samples). Rather than minimizing the scatter in the relation between photometric and spectroscopic redshift using the spectroscopic sample as a training set, a user-defined template error function is introduced to account for wavelength-dependent template mismatch. The code combines features from various existing programs: the possibility of fitting linear combinations of templates (as done in GREGZ), the use of priors (as first done in BPZ), and a user-friendly interface based on HYPERZ. The default template set and the redshift-magnitude priors are derived from semi-analytical models. These models are, of course, only an approximation of reality, but their ``perfect'' completeness down to very faint magnitudes outweighs their imperfect representation of the real Universe.  The outline of this paper is as follows.  In \\S\\ref{s:implementation} we describe the implementation of the code, including the optimized template set and redshift priors derived from semi-analytical models and the template error function derived from the GOODS-CDFS photometric catalog.  In \\S\\ref{s:application} we test the code on a combined photometric catalog from a variety of deep multi-wavelength surveys and compare the photometric redshifts to spectroscopic redshifts of nearly 2000 galaxies at $0 < z < 4$.  In \\S\\ref{s:reliability} we discuss the reliability of the photometric redshift estimates and provide cautionary examples for relying solely on spectroscopic samples to estimate the photometric redshift quality. Finally, in \\S\\ref{s:summary} we summarize the features and performance of the photometric redshift code and discuss future avenues for improvement. ",
        "conclusions": ""
    },
    "0807/0807.0791_arXiv.txt": {
        "abstract": "% We review VLBI observations of supernovae over the last quarter century and discuss the prospect of imaging future supernovae with space VLBI in the context of VSOP-2. From thousands of discovered supernovae, most of them at cosmological distances, $\\sim$50 have been detected at radio wavelengths, most of them in relatively nearby galaxies. All of the radio supernovae are Type II or Ib/c, which originate from the explosion of massive progenitor stars. Of these, 12 were observed with VLBI and four of them, SN 1979C, SN 1986J, SN 1993J, and SN 1987A, could be imaged in detail, the former three with VLBI.  In addition, supernovae or young supernova remnants were discovered at radio wavelengths in highly dust-obscured galaxies, such as M82, Arp 299, and Arp 220, and some of them could also be imaged in detail. Four of the supernovae so far observed were sufficiently bright to be detectable with VSOP-2.  With VSOP-2 the expansion of supernovae can be monitored and investiated with unsurpassed angular resolution, starting as early as the time of the supernova's transition from its opaque to transparent stage.  Such studies can reveal, in a movie, the aftermath of a supernova explosion shortly after shock breakout. ",
        "introduction": "A supernova (SN), the explosion of a star, is one of the most energetic single events in the universe.  Thousands of optical SNe are now known but most of them are at cosmological distances.  Only $\\sim$50 have been detected at radio wavelengths and 12 observed with VLBI, the most distant at almost 100 Mpc. In addition, several SNe and supernova remnants (SNRs) were discovered at radio wavelengths in dust-obscured galaxies, such as M82, Arp 299, and Arp 220, and observed with VLBI. All of the SNe observed with VLBI are thought to have resulted from the core collapse of a massive progenitor and emit synchrotron radiation generated from the electrons accelerated in the region where the shock front interacts with the circumstellar medium (CSM) left over from the wind of the mass-losing progenitor.  In addition, some of the radiation may come from the environment of the stellar corpse, a neutron star or a black hole left over from the explosion.  During the early stage of the expansion of a radio SN, the radio lightcurve rises quickly because of the decreasing optical depth due to synchrotron self-absorption or to thermal absorption in the ionized CSM along the line of sight. It reaches its peak, first at high frequencies later at lower frequencies, when the optical depth decreases below unity, and it declines thereafter during the optically thin stage of the SN while the SN is expanding \\citep[e.g.,][]{Chevalier1982a}.  For a typical shock front expansion velocity of 10,000 km s$^{-1}$, the angular size of a SN at a distance of 4 Mpc expands at a rate of 1 mas yr$^{-1}$. With global VLBI and a wavelength of 4 cm, an angular resolution of $\\sim$0.6 mas can be obtained, allowing detailed investigations of the expanding SN and the making of a movie. Global VLBI with VSOP-2 promises to increase the angular resolution at this wavelength to 0.2 mas and to a correspondingly higher resolution at shorter wavelengths, enabling us to witness the aftermath of the explosion shortly after shock breakout. ",
        "conclusions": ""
    },
    "0807/0807.0758_arXiv.txt": {
        "abstract": "Models of magnetospheric accretion on to classical T Tauri stars often assume that stellar magnetic fields are simple dipoles.  Recently published surface magnetograms of BP~Tau and V2129~Oph have shown, however, that their fields are more complex.  The magnetic field of V2129~Oph was found to be predominantly octupolar.  For BP~Tau the magnetic energy was shared mainly between the dipole and octupole field components, with the dipole component being almost four times as strong as that of V2129~Oph.  From the published surface maps of the photospheric magnetic fields we extrapolate the coronal fields of both stars, and compare the resulting field structures with that of a dipole.  We consider different models where the disc is truncated at, or well-within, the Keplerian corotation radius. We find that although the structure of the surface magnetic field is particularly complex for both stars, the geometry of the larger scale field, along which accretion is occurring, is somewhat simpler.  However, the larger scale field is distorted close to the star by the stronger field regions, with the net effect being that the fractional open flux through the stellar surface is less than would be expected with a dipole magnetic field model.  Finally, we estimate the disc truncation radius, assuming that this occurs where the magnetic torque from the stellar magnetosphere is comparable to the viscous torque in the disc. ",
        "introduction": "Classical T Tauri stars (cTTSs) are young solar analogs which are accreting material from circumstellar discs.   Many observations are consistent with the scenario of the stellar field truncating the inner disc and channelling gas onto discrete regions of the stellar surface.   The shapes of near-IR spectral energy distributions (e.g. \\citealt{rob07}), and the kinematics of CO lines formed in the disc \\citep*{naj03}, are consistent with the disc having been truncated at a distance of a few stellar radii.  Average surface fields of order a kG have been detected on a number of cTTSs \\citep{joh07}, which models suggest will be strong enough to disrupt the inner disc \\citep{kon91}.  The detection of inverse P-Cygni profiles, with widths of several hundred kms$^{-1}$, can also be explained by material essentially free-falling along the field lines of the stellar magnetosphere from the location of the inner disc \\citep{edw94}.  Furthermore, the excess continuum emission (veiling) in the optical and UV is likely to arise from shocks at the base of accretion funnels, with emission lines with high excitation potentials, e.g. HeI 5876{\\AA }, forming mainly in such regions \\citep{ber01}.  The magnetic interaction between the stellar field and the disc may also have important consequences for the formation of planets.  Simulations by \\citet{rom06}, and analytic work by \\citet*{lin96} and \\citet{fle08}, suggest that the inner disc hole, cleared by the star-disc interaction, may provide a natural barrier that decreases the rate of inward migration of forming planets.  There is also evidence that the large scale magnetosphere may directly disrupt the inner disc, producing warps which in some systems cross the observer's line-of-sight to the star (e.g. AA Tau, \\citealt{bou07}).  Indeed 3D MHD simulations have demonstrated that complicated warping effects in the disc truncation region may be common for T Tauri stars \\citep{rom03,rom04}.  The star-disc interaction may also explain the slower rotation of cTTSs compared to the typically older weak-line T Tauri stars, whose discs have largely dispersed (see e.g. the review by \\citealt{bou07IAU}).  Accretion of material from the inner disc would act to spin-up the central star in the absence of some physical mechanism to remove angular momentum from the system. Various magnetospheric accretion models have been developed to explain this mechanism, which differ in their assumed location of the inner disc, the details of how angular momentum is removed, and how the magnetic field topology controls both accretion and outflows (\\citealt{kon91}; \\citealt{col93}; \\citealt{shu94}; \\citealt*{fer00}; \\citealt*{fer06}; \\citealt*{kuk03}; \\citealt{mat05a,mat05b,mat08a,mat08b}; \\citealt*{lon05}; \\citealt{bes08}).  Although observations indicate that the magnetic field topologies of cTTSs are complex (e.g. see the discussion in \\citealt{gre06a}), the majority of accretion models simply assume that the star has a strong dipolar magnetic field.  Recently however a few studies have dropped this assumption, with \\citet{gre05,gre06a} and \\citet{jar06} being the first to consider how magnetic fields with a realistic degree of complexity would affect the accretion process. Other field geometries have since been considered by \\citet*{lon07,lon08} who present the results of a 3D MHD simulation of accretion to stars with composite fields that consist of dipole and quadrupole components tilted at various angles relative to the stellar rotation axis.  It is only recently, however, that instrumentation has advanced to the stage where it is possible to map the magnetic topologies of classical T Tauri stars.  \\citet{don07,don08a} have recently published magnetic surface maps derived from Zeeman-Doppler imaging of the young solar-analogs V2129~Oph and BP~Tau, using data from the ESPaDOnS and NARVAL spectropolarimeters at the Canada-France-Hawai'i Telescope and T\\'elescope Bernard Lyot respectively. Using the pre-main sequence evolutionary model of \\citet*{sie00}, \\citet{don07} determined that V2129~Oph, a high-mass T Tauri star of $1.35\\,{\\rm M}_{\\odot}$, had developed a radiative core.  In contrast, BP~Tau, at $0.7\\,{\\rm M}_{\\odot}$, should be completely convective \\citep{don08a}.  As the internal structure, and therefore the magnetic field generation process, is different in both V2129~Oph and BP~Tau, it is of particular interest to compare the structure and properties of their magnetic fields, as well as comparing both to a dipole.  In \\S2 we summarise the stellar parameters and magnetic field measurements of BP~Tau and V2129~Oph. In \\S3 we describe how three-dimensional coronal fields can be extrapolated from Zeeman-Doppler maps of photospheric magnetic fields, and compare the resulting field topologies with a large scale dipole.  In \\S4 and \\S5 we discuss the structure of the open stellar, and accreting, field, while \\S6 contains our conclusions. ",
        "conclusions": "By extrapolating the fields of two T Tauri stars from observationally derived surface magnetograms we have compared the resulting magnetic field topologies with a simple dipole.  We have found that although the surface magnetic fields of both stars are particularly complex, the larger scale field is simpler and more well ordered. This is consistent with previously published spectropolarimetric studies of accreting T Tauri stars, which suggested that bundles of accreting field lines connect to the stellar surface in single polarity regions, and were therefore likely well-ordered (e.g. \\citealt{val04}). The larger scale field of BP~Tau in particular is closely matched to a dipole magnetic field, although V2129~Oph still shows departures from a dipole field even on the largest scales.  However, for both stars, the footpoint separation at the stellar surface of the largest scale field lines is greater than for a purely dipolar magnetic field.  In other words as the largest scale field lines arrive at the stellar surface, their structure is influenced by the much stronger field regions. The effect of this is that the largest scale field connects to the star at higher latitudes than would be expected when using a dipole magnetic field model.  Thus by considering complex magnetic field topologies we find that the fractional open flux through the stellar surface is less than would be expected from dipole field models.  The reduction in the fractional open flux is mainly attributable to the larger closed flux for the complex fields, which arises due to the numerous small magnetic loops close to the stellar surface.  For both V2129~Oph and BP~Tau there is less total open flux than for a dipole field, however, the difference is small enough (see Fig. \\ref{fluxes}) that there is unlikely to be any significant implications for the amount of angular momentum that can be carried away from such systems by a stellar wind \\citep{mat05b,mat08a,mat08b}.  However, more spectropolarimetric data for many different T Tauri stars will be required to confirm this.  If the extent of stellar magnetospheres are limited to well within corotation then accretion may proceed along the more complex regions of the stellar surface field.  The result of this would be the formation of many low-latitude accretion hotspots \\citep{gre05,gre06a}.  There is at least one star which shows evidence for such equatorial spots, GQ~Lup \\citep{bro07}.  However, for V2129~Oph and BP~Tau the spectroscopic data of \\citet{don07,don08a} clearly showed that the bulk of accreting material is carried into high latitude regions of the stellar surface. In order to explain such high latitude hotspots, \\citet{don07} and \\citet{jar08} demonstrated that for V2129~Oph the source surface must be set to at least 6.0$\\,{R_{\\ast}}$, while for BP~Tau the stellar magnetosphere must extend to at least 4$\\,{R_{\\ast}}$ \\citep{don08a}.  Selecting the location of the source surface sets the structure of the stellar magnetic field, although the inner disc may still be truncated closer to the stellar surface.  Using the assumption that inner disc is located at the point where the torque due to viscous processes in the disc is comparable to the magnetic torque due to the stellar magnetosphere we estimated the disc truncation radius using the extrapolated fields.  For BP~Tau we found that the inner disc would be truncated close to the radius predicted for a dipole field, whereas for V2129~Oph, the inner disc is likely to be closer to the star. However, with large uncertainties in the assumed mass accretion rate (which influences the torque due to viscous processes in the disc), and with the static field structures considered here, such calculations are only of qualitative value.  Further 3D MHD simulations, such as those already performed by, for example \\citet{lon08}, are required in order to model how magnetic fields with an observed degree of complexity can disrupt and influence the structure of planet forming discs, and will represent a major future development in this field.  Although the radius of the source surface is essentially a free parameter of our model, the observed hotspot locations provide constraints.  Once the source surface has been selected, however, this then also sets the structure of the X-ray emitting surface field, allowing predictions to be made regarding the stellar X-ray emission properties \\citep{jar08}.  Such predictions will allow direct testing of our field extrapolation model with future independently obtained X-ray observations. A further test of the model will be the simulation of accretion related emission line profiles, and are currently being undertaken.  V2129~Oph, despite its young age, is massive enough to have developed a radiative core.  In contrast to this BP~Tau is likely to be completely convective.  Although the field of BP~Tau is more complex than a dipole, it is much simpler than that of V2129~Oph.  The differences in the field structure of both stars likely reflect the different internal structure and the process by which their magnetic fields are generated and maintained.  For stars with radiative cores the magnetic field is likely generated in the shear layer between the core and the outer convective envelope.  In the absence of such a shear layer, it is difficult to explain how fully convective stars can generate and maintain large-scale almost axisymmetric kilo-Gauss fields, although recent theoretical models suggest that this is possible \\citep{bro08}.  With only three magnetic surface maps available to date, V2129~Oph and two of BP~Tau \\citep{don07,don08a}, it is not possible to determine whether this is a general feature of all T Tauri stars, however it is similar to what it found for low mass main sequence stars \\citep{don08b}. More magnetic surface maps of T Tauri stars with varying stellar parameters are now required to test these ideas fully, and are an essential requirement to advance our understanding of stellar magnetism on the pre-main sequence."
    },
    "0807/0807.2088_arXiv.txt": {
        "abstract": "We present a new survey for pulsars in the error boxes of the low-latitude EGRET sources 3EG J1027$-$5817, 3EG J1800$-$2338 and 3EG J1810$-$1032. Although all of these sources have been covered by previous pulsar surveys, the recent discovery of the young, energetic pulsar PSR~J1410$-$6132 at 6.7~GHz has shown that pulsars of this type can be hidden from low frequency surveys. Using an observing frequency of 3.1~GHz we discovered a 91-ms pulsar, PSR J1028$-$5819, which observations made at the Parkes telescope and the Australia Telescope Compact Array have shown to be young and energetic. We believe this pulsar is likely to be powering the unidentified EGRET source 3EG J1027$-$5817. Like other energetic pulsars, PSR J1028$-$5819 is highly linearly polarised, but astonishingly has a pulse duty cycle of only 0.4\\%, one of the smallest in the entire pulsar catalogue. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0807/0807.0334_arXiv.txt": {
        "abstract": " ",
        "introduction": "\\noindent Evolution of galaxies is directed both by external and internal factors. Internal factors include mainly dynamical instabilities and depend on mass and angular momentum of a galaxy. External factors are determined by the galaxy environments -- other galaxies, intergalactic medium -- and can be either gravitational or hydrodynamical mechanisms. Galactic groups are the best place to study the external factors of galaxy evolution. Many of them have X-ray halo of hot gas; this high-pressure intergalactic medium can strip the own gas reservoirs of spiral galaxies and transform them into lenticulars. Gravitational interactions are provided by the tight neighborhood.  Galaxy velocity dispersions inside the groups are not so large, of order of the intrinsic stellar velocity dispersions of early-type galaxies, and cannot prevent development of tidal effects. Numerical simulations proved that tidal interactions affect not only external parts of galaxies but also stimulate bar formation in the centers which in turn re-distribute matter along the radius. Secular evolution can even change the morphological type of a galaxy by causing a growth of a bulge by resulting in gas inflow and subsequent star formation in the center; a whole class of bulges, named `pseudobulges', may be formed rather recently by such a mechanism (Kormendy and Kennicutt 2004).  If several galaxies are settled within similar environments in a center of a group and mutually interact gravitationally, it seems probable that the re-shaping of their structures and formation of new central stellar components should be synchronous. Observationally, it means that the stellar ages would be the same, and the spatial star distributions would be similar. If observations reveal synchronous evolution of the central parts of the group galaxies, it would be an argument in favour of dominance of external mechanisms; if no -- we should put attention to variance of internal conditions in the galaxies under consideration. We have already studied 5 nearby groups with the Multi-Pupil Fiber Spectrograph (MPFS) of the Russian 6m telescope; in each group we have observed central parts of 2--3 centrally located galaxies. In Leo I (Sil'chenko et al. 2003a) and in the groups around NGC~5576 (Sil'chenko et al. 2002) and NGC~3169 (Sil'chenko and Afanasiev 2006) we have found similar properties of the stellar populations in the circumnuclear stellar disks: they all have been formed rather recently, 1--3 Gyr ago; and in 3 dominant galaxies of Leo I they are almost co-spatial. In the Leo Triplet, on the contrary, the stellar ages and kinematics in the centers of NGC~3623 and NGC~3627 were quite different (Afanasiev and Sil'chenko 2005), and we concluded that the galaxies of the Triplet had met recently, not earlier than 1 Gyr ago. The only group with the X-ray halo among those studied by us is Leo II (Afanasiev and Sil'chenko 2007); its two central galaxies, NGC~3607 and NGC~3608, have both old kinematically decoupled stellar structures, but their mean stellar ages are different, 10 and 6 Gyr respectively. In this paper we present results of our study of stellar populations in galaxies of another massive group with X-ray hot gas -- the group of NGC~80.  A group of galaxies around the giant lenticular galaxy NGC 80 is a rich and massive one. Ramella et al. (2002) in their catalog of galactic groups mentioned 13 members within 2 magnitudes from the brightest one, Dell'Antonio et al. (1994) estimated the number of galaxies in the group as 21, and Mahdavi et al. (2004), with more data, -- even as 45. The X-ray luminosity of the group, according to Mahdavi et al. (2000), is $\\log L_x = 42.56 \\pm 0.09$ ($L_x$ is expressed in erg/s); it is rather high for a group. The center of the X-ray brightness distribution coincides practically with NGC 80, the most luminous galaxy in the group. According to Mahdavi et al. (2004), the velocity dispersion of the galaxies in the NGC 80 group is 336 km/s and the systemic velocity is $5771 \\pm 48$ km/s; Dell'Antonio et al. (1994) rejected some non-member galaxies and after that obtained 246 km/s and $5663 \\pm 51$ km/s, respectively. The central galaxy of the group, the S0 NGC 80, has its own line-of-sight velocity of 5698 km/s which is close to the systemic velocity of the group; it is consistent with NGC 80 being a dynamical center of the group. Interestingly, another giant galaxy, a E0 NGC 83, having the same luminosity as NGC 80 and being located very nearby the group center, has a line-of-sight velocity exceeding the systemic group velocity by more than 500 km/s. One more luminous galaxy in the group is a spiral galaxy NGC 93. Bothun and Schommer (1983) observed NGC 93 at the 21 cm, in the neutral-hidrogen emission line, and they found that the galaxy is very massive, with the rotation velocity of 317 km/s, and `anemic', that means very red and with the low $M_{HI}/L_B$ ratio -- just as massive spiral galaxies of the Virgo cluster. Spiral galaxies in clusters are thought to be `anemic' because of effect of the surrounding hot intergalactic medium -- the most probable cause is the stripping of neutral hydrogen from the outer disks of galaxies by ram pressure. Since the group of NGC 80 is rich by hot gas, NGC 93 may be also stripped in such a manner.  We have made integral-field spectroscopic observations of the central parts of the most luminous galaxies in the group, NGC 80, NGC 83, and NGC 93, as well as of E-galaxy NGC 79 and S0-galaxies NGC 86, IC 1541, and IC 1548; the stellar population properties and stellar and gaseous (if the gas is present) kinematics have been studied. The galaxies studied are distributed from the very center (NGC 80 and NGC 83) to the most peripheric parts, at 0.5 Mpc from the group center (IC 1548 and IC 1541). Their global parameters taken from the NED and HYPERLEDA are given in the Table 1.  {\\normalsize \\begin{table*} \\caption[ ] {Global parameters of the galaxies studied} \\begin{flushleft} \\begin{tabular}{lccccccc} \\hline\\noalign{\\smallskip} NGC(IC) & 80 & 83 & 79 & 93 & 86 & (1541) & (1548) \\\\ Morph. type (LEDA$^1$) & SA$0-$ & E & E & S? Sab & Sa & S0 & S0 \\\\ $D_{25},\\, ^{\\prime}$  (LEDA) & 1.82 & 1.29 & 0.81 & 1.35 & 0.76 & 0.76 & 0.69 \\\\ $B_T^0$ (LEDA) & 12.98 & 13.22 & 14.66 & 13.66 & 14.35 & 15.15 & 15.20\\\\ $M_B$ (LEDA) & --21.6 & --21.6 & --19.8 & --20.8 & --20.2 & --19.5 & --19.4 \\\\ $(B-V)_e$ (LEDA) & 1.07 & 1.12 & -- & 1.16 & -- & -- & --  \\\\ $(U-B)_e$ (LEDA) & 0.66 & 0.60 & -- & 0.61 & -- & -- & -- \\\\ $V_r$, km/s (NED$^2$) & 5698 & 6227 & 5485 & 5380 & 5591 & 5926 & 5746 \\\\ Distance, Mpc & \\multicolumn{7}{c}{77 (Mahdavi \\& Geller 2004)} \\\\ $PA_{phot}$ (LEDA) & $4.5^{\\circ}$ & $112^{\\circ}$ & $163^{\\circ}$ & $49^{\\circ}$ & $8^{\\circ}$ & $36^{\\circ}$ & $78^{\\circ}$  \\\\ $\\sigma _*$, km/s (LEDA) & 260 & 262 & 201 & -- & -- & -- & 154 \\\\ \\hline \\multicolumn{8}{l}{$^1$\\rule{0pt}{11pt}\\footnotesize NASA/IPAC Extragalactic Database}\\\\ \\multicolumn{8}{l}{$^2$\\rule{0pt}{11pt}\\footnotesize Lyon-Meudon Extragalactic Database} \\end{tabular} \\end{flushleft} \\end{table*} } ",
        "conclusions": "\\noindent We have studied 7 early-type (E--Sab) galaxies which are members of the massive X-ray galaxy group NGC 80, by means of 2D spectroscopy at the 6m telescope. We have searched for consequences of their synchronous secular evolution. Five of seven galaxies have old bulges, with the mean stellar ages of 10--15 Gyr. Concerning the central galaxy of the group, NGC~80, we have found earlier (Sil'chenko et al. 2003b) that it has circumnuclear structures of intermediate age, namely, the nucleus and the ring with the radius of about 2 kpc, but they are not younger than 5 Gyr. The `anemic' giant spiral galaxy NGC~93, which neutral gas has been probably stripped away partly by ram pressure of the hot intragroup medium at the moment of galaxy infall, has the mean stellar age of the chemically distinct nucleus of 4 Gyr. If to relate its nuclear star formation burst having produced the chemically distinct nucleus with the galaxy infall into the group, we would obtain that the moment of the main group gathering occured about 4--5 Gyr ago.  However there are some early-type galaxies in this group that possess rather young stellar populations. The S0-galaxy IC~1548 reveals clear signatures of the recent central star formation burst: the mean stellar age in the bulge is 3 Gyr, and that in the nucleus is 1.5 Gyr. In the same galaxy we have found a circumnuclear polar gaseous ring which warps smoothly into the counterrotating gaseous disk coplanar with the stellar one in more outer parts. Probably, IC~1548 suffered strong tidal interaction with the nearby late-type spiral galaxy and perhaps even accreted gas from this neighbor; once acquiring the gas, IC 1548 had to be provoked to the central star formation by the same tidal interaction.  The giant E0 galaxy NGC~83 possesses a compact massive stellar-gaseous disk with the radius of some 2 kpc, very fast rotating, with the noticeable current star formation. Perhaps, all these properties are consequences of merging a spiral or an irregular gas-rich galaxy, so called `minor merger'. It is interesting that the line-of-sight systemic velocity of NGC~83, a giant galaxy in the center of the group, exceeds far the systemic LOS velocity of the group, by more than two group velocity dispersion values. Only a few nearby group members -- NGC 81, NGC 85, PGC 1327 -- have LOS velocities close to that of NGC~83. The majority of the group members have LOS velocities that differ from that of NGC~80 by no more than 200--300 km/s. We may suggest that the poor subgroup of NGC~83 has been accreted by the group of NGC~80 only recently. Was NGC~83 an elliptical galaxy before the accretion event? Did its central gas belong to NGC~83 before the subgroup accretion, or now it is a result of merging? Several variants of answers are possible.  \\medskip We thank Dr. A. V. Moiseev of SAO RAS for supporting the observations at the 6m telescope. The 6m telescope is operated under the financial support of Science Ministry of Russia (registration number 01-43). During the data analysis we have used the Lyon-Meudon Extragalactic Database (LEDA) supplied by the LEDA team at the CRAL-Observatoire de Lyon (France) and the NASA/IPAC Extragalactic Database (NED) which is operated by the Jet Propulsion Laboratory, California Institute of Technology, under contract with the National Aeronautics and Space Administration. Our study  is supported by the grant of the Russian Foundation for Basic Researches no. 07-02-00229a  \\clearpage"
    },
    "0807/0807.2331_arXiv.txt": {
        "abstract": "We analyse the signatures of brane inflation models with moduli stabilisation. These are hybrid inflation models with a non-trivial field-space metric which can induce complex trajectories for the fields during inflation. This in turn could lead to observable features on the power spectrum of the CMB fluctuations through departures from near scale invariance or the presence of isocurvature modes. We look specifically at multi-brane models in which the volume modulus also evolves. We find that the signatures are highly sensitive to the actual trajectories in field space, but their amplitudes are too small to be observable even for future high precision CMB experiments. ",
        "introduction": "The early universe gives us the best possibility to test string theory models. Brane-anti-brane \\cite{Dvali:1998pa,Burgess:2001fx,GarciaBellido:2001ky,Jones:2002cv} and modular inflation models \\cite{BlancoPillado:2004ns,Conlon:2005jm} give predictions that are largely compatible with the present cosmological data. Therefore, in order to differentiate between the various models we have to look at their predictions beyond the usual set of experimentally determined parameters (the amplitude of the density perturbations, the spectral index and the running of the spectral index).  With the advent of high precision cosmological Cosmic Microwave Background (CMB) and Large Scale Structure (LSS) data the possibility of characterising the initial perturbation spectrum beyond these simplest spectral parameters is now a reality. Efforts to reconstruct the initial perturbation spectrum with model-independent 'inversion' techniques are already under way (see e.g. \\cite{reconstruct}) and have yielded some tantalizing hints of structure in the spectrum. There is a lack of power on the largest scales \\cite{quadrupole} and some indications of oscillations on intermediate scale \\cite{oscillations}. Arguably, the statistical significance of these effects are not sufficiently strong to motivate a departure from the simplest phenomenological models of single-field, slow roll inflation but an accurate analysis of observable features in theoretically driven scenarios is certainly warranted.   A generic feature of models of inflation within warped compactification mechanisms is the presence of many extra degrees of freedom in addition to the field driving inflation. Depending on the masses and couplings of the extra fields their presence can induce non-trivial trajectories in the field space. This can result in observable effects such as radically broken scale invariance, presence of isocurvature (or entropy) perturbations and enhanced non-Gaussianity of the perturbations close to points of high acceleration in the trajectories \\cite{ng}.  A non-vanishing isocurvature contribution to the total initial perturbations is a particularly interesting possibility since tight limits on any such contribution will be available when CMB polarization measurements becomes precise. The polarization data will also help break fundamental degeneracies in our ability to accurately constrain any broken scale invariance beyond a simple running of the spectral index.  In this work we specifically look at multi-brane models \\cite{Cline:2005ty}. These models are very promising since they naturally avoid a number of significant fine-tuning problems present in the simplest brane inflation models. Firstly they naturally produce a larger number of $e$-folds via the assisted inflation mechanism \\cite{Liddle:1998jc}, and secondly, they offer the possibility that the inflaton potential is flattened dynamically.  In these models, all branes are initially initially in a local minimum and as branes tunnel through the potential barrier and annihilate with anti-branes the barrier is lowered dynamically until it completely disappears. The potential becomes monotonic and very close to flat once a sufficient number of branes-anti-brane pairs have annihilated and the remaining branes will roll down the potential before annihilating. Quantum fluctuations will cause the rolling branes to start from slightly different initial positions and therefore follow different trajectories. We study the signatures that these models give which may be detected through cosmological observations.  All models under consideration feature moduli stabilisation for both the shape \\cite{Giddings:2001yu} and volume \\cite{Kachru:2003aw} moduli and anti-branes needed to lift the anti-deSitter vacuum generated by the stabilising mechanism to a deSitter one during inflation. The bulk contains a warped throat, the anti-branes being located at the bottom of it, and the mobile branes roll down this throat while inflation takes place.  Specifically, we follow the evolution of a number of perturbation modes for a set of scalar fields (in this case the positions of the mobile branes) coupled to gravity both inside and outside the Hubble horizon. The non-trivial shape of the potential and K\\\"ahler metric leads to a residual evolution of the scalar fields after horizon crossing. We evolve the perturbations of the scalars coupled to gravity using the Mukhanov-Sasaki \\cite{Mukhanov:1990me,Sasaki} variables and decompose the perturbations into adiabatic and entropy ones. We finally compute the separate spectra for the adiabatic and entropy perturbations and we find that while the adiabatic perturbations are insensitive to the trajectories followed by the many inflaton fields, the entropy ones are highly dependent on the trajectory. However, their amplitude is much too small to lead to observable features in the measured CMB spectrum.  The study performed here applies to a larger class of models, namely the ``inflection point'' models \\cite{Baumann:2007np}, where the mass of the inflaton(s) is large, except for a very small region  around an inflection point of the potential seen as a function of one inflaton field at a time (keeping all other fields constant). Most of the $e$-folds are coming from this small region of the field space where the inflaton trajectory can be very well approximated by a straight line therefore leading to a very small amplitude for the density perturbations. ",
        "conclusions": "We analysed the possible signatures that brane-inflation models could give, that may be detectable in observations of the CMB spectrum. To do so we derived the evolution equation for the perturbations of the inflatons and metric and numerically evolved a large number of modes starting inside the horizon, all the way to the end of inflation.  We then separated the fluctuations into adiabatic and entropy ones and derived their respective spectra. We found that the adiabatic spectrum is insensitive to the inflaton trajectories and has the has the same shape as was determined in earlier work \\cite{Burgess:2004kv,Cline:2005ty} using the Sasaki-Stewart formalism \\cite{Sasaki:1995aw}. The entropy one is highly sensitive to the inflaton trajectories but its amplitude is too small to be observable. We also offered an analytical argument that in a large class of models where most of the $e$-folds come from a small region around an inflection point of the potential, the inflaton trajectory is almost a straight line and therefore the entropy perturbations are expected to be small.  Thus multi-brane models solve a number of fine-tuning problems but due to their $e$-foldings coming from a reduced section of the field trajectory they display less phenomenology than single field models. This argument would lead us to expect that other potentially observable effects such as non-Gaussianities would also have a low amplitude in these models. Typically large values of the non-Gaussian $f_{\\rm NL}$ parameter are produced during periods of high acceleration in the field trajectory which are not present in these models \\cite{ng}.  Interestingly, the old models of brane inflation using branes oriented at angles and toroidal compactifications \\cite{Jones:2002cv,Burgess:2001fx,GarciaBellido:2001ky,Shandera:2003gx} have highly non-trivial inflaton trajectories as a result of the brane interaction inside a compact manifold. This leaves open the possibility that more complicated models, like multi-throat models \\cite{Iizuka:2004ct} will feature non-trivial inflaton trajectories and therefore observable features in the power spectrum and non-Gaussianity.  As we have shown the models have no particular problem in reproducing the observed tilt in the power spectrum but as expected in models that probe short trajectories, the running is small. If the observed running is confirmed to be at the current best-fit value of $dn_s/d\\ln k \\approx -0.05$ with increasing significance in future measurements it may turn out to be the strongest constraint on these models. As with most multi-field inflation models a detection of a gravitational wave background would also be in conflict with the models since the tensor perturbation are not boosted by the number of fields contributing to inflating the universe."
    },
    "0807/0807.4983_arXiv.txt": {
        "abstract": "{ Previous surveys for \\HI\\ 21-cm absorption in $z > 0.1$ radio galaxies and quasars yield a $\\approx40$\\% detection rate, which is attributed to unified schemes of active galactic nuclei (AGN). In this paradigm absorption is only witnessed in (close to) type-2 objects, where the central obscuration is viewed (nearly) edge-on and thus absorbs the rest frame 1420 MHz emission along our sight-line. However, we find this mix of detections and non-detections to only apply at low redshift ($z < 1$): From a sensitive survey of eight $z\\gapp3$ radio sources we find no 21-cm absorption, indicating a low abundance of cold neutral gas in (the sight-lines searched in) these objects. Analysing the spectral energy distributions of these sources, we find that our high redshift selection introduces a bias where our sample consists exclusively of quasars with ultra-violet luminosities in excess of $L_{\\rm UV}\\sim10^{23}$ W Hz$^{-1}$. This may suggest that we have selected a class of particularly UV bright type-1 objects. Whatever the cause, it must also be invoked to explain the non-detections in an equal number of $z < 0.7$ sources, where we find, for the first time, the same exclusive non-detections at $L_{\\rm UV}\\gapp10^{23}$ W Hz$^{-1}$. These objects also turn out to be quasars and, from these exclusive high UV luminosity--21-cm non-detections, it is apparent that orientation effects alone cannot account for the mix of 21-cm detections and non-detections at any redshift. ",
        "introduction": "Observations of the redshifted 21-cm transition of neutral hydrogen (\\HI) provide a powerful probe of the nature and contents of the early Universe. For example, this line can be used to: \\begin{enumerate} \\item Measure the baryonic content of the young Universe, when, before its consumption by star formation, neutral gas outweighed the stars.  \\item Probe the evolution of large-scale structure and galaxy formation. At high redshifts, interactions occur more frequently and \\HI\\ observations provide an indispensable means to studying the dynamics of galactic mergence and accretion.  \\item Provide a lower limit the time of the epoch of reionisation, when neutral hydrogen collapsed forming the first galaxies and igniting the first stars.  \\item Measure any changes in the values of the fundamental constants of nature: Comparing the redshifted frequencies of the 21-cm transition with those of metal-ion and molecular lines against laboratory values, can in principle yield measures of various combinations of fundamental constants at large look-back times (see \\citealt{cdk04}). \\end{enumerate} However, redshifted \\HI\\ 21-cm absorption systems are currently very rare with only 67 known at $z\\gapp0.1$ (see table 1 of \\citealt{cww+08}). We have therefore embarked upon a large survey to search for this transition in both systems intervening the lines-of-sight to background radio sources \\citep{cmp+03,ctp+07} and within the sources themselves (i.e.  for absorption ``associated'' with the host, \\citealt{cwm+06,cww+08}). From a recent survey for associated absorption in $z = 2.9 - 3.8$ radio sources with the Giant Metrewave Radio Telescope (GMRT), 21-cm was not detected in any of the eight targets for which good data were obtained. We discuss the reasons for this here. ",
        "conclusions": "Our observations are discussed in detail in \\citet{cwm+06,cww+08} and in Fig. \\ref{N-z} we show the derived 21-cm line strengths, indicating the upper limits of our and the previous surveys. \\begin{figure*}[t!] \\resizebox{\\hsize}{!}{\\includegraphics[angle=-90,clip=true]{N-z-shrink.eps}} \\caption{\\footnotesize The scaled velocity integrated optical depth of the \\HI\\ line ($1.823\\times10^{18}.\\int \\tau dv$) versus the host redshift for the published $z\\gapp0.1$ searches for associated 21-cm absorption. The filled symbols represent the 21-cm detections and the unfilled symbols the non-detections, with stars designating quasars and circles galaxies. The hatched region shows the range of our recent survey. The detections in this range are from \\cite{ubc91,mcm98} and the two $z_{\\rm em}>5$ non-detections are from \\cite{cwh+07}.} \\label{N-z} \\end{figure*} As seen from the figure, our limits are comparable to the vast majority of detections (which are primarily at low redshift)\\footnote{Note that most of the low redshift non-detections have been searched sufficiently deeply to detect 21-cm in most of the known absorbers.}. Currently, the known mix of detections and non-detections (until our survey, mostly at $z\\lapp1$) are attributed to unified schemes of active galactic nuclei (AGN), where the observed properties of the active galaxy or quasar are due to the orientation at which the active nucleus is viewed (see \\citealt{ant93,up95}): In type-1 objects, the AGN is viewed directly and in type-2 objects through a large column of dense obscuring gas, thus giving rise to the 21-cm absorption by the cold neutral gas located along our line-of-sight (e.g. \\citealt{jm94,cb95}, Fig. \\ref{agn}). \\begin{figure*}[t!] \\begin{center} % \\includegraphics[scale=0.32]{agn3-rot.eps} \\end{center} % \\caption{\\footnotesize Schematic showing the lines-of-sight to an AGN, where according to unified schemes the object type depends upon the orientation at which the nucleus is observed: In type-1 objects we view the AGN directly and in type-2s this is obscured by a circumnuclear torus of dense neutral gas. The 21-cm absorption detection shown is in PKS 1555--140, a known type-2 object, and the non-detection is in PKS 2300--189, a known type-1 object (the spectra are taken from \\citealt{cwm+06}, where the ordinate is the flux density in Jy and the abscissa the observed frequency in GHz). } \\label{agn} \\end{figure*}  However, for our sample we obtain exclusive non-detections, where, according to unified schemes, we may expect the $\\approx40$\\% (31 out of 73) mix seen at $z\\lapp1$\\footnote{Perhaps more, since the density of \\HI\\ at $z \\sim 3$ is expected to be higher than it is presently (e.g. \\citealt{psm+01}).}. With two detections already obtained at $z = 2.64$ and $3.40$ (Fig. \\ref{N-z}), it is clear that our high redhsift selection (alone, at least) cannot be responsible for the lack of 21-cm absorption in our targets. If we consider the ultra-violet ($\\lambda\\sim1000$ \\AA) luminosities, however (which we have estimated from the available optical and near-infrared photometry, \\citealt{cww+08}), we see that all our targets have luminosities of $L_{\\rm UV}\\gapp10^{23}$ W Hz$^{-1}$ (Fig. \\ref{lum-z}). \\begin{figure*}[t!] \\resizebox{\\hsize}{!}{\\includegraphics[angle=-90,clip=true]{lum-z-shrink.eps}} \\caption{\\footnotesize As per Fig. \\ref{N-z}, but with the calculated ultra-violet luminosities on the ordinate. } \\label{lum-z} \\end{figure*} Furthermore, although there is a roughly equal mix of detections (33 objects) and non-detections (36 objects) at $L_{\\rm UV}\\lapp10^{23}$ W Hz$^{-1}$, the exclusive non-detections at high UV luminosities relation is also seen for the low redshift sources.  We therefore suspect that the high UV fluxes are ionising the neutral gas, rendering this undetectable through the 21-cm transition of hydrogen. This is the first time such a correlation has been noted and, if indeed the case, perhaps exclusive 21-cm non-detections should have been expected for our sample, which at luminosity distances of 26 to 34 Gpc are severely flux limited. That is, at such large distances only the most UV luminous sources are known\\footnote{The quasar frame UV flux being redshifted into the optical band in these cases.}. It should be noted, however, that our main selection criterion was choosing radio-loud sources with $B\\gapp19$. This introduces a luminosity ``ceiling'' of $L_{\\rm UV}\\lapp3\\times10^{24}$ W Hz$^{-1}$, thus causing us to select the {\\it dimmer} ultra-violet sources known at these redshifts.  With hindsight, a high redshift selection yielding high luminosity objects undergoing a large degree of ionisation may be expected, although, what is surprising is the discovery of the eight 21-cm non-detections at low redshift (all at $z\\leq0.7$, Fig. \\ref{lum-z}) for which the UV luminosity also exceeds $L_{\\rm UV}\\sim10^{23}$ W Hz$^{-1}$. These objects could be understood in terms of unified schemes (Fig. \\ref{agn}), where we may expect high UV luminosities to go hand-in-hand with 21-cm non-detections (i.e. all are type-1 objects), if it were not for the non-detections at $L_{\\rm UV}\\lapp10^{23}$ W Hz$^{-1}$ (in fact down to $L_{\\rm UV}\\approx4\\times10^{18}$ W Hz$^{-1}$). Again, this is the first time that such a segregation has been noted and suggests that the $L_{\\rm UV}\\gapp10^{23}$ W Hz$^{-1}$ targets (which are all flagged as quasars) are different from their low luminosity counterparts in which 21-cm also remains undetected (which comprise of a mix of galaxies and quasars)."
    },
    "0807/0807.1274_arXiv.txt": {
        "abstract": "{Recent observations of two black hole candidates (GX 339-4 and J1753.5-0127) in the low-hard state ($L_{\\rm{X}}/L_{\\rm{Edd}} \\simeq 0.003-0.05$) suggest the presence of a cool accretion disk very close to the innermost stable orbit of the black hole. This runs counter to models of the low-hard state in which the cool disk is truncated at a much larger radius. We study the interaction between a moderately truncated disk and a hot inner flow. Ion-bombardment heats the surface of the disk in the overlap region between a two-temperature advection-dominated accretion flow and a standard accretion disk, producing a hot ($kT_{\\rm{e}}\\simeq 70$keV) layer on the surface of the cool disk. The hard X-ray flux from this layer heats the inner parts of the underlying cool disk, producing a soft X-ray excess. Together with interstellar absorption these effects mimic the thermal spectrum from a disk extending to the last stable orbit. The results show that soft excesses in the low-hard state are a natural feature of truncated disk models. ",
        "introduction": "The geometry of the low-luminosity (``low-hard'') state of Galactic Black Hole Candidates (GBHC), in which the spectrum is dominated by a power law X-ray flux extending to high energies, has been an open question for several decades. While it is generally believed that the power law spectrum is formed by inverse Compton scattering, there is no consensus about the geometry of the flow, source of seed photons or energy distribution for the Comptonizing electrons.  Broadly speaking, there are two classes of model to explain the spectrum in the low-hard state. The first is the ``corona'' model, in which the disk remains untruncated or nearly untruncated at luminosities $L_{\\rm{X}} \\simeq 10^{-3} L_{\\rm{Edd}}$. The hard power law spectrum comes from a hot and patchy corona (perhaps powered by magnetic flares \\citep{1998MNRAS.299L..15D, 1999ApJ...510L.123B, 2001MNRAS.321..549M}) on top of the disk, while the surrounding region is bombarded with high energy photons, producing the observed reflection and Fe-K fluorescence components. In the alternate, ``truncated disk'' model the thin disk is truncated at some distance from the black hole and the inner region is filled with a hot, radiatively-inefficient flow, which produces the hard spectrum. The reflection spectrum and Fe-K fluorescence is then produced by the interaction of the hard X-rays with the inner part of the truncated disk, or in some cool outflow moving away from the disk. For a recent discussion of the low-hard state see sect. 4 of \\cite{2007A&ARv..15....1D}.  In theory, the presence or absence of a cool disk should be confirmable through direct detection of a soft X-ray blackbody component at low energies. In practice however, this is made difficult by the fact that at low accretion rates the temperature of even an untruncated disk will drop from about 1-2keV in the high soft state to $\\sim$ 0.1-0.3keV, which puts it out of the range of most X-ray detectors. Additionally, the effects of interstellar absorption become very strong at around 0.1keV, so that detecting a soft excess and accurately measuring its parameters will depend somewhat on how accurately the interstellar absorption can be determined.  Even with these challenges, a soft excess in the low-hard state has previously been reported in several sources. The first was Cyg X-1 \\citep{1995A&A...302L...5B, 2001ApJ...547.1024D}, although its association with an accretion disk is complicated by the fact that Cyg X-1 is a high mass X-ray binary accreting from a wind. This question was also the focus of two recent papers, \\cite{2006ApJ...652L.113M, 2006ApJ...653..525M}, in which the authors studied long-exposure {\\it{XMM-Newton}} spectra of two different GBHCs, SWIFT J1753.5-0127 and GX 339-4, at low luminosities ($L_{\\rm{X}}/L_{\\rm{Edd}} \\sim 0.003-0.05$). Soft excesses at similar luminosities in these two sources have also been reported in \\cite{2007MNRAS.378..182R} (J1753.5-0127) and \\cite{2008arXiv0802.3357T} (GX 339-4). Since these two observations, there have also been observations of soft excesses in several other sources. \\cite{2007ApJ...666.1129R} made several observations of the soft component of XTE J1817-330 with $Swift$ during the outburst decline of that source down to a luminosity of $L_{\\rm{X}}/L_{\\rm{Edd}} \\sim 0.001$, while a soft component in GRO J1655-40 has been reported by both \\cite{2006MNRAS.365.1203B} and \\cite{2008PASJ...60S..69T} using different telescopes.  To interpret the soft excesses in SWIFT J1753.5-0127 and GX 339-4, \\cite{2006ApJ...652L.113M, 2006ApJ...653..525M} fit the data with a several {\\it XSPEC} models, trying various black-body disk models and simple hard X-ray components (both a power law and various Comptonization models). In GX 339-4 a broad Fe-K line was also observed and fit with a relativistically broadened reflection model. Using blackbody models for a standard accreting disk, the authors found disks with maximum temperatures of {\\it{kT}} $\\sim$ 0.2-0.4 keV, and inner radii consistent with the innermost stable circular orbit of a black hole.  At the inferred low accretion rates in the hard state, a disk extending to the last stable orbit would produce a soft X-ray component with peak close to the cutoff due to interstellar absorption. Unless an accurate independent measure of the interstellar absorption column is available, spectral fitting procedures cannot reliably distinguish between a thermal peak at $kT = 0.3$ keV with one interstellar absorption column and a cooler component with a lower energy component cutoff by a slightly higher interstellar absorption column.  For energetic reasons the hard X-ray component which dominates the luminosity in the hard state must originate near the black hole, the same region as the proposed cool disk. Some form of interaction of hard X-rays with the cool disk must take place, and this implies that the isolated cool disk models used as `components' in fits to observed spectra are unrealistic. In fact, most models for the hard X-ray component include some prescription for the reprocessing of hard into soft radiation, whether these be truncated disks or extended disk models. As shown by Haardt and Maraschi (1991), such models generically produce a similar energy flux in soft and hard X-rays. A strong soft component is thus a natural consequence in truncated as well as extended disk models for the hard state.  The main difference in a truncated disk model is that the soft flux originates from a larger surface area and consequently has a lower temperature, putting its spectral peak below 0.5 keV. After interstellar absorption the soft component has a peak around 0.5 keV that can be mistaken for an apparent thermal peak with the temperature of a disk near the last stable orbit.  In this paper we examine this with a more quantitative model for truncated disks. At the inner edge of a truncated disk the accretion flow must change in nature from a relatively cool, thin disk into a much hotter, vertically-extended inner flow. There will thus necessarily be some interaction between the two, either through radiation (e.g. \\citet{1991ApJ...380L..51H}) or matter exchange \\citep{1997LNP...487...67S}, or both. Our goal is to determine whether such a model could reproduce the soft spectral components reported by \\cite{2006ApJ...652L.113M,2006ApJ...653..525M}. We will find that disks truncated at 15--20 Schwarzschild radii can in fact produce soft components of the observed strength and shape. An alternative model investigating re-condensation fron am ADAF is considered by \\cite{2008arXiv0807.3402T}. ",
        "conclusions": "\\label{sec:conclusions}  From energetic considerations, the hard spectra observed in the low-hard state of LMXBs must be produced by hot ($kT_{e}\\sim$ 100 keV) matter in the inner regions surrounding the black hole. If there is also a much cooler disk present, there will necessarily be some degree of interaction between the two components, and the disk will be somewhat heated by irradiation from the hot Comptonizing component. The fits reported in \\cite{2006ApJ...652L.113M} and \\cite{2006ApJ...653..525M} neglect this interaction by fitting the disk and hard component separately. In this paper we have shown that incorporating the effects of this interaction heats the inner regions of a moderately truncated disk so that, when coupled with the effects of interstellar absorption, the size of the soft excess matches observations. Our work also highlights the potential pitfalls of using simple power law or analytic Comptonization fits at low energies, which can provide significant deviations in the soft X-rays, thus changing the shape and intensity of the observed soft excess.  In the case of GX 339-4, our model predicts an Fe-K component of comparable strength to that observed, although we did not do a detailed comparison. However, work by others has suggested that part of the broadening in the Fe-K line that was observed for GX 339-4 can be attributed to a large outflow, and detailed models of Fe-K fluorescence in galactic black holes show lines that are much broader than is found in AGN models (and which are normally used to fit spectra).  The model we have envisioned presents several opportunities for further improvement, in order to better constrain the introduced fitting parameters, $C$ and $\\eta$. The spectrum from the hot ring and ADAF are particularly uncertain, and dependent on a more detailed model for the radiative transfer through this region, as well as the source and number of seed photons (which will set the electron temperature in both regions). The model's global accretion rate (which is limited by the rate at which the hot layer spills over into the hot ring and then evaporates into the ADAF) is also very low, although this can be increased if the viscosity in the warm layer can be increased, perhaps as a result of accretion through an ordered magnetic field.  CD'A acknowledges financial support from the National Sciences and Engineering Research Council of Canada."
    },
    "0807/0807.2985_arXiv.txt": {
        "abstract": "We consider black holes resulting from binary black hole mergers. By fitting to numerical results we construct analytic formulas that predict the mass and spin of the final black hole. Our formulas are valid for arbitrary initial spins and mass ratios and agree well with available numerical simulations. We use our spin formula in the context of two common merger scenarios for supermassive galactic black holes. We consider the case of isotropically distributed initial spin orientations (when no surrounding matter is present) and also the case when matter closely aligns the spins with the orbital angular momentum. The spin magnitude of black holes resulting from successive generations of mergers (with symmetric mass ratio $\\eta$) has a mean of $1.73\\eta + 0.28$ in the isotropic case and $0.94$ for the closely aligned case. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0807/0807.2511_arXiv.txt": {
        "abstract": "We study the ejecta chemistry of a zero-metallicity progenitor, massive, supernova using a novel approach based on chemical kinetics. Species considered span the range of simple, di-atomic molecules such as CO or SiO to more complex species involved in dust nucleation processes. We describe their formation from the gas phase including all possible relevant chemical processes and apply it to the ejecta of a primordial 170  \\Ms~supernova. Two ejecta cases are explored:  full mixing of the heavy elements, and a stratified ejecta reflecting the progenitor nucleosynthesis. Penetration of hydrogen from the progenitor envelope is considered. We show that molecules form very efficiently in the ejecta of primordial supernovae whatever the level of mixing and account for 13 to 34\\% of the total progenitor mass, equivalent to 21 to 57 \\Ms~of the ejecta material in molecular form. The chemical nature of molecules depends on mixing of heavy elements and hydrogen in the ejecta. Species produced include O$_2$, CO, CO$_2$, SiS, SO, SiO and H$_2$. Consequently, molecules can be used as observational tracers of supernova mixing after explosion. We conclude that primordial massive supernovae are the first molecule providers to the early universe. ",
        "introduction": "Large amounts of dust have been conjectured to explain the reddening of background quasars and damped Ly$\\alpha$ systems in the early universe (Pettini et al. 1994, Pei \\& Fall 1995). Possible dust makers could be primordial, very massive stars exploding as supernovae, for the time scales at redshifts $\\ge$ 6 imply short stellar evolution times, thus excluding low-mass, evolved stars. Dust formation in such massive objects has been studied (Nosawa et al. 2003, Schneider et al. 2004) using a classical nucleation theory and excluding details on nucleation processes of dust clusters from the gas phase. However, it is now well accepted that dust forms in other evolved stellar environments under non-equilibirum conditions close to those encountered in laboratory dust condensation experiments where chemical kinetics commands dust nucleation from the gas phase (Donn \\& Nuth 1985, Cherchneff et al. 1992). In supernovae like in other environments, the nucleation will take place via the formation of a molecular phase in the ejecta. Detection of CO, SiO and H$_3^+$ in SN1987A, a core-collapse supernova in the Large Magellanic Cloud (Spyromilio et al. 1988, Roche et al. 1991, Miller et al. 1992), triggered a few theoretical studies on modeling CO and SiO observations (Petuchowski et al. 1989, Lepp et al. 1990, Liu \\& Dalgarno 1994, Gearhart et al. 1999). Those models often used incomplete chemical networks at steady state. No complete chemical description of a primordial, massive supernova ejecta has been so far attempted. We present here the first results of such an endeavor and consider as a surrogate a primordial massive supernova (hereafter PMSN) with zero-metallicity 170  \\Ms~progenitor. We study the formation of chemical species, some of which being dust precursors, in its ejecta. We show that chemistry is usually not at steady state and fosters the formation of complex molecules. Up to one third of the ejecta is found to be in molecular form depending on the level of mixing after explosion. ",
        "conclusions": "Molecular abundances with respect to the total gas number density are displayed in Figure 2 for the fully-mixed case. Here, we consider two sub-cases: the extreme case where the entire H envelope is microscopically mixed to the He core (hereafter referred as H-rich case), and a case for which only 1\\% of the H envelope mixes with the He-core (i.e. H-poor case). In both cases several molecules do form but differences exist among the two situations.  For the H-rich case, the dominant ejected species at 1000 days are H$_2$, O$_2$, SO, CO$_2$, and N$_2$. Their corresponding masses are respectively 19.3 \\Ms, 15.3 \\Ms, 13.0 \\Ms, 8.4 \\Ms, and 1.0 \\Ms, thereby implying a total molecular mass of $\\sim$ 57 \\Ms~equivalent to 34\\% of the PMSN progenitor mass. SiO is abundant up to t = 550 days, but is rapidly depleted due to silica/quartz precursor formation. Inspection of Figure 2 shows that, chemically speaking, the steady state assumption does not hold for molecular formation and destruction at t $\\le$ 1000 days. At early times, the dominant chemical processes at play are neutral-neutral and RA reactions. The OH radical is a key species to molecular formation. Destruction occurs mainly via He$^+$ attack. For t$\\ge$ 600 days, neutral-neutral processes without activation barrier and ion-molecule reactions are dominant. As for dust precursors, (SiO$_2$)$_2$, a ring precursor to silica/quartz nucleation, forms in large amount --  33.3 \\Ms, as early as 500 days after explosion, followed by AlO, the gas phase precusor to corumdum (Al$_2$O$_3$) -- 0.04~\\Ms~at 530 days, and finally (FeO)$_2$ -- 0.001~\\Ms~at 750 days. For the H-poor case, the dominant molecular species are SiO - 15.8~\\Ms, CO - 4.6~\\Ms~and CO$_2$ - 1~\\Ms, resulting in a total molecular content of 21.5~\\Ms, or 12.6\\% of the progenitor mass. When H is absent, the formation chemistry is powered at early times by RA reactions as opposed to neutral-neutral reactions involving OH. The RA forming SiO has a reaction rate 50 times greater than that for O$_2$ formation, resulting in large amounts of SiO compared to O$_2$. Molecular  formation by  neutral-neutral processes without activation barrier is postponed to later times, as seen in Figure 2. A smaller molecular content is thus produced due to lower gas temperature and density. As for dust, (SiO$_2$)$_2$ again is the dominant precursor to form  at level up to 10.7~\\Ms, while (MgO)$_2$, AlO and (FeO)$_2$ contents are negligible.  The above dust nucleation sequences are different from those of existing classical nucleation studies which rely on condensation temperature analysis for solids (Todini \\& Ferrrara 2001, Nozawa et al. 2003, Schneider at al. 2004). Such studies predict the condensation of corundum, forsterite (Mg$_2$SiO$_4$), quartz (SiO$_2$), amorphous carbon and magnetite (Fe$_3$O$_4$). The discrepancy resides in the approach used: in the present model, molecules form from the gas simultaneously to dust precursors, thereby depleting some of the available elements from the gas phase. Dust precursor formation is then commanded by chemical kinetics at play in an altered gas phase whose chemical composition differs drastically from the initial elemental composition. The absence of carbon dust precursors in our fully-mixed case illustrates this point. Indeed, carbon is locked up in CO and CO$_2$ despite the inclusion of Compton electron destruction reactions (Clayton et al. 1999), and is not available for further build-up of C-rich molecules and amorphous carbon precursors. Hence chemistry is acting as a bottleneck to dust nucleation. If we assume that all gas-phase precursors formed in our fully-mixed ejecta are included into dust during condensation, we get an upper limit for freshly formed dust of roughly 33.4~\\Ms~in the H-rich case and 10.7~\\Ms~in the H-poor case, equivalent to $\\sim 19\\%$ and 6\\% of the PMSN progenitor mass, respectively. However, such an assumption overestimates the dust content as dust condensation efficiencies in pyrolysis or flame experiments in the laboratory are usually less than one, with a residual population of dust precursors in the gas phase (J{\\\"a}ger et al. 2006). Furthermore, our H-rich case is somehow extreme and the total H envelope mixing with the He core unlikely, though H mixing was observed in SN1987A and should occur in PISNe as well. Nozawa et al. (2003) and Schneider et al. (2004) find respectively  $\\sim$ 33 \\Ms~and $\\sim$ 39 \\Ms~of dust for the fully-mixed He core of a 170 \\Ms~progenitor mass PISN and similar ejecta thermodynamics than that in the present study. These results are to be compared to the 10.7 \\Ms~upper limit derived for our H-poor case. Our two mixed cases highlight trends in dust formation scenarios and content in PMSNe and our limits for freshly formed dust are a few \\% of the progenitor mass only. This is certainly lower than existing predicted dust amounts.  For the unmixed ejecta, molecules and their derived mass are listed in Table 2. Again, molecules do form efficiently and account for $\\sim$ 42~\\Ms, equivalent to 25\\% of the PMSN progenitor mass. Their chemical nature now traces their location in the He core and the amount of H mixing from the progenitor envelope. In Zone 1, devoid of oxygen, SiS forms from reaction between atomic Si and S$_2$ and traces a Si/S-rich gas. In Zone 4, we allowed for 18 \\% of total mass hydrogen penetration in the C/He--rich layer as in Nosawa et al. (2003), with subsequent formation of H$_2$, C$_2$H$_2$, and CO. Carbon monoxide is preferably formed over CO$_2$ in this layer as CO$_2$ is rapidly destroyed by reactions with ions like C$^+$ and He$^+$ to re-form CO. Thus, CO presence in the ejecta indicates a C/O ratio greater than 1, a large free carbon atom content, and some mixing with He and H. This result is supported by the simultaneous detection of CO and amorphous carbon dust in SN1987A, implying carbon-rich inhomogeneities in the ejecta. In Zone 2, CO$_2$ dominates over CO, because of the rapid CO conversion to CO$_2$ via reaction with O$_2$, which is very abundant in those regions where the C/O ratio is less than 1. CO$_2$ is thus a tracer of oxygen-rich ejecta regions whereas the molecular composition of Zone 4 is typical of a C/H/He-rich environment. Similar chemical processes and species are encountered in the inner shocked winds of O-rich and C-rich, evolved low-mass stars, where the gas parameters resemble those of SN ejecta (Cherchneff 2006).  As for dust molecular precursors, we find that Zone 1 produces $\\sim$ 1.6~\\Ms~of (Si)$_4$ and $\\sim$ 1.3~\\Ms~of (FeS)$_2$, while Zone 2 produces 3~\\Ms~of (SiO$_2$)$_3$. No dust precursors are formed in Zone 3, and Zone 4 produces 0.5~\\Ms~of C$_3$. Again, grain types and nucleation sequences are different from the study of Nosawa et al. (2003). An upper limit to dust production in the unmixed case is $\\sim$ 6.5 \\Ms~or $\\sim$ 3.8$\\%$ of the total PMSN progenitor mass, which is much less than the value of $\\sim$ 18 \\% derived by Nozawa et al. for their 170 \\Ms~PISN unmixed case. This points again to the crucial role of chemical kinetics as bottleneck to dust formation. Chemically speaking, we notice that Zone 4 forms carbon dust precursors via pure carbon chains like C$_3$ despite the large quantities of C$_2$H$_2$ available and the aromatic formation reactions included in our chemical network. Aromatic chemical pathways to carbon dust are inhibited even in the presence of hydrogen because of the high temperatures encountered in the ejecta up to 600 days after explosion, and the competitive RA pathways to carbon chains combined to large amounts of free C atoms. This result may be specific to supernova ejecta and is supported by the non-detection of polycyclic aromatic hydrocarbon infrared emission lines in SN1987A spectra (Wooden et al. 1993).  We conclude that Pop III, massive supernovae are efficient molecule providers to the early universe, with 13\\%  to 34$\\%$ of their ejecta in molecular form, corresponding to 22 \\Ms~to 57~\\Ms~of molecular material released to the local, pristine gas after explosion. Formed species depend on mixing in the ejecta and include O$_2$, CO$_2$, SO, CO, SiS, OH, C$_2$H$_2$ and H$_2$. H mixing boosts molecular formation at early times via neutral-neutral processes, resulting in a large molecular component in the ejecta. In a more general context, molecules could be used as observational tracers of mixing in nearby core-collapse supernovae. The substantial amounts of molecular material imply some impact on the local gas cooling if cooling time scales are comparable to those for the ejecta adiabatic expansion. It was recently shown that some dust grains could survive the passage of the PMSN reverse shock some 10$^4$ years after explosion (Nozawa et al. 2007, Bianchi \\& Schneider 2007). The survival of molecules in cool, dense, inhomogeneities passing the reverse shock must then be studied. At later times, it is conjectured that Pop.~II.5 stars can form in the PMSN dense shell if other cooling than that of H$_2$ is provided (MacKey et al. 2003, Salvaterra et al. 2004). If surviving the reverse shock, our predicted  molecules could provide part of the necessary cooling and their subsequent impact on second generation star formation needs further investigation. Finally, under present chemical kinetic conditions, our results show that PMSNe do form dust efficiently as the amounts of dust precursors produced are significant. However, the present study points to lower dust contents formed in PMSNe than those predicted by existing, classical nucleation studies."
    },
    "0807/0807.4930_arXiv.txt": {
        "abstract": "We study the phenomenology of hybrid scenarios of neutrino dark energy, where in addition to a so-called Mass Varying Neutrino (MaVaN) sector a cosmological constant (from a false vacuum) is driving the accelerated expansion of the universe today. For general power law potentials we calculate the effective equation of state parameter $w_{eff}(z)$ in terms of the neutrino mass scale. Due to the interaction of the dark energy field (``acceleron'') with the neutrino sector, $w_{eff}(z)$ is predicted to become smaller than $-1$ for  $z>0$, which could be tested in future cosmological observations. For the considered scenarios, the neutrino mass scale additionally determines which fraction of the dark energy is dynamical, and which originates from the ``cosmological constant like'' vacuum energy of the false vacuum. On the other hand, the field value of the ``acceleron'' field today as well as the masses of the right-handed neutrinos, which appear in the seesaw-type mechanism for small neutrino masses, are not fixed. This, in principle, allows to realise hybrid scenarios of neutrino dark energy with a ``high-scale'' seesaw where the  right-handed neutrino masses are close to the GUT scale. We also comment on how MaVaN Hybrid Scenarios with  ``high-scale'' seesaw might help to resolve stability problems of dark energy models with non-relativistic neutrinos. ",
        "introduction": "The evidence for the existence of a dark sector in the universe has made the present era of cosmology fascinating and challenging. At present, the microscopic natures of dark energy and dark matter are still an open question, with both components only probed gravitationally \\cite{Colless:2001gk,Riess:1998cb,Perlmutter:1998np,Spergel:2006hy}. Regarding dark matter, there are various particle physics candidates which may belong to the class of Weakly Interacting Massive Particles (WIMPS) or which may interact only gravitationally. On the other hand, for particle physics explanations of the observed dark energy, the main fundamental question is whether dark energy is a cosmological constant or a dynamical field. While in the former case the equation of state parameter $w$ of dark energy is constant and equal to $-1$, in the latter case it is a function of redshift $z$ and in general differs from $-1$, providing a way to distinguish both fundamental types of dark energy experimentally. Such a deviation of $w$ from $-1$ will be searched for in various future surveys, studying supernovae, baryonic acoustic oscillations, weak gravitational lensing effects, galaxy clusters, or other techniques \\cite{Albrecht:2006um}. One obstacle to relate dark energy more explicitly to particle physics is the fact that the amount of the observed energy density $\\rho_{DE}^{1/4} \\sim 0.003$ eV is {\\em much} smaller than the ``theoretical expectation''. In this paper we will not address this ``cosmological constant'' problem, i.e.\\ the suppression/cancellation of the various generic contributions to dark energy which are too large, but assume that it is resolved by some other  mechanism.  Our work is motivated by the intriguing observation that the energy density of dark energy is close to another very small scale in particle physics, namely the one of neutrino masses. The discovery of flavour conversion of neutrinos from various sources, interpreted within the framework of neutrino oscillations, points to two mass eigenvalues of the light neutrinos above about $0.01$ eV and $0.05$ eV \\cite{Mohapatra:2005wg}, while searches for neutrino masses from Tritium $\\beta$-decay and neutrinoless double $\\beta$-decay yield an upper bound for each mass eigenvalue of roughly $0.5$ eV \\cite{Mohapatra:2005wg}. The proximity of the scale of dark energy to that of neutrino masses has inspired the proposal of the so-called Mass Varying Neutrino (MaVaN) scenario \\cite{Fardon:2003eh}. In this scenario, the relic neutrino density contributes to the effective potential for the dynamical field which is called ``acceleron''. One consequence is that neutrino masses vary on cosmological time scales and, since the relation between neutrino masses and $\\rho_{DE}^{1/4}$ remains valid in earlier cosmic epochs, also dark energy behaves differently from a cosmological constant. Compared to other dynamical models of dark energy (e.g.\\ quintessence \\cite{Wetterich:1987fm,Peebles:1987ek,Caldwell:1997ii}), in the MaVaN scenario one does not have to choose a mass for the dynamical field of the order of the present Hubble parameter $H_0$. In fact it can be much larger, even of the order of the neutrino mass scale, due to the stabilising effect of the contribution of the relic neutrino density to the potential. As a result, the field adiabatically tracks the minimum of the effective potential whose evolution is controlled by the time evolution of neutrino number density.  The MaVaN models have been investigated in many studies, e.g.\\ regarding possible experimental signatures and constraints \\cite{Kaplan:2004dq,Zurek:2004vd,Li:2004tq,Barger:2005mn,Schwetz:2005fy,Ringwald:2006ks, Cirelli:2005sg, Weiner:2005ac, Brookfield:2005bz,Bernardini:2008pn} as well as model-building issues \\cite{Fardon:2005wc, Takahashi:2005kw, Ma:2006mr, Takahashi:2007ru, Bjaelde:2008yd}. It has turned out that while there are many interesting possible signatures and attractive features, the  scenario is tightly constrained by the requirement of consistency with late time structure formation. In particular, it has been pointed out \\cite{Afshordi:2005ym} that if neutrinos are non-relativistic today, due to fifth force effects the neutrinos would cluster at late time and finally form ``neutrino nuggets'' which would spoil the dark energy behavior of the neutrino fluid.  The goal of this paper is to investigate the phenomenology of a generalised scenario which we will refer to as MaVaN Hybrid Scenario, where in addition to a MaVaN sector a cosmological constant (from a false vacuum) is driving the accelerated expansion of the universe today. As we will show, for a generalised power law potential in the MaVaN sector the effective equation of state parameter $w_{eff}(z)$ as well as the fraction to which the dark energy is of dynamical nature is determined by the neutrino mass scale. We will also analyse the possibility to realise neutrino dark energy with a ``high-scale'' seesaw mechanism, where the right-handed neutrino masses are close to the GUT scale, and comment on how such models might help to suppress the formation of ``neutrino nuggets'' and resolve stability problems of MaVaN models with non-relativistic neutrinos. For the main part of the paper we will focus on neutrinos that are non-relativistic today.  The paper is organised as follows: In Section 2 we will review the basics of the MaVaN scenario. Section 3 contains the definition of the phenomenological framework and the analysis is performed in Section 4. In Section 5 we discuss the possibility of realising neutrino dark energy with a ``high-scale'' seesaw mechanism and in Section 6 we comment on how this might help to solve stability issues for non-relativistic neutrinos. Section 7 contains our Summary and Conclusions. ",
        "conclusions": "Motivated by the intriguing proximity of the energy density of dark energy and the neutrino mass scale we have studied the phenomenology of hybrid scenarios of neutrino dark energy, where in addition to a so-called Mass Varying Neutrino (MaVaN) sector, a cosmological constant (from a false vacuum) is driving the accelerated expansion of the universe today. Within the generalised framework we have focused on phenomenological issues such as on the connection to the neutrino mass scale and on its consequences for the dynamical nature of dark energy.  We have therefore calculated the effective equation of state parameter $w_{eff}(z)$ in the MaVaN Hybrid Scenario where the effective potential for the dynamical real scalar field (the ``acceleron'' field $A$) has the following form \\begin{equation} V(A)_{eff} =  M^{4- \\alpha}A^{\\alpha} + \\frac{\\rho_{\\nu}^{(0)}}{a^3} \\left(\\frac{A_0}{ A}\\right)^{\\beta} +  V_{0} \\;. \\end{equation} We found that, for the case of a power law potential in the MaVaN sector, $w_{eff}(z)$ is determined by the neutrino mass scale and by the parameters $\\alpha$ and $\\beta$ (c.f.\\ Eq.~(\\ref{Eq:weff_mnu})). Due to the interactions of the dark energy field with the neutrino sector, $w_{eff}(z)$ is predicted to become smaller than $-1$ for increasing $z>0$ (c.f.\\ Fig.~\\ref{weffective_plot}), which could be tested in future cosmological observations. For the considered scenarios, we have also calculated how the neutrino mass scale determines which fraction of the dark energy is dynamical, and which originates from the ``cosmological constant like'' vacuum energy. In particular, for the case of a mass term potential and a standard seesaw relation (i.e.\\ $\\alpha=2$ and $\\beta=1$) we found that compatibility with the terrestrial neutrino mass bounds requires a large contribution of constant vacuum energy (c.f.\\ Eq.~(\\ref{Eq:MnuvsGamma})).   Another interesting question, which we have investigated in the MaVaN Hybrid Scenario with power law potentials is whether it is possible to realise neutrino dark energy with a ``high-scale'' seesaw mechanism, where the right-handed neutrino masses are close to the GUT scale. We found that the field value of the ``acceleron'' field as well as the masses of the right-handed neutrinos can indeed be large and in principle a hybrid scenario of neutrino dark energy might be realised with a ``high-scale'' seesaw. We have also commented on how the Hybrid MaVaN Scenarios with  ``high-scale'' seesaw might allow to suppress the formation of ``neutrino nuggets'' and resolve stability problems of dark energy models with non-relativistic neutrinos.  In summary, we have found that the considered hybrid scenarios of neutrino dark energy have several attractive features, in particular the close connection to the neutrino mass scale and an effective dark energy equation of state parameter $w_{eff}(z)$ which depends only on the parameters $\\alpha$ and $\\beta$ of the potential and on the neutrino mass scale. The prediction that $w_{eff}(z) < -1$ for $z>0$ provides a ``smoking gun'' signal for such interacting dark energy scenarios, which could be observed in future surveys. Issues which are still open, and which are left for further studies, are whether a MaVaN Hybrid Scenario with ``high-scale'' seesaw can indeed solve the stability problems of conventional MaVaN scenarios with non-relativistic neutrinos and whether a consistent model can be constructed where this is realised."
    },
    "0807/0807.1724_arXiv.txt": {
        "abstract": "Achieving maximum scientific results from the overwhelming volume of astronomical data to be acquired over the next few decades will demand novel, fully automatic methods of data analysis.  Artificial intelligence approaches hold great promise in contributing to this goal.  Here we apply neural network learning technology to the specific domain of eclipsing binary (EB) stars, of which only some hundreds have been rigorously analyzed, but whose numbers will reach millions in a decade. Well-analyzed EBs are a prime source of astrophysical information whose growth rate is at present limited by the need for human interaction with each EB data-set, principally in determining a starting solution for subsequent rigorous analysis.  We describe the artificial neural network (ANN) approach which is able to surmount this human bottleneck and permit EB-based astrophysical information to keep pace with future data rates. The ANN, following training on a sample of 33,235 model light curves, outputs a set of approximate model parameters ($T_2/T_1$, $(R_1+R_2)/a$, $e \\sin \\omega$, $e \\cos \\omega$, and $\\sin i$) for each input light curve data-set. The whole sample is processed in just a few seconds on a single 2GHz CPU. The obtained parameters can then be readily passed to sophisticated modeling engines. We also describe a novel method \\emph{polyfit} for pre-processing observational light curves before inputting their data to the ANN and present the results and analysis of testing the approach on synthetic data and on real data including fifty binaries from the Catalog and Atlas of Eclipsing Binaries (CALEB) database and 2580 light curves from OGLE survey data.  The success rate, defined by less than a 10\\% error in the network output parameter values, is approximately 90\\% for the OGLE sample and close to 100\\% for the CALEB sample -- sufficient for a reliable statistical analysis. The code is made available to the public. Our approach is applicable to EB light curves of all classes; this first paper in the Eclipsing Binaries via Artificial Intelligence (\\ebaints) series focuses on detached EBs, which is the class most challenging for this approach. ",
        "introduction": "Over the past decade advances in observational technologies, computers and eclipsing binary (EB) analysis codes (e.g., WD code: \\citet{wd1971,wd1993,wd2007}) have enabled the accumulation of an impressive sample of fundamental stellar data.   Careful analysis of EB light curves has produced fundamental stellar properties, tests of stellar evolution theories, accurate distances within the Galaxy and to external galaxies, as well as providing tests of stellar structure models and general relativity \\citep[see e.g.][]{guinan2007}.  Despite the importance of these astrophysical results, only some hundreds of EBs have been subject to the requisite analysis and the cumulative results populate the astrophysical parameter space sparsely. Current standard practice requires significant human interaction with EB light curve data, particularly in the initial stages of analysis, which defines a \"rate-determining step\" in the process generating these astrophysical results.  By 2020, the observational bounty from ground- and space-based programs such as Optical Gravitational Lensing Experiment \\citep[OGLE;][]{udalski1997}, Exp\\'erience de Recherche d'Objets Sombres \\citep[EROS;][]{pd1998}, All Sky Automated Survey \\citep[ASAS;][]{pojmanski2002}, Pan-Starrs \\citep{kaiser2002}, Large Synoptic Survey Telescope \\citep[LSST;][]{tyson2002}, and their space counterparts Hipparcos \\citep{perryman1997}, Kepler \\citep{borucki2004} and Gaia \\citep{perryman2001} will include millions of new EB light curves, even catching some EBs in fleeting stages of stellar evolution. Powerful and mature codes for light curve analysis stand ready to mine this enormous and rich vein of new astrophysical information, and pioneering steps in automating these tools have already been taken by \\citet{devor2005,tamuz2006,mazeh2006} (see \\citet{prsa2007} for a concise overview of current automated techniques). A key issue is to efficiently determine an initial set of model parameters for every light curve as input to the automated analysis process. Starting values have been conventionally supplied by the analyst/astronomer using expert knowledge typically guided by checks with light curve synthesis codes, a time-intensive process that is certainly out of the question for the coming fire-hose of EB data.  Approaches to automating light curve solutions have taken various forms to date. \\citet{ww2001,ww2002}, in their work to establish the best distance indicators among detached and semi-detached binaries in the Small Magellanic Cloud, obtained starting parameters for the rigorous WD model by comparing each candidate light curve with a set of template model light curves, sending the best match to an automated version of the WD differential corrector program DC. This could be computationally prohibitive to apply to the expected large future data-sets, even if clever pruning reduced the number of comparison templates.  Employing less rigorous physical models, of course, is one approach to computational efficiency. \\citet{devor2005} provides a critical discussion of an automated pipeline employing a simple model of spherical stars without tidal or reflection physics, whose starting values are obtained from an initial guess and then refined using a downhill simplex method with simulated annealing. \\citet{tamuz2006} employ the EBOP ellipsoidal model \\citep{popper1981}.  Using this engine, they arrive at initial solutions after a combination of grid search, gradient descent and geometrical analysis of the LC.  The challenge is therefore to gain the advantages of a sophisticated model yet keep processing time limited. The Eclipsing Binaries via Artificial Intelligence (\\ebaints) project introduces artificial neural networks (ANNs) that are trained on the rigorous WD physical model and are computationally extremely efficient, towards fully automating the solution process for EBs.  In this first paper in the \\ebai series we describe the basic ANN concepts and procedures for applying ANNs to detached EBs.  We present results of applying the trained ANN to a set of 10,000 synthetic detached EBs, to 50 detached real world binaries from the Catalog and AtLas of Eclipsing Binaries (CALEB\\footnote{CALEB is maintained by D.~Bradstreet at Eastern University -- see {\\tt http://caleb.eastern.edu}.}), and to the set of 2580 OGLE LMC binaries \\citep{wyrzykowski2003} classified as detached. Subsequent papers will deal similarly with semi-detached systems and overcontact EBs, and address automated light curve classification. ",
        "conclusions": "ANNs have proved successful in problems of classification, real-time control, data mining and many other tasks in a variety of scientific and technical applications. This paper demonstrates their utility in parameter estimation for detached binaries, even in the face of parameter degeneracy, the correlation of certain pairs of parameters; e.g., the compensating effect of increasing a system's sum of radii vs. increasing its inclination.  A task-optimal ANN topology and learning rate was determined and the network was trained on the rigorous WD model.  The trained network was tested on a previously unseen test set of 10,000 WD model light curves, yielding model parameters to better than 90\\% accuracy for 90\\% of the test set.  In tests with 50 real-world detached binary systems from the CALEB database, the ANN alone achieved sufficiently good fits for 22 systems that differential corrections were not needed.  Those requiring differential corrections all converged to produce good quality fits, confirming the utility of the ANN-produced starting parameters.  Overall, the CALEB test produced a 100\\% success rate for the ANN.  We have shown that a suitable ANN can be successfully trained on a sophisticated EB model, and that the trained network produces quite satisfactory approximate light curve solutions with high computational efficiency.  In addition to these advantages, ANNs have favorable properties of interpolation; they are well-behaved in regions around their multi-dimensional training points.  Thus the network will interpolate reasonable starting parameters for light curves with asymmetry or spots, and in fact for any (detached) light curve that it has not seen before.  It is quite remarkable that the network is able to memorize to such high quality all the light curve solutions in the space of detached binaries with only 8240 numbers ($201\\times40 + 40 \\times 5$ weights).  Clearly, this network has sufficient information storage capacity, but it is worth noting that the information capacity corresponding to a given network topology is only known in heuristic approximation \\citep{neelakanta1994}. This confirms that much is yet to be learned about ANNs, and in particular it explains the necessity to explore network topologies and learning rates for their appropriateness to the given problem.  Such exploration, in addition to the ability to rapidly train many iterations, further emphasizes the need for speed in the learning phase.  This work was greatly facilitated by two developments.  The \\emph{polyfit} algorithm produced excellent analytic approximations to real-world light curves sampled at (often quite) unequal phase intervals, where splines or other interpolating methods behave badly \\citep{emery2001}, making equal-phase interpolated data points available for input to the ANN.  This algorithm has been found to be an excellent interpolator for EBs of all classes, as well as, for example, Cepheid variables.  Since well-behaved interpolation is a problem common to many fields, \\emph{polyfit} could well see further application beyond astronomy.  The scheme for parallelization of ANN training over multiple processors was key to quickly exploring topologies and training parameters, and churning through a half-million training iterations in days instead of weeks.  This scheme had the beneficial side effect of lending simulated annealing behavior to the training phase, improving the search for the global minimum.  The source code for \\emph{polyfit} and ANN is released under the GNU General Public License, which grants users the right to freely use, distribute and modify the code. The programs may be downloaded from the IDP project homepage: {\\tt http://www.eclipsingbinaries.org} or \\phoebe homepage: {\\tt http://phoebe.fiz.uni-lj.si}.  This paper focuses on detached EBs.  Systems with large proximity effects which are fully treated in the WD model will be discussed in a subsequent paper.  We will thus be able to apply ANNs in a unified, automated approach to provide starting solutions for eclipsing binary light curves of all classes."
    },
    "0807/0807.2216.txt": {
        "abstract": "We perform a detailed photometric analysis (bulge-disk-bar decomposition and Concentration-Asymmetry-Clumpiness - CAS parametrization) for a well defined sample of isolated galaxies, extracted from the Catalog of Isolated Galaxies \\citep{karachentseva73} and reevaluated morphologically in the context of the AMIGA project (\\textbf{A}nalysis of the interstellar \\textbf{M}edium of \\textbf{I}solated \\textbf{GA}laxies). We focus on Sb-Sc morphological types, as they are the most representative population among the isolated spiral galaxies. Our analysis yields a large number of important galactic parameters and various correlation plots are used to seek relationships that might shed light on the processes involved in determining those parameters. Assuming that the bulge S\\'{e}rsic index and/or Bulge/Total luminosity ratios are reasonable diagnostics for pseudo- versus classical bulges, we conclude that the majority of late-type isolated disk galaxies likely host pseudobulges rather than classical bulges. Our parametrization of galactic bulges and disks suggests that the properties of the pseudobulges are strongly connected to those of the disks. This may indicate that pseudobulges are formed through internal processes within the disks (i.e. secular evolution) and that bars may play an important role in their formation. Although the sample under investigation covers a narrow morphological range, a clear separation between Sb and Sbc-Sc types is observed in various measures, e.g. the former are redder, brighter, have larger disks and larger bars, more luminous bulges, are more concentrated, more symmetric and clumpier than the latter. A comparison with samples of spiral galaxies (within the same morphological range) selected without isolation criteria reveals that the isolated galaxies tend to host larger bars, are more symmetric, less concentrated and less clumpy. ",
        "introduction": "The properties of galaxies and their evolution are thought to be strongly related to their environment. The empirical quantification of environmental influence (``nurture'') on morphology, structure, nuclear activity, star formation properties, etc. requires a robust definition of a sample of galaxies that are minimally perturbed by other galaxies. Such a sample could serve as a ``pure nature'' baseline. In this sense, perhaps the best compilation of isolated galaxies available at this time is the Catalog of Isolated Galaxies \\citep[CIG;][]{karachentseva73}. Both the size (n=1050 galaxies) and the restrictive isolation criteria in the catalog contribute to its statistical value. The definition of isolation requires that, for a galaxy of diameter D, there is no companion/neighbor with a diameter d in the range D/4 to 4D within a distance of 20d. The isolation criteria used to construct the CIG suggest that a typical galaxy of 25 kpc diameter has not been visited by a similar mass perturber in the past $\\sim$ 3Gyr \\citep[assuming a typical field velocity of $\\sim$ 150 km s$^{-1}$;][]{verdes05}. Thus, the evolution of such isolated galaxies is mostly driven by internal processes and to a much lesser degree by environment, at least for the last $\\sim$ 3Gyr of their existence.  A recent morphological reevaluation of the CIG galaxies in the context of the AMIGA project (\\textbf{A}nalysis of the interstellar \\textbf{M}edium of \\textbf{I}solated \\textbf{GA}laxies) revealed that the bulk ($\\sim$ 63\\%) show morphological types in the range Sb-Sc \\citep{sulentic06}. In this study we present the results of a photometric characterization for a representative subsample of n $\\sim$ 100 CIG galaxies classified as Sb-Sc in this latter reference. We perform multicomponent decomposition (bulge/disk/bar) using the BUDDA code\\footnote{http://www.mpa-garching.mpg.de/~dimitri/budda.html} \\citep{deSouza04}. Additionally, we evaluate CAS parameters Concentration(C)-Asymmetry(A)-Clumpinesss(S) \\citep[e.g.][]{conselice00,bershady00,conselice03,taylor07}. Assembling a set of parameters combining a model-dependent description of the main components of galaxies (BUDDA) with global structural measures (CAS) could provide valuable hints into the formation and evolution of galaxies.  This is the first attempt to date to present a detailed examination of this kind (bulge-disk-bar decomposition combined with CAS parameters) for a well defined sample of isolated galaxies. This study is an integral part of the AMIGA project, which is a dedicated multiwavelength study of the revised CIG catalog. The goal of AMIGA is to quantify the fundamental properties of a statistically meaningful sample of isolated galaxies which can then be used as a baseline for comparison and for estimation of the effects of environment in other less isolated samples of galaxies. The CIG catalog has recently been reevaluated in terms of galaxy positions \\citep{leon03}, isolation \\citep{verley07a,verley07b} and morphology \\citep{sulentic06}. A series of studies were produced in the context of the AMIGA project: 1) an optical characterization of the refined sample \\citep{verdes05}, 2) an analysis of mid- and far-infrared properties \\citep{lisenfeld07}, 3) a study of the neutral CO and HI gas \\citep{espada05,espada06}, 4) radio continuum emission \\citep{leon08} and 5) nuclear activity \\citep{sabater08}. Another recent study used a subsample of isolated AMIGA galaxies to investigate the role of bars in star formation processes \\citep[e.g.][]{verley07c}. Our present study offers a detailed photometric analysis of a representative sample of the core AMIGA population of Sb-Sc morphological types. We should note that all data produced within the AMIGA project are periodically updated and made publicly available at http://www.iaa.es/AMIGA.html.  Theoretical models and numerical simulations exploring the formation and evolution of galaxies rely on empirical results that could separate and quantify the relative roles of internal secular processes (that develop on time scales much longer than the galaxy formation/collapse process itself) and slow or fast external perturbations (environment) in defining the structural properties of galaxies. In this sense, our present study has a twofold importance: a) it explores a representative and well defined sample of \\textit{the most isolated galaxies} in the local Universe and b) provides an \\textit{extensive photometric structural analysis} of these galaxies. Our main goal is to identify potential scaling relations and correlations: i) between parameters describing the same structural component (bulge, disk or bar), ii) between components, iii) between components and global properties of the galaxy (morphological type, color, luminosity, concentration, asymmetry, clumpiness, etc.). With such correlations available one could explore for example the nature of bulges in isolated galaxies and how they are formed, the role of bars (if any) in the formation/evolution of bulges, whether the isolated spiral galaxies are different relative to spirals in richer environments in terms of global properties and/or in terms of properties of their components (bulge, bar, disks).  This paper is organized as follows: \\S~2 presents the selection and basic properties of the sample, \\S~3 offers a concise view on data reduction, \\S~4 and \\S~5 present the results of BUDDA decomposition analysis and CAS parametrization, respectively. \\S~6 combines various measures obtained from the BUDDA code with CAS parameters. \\S~7 is dedicated to discussion and conclusions. Throughout the paper we use H$_{o}$ = 75 km s$^{-1}$ Mpc$^{-1}$. ",
        "conclusions": "We have presented a detailed structural analysis for a well defined sample of $\\sim$ 100 late-type isolated galaxies. If a {\\it bona fide} isolated (pure ``nature'') population of galaxies exists then our previous work \\citep{sulentic06} suggests that it is dominated by systems with spiral morphology ($\\sim$ 84\\%) with the bulk in the range Sb-Sc (63\\%). We assume that the galaxies we investigate here are best described as ``minimal nurture and maximal nature'' systems because they are as isolated as individual galaxies can be. This hypothesis does not imply that these isolated galaxies have undergone no merger activity since their epoch of formation but rather that major mergers are probably absent from their past $\\sim$ 3Gyr history. We do note that the AMIGA sample includes 14\\% early-type galaxies and those are systems of such low luminosity as to suspect little or no major merger activity over their entire history \\citep{sulentic06}.  One might reasonably expect the tightest correlations between various intrinsic properties from a sample of isolated galaxies, where it is assumed that nurture (i.e. interactions) would increase the scatter (e.g. UBV-colors; \\citealt{larson78}). The strength of this study is manifold: the large size of the sample, the uniformity of the SDSS data, the robustness of the BUDDA code and the stringent isolation criteria underlying the definition of the parent AMIGA sample. In this study we have retained subjective morphological classifications and investigate morphological type dependence of various properties even though the typical range is narrow. This narrowness coupled with our ``nurture-free'' assumption raises the possibility that Hubble type T=4$\\pm$1 may represent the seed population for all spiral galaxies.  \\subsection{Pseudobulges in Isolated Galaxies}  We present evidence favoring the hypothesis that most or all late-type isolated galaxies host pseudobulges (\\S\\S~4.1) rather than classical bulges:  A. A large majority of our isolated systems host relatively ``unevolved'' bulge structures (as hypothesized by \\citealt{hunt04}); most S\\'{e}rsic indices (n$_{bulge}$) are smaller than 2.0-2.5 (see Table 5 and Figure 4) with the largest concentration around n$_{bulge}$ $\\sim$ 1.3-1.4. Such bulges are probably not as relaxed as larger bulges in earlier spiral types. They are likely dominated by rotation unlike higher S\\'{e}rsic index bulges (for a detailed discussion on this subject see section 4.6 in \\citealt{kormendy04}, section 4.2 in \\citealt{laurikainen07} and references therein).  B. We observe a large range of effective surface brightness $\\mu_{e}$ for a rather narrow range of r$_{e}$ (these two last parameters defining in part the fundamental plane) - see Figure 5a. The locus occupied by the bulges of our Sb-Sc galaxies in this plane is similar to that of disky bulges of galaxies at the end of dissipative collapse \\citep{capaccioli92}. The lack of correlation between $\\mu_{e}$ and r$_{e}$ supports the case of ``pseudobulges'' for isolated spiral galaxies in our sample. As pointed out in \\citet{macarthur03}, these results support an ``iceberg'' scenario, i.e. late-type spiral bulges are ``more deeply embedded in their host galaxy disk than earlier type bulges''. This idea is further complemented by the fact that the size of the bulge (r$_{e}$) scales with the scalelength of the disk (Figure 10a) \\citep[see also e.g.][]{khosroshahi00,macarthur03,mendez08,fisher08}.  We observe a larger dispersion in r$_{e}$/h$_{R}$ for larger Bulge/Total luminosity ratios (Figure 10b). However, in Figures 10b-c one can clearly see that Sbc-Sc galaxies do show a clear increasing trend for r$_{e}$/h$_{R}$ and r$_{e}$/a$_{25}^{i}$ with Bulge/Total and Bulge/Disk luminosity ratio, respectively. In contrast, Sb galaxies appear detached from the Sbc-Sc population. Assuming that the bulge S\\'{e}rsic index and/or Bulge/Total luminosity ratios are reasonable discriminators of pseudo- versus classical bulges (\\S\\S~4.1), then amongst our sample Sb galaxies have the greatest chance of hosting classical bulges. Thus, in Figures 10b-c we may have yet another indication that the pseudobulges and the galactic disks are clearly connected, while the classical bulges do not show similar scaling relations.  Some studies (e.g. \\citealt{thomas06}) argue that ``secular evolution through the disk and the phenomenon of pseudobulge formation are most likely restricted to spirals of types Sc and later''. Our results (but see also \\citealt{laurikainen07}) find a large fraction of pseudobulges among spiral types earlier than Sc (see \\S\\S~4.1). This may be telling us that the formation of pseudobulges does not appear exclusively restricted to Sc types or later. Our results suggest that if one considers only morphological types later than Sc, one may identify an almost pure pseudobulge population of galaxies. A fundamental question mentioned earlier is whether the isolated Sb-Sc spiral galaxies constitute the seed population of unnurtured spirals? If so then isolated galaxies might be expected to host a pure pseudobulge population. In this context Sb types in our sample have the greatest chance of bulge building via nurture and may involve a mixed classical and psuedobulge population. In Figure 4 it is interesting that a linear correlation emerges only for galaxies of Sb type which bridges the classical and pseudobulges domains. Alternatively, the trend may be telling us that all/most Sb galaxies contain a real (classical) bulge. This would suggest that some large bulges are natural or that all Sb spirals in the sample are a product of nurture. The latter interpretation is disfavored by the extreme isolation of our sample.  \\subsection{The Role of Bars in the Formation of Pseudobulges}  The results of the present study could set constraints for various galaxy formation and evolution models. Two important galaxy formation scenarios have been proposed and advocated: 1) spheroidal component (bulge) forms prior to the disk component in a monolithic collapse or via early mergers (so called ``inside out'' formation, e.g. \\citealt{eggen62,baugh96,kauffmann96,vandenBosch98,cole00,merlin06}) and 2) bulges form after the disk component as a result of secular dynamics/evolution driven by a disk instability \\citep[e.g.][]{courteau96,zhang04} possibly triggered by external satellite accretion \\citep[e.g][]{aguerri01,eliche06}. The former mechanism may be dominant for elliptical galaxies and in early spiral galaxies with large bulges (as they all appear to share similar properties and scaling relations within the fundamental plane; e.g. \\citealt{kormendy85,djorgovski87,faber87}). The latter mechanism may be more plausible for late type spiral systems \\citep[e.g.][]{carollo99,hunt04,debattista04}, as they largely harbor pseudobulges.  Some authors proposed that bulges of late type spiral galaxies are formed primarily through secular evolution of bars \\citep[e.g.][]{kormendy79,kormendy93,norman96,hasan98,fathi03,kormendy04,athanassoula05,jogee05,debattista06}. Others have suggested that bars can help the process of ``pseudobulge'' formation (making it faster and more efficient), but is not a necessary requirement for that process \\citep[e.g.][and references therein]{laurikainen07}. Bars can transport gas inward \\citep[e.g.][]{sakamoto99, sheth05} potentially contributing to the formation of a bulge. On the other hand it has been proposed that even without a bar the stellar disk component could be redistributed due to a secular torque action \\citep[e.g.][]{zhang07}.  We find a larger fraction of barred galaxies among Sb types relative to Sbc-Sc types (\\S\\S~4.2). Sb galaxies also appear to host the largest bars (Table 6a) within the morphological sequence Sb-Sbc-Sc. If bars are assumed as necessary precursors of all pseudobulges, then the smaller bars in later type galaxies ``dissolve'' more efficiently in the process of bulge formation. It is interesting to mention that for Sb and Sbc types in our sample of isolated galaxies we find systematically larger values of the index n$_{bulge}$ for barred galaxies compared to the non-barred galaxies (Tables 6a-b). The difference almost vanishes for Sc barred and non-barred. \\citet{laurikainen07} report a rather opposite result for Sb type (see their Figure 3). If n$_{bulge}$ is one of the empirical discriminators between classical and pseudobulges then any connection with the presence/absence of bars merits further attention. In this context it is relevant to review our Figures 8a-b. We note the ``disappearance'' of objects with n$_{bulge}$ above 1.7 for non-barred galaxies (Figure 8b in contrast to Figure 8a). We tested whether this may be caused by the resolution limitation in SDSS images, thus the BUDDA code's inability to identify the presence of a bar. First, we analyzed the distribution of n$_{bulge}$ values of non-barred galaxies with V$_{R}$ lower and higher than the median V$_{R}$ of the full non-barred sample ($\\sim$ 5700 km s$^{-1}$), respectively. We found no significant difference. Secondly, for our galaxies, the typical seeing FWHM is better than 1\\arcsec, with very few cases at 1.5\\arcsec. Considering the most extreme case, for a galaxy showing V$_{R}$ $\\simeq$ 10,000 km s$^{-1}$ a 1.5\\arcsec seeing would translate into a spatial resolution of $\\sim$ 1.0 kpc, which is well within the capability of the BUDDA code to provide reliable structural measures \\citep{gadotti08}. The scarcity of non-barred galaxies with n$_{bulge}$ above 1.7 is consistent with the scenario that bars could transform by dissolution into pseudobulges. The presence of bars may influence the degree of relaxation of bulges in the sense that n$_{bulge}$ decreases from Sb through Sc only for barred galaxies, but not for non-barred spirals (\\S\\S~4.2).  The formation and lifetime of bars may be sensitive to environment \\citep[e.g.][]{gerin90}. It has been suggested that bars in early type spiral galaxies are formed by tidal interactions with other galaxies and those in late types have intrinsic origin \\citep{noguchi96}. The connection bars-environment may be different for early and late type spirals \\citep{noguchi96, noguchi00}, being proposed a ``bimodality'' of bars in this sense. Moreover, numerical simulations have shown that for Sb-Sc galaxies bars are transient features and dissolve progressively in $\\sim$ 1-2 Gyr \\citep{bournaud05}. As we pointed out earlier, the AMIGA/CIG isolated galaxies have been basically nurture-free for at least a comparable time. We find that $\\sim$ 50-60\\% of our present sample are barred galaxies. The conclusion here could be that the bars we observe in these late type isolated spiral galaxies have been likely renewed or reformed through internal processes and not by external accretion or interactions \\citep[e.g.][]{block02,berentzen04}. It is also interesting to mention we find that the largest bars lie in disks with the lowest central surface brightness $\\mu_{o}$ (Figure 9b). This is consistent with the idea that bars build up from the material in the central parts of disks and they are products of secular dynamical evolution within the disk.  We find that our isolated galaxies tend to host large bars\\footnote{However, one must be aware that there is no standard definition for the length of a bar in a galaxy \\citep{erwin05} and scaling parameters for galaxy components may be sensitive to the filter that is used for photometry.}. Our Figure 6 shows that most bar radii are clustered in the range 2-6 kpc. \\citealt{erwin05} (based on \\citealt{martin95}) reports typical bar sizes in the range 1-3 kpc (B-band) for morphological types Sb-Sc having absolute magnitudes similar to our sample. A more recent study \\citep{marinova07} presents a characterization of bars in optical (B-band) and near-IR (H-band) for the OSUBSGS sample of galaxies. In order to compare the bar sizes with their estimates, we restrict their sample to Sb-Sc morphological range, based on the RC3 catalog \\citep{deVaucouleurs91}. The OSUBSGS-based sample has a similar distribution of absolute magnitudes as our sample. In terms of l$_{bar}$, our sample of barred galaxies (n=48) is characterized by a mean $\\sim$ 5.0 kpc and a median $\\sim$ 4.8 kpc. For The OSUBSGS sample of n=49 barred galaxies, the mean and median values (H-band) of l$_{bar}$ $\\sim$ 3.8 kpc and $\\sim$ 3.4 kpc, respectively. The conclusion is that the size of bars may be related to the environment, isolation favoring larger bars.  Moreover, this conclusion seems to be consistent with reports that the disk scalelength h$_{R}$ of spiral galaxies in rich environments is typically smaller than that of field (i.e. isolated) galaxies \\citep[e.g.][]{aguerri04}. We find that the bar size scales with the disk scalelength h$_{R}$ (our Figure 9; see also \\citealt{laine02}). In extreme environments (e.g. compact groups) spiral galaxies tend to lose their disk components by dissolution into a stellar halo. The size of the disks in what we assumed were initially late type spiral galaxies in Seyfert's Sextet for example (estimated by the last concentric isophote) is less than 10 kpc diameter, comparable to the smallest disks in our present sample \\citep{durbala08}. Comparing our $l_{bar}/a_{25}^{i}$ estimates with the similar quantities reported in \\citet{erwin05} we observe the same declining trend from Sb through Sc; for our sample we do obtain systematically larger $l_{bar}/a_{25}^{i}$ ratios relative to that study, although we note that \\citet{erwin05} measures are based on B-band data from \\citet{martin95}.  \\subsection{CAS Structural Measures in Isolated Galaxies}  The minimal environmental influence on AMIGA/CIG galaxies investigated here is revealed also by an analysis of the structural properties in terms of CAS parameters (\\S~5). Due to the narrow morphological range represented in our sample of isolated galaxies, any attempt at comparison with other studies must be cautiously explored. Nonetheless, the size of the sample examined in the present study allows a meaningful comparison of the 96 galaxies as a whole (i.e. the full set of Sb-Sc galaxies) with galaxies of same morphological types selected without isolation constraints. Table 11 offers such a comparison with the subsample of Sb-Sc galaxies (n=49) examined in \\citet{conselice03}, extracted from the \\citet{frei96} sample, assumed representative for the population of nearby normal galaxies. The general conclusion is that the isolated galaxies are less concentrated, less asymmetric and less clumpy than other galaxies of same morphological type selected without isolation criteria. Thus, we may have clear indications of environmental influence on the structure of galaxies. This may be telling us that the formation of large central concentrations and large clumps within disks are disfavored in the absence of comparable sized neighbors.  \\subsection{Describing the Morphological Classification}  Although our study involves a narrow range of morphological types all plots that involve exclusively bulge measures show clear morphological separation (e.g. Figure 5). When we combine disk measures (e.g. Figure 7a), the morphological segregation is less clear or absent suggesting some commonality among disk properties over the Sb-Sc range. Thus, it appears that the morphological separation may be associated with a change in the luminosity profile of bulges as indicated by their S\\'{e}rsic indices. \\citet{hunt04} proposed that spiral galaxies may begin with low bulge S\\'{e}rsic index. As they age they change into structurally more evolved systems (toward n$_{bulge}$ = 4 or higher) also  characterized by higher surface brightness (see Figures 5a, 12d) and an increased absolute magnitude (see Figures 5b-d). However, \\citet{carollo99} argue that pseudobulges cannot evolve into denser r$^{1/4}$ (i.e., n$_{bulge}$=4) bulges just by repeated cycles of bar formation/disruption.  At the same time one should keep in mind that the bulge S\\'{e}rsic index is associated with rather large uncertainties (\\S\\S~4.1), which complicates its use for a quantitative morphological classification \\citep{gadotti08}. It appears that the concentration index C, the Bulge/Total (or Bulge/Disk) luminosity ratio and the bulge S\\'{e}rsic index are relevant parameters when one describes the morphological sequence of spiral galaxies from earlier to later types. Nonetheless, Figure 12 suggests that the morphological diversity of spiral galaxies is deeply connected to the structure of their bulges. The concentration index C is not a good tracer of Bulge/Total ratio for Bulge/Total $>$ 0.1. This is true because the bulge light is no longer concentrated within the radius that includes 20\\% of the total light (section \\S~3; see also \\citealt{graham01b}). It is not obvious why the bulge surface brightness shows a plateau in its trend versus C (Figure 12d). However, one may speculate that the fact that some of the Sb galaxies curve away from the main trend (described largely by Sbc and Sc types) toward larger C values could be due to a different type of bulges they host.  \\subsection{Final Remarks}  This present study could be complemented by an extension of a similar type of analysis to the whole set of isolated spiral galaxies, which would include the whole sequence of Hubble morphological types. This would provide a more general and a more clear picture on the morphological type dependence of various structural properties and scaling relations presented and discussed here. Measures of bulge colors and kinematics would both provide strong tests of our hypothesis that most isolated spirals involve pseudobulges. Another complementary approach is a Fourier analysis of our images, which would provide a quantitative description of the spiral structure, intimately connected to galactic morphology as well. This is part of an ongoing project we are working on at this time and the results will be presented in a future paper."
    },
    "0807/0807.1348_arXiv.txt": {
        "abstract": "{ We present the results from an ESO/VLT campaign aimed at studying the afterglow properties of the short/hard gamma ray burst GRB\\,070707. Observations were carried out at ten different epochs from $\\sim 0.5$ to $\\sim 80$ days after the event. The optical flux decayed steeply with a power-law decay index greater than 3, later levelling off at $R\\sim 27.3$~mag; this is likely the emission level of the host galaxy, the faintest yet detected for a short GRB. Spectroscopic observations did not reveal any line features/edges that could unambiguously pinpoint the GRB redshift, but set a limit $z < 3.6$. In the range of allowed redshifts, the host has a low luminosity, comparable to that of long-duration GRBs. The existence of such faint host galaxies suggests caution when associating short GRBs with bright, offset galaxies, where the true host might just be too dim for detection. The steepness of the decay of the optical afterglow of GRB\\,070707 challenges external shock models for the optical afterglow of short/hard GRBs. We argue that this behaviour might results from prolonged activity of the central engine or require alternative scenarios.} ",
        "introduction": "Gamma-ray bursts (GRBs) are among the most powerful explosions in the universe. They are revealed in the hard X-ray/soft gamma-ray band and are followed in many cases by a fading afterglow observable from radio to X-ray wavelengths. GRBs are empirically classified in two groups \\citep{Maz81, Nor84, Kou93,Tav96}: short GRBs last less than 2\\,s and have a hard spectrum; long GRBs have longer durations (typically tens to hundreds seconds) and somewhat softer spectra.  The emergence of a typical supernova (SN) spectrum superposed on the rapidly decaying non-thermal afterglow weeks after the events and the association with blue, highly star-forming galaxies provided strong evidence that a significant fraction of long GRBs originates in the gravitational collapse of massive stars (but see \\citealt{MDV06, Fyn06, Gal06a}).   Short GRBs are revealed less frequently than long GRBs \\citep[they comprise about 1/4 and 1/10 of the  BATSE and \\textit{Swift} samples, respectively;][]{Kou93,Ber07}; moreover their afterglows are weaker and, thus, more difficult to detect and follow up. These are among the reasons why the origin of short GRBs is still under debate, despite the important progresses made in the \\textit{Swift} era. The tight upper limits on any associated SN \\citep{Hjo05,Cov06,Kann08} as well as the association with a broad variety of Hubble types hosts, from elliptical \\citep{Berger05} to moderate star forming galaxies \\citep[e.g.][]{Cov06, Fox06, Berger08}, rules out the core-collapse mechanism as the main channel for short-GRBs production and strongly suggest that the explosion mechanism and/or progenitors of short GRBs are different from those of long GRBs \\citep[for a recent review, see][]{LeRa07}.  The leading model for short GRBs involves the merging of a system composed of two collapsed objects, a double neutron star (DNS) or a black hole/neutron star binary. In those systems that evolve out of  massive stars that were born in a binary system (we term these ``primordial binaries''), the delay between formation and merging is dominated by the gravitational wave inspiral time, ranging from tens of Myr to a few Gyr \\citep{PeBe02}, strongly dependent on the initial system separation. Short GRBs that result from them are expected to: (a) have a redshift distribution which broadly follows that of star formation and (b) drift away in some cases from the star-forming regions in which they were born, and merge outside, or in the outskirts, of galaxies \\citep{Bel02}.  A fraction of two collapsed object binaries may also form dynamically through binary exchange interactions in the core of globular clusters \\citep{Gri06}. For such a formation mechanism, the delay between star-formation and merging is driven by the cluster core collapse time, which is comparable to the Hubble time \\citep{Hop06}. Therefore, short GRBs originating from dynamically formed double collapsed object binaries should go off at lower redshifts than short GRBs from primordial binaries \\citep{GuPi05,GuPi06,Gal06,Hop06,Salva08}. In another scenario, a fraction of the short GRBs is due to hyperflares from soft gamma-ray repeaters in the local  universe (distances up to $\\sim 100$\\,Mpc; \\citealt{Hur05,Tan05,Fre07,Maz08}).  The above summary emphasises that redshift determination, GRB position relative to the host galaxy and properties of the host galaxy are all crucial pieces of information for understanding short GRBs and discriminating among different models \\citep{Bel06}. One key issue in the study of short GRBs  is the secure identification of the host galaxy. Indeed, several short GRBs afterglows have been detected only in X-rays and thus localized with a precision of a few arcseconds (whereas sub-arcesecond localizations are required to unambiguously reveal their host). The possibility that  a sizeable fraction of short GRBs lie outside the light of their hosts, as predicted for both primordial and dynamically formed NS-NS/BH binaries, further complicates the identification process. Indeed some proposed short GRB associations with bright, nearby galaxies, based on some angular separation,   might result from by-chance alignment.   So far, a dozen short GRBs were localized with sub-arcsecond precision and host galaxies were firmly detected with small offset. Only in a few cases (e.g. \\object{GRB\\,061201}: \\citealt{Str07}; \\object{GRB\\,080503}: \\citealt{Perley08}) no host galaxy was found down to $R \\geq 26$--28 after the optical afterglow had faded. For other short GRBs at unknown redshift, a putative host galaxy was proposed with magnitude $R \\sim 23$--26. Spectroscopy of the brightest four of these galaxies indicates that they lie at $0.4 \\leq z \\leq 1.1$. A comparison with field galaxy magnitudes suggests that the rest of the sample lies at $z \\geq 1$ \\citep{Ber07b}. The unambiguously localized hosts are both early- and late-type galaxies, with very different star formation rates and masses \\citep{Nak07}.   The association with early-type galaxies has provided clear evidence that at least a fraction of the short GRBs have progenitors related to an older stellar population than that of long GRBs, as expected from the compact binary system models. The nature of the progenitors of short GRBs that go off in star forming galaxies is still under debate. The possibility that short GRBs comprise different subclasses cannot be confirmed yet but neither excluded.  In this paper we present the results from an extensive campaign aimed at studying the optical afterglow of \\object{GRB\\,070707}. This campaign monitored the decay of the optical afterglow evolution from about 0.45\\,days to more than one month after the burst. This is one of the best optical light curves for a short/hard GRB so far obtained. In Sect.\\,\\ref{sec:grb} we report on previous results on \\object{GRB\\,070707}. In Sect.\\,\\ref{sec:obs} our observations and data analyses are discussed and in Sect.\\,\\ref{sec:res} our results are presented. In Sect.\\,\\ref{sec:disc} we discuss our findings.  \\begin{table*} \\caption{VLT observation log for GRB\\,070707. Magnitudes are not corrected for Galactic absorption. Upper limits are given at $3\\sigma$ confidence level.} \\centering \\begin{tabular}{llrcccc} \\hline Mean time          &  Exposure time             & $t-t_0$    &  Seeing     & Instrument &  Magnitude   & Filter / Grism \\\\ (UT)               &  (s)                       &  (days)    & ($\\arcsec$) &            &              &                \\\\ \\hline 2007 Jul  08.12988 &  $10 \\times 120$           & 0.45742    & 0.9  & VLT/FORS1 & $23.05 \\pm 0.02$   & $R$       \\\\ 2007 Jul  09.07984 &  $10 \\times 120$           & 1.40738    & 1.0  & VLT/FORS1 & $23.86 \\pm 0.05$   & $R$       \\\\ 2007 Jul  09.23005 &  $ 1 \\times 180$           & 1.55759    & 1.0  & VLT/FORS1 & $24.07 \\pm 0.12$   & $R$       \\\\ 2007 Jul  10.14491 &  $10 \\times 120$           & 2.47245    & 0.7  & VLT/FORS1 & $25.33 \\pm 0.08$   & $R$       \\\\ 2007 Jul  10.21523 &  $20 \\times 3 \\times 30$   & 2.54277    & 0.7  & VLT/ISAAC & $> 23.6        $   & $J$       \\\\ 2007 Jul  11.13697 &  $20  \\times 180$          & 3.46721    & 0.5  & VLT/FORS1 & $26.62 \\pm 0.18$   & $R$       \\\\ 2007 Jul  12.17370 &  $17 \\times 300$           & 4.50124    & 0.8  & VLT/FORS1 & $26.81 \\pm 0.22$   & $R$       \\\\ 2007 Jul  15.21785 &  $18 \\times 300$           & 7.54539    & 0.6  & VLT/FORS1 & $27.39 \\pm 0.22$   & $R$       \\\\ 2007 Jul  19.12752 &  $20 \\times 300$           & 11.45506   & 0.6  & VLT/FORS1 & $27.21 \\pm 0.20$   & $R$       \\\\ 2007 Aug  15.07151 &  $20 \\times 300$           & 38.38355   & 0.6  & VLT/FORS1 & $27.15 \\pm 0.29$   & $R$       \\\\ 2007 Sep  28.03125 &  $20 \\times 3 \\times 60$   & 81.95974   & 0.5  & VLT/NACO  & $> 22.7$           & $K$       \\\\ \\hline 2007 Jul  09.32508 &  $5 \\times 2400$           & 1.65262    & 1.0  & VLT/FORS1 & ---       & $300V+$GG375   \\\\ 2007 Jul  13.15960 &  $2 \\times 900$            & 5.48714    & 0.8  & VLT/FORS1 & ---       & $300V+$GG375   \\\\ \\hline \\end{tabular} \\label{tab_log1} \\end{table*} ",
        "conclusions": "\\label{sec:disc}  \\begin{figure} \\includegraphics[width=\\columnwidth]{Amati_2.ps} \\caption{Location of GRB\\,070707 in the plane peak energy vs. isotropic-equivalent energy. The thick solid line shows the position of GRB\\,070707 as a function of redshift, with the diamonds indicating specific values discussed in this paper. Filled circles represent long-duration GRBs (from \\citealt{Amati08}), and the diagonal lines indicate the best-fit Amati relation (dashed) and the 2$\\sigma$ contours (dotted). Empty squares indicate other short-duration events with known redshift and peak energy \\citep{Ama06,Gole06,Ohno07,Gole07b}.} \\label{Amati} \\end{figure}  As with a number of other short GRBs, the nearly unconstrained redshift of GRB\\,070707 remains an important limiting factor. In the peak energy vs. isotropic-equivalent energy diagram (Fig. \\ref{Amati}), \\object{GRB\\,070707} should lie on the thick solid curve marked with selected values of the redshift. For the range of possible redshifts, the position of \\object{GRB\\,070707} would be in broad agreement with that of other short GRBs with known redshift and peak energy, which notoriously do not follow the so-called ``Amati-relation'' of long GRBs \\citep{Ama06}.  Whichever its redshift, the host of \\object{GRB\\,070707} ($R \\sim 27.3$) is the faintest ever detected for a short GRB, with a magnitude comparable to that of long GRB hosts at high redshift (e.g. \\citealt{Wain07,Fruchter06}). In past works \\citep[e.g.][]{Blo07,Str07,Lev07}, there has been discussion about the possibility to find short GRBs not spatially coincident with a host galaxy. In some cases, nearby, bright objects were proposed to be the GRB host based just on angular proximity. The case of GRB\\,070707 shows that caution is needed. Without deep VLT images, it might have been tempting to associate GRB\\,070707 with the closeby galaxy G6, which is certainly not the GRB host. With such a large magnitude contrast between the afterglow and the host, many more galaxies of short GRBs may be just fainter than the fluxes probed by shallower exposures. While this is consistent with the suggestion that a sizeable fraction of short GRBs reside at redshift larger than $z \\sim 1$ \\citep{Ber07b}, this may also indicate that some short GRBs go off inside low-luminosity galaxies at low redshift.   The sparse data on the host galaxy of GRB\\,070707 do not allow for a detailed analysis. At the maximum allowed redshift $z = 3.6$, the host would have an absolute magnitude $M_{\\rm AB} = -18.6$ (at rest-frame wavelength $\\lambda \\approx 1500$~\\AA). This is more that 2 magnitudes fainter than the Schechter luminosity at this redshift ($M^*_{\\rm AB} = -20.7$; \\citealt{Gabasch04}). We can thus firmly set $L < (1/6)L_*$ at any redshift, implying that the host was intrinsically faint. Our late-time NIR upper limit ($K > 22.7$) implies $R - K < 4.3$ and hence can rule out a bright, red host, such as an extremely red object (ERO) or a moderate-redshift elliptical. GRB\\,070707 hence confirms that short GRBs can explode inside faint and possibly extremely faint systems. Short GRB hosts indeed exhibit a wide range of luminosities.  Our spectrum did not show any clear absorption feature, which is quite possibly due to the low signal-to-noise and/or limited covered wavelength range. It is however interesting to note that a featureless spectrum has been reported also for the afterglow of the short GRB\\,061201 \\citep{Str07}. If confirmed by further, better-quality observations, this fact may provide hints for a lower density of short GRB environments compared to those of long-duration events.  Further information on the GRB progenitor comes from the optical light curve of the GRB\\,070707 afterglow, which decayed very steeply starting $\\sim 1.5$~days from the burst. The power law decay index $\\alpha$ is constrained by our fits to be steeper than 3, or may even be exponential. Assuming that the optical afterglow emission came from the forward shock, as commonly supposed in the fireball model for GRB afterglows, the steepening of the optical light curve could be interpreted as a jet break. However the index measured for GRB\\,070707 is too steep to be explained in terms of jetted emission. In fact this would require $\\alpha_2 = p$ \\citep{Rhoads99}, where $p$ is the electron energy distribution index. Such a steep decay could be marginally consistent with the post-break phase only by adopting a very soft electron energy distribution ($p > 3$). Such large $p$ values have never been found in GRB afterglows, both from theoretical and empirical investigations (e.g. \\citealt{Pana05,She06,Tag06, Kann06}). In particular, values of $p < 3$ are inferred based on afterglow spectra, in which case the analysis is more robust than when relying on the temporal behaviour. Moreover, the magnitude of the steepening from the pre- to the post-break decay would be too pronounced for a jet-break interpretation in GRB\\,070707.   The steepness of the  observed decay might be difficult to reconcile with forward shock emission even just for causality reasons. The fastest possible decay is the so-called high-latitude emission, which occurs when the fireball emission stops abruptly and the observers see photons coming from the wings of the emitting surface. In this case, $\\alpha = 2 + \\beta$, where $\\beta$ is the afterglow spectral index \\citep{Kumar00}. The observed upper limit in the $J$ band (Table\\,\\ref{tab_log1}) allows us to derive  $R-J < 2.1$ on Jul 10.2 UT, which corresponds to $\\beta < 1.95$. This would still be consistent with the high-latitude interpration. A stronger limit on $\\beta$ can however be inferred by using the X-ray data. Assuming a synchrotron spectrum, the optical spectral index can never be softer than the optical-to-X-ray spectral index $\\beta_{\\rm OX}$. For GRB\\,070707, the observed X-ray flux \\citep{Bea07b} implies $\\beta_{\\rm OX} = 0.75^{+0.13}_{-0.09}$. High-latitude emission from a source with such a spectrum cannot decay faster than $t^{-2.75}$, hence effectively ruling out such possibility for GRB\\,070707.    A viable alternative to explain the steep decay of GRB\\,070707 is that of a long-lived central engine. In this case, as discussed by several authors \\citep[e.g.][]{Zhang06}, the decay index should be computed after setting the zero time $t_0$ to the end of the extended activity phase. The intrisic decay slope would therefore be shallower than the observed value, eliminating the causality problem. The very steep decay of \\object{GRB\\,070707} thus would provide further evidence that the inner engine powering short GRBs is working for a much longer time than the observed gamma-ray emission. The optical light curve of \\object{GRB\\,070707} might be produced, for example, by a large flare, as also argued for \\object{GRB\\,050724} \\citep{Barthelmy05,Malesani07}.  Finally, we mention that the optical light curve of GRB\\,070707 could also be explained within the context of cannonball model (\\citealt{Dado08a}; \\citealt{Dado08b} and references therein). The emitting region in this case has angular dimension well below $1/\\Gamma$, and therefore the causality constraints are much more relaxed."
    },
    "0807/0807.3667_arXiv.txt": {
        "abstract": "Absorption lines of the $^{12}$CO and $^{13}$CO molecular bands ($\\Delta v$ = 2) at 2.29 -- 2.45 micron are modelled in spectrum of Arcturus (K2III). We compute a grid of model atmospheres and synthetic spectra for giant of Teff = 4300, log g = 1.5, and the elemental abundances of Peterson et al. (1993), but abundances of carbon, oxygen and the carbon isotopic ratio, \\CC are varied in our computations. The computed spectra are fitted to the observed spectrum of Arcturus from the atlas of Hinkle et al. (1995). The best fit to observed spectrum is achieved for log N(C) = -3.78 $\\pm$ 0.1, \\CC = 8 $\\pm$ 1. A dependence of the determined \\CC vs. log N(C) and log N(O) in atmospheres of red giants is discussed. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0807/0807.3445_arXiv.txt": {
        "abstract": "We study the $f(R)$ theory of gravity using metric approach. In particular we investigate the recently proposed model by Hu-Sawicki, Appleby $-$ Battye and Starobinsky. In this model, the cosmological constant is zero in flat space time. The model passes both the Solar system and the laboratory tests. But the model parameters need to be fine tuned to avoid the finite time singularity recently pointed in the literature. We check the concordance of this model with the $H(z)$ and baryon acoustic oscillation data. We find that the model resembles the $\\Lambda$CDM at high redshift. However, for some parameter values there are variations in the expansion history of the universe at low redshift. ",
        "introduction": "\\vskip 0.5cm \\noindent It is remarkable that different data sets of complementary nature such as supernovae, baryon oscillations, galaxy clustering, microwave back ground and weak lensing all taken together strongly support the late time acceleration of universe. In the standard lore, one assumes that the history of universe is described by the general relativity (GR). The late time acceleration can easily be captured in this frame work by introducing a scalar field with large negative pressure known as {\\it dark energy} \\cite{review}. In view of the fine tuning problem, the scalar field models, specially those with tracker like solutions, are more attractive compared to the models based on cosmological constant. At present, observations are not in a position to reject or to establish the dark energy metamorphosis. A host of scalar field models have been investigated in the literature. The scalar field models can fit the data but lack the predictive power. It then becomes important to seek the support of these models from a fundamental theory of high energy physics.  \\vskip 0.5cm \\noindent One can question the standard lore on fundamental grounds. We know that gravity is modified at small distance scales; it is quite possible that it is modified at large scales too where it has never been confronted with observations directly. It is therefore perfectly legitimate to investigate the possibility of late time acceleration due to modification of Einstein-Hilbert action. It is tempting to study the string curvature corrections to Einstein gravity amongst which the Gauss-Bonnet correction enjoys special status. A large number of papers are devoted to the cosmological implications of string curvature corrected gravity\\cite{fr,NOS,KM06,TS,CTS,Neupane,Cal,Sami06,Annalen,Sanyal}. These models suffer from several problems. Most of these models do not include tracker like solution and those which do are heavily constrained by the thermal history of universe. For instance, the Gauss-Bonnet theory with dynamical dilaton might cause transition from matter scaling regime to late time acceleration allowing to alleviate the fine tuning and coincidence problems. However, it is difficult to reconcile this model with nucleosynthesis\\cite{TS,KM06}. Another possibility of large scale modification is provided by non-locally corrected gravity which typically involves inverse of d'Alembertian of Ricci scalar. The non-local construct might mimic dark energy; the model poses technical difficulties and there has been a little progress in this direction\\cite{Wood}. The large scale modification may also arise in extra dimensional theories like DGP model which contains self accelerating brane. Apart from the theoretical problems, this model is heavily constrained by observation.  \\vskip 0.5cm \\noindent On purely phenomenological grounds, one could seek a modification of Einstein gravity by replacing the Ricci scalar by $f(R)$. The $f(R)$ gravity models have been extensively investigated in past five years\\cite{Nojiri:2003ft,review1,FRB,FRB1}. The $f(R)$ gravity theories  giving rise to cosmological constant in low curvature regime are plagued with instabilities and on observational grounds they are not distinguished from cosmological constant. The recently introduced models of $f(R)$ gravity by Hu-Sawicki and Starobinsky (referred as HSS models hereafter) with disappearing cosmological constant\\cite{HS,star} have given rise to new hopes for a viable cosmological model within the framework of modified gravity ($f(R)$ gravity model with similar properties is proposed in Ref.\\cite{ABTT}). These models contain Minkowski space time as a solution in the low curvature regime which is an unstable solution. In high curvature regime these models reduce to cosmological constant. Both the first and the second derivatives of $f(R)$ with respect to R are positive. The positivity of the first derivative ensures that the scalar degree of freedom, a characteristic of any f(R) theory, is not tachyonic where as the positivity of second derivative tells us that graviton is not ghost thereby guaranteing the stability. In Starobinsky parametrization\\cite{star}, $f(R)$ is given by, $f(R) = R + \\Lambda \\left \\lbrack \\left ( 1 + R^2/R_0^2 \\right )^{-n} - 1 \\right]$. The HSS models can evade solar physics constraints provided that the model parameters are chosen properly. An important observation has recently been made by by Appleby $-$ Battye and Forolov\\cite{AppBatt,frolov} (see also \\cite{Nojiri:2008fk}). The minimum of scalaron potential which corresponds to dark energy can be very near to $\\phi=0$ or equivalently $R=\\infty$. As pointed out in Ref.\\cite{Tsujikawa}, the minimum should be near the origin for solar constraints to be evaded. Hence, it becomes most likely that we hit the singularity if the parameters are not fine tuned.  \\vskip 0.5cm \\noindent In order to check whether the $f(R)$ gravity theory is cosmological viable or not, it is necessary that this theory  must be compatible with the observations. In this work, we study the Starobinsky model using the data from the recent observations which include $H(z)$ , Hubble parameter  at various red-shifts and the Baryon Acoustic Oscillation (BAO) peak from Sloan Digital Sky Survey (SDSS).  \\vskip 0.5cm \\noindent This paper is organised as follows. In Section II, we describe the general properties of the model highlighting the fine tuning problem. The Friedmann equation and the special cases of trace equation is studied in Section III. The cosmological constraints from the recent observations are described in Section IV. Finally, Section V contains the results and discussions. ",
        "conclusions": "\\no The observed late time acceleration of the universe is one of the major unsolved problem in the cosmology. It may hint at the breakdown of Einstein GR . This has lead to modification of the Einstein theory of gravity. One of the attractive possibility to modify this theory is to replace the Ricci scalar $R$ with the generic function $f(R)$ in the Hilbert action.  \\vskip 0.5 cm \\no $f(R)$ theories are usually studied by two methods: {\\it metric} and {\\it Palatini} approach. The Palatini approach is explored extensively in the literature both theoretically and observationally \\cite{pal}( see the other ref's given in these papers). This formulation gives second order differential field equation which can explain the late time behavior of the universe.  \\vskip 0.5 cm \\no The metric approach of $f(R)$ theory  leads to a fourth order non linear differential equation in terms of scale factor. This equation is difficult to solve both analytically and numerically even for the special cases. Therefore, not much observational tests have been performed on $f(R)$ theories based on metric formulation. Our work is an attempt to check the concordance of $f(R)$ theory of gravity using the metric approach  with some of the cosmological observations. In particular, we have explored  the Starobinsky model in which function $f(R)$ is analytic, satisfying the condition $ f(0) = 0$. This model also passes the Solar system and laboratory tests successfully for large values of $n$. In this work we study the $f(R)$ model for $n= 1$ and $n=2$. \\vskip 0.5 cm \\noindent In order to study the cosmological viability of this theory, we investigate this class of model with two observational tests. The first method is based on the Hubble parameter versus redshift data, $H(z)$. The Hubble parameter  is related to the differential age of the universe through this form $$ H(z) = - \\frac{1}{1 +z} \\frac{dz}{dt}. $$ By estimating the $dt/dz$, one can obtain directly the Hubble parameter, $H(z)$ at different redshifts. The H(z) data has one major advantage, unlike in the standard candle approach (SNe Ia): the Hubble function is not integrated over. The other important feature of this test is that differential ages are less sensitive to systematic errors as compared to the absolute ages \\cite{svj}. This observational $H(z)$ have been used earlier also to constrain various other dark energy models \\cite{other}.  \\vskip 0.5 cm \\noindent In this work, we check the compatibility of $f(R)$ model with $n =1$ and $n=2$ with the $H(z)$ data. Since $H_0$ is a nuisance parameter, we marginalise over $H_0$. We further combine the results obtained from $H(z)$ data set with the BAO data. To perform the joint test, we define the quantity: \\begin{equation} \\chi^2_{\\mathrm{total}} = \\chi^2_{\\mathrm{H(z)}} + \\chi^2_{\\mathrm{BAO}}. \\end{equation}  \\no The results are given in the Table 1. It appears that $n = 2$ is favored by the observations. This fact is also in agreement with solar and laboratory test. The other important conclusions are as follow:  \\begin{itemize} \\item \\no The variation of $H(z)$ with redshift becomes independent of $n$ after the redshift around $ z =1.8$, see Fig. 7. Before this redshift ( $ z < 1.8$),  there is a small difference between the behavior of $H(z)$ with $z$ for $ n =1$ and $n =2$ models. This variation is still within the error bars of the data points. At higher redshift ( i.e when $R >> R_0$) we recover the standard $\\Lambda$CDM universe for $\\lambda =2$, for all the values of $n$. Hence the thermal history of the universe is correctly reproduced by this model.   \\item \\no The expansion history for this model with $ n = 2, \\lambda = 2$ matches exactly with the standard cosmological model ($\\Lambda$CDM) with equation of state parameter $\\omega = -1, \\Omega_m = 0.3$ and $\\Omega_{\\Lambda} = 0.7$ (see fig. 7). The variation in parameters $\\Omega_m$ and $\\Omega_{\\Lambda}$ (keeping $\\Omega_{m0} + \\Omega_{\\Lambda} = 1$) leads to very small changes in values of $\\lambda$, keeping it close to 2.   \\item \\no As shown in Table 1, it seems that the present observational data used in this work prefer $n = 2$ over $n = 1$, as indicated by $\\chi^{2}$ per degree of freedom, $\\chi^2_{\\nu}$. However, $ n=1$ cannot be ruled out.  \\item \\no As stated earlier, the algebraic expressions become cumbersome for larger values of $n$. It is quite possible that observational tests discussed in this paper may be compatible with values of $n$ larger than two which are also consistent with solar and laboratory tests. However, it should be emphasized that for a given value of $\\lambda$, the minimum of the scalaron potential gets closer to zero for larger values of $n$ making the model vulnerable to singularity, see Fig. 3.  \\end{itemize}  \\vskip 0.5cm \\no Due to large error bars in the data sets used, the variation of the expansion history of the universe studied in this paper can easily be accommodated by the observations, thereby making the models close to $\\Lambda$CDM at the background level. \\no At present the sample of $H(z)$ data is too small  and the error bars are large. In the future, large amount  of precise $H(z)$ data is expected to become available. This will not only  reveal the fine features of the expansion history of the universe but also tightly constrain the cosmological parameters.   \\vskip 0.5cm \\no There is a further need to explore this model with time based observational tests, like age of the universe and high redshift objects. It is  known that the evolution of the age of the universe with the redshift vary from model to model. So it is possible that the model of the universe which are able to reproduce the total age of the universe at $z= 0$, may not accommodate the objects at high redshift\\cite{dj}. Therefore the time based observational test may  play the key role in short-listing the viable dark energy model in the near future. It is also expected that the study of matter power perturbations would allow to distinguish this model from the standard cosmology."
    },
    "0807/0807.4579_arXiv.txt": {
        "abstract": "We present 1.4 and 2.5 GHz Australia Telescope Compact Array (ATCA) observations of the galaxy cluster A3158 (z=0.0597) which is located within the central part of the Horologium-Reticulum Supercluster (HRS). Spectroscopic data for the central part of the HRS suggest that A3158 is in a dynamically important position within the supercluster and that it is moving toward the double cluster system A3125/A3128 which marks the centre of the HRS. A total of 110 radio galaxies are detected in a 35 arcminute radius about the cluster at 1.4 GHz, of which 30 are also detected at 2.5GHz. We examine the source counts and compute the Radio Luminosity Function (RLF) at 1.4 GHz from the subset of 88 sources found within the full-width half-power area of the ATCA beam. Comparison of the source counts in the area over the background, as computed by Prandoni et al.~(2001), shows some evidence of an excess of galaxies with L$_{1.4 GHz} \\leq 2 \\times 10^{22} W Hz^{-1}$. This result seems to indicate a star forming population and is a result similar to that found recently by Owen et al.~(2005) for the merging cluster A2125. In addition we find that the radio luminosity function for early-type galaxies (E and S0) below log P$_{1.4} \\sim$ 22.5 is lower than that found for a composite cluster environment (Ledlow \\& Owen, 1996) but is similar to the early-type RLF for clusters in the centre of the Shapley Super cluster (Venturi et al.~ 2000) which are believed to be in the latter stages of merging. This result implies that the cores of superclusters are environments where radio emission, particularly resultant from AGN, is suppressed in the later stages of merging. Thus, radio observations of clusters might be sensitive indicators of the precise merger stage of the cluster but more observational evidence is still required to establish this trend. ",
        "introduction": "Observational and numerical studies suggest that dynamical interactions are frequent occurrences in clusters of galaxies and that these events are the most energetic in the recent history of the Universe. Matter is believed to stream along preferred axes (so-called filaments and sheets) often accelerating in-falling galaxies to velocities of the order of 10$^3$ km s$^{-1}$. Rich clusters form at the intersection of the filament and sheets, and larger mass congregations, such as superclusters, are often detected as a collection of several clusters and filaments. These dynamical processes cause shocks, turbulence and bulk flows along the filamentary axes resulting in disturbance of the intracluster medium (ICM). Such perturbations are frequently detected with X-ray observations which show a variety of features including sharp temperature jumps, distorted X-ray isophotes and regions of enhanced pressure and entropy. At radio wavelengths, cluster mergers are thought to generate large areas of diffuse synchrotron emission either as a central, unpolarised halo or one or occasionally two peripheral ``radio relics''. Cluster ``weather'' and bulk flows have also been cited as causes of the bending of the jets in head-tail galaxies (Burns et al.~1998, Johnston-Hollitt et al.~2004), although recent results suggest that more local conditions might play a more important role in their formation (Mao et al.~2008).  \\begin{figure*} \\begin{center} \\includegraphics {mjh_fig1.eps} \\caption{Plot of all spectroscopic redshifts available in the region between A3158 (shown on the left, marked with a circle) and the A3125/A3128 complex on the right, again with the clusters denoted by large circles. Filled circles represent objects which have velocities in the range of 17425 km s$^{-1}$ and 18575 km s$^{-1}$ which encompasses those in the bridge investigated by Fleenor (2006).} \\label{bridge} \\end{center} \\end{figure*}   In addition to the large-area phenomena discussed above, the dynamical history of galaxy clusters should also be imprinted on the constituent galaxies. In particular, increased secondary star formation either observed in the optical (Burns et al.~1994; Caldwell \\& Rose 1997) or through increased low-power radio emission (Owen et al.~1999) is one possible signature of the history of a dynamical cluster. However, the balance between mergers driving gas towards the centre of the galaxies and hence feeding the central AGN to induce a starburst (Bekki 1999) and an increase in the ram pressure which would strip gas from galaxies and hence curtail the feeding of the AGN (e.g. Balogh et al.~1998; Fujita et al.~1999) still needs to be addressed observationally. To date, few studies have been conducted with both sufficient sensitivity in the radio to probe the low-powered radio sources, and good optical observations to confirm cluster membership. The most notable progress to date has been made on sections of the Shapley Supercluster with several areas analysed in both the radio and optical, including radio observations of A3556 (Venturi et al.~1997), A3571 (Venturi et al. 2002), the A3528 complex (Venturi et al.~2001) and the A3558 complex (Venturi et al.~2001; Miller 2005) and its outskirts (Giacintucci et al.~2004, Miller 2005). In addition to clusters in superclusters, the host radio population for several individual clusters have also been examined including A2255 (Miller \\& Owen, 2005), A2256 (Miller et al.~2003), A2125 (Owen et al.~2005) and A2111 (Miller et al.~2006). However, at present, those observational studies that have been undertaken give incomplete results with both instances in which effects are seen on the radio population, and cases where they are not observed. Sometimes neighbouring clusters in the same supercluster show markedly different results in terms of radio population. This may well be due to the radio galaxy population only being affected in certain phases of dynamical activity (Venturi et al.~2002) however, the exact point at which a merger might affect the radio population is, as yet, unclear. Nevertheless, large-area observations of superclusters, such as Shapley, present ideal datasets with which to investigate the way in which environment affects the galaxy population.  In this paper, which will be the first in a series, we present a study of the radio population of A3158 which lies within the kinematic core of the massive Horologium-Reticulum Supercluster (HRS). The HRS spans around 180 square degrees on the sky (Fleenor et al.~2006), contains upwards of 20 galaxy clusters (Zucca et al.~1993), is comprised of at least two major filaments (Einasto et al.~2003), and is the second largest mass concentration in the local 300 Mpc (Hudson et al.~1999). A3158 lies adjacent to the so-called center of the HRS, defined by the double cluster complex A3125/A3128 (Rose et al.~2002). The paper will be laid out as follows: Section 2 will discuss A3158 in detail; Section 3 will give details of the radio observations of the cluster; Section 4 will present the radio source population and in Sections 5 and 6 we will discuss the resultant radio source counts and radio luminosity functions; Sections 7 \\& 8 will present the discussion and conclusions. Throughout this paper we will assume H$_0$=71 km s$^{-1}$ Mpc$^{-1}$ which at the average redshift of A3158 (z=0.0597) means that 1 arcsecond $\\sim$ 1.1 kpc.  \\begin{table} \\begin{center} \\label{optp} \\caption{Known cluster parameters for A3158. Column 1 and 2 give the parameter and the most recent value from the literature. The reference is given in column 3 where 1 means Havlen \\& Quintana~(1978), 2 is Ebeling et al.~(1996), 3 is  Lokas et al.~(2006) and 4 is Fleenor et al.~(2006).} \\begin{tabular}{lll} \\hline Parameter & Value & Reference\\\\ \\hline redshift & 0.0597 & 3\\\\ ellipticity & 0.3 & 1\\\\ X-ray Luminosity & 5.31 $\\times 10^{44}$ erg s$^{-1}$ & 2\\\\ velocity & 17,910 kms$^{-1}$ & 4\\\\ Redshift n=145 & 0.0597 & 3\\\\ velocity disp. (n=145)& 970$\\pm$57 kms$^{-1}$ & 3 \\\\ Viral Mass & $15.4^{+7.6}_{-5.4} 10^{14}$M$_{\\sun}$ & 3\\\\ \\hline \\end{tabular} \\end{center} \\end{table} ",
        "conclusions": "We present the the first of a series of papers investigating radio emission in the Horologium-Reticulum Supercluster. In this paper we present radio imaging at 1.4 and 2.5 GHz of the cluster A3158, from which we detect 109 radio sources of which 64 are found to have optical counterparts. By considering a reduced subsample of 87 sources within a $23\\farcm05$ radius, over which the primary beam attenuation of the ATCA is negligible, we are able to investigate the statistical properties of radio emission in the cluster. Our results can be summarised as follows.  (1) The radio source counts in A3158 are unlikely to be consistent with the background source counts and in particular there appears to be an excess of galaxies with S$_{1.4} \\leq$ 3.96 mJy; a KS test of the entire sample against the background only returns a $\\sim$3\\% - 8\\% chance that the distributions are drawn from the same population, while a KS test over the sample with S$_{1.4} \\leq$ 3.96 mJy gives a much higher chance ($\\sim$60\\%) that the distributions are drawn from the same population suggesting that the low-powered excess is the significant difference.  (2) The percentage of radio galaxies associated with the cluster as determined by spectroscopic redshifts is highly correlated to radio flux/power which is consistent with the excess source counts being the result of low-powered radio galaxies associated with A3158; in addition colours of these galaxies suggest they are blue which is likely to indicate starformation as a driving mechanism.  (3) The alignment of the candidate starforming galaxies associated with the increased source counts follows the bridge of galaxies connecting A3158 to the A3125/A3128 complex suggestive of merger-induced formation.  (4) The radio luminosity function for early-type galaxies is significantly lower than for other composite cluster environments and the field; this result is similar to the results for the late merger in A3558 obtained by Venturi et al.~(2000).  Thus, if we can use the radio galaxy population as an indicator of merger history it is likely that A3158 is observed in a late merger state with a significant excess of low-powered, blue galaxies aligned along the axis connecting this cluster to the A3125/A3128 complex. These galaxies are likely the result of merger-induced starformation."
    },
    "0807/0807.4323_arXiv.txt": {
        "abstract": "Astrometric and radial-velocity planet detections track very similar motions, and one generally expects that the statistical properties of the detections would also be similar after they are scaled to the signal-to-noise ratio of the underlying observations.  I show that this expectation is realized for periods small compared to the duration of the experiment $P/T\\ll 1$, but not when $P/T\\ga 1$.  At longer periods, the fact that models of astrometric observations must take account of an extra nuisance parameter causes the mass error to begin deteriorating at $P/T\\sim 0.8$, as compared to $P/T\\sim 1.0$ for RV.  Moreover, the deterioration is much less graceful. This qualitative difference carries over to the more complicated case in which the planet is being monitored in the presence of a distant companion that generates an approximately uniform acceleration.  The period errors begin deteriorating somewhat earlier in all cases, but the situation is qualitatively similar to that of the mass errors. These results imply that to preserve astrometric discovery space at the longest accessible orbits (which nominally have the lowest-mass sensitivity) requires supplementary observations to identify or rule out distant companions that could contribute quasi-uniform acceleration. ",
        "introduction": "\\label{sec:intro}}  Astrometric and radial-velocity (RV) planet detections are, from a mathematical standpoint, extremely similar.  In each, one models 1-dimensional projections of Kepler orbits and attempts to fit Kepler parameters.  In the limit of circular orbits, the planet signature is encoded in the simple form \\begin{equation} F(t;a_1,a_2,a_3) = a_1\\sin(a_2 t + a_3), \\label{eqn:basicform} \\end{equation} where \\begin{equation} a_1 = \\alpha \\quad {\\rm (astrometry)},\\qquad a_1 = K \\quad {\\rm (RV)}, \\label{eqn:a1eq} \\end{equation} are the astrometric and velocity semi-amplitudes for the two cases, and where \\begin{equation} a_2 = {2\\pi\\over P}, \\quad a_3 = \\phi \\qquad {\\rm (astrometry + RV)} \\label{eqn:a23eq} \\end{equation} designate the period $P$ and phase $\\phi$, respectively. One then goes on to combine information about the star's mass and (for astrometry) its distance, to infer the planet mass $m$ (astrometry) or $m\\sin i$ (RV), where $i$ is the inclination. In the astrometric case, there are of course two such equations, one for each direction in the plane of the sky, whose ratio gives $\\cos i$, and so permit one to break the $m\\sin i$ degeneracy that plagues the intrinsically 1-dimensional RV measurement.  Because the form of equation (\\ref{eqn:basicform}) is essentially identical in the two cases, it is generally assumed that the error properties, i.e., the relation between the measurement errors and the derived-parameter errors, is also the same. Of course, it is well known that the parameter errors have a different dependence on semimajor axis $a$, system distance $D$, etc. Most notably, with other parameters held fixed, astrometric sensitivity increases linearly with $a$ whereas RV sensitivity declines as $a^{-1/2}$.  But here I am referring to something else.  The differences just mentioned all impact the final result because they change the characteristic signal-to-noise ratio of the experiment, \\begin{equation} {\\rm SNR}\\equiv {a_1\\over\\sigma}\\sqrt{N}, \\label{eqn:snrdef} \\end{equation} where $N$ is the number of measurements and $\\sigma$ is the error in each measurement (assumed for simplicity to be all the same).  Here I will show that the astrometric and RV measurements have substantially different error properties even when SNR is identical. ",
        "conclusions": "\\label{sec:discussion}}  The results derived here are primarily of interest in regard to future astrometric missions such as GAIA and SIM.  The first point is that planets with periods that are even slightly longer than the mission cannot be reliably detected unless they are many times more massive than the nominal thresholds of detection.  And second, if one is forced to allow for uniform acceleration, then the same statement applies to $P/T\\ga 0.7$.  Hence, to preserve the discovery space at the longest accessible periods, it is critically important to identify or rule out distant perturbers by supplementary data, such as longer-term RV observations to look for large, distant planetary or stellar companions, or perhaps AO observations to find stellar companions.  This should be relatively straightforward for SIM, with its limited number of high-precision targets, but may be more difficult for GAIA, which is a survey instrument."
    },
    "0807/0807.3029_arXiv.txt": {
        "abstract": "We study an extension of the MSSM with a singlet supermultiplet $S$ with coupling $SH_1H_2$ in order to solve the $\\mu$ problem as in the NMSSM, and right-handed neutrino supermultiplets $N$ with couplings $SNN$ in order to generate dynamically electroweak-scale Majorana masses. We show how in this model a purely right-handed sneutrino can be a viable candidate for cold dark matter in the Universe. Through the direct coupling to the singlet, the sneutrino can not only be thermal relic dark matter but also have a large enough scattering cross section with nuclei to detect it directly in near future, in contrast with most other right-handed sneutrino dark matter models. ",
        "introduction": "Weakly interacting massive particles (WIMPs) are among the best motivated candidates for explaining the cold dark matter in the Universe. WIMPs appear in many interesting extensions of the standard model providing new physics at the TeV scale. Such is the case of supersymmetric models, in which imposing a discrete symmetry (R-parity) to avoid rapid proton decay renders the lightest supersymmetric particle (LSP) stable and thus a good dark matter (DM) candidate.   The minimal supersymmetric extension of the standard model (MSSM) provides two natural candidates for WIMPs, the neutralino and the (left-handed) sneutrino. The neutralino is a popular and extensively studied possibility. On the contrary, the left-handed sneutrino in the MSSM~\\cite{Ibanez} is not a viable dark matter candidate. Given its sizable coupling to the $Z$ boson, they either annihilate too rapidly, resulting in a very small relic abundance, or give rise to a large scattering cross section off nucleons and are excluded by direct DM searches~\\cite{Falk:1994es} (notice however that the inclusion of a lepton number violating operator can reduce the detection cross section~\\cite{Hall:1997ah}).   However, there is a strong motivation to consider an extension of the MSSM, the fact that neutrino oscillations imply tiny but non-vanishing neutrino masses. These can be obtained introducing right-handed neutrino superfields. Several models have been proposed to revive sneutrino DM by reducing its coupling with Z-boson. This can be achieved by introducing a mixture of left- and right-handed sneutrino~\\cite{ArkaniHamed:2000bq,Arina:2007tm,valle}, or by considering a purely right-handed sneutrino~\\cite{Asaka:2005cn,Gopalakrishna:2006kr,McDonald:2006if,Lee:2007mt}. In the former, a significant left-right mixture is realized by adopting some particular supersymmetry breaking with a large trilinear term~\\cite{ArkaniHamed:2000bq}. Such a mechanism is not available in the standard supergravity mediated supersymmetry breaking, where trilinear terms are proportional to the small neutrino Yukawa couplings. Recently, another realization of large mixing was pointed out~\\cite{valle} by abandoning the canonical see-saw formula for neutrino masses. On the other hand, pure right-handed sneutrinos cannot be thermal relics, since their coupling to ordinary matter is extremely reduced by the neutrino Yukawa coupling~\\cite{Asaka:2005cn,Gopalakrishna:2006kr,McDonald:2006if}. Furthermore, these would be unobservable in direct detection experiments. Another possibility to obtain the correct thermal relic density would consist of coupling the right-handed sneutrino to the observable sector, e.g., via an extension of the gauge \\cite{Lee:2007mt} or Higgs \\cite{pilaftsis,dp} sectors   There is one more motivation to consider another extension of the MSSM, % the so-called ``$\\mu$ problem''~\\cite{Kim:1983dt}. The superpotential in the MSSM contains a bilinear term, $\\mu H_1 H_2$. Successful radiative electroweak symmetry breaking (REWSB) requires $\\mu$ of the order of the electroweak scale. The next-to-minimal supersymmetric standard model (NMSSM) offers a simple solution % by introducing a singlet superfield $S$ and promoting the bilinear term to a trilinear coupling $\\l S H_1 H_2$. After REWSB, % $S$ develops a vacuum expectation value (VEV) of order of the electroweak scale thereby providing an effective $\\mu$ term, $\\mu=\\lambda \\langle S\\rangle$. Furthermore, the NMSSM also alleviates the ``little hierarchy problem'' of the Higgs sector in the MSSM~\\cite{BasteroGil:2000bw} and has an attractive phenomenology, featuring light Higgses and interesting consequences for neutralino DM \\cite{Cerdeno:2004xw}. Although the $Z_3$ symmetry of the NMSSM may give rise to a cosmological domain wall problem, this can be avoided with the inclusion of non-renormalisable operators.   Motivated by the above two issues, we study an extension of the MSSM where singlet scalar superfields are included % \\cite{ko99,pilaftsis}. A singlet $S$  in order to solve the $\\mu$ problem as in the NMSSM (and which accounts for extra Higgs and neutralino states) and right-handed neutrinos $N$ to obtain non-vanishing neutrino Majorana masses with the canonical, but low scale, see-saw mechanism. Terms of the type $SNN$ in the superpotential can generate dynamically Majorana masses through the VEV of the singlet $S$. In addition, the presence of right-handed sneutrinos, $\\tilde N$, with a weak scale mass provides a new possible DM candidate within the WIMP category.   In this letter we analyse the properties of right-handed sneutrinos, showing that not only they can be thermally produced in sufficient amount to account for the DM in the Universe because of the direct coupling between $S$ and $N$, but also that their elastic scattering cross section off nuclei is large enough to allow their detection in future experiments. ",
        "conclusions": "We propose the right-handed sneutrino as a viable thermal DM candidate in an extension of the MSSM where the singlet superfields, $S$ and $N$, are included to solve the $\\mu$ problem and account for neutrino masses. A direct coupling between $S$ and $N$ provides a sufficiently large annihilation cross section for the right-handed sneutrino, as well as a detection cross section in the range of future direct DM searches.   DGC was supported by the program ``Juan de la Cierva''. CM was supported by the MEC project FPA2006-05423 and the EU  program MRTN-CT-2004-503369. DGC  and CM were also supported by the MEC project FPA2006-01105, and  the EU network MRTN-CT-2006-035863. OS was supported by the MEC project FPA 2004-02015. We thank the ENTApP Network of the ILIAS project RII3-CT-2004-506222 and the project HEPHACOS P-ESP-00346 of the Comunidad de Madrid."
    },
    "0807/0807.2733_arXiv.txt": {
        "abstract": "s{ M\\,87 is the first extragalactic source detected in the Very High Energy (VHE; $E > 100$\\,GeV) $\\gamma$-ray domain that is not a blazar, its large scale jet not being aligned to the line of sight. Slight modification of standard emission models of TeV blazars allows to account for the $\\gamma$-ray spectra obtained with H.E.S.S. We present a multi-blob synchrotron self-Compton model taking explicitly into account large viewing angles and moderate values of the Lorentz factor as inferred from MHD simulations of jet formation. Predictions of the VHE emission for the nearby radiogalaxy Cen\\,A and an interpretation of the broadband radiation of M\\,87 are presented. } ",
        "introduction": " ",
        "conclusions": "Our scenario shows the ability to extend the interpretation of TeV blazars by leptonic models to active galactic nuclei with intermediate beaming which jet is not strictly aligned with the line of sight.  The {\\it GLAST} mission and H.E.S.S.\\,II project, and CTA in a near future, will allow to probe the sub-TeV spectral region, which is still poorly known and yet essential to constrain the inverse Compton bump. This is required to distinguish between leptonic or hadronic origin of the high energy emission of the jet in these objects."
    },
    "0807/0807.2443_arXiv.txt": {
        "abstract": "We present and discuss the bounds from the energy conditions on a general $f(R)$ functional form in the framework of metric variational approach. As a concrete application of the energy conditions to locally homogeneous and isotropic $f(R)-$cosmology, the recent estimated values of the deceleration and jerk parameters are used to examine the bounds from the weak energy condition on the free parameter of the family of $f(R)=\\sqrt{R^2 - R_0^2}\\;$ gravity theory. ",
        "introduction": "The observed late-time acceleration of the Universe poses a great challenge to modern cosmology, which  may be the result of unknown physical processes involving either modifications of gravitation theory or the existence of new fields in high energy physics. This latter route is most commonly used, however, following the former, an attractive approach to this problem, known as $f(R)-$gravity,\\cite{Kerner} examines the possibility of modifying Einstein's general relativity (GR) by adding terms proportional to powers of the Ricci scalar $R$ to the Einstein-Hilbert Lagrangian (see also Refs.~\\refcite{Francaviglia} for recent reviews). Although these theories provide an alternative way to explain the observed cosmic acceleration without dark energy, the freedom in the choice of different functional forms of $f(R)$ gives rise to the problem of how to constrain on theoretical and/or observational grounds, the many possible $f(R)-$gravity theories. Theoretical limits have long been discussed in Refs.~\\refcite{Amendola}, while only recently observational constraints from several cosmological datasets have been explored for testing the viability of these theories.\\cite{Fabiocc}  Additional constraints to $f(R)$ theories may also arise by  imposing the so-called energy conditions.\\cite{Kung}  It is well known that these conditions, initially formulated in GR context,\\cite{Hawking} have been used in different contexts to derive general results that hold for a variety of physical situations.\\cite{EC} More recently, several authors have employed the GR classical energy conditions to investigate cosmological issues such as the phantom fields\\cite{Santos} and the expansion history of the universe.\\cite{Nilza} While they are well founded in the context of GR, one has to be cautious when using the energy conditions in a more general framework, such as the $f(R)-$gravity. In this regard, in a recent work Santos \\emph{et al.}\\cite{SARC} have used Raychaudhuri's equation along with the requirement that gravity is attractive, to derive the energy conditions for a general $f(R)-$gravity in the metric formulation.  They have shown that, although similar, the energy conditions differ from their formulation in GR context. In this work, we use estimated values of the deceleration and jerk parameters to examine the bounds from these newly derived $f(R)-$energy-conditions on the one-parameter family of a recently proposed $f(R)-$gravity.\\cite{Baghram-a,Baghram-b} ",
        "conclusions": ""
    },
    "0807/0807.1056_arXiv.txt": {
        "abstract": "The GMRT reionization effort aims to map out the large scale structure of the Universe during the epoch of reionization (EoR).  Removal of polarized Galactic emission is a difficult part of any 21 cm EoR program, and we present new upper limits to diffuse polarized foregrounds at 150 MHz.  We find no high significance evidence of polarized emission in our observed field at mid galactic latitude (J2000 08h26m+26). We find an upper limit on the 2-dimensional angular power spectrum of diffuse polarized foregrounds of $[l^2C_l/2\\pi]^{1/2}<3\\,$K in frequency bins of width $\\delta\\nu=1\\,$MHz at $300<l<1000$. The 3-dimensional power spectrum of polarized emission, which is most directly relevant to EoR observations, is $[k^3P_P(k)/2\\pi^2]^{1/2}<2\\,$K at $k_\\perp>0.03h\\,$Mpc$^{-1}$, $k<0.1h\\,$Mpc$^{-1}$.  This can be compared to the expected EoR signal in total intensity of $[k^3P(k)/2\\pi^2]^{1/2}\\sim 10\\,$mK. We find polarized structure is substantially weaker than suggested by extrapolation from higher frequency observations, so the new low upper limits reported here reduce the anticipated impact of these foregrounds on EoR experiments. We discuss  Faraday beam and depth depolarization models and compare predictions of these models to our data. We report on a new technique for polarization calibration using pulsars, as well as a new technique to remove broadband radio frequency interference.  Our data indicate that, on the edges of the main beam at GMRT, polarization squint creates $\\sim$ 3\\% leakage of unpolarized power into polarized maps at zero rotation measure. Ionospheric rotation was largely stable during these solar minimum night time observations. ",
        "introduction": "A current frontier in observational and theoretical cosmology is the search for and understanding of large scale structure during the epoch of reionization (EoR).  Detection of reionization at $z\\sim 10$ would allow study of the emergence of the first abundant luminous objects. Such observations holds the promise of examination of astrophysical and cosmological processes at epochs as early as a few hundred million years after the start of the expansion of the Universe.   The WMAP satellite has measured polarization in the Cosmic Microwave Background (CMB) at large angular scales.  This polarization is believed to arise from Thomson scattering of the CMB photons near the EoR \\citep{2008arXiv0803.0593N,2008arXiv0803.0547K}. The observed optical depth $\\tau\\sim 0.089\\pm 0.016$ corresponds to an instantaneous reionization redshift of $z_{\\rm reion}=10.8\\pm 1.4$.  One way to study the reionization transition in detail is by imaging redshifted 21cm emission. At redshifts above the EoR transition the gas is neutral and is predicted to glow with about  a 25 mK sky brightness temperature. After reionization is complete this glow is absent. At redshifts close to the transition a patchy sky is expected. Simulations \\citep{2008MNRAS.tmp...77I} suggest that with existing telescopes a measurement near 150 MHz may allow for statistical detection of $\\sim$ 20 Mpc patchiness in the neutral hydrogen.  This detection would pin down the reionization redshift and begin the process of more detailed study of the transition.  The 21 cm EoR signal, in the 10 milliKelvin range, is much smaller than the expected $\\sim 10$K patchiness due to Galactic synchrotron emission and extragalactic foregrounds. One of the most difficult challenges of 21 cm EoR astronomy is the removal of these bright foregrounds.  The proposed strategy makes use of the very different spatial and spectral character of the two components of sky structure. Since the received frequency of the 21 cm line emission encodes the distance of the source, the EoR signal consists of a series of source screens that are largely independent. For example, if one detects a patchy layer of cosmic material at 150 MHz due to the EoR transition, another patchy layer must be present at 151 MHz, but these two images come from distinct radial shells and the bright and dim patches in the two images are not expected to be aligned on the sky.  The characteristic size of the ionization structures near the epoch of ionized-cell overlap is approximately 10 $h^{-1}$ Mpc, which corresponds to 5 arc minutes in angular size, or about 1 MHz in radial distance for the redshifted 21cm line at $z \\sim 9$.  In contrast, the foreground emission is highly correlated from one frequency to the next, since most radio-bright structures emit with a very smooth synchrotron spectrum. To remove the foregrounds some type of frequency-differencing technique will be needed.  At frequencies above 300 MHz many radio sources with a synchrotron spectrum have been shown to be linearly polarized, and this raises a concern for the frequency differencing scheme.  In propagation through the interstellar medium the polarization rotates on the sky because of the Faraday effect.  In the polarization basis of the telescope this means that each linear polarization component can oscillate with frequency. If the oscillation period falls close to the frequency separation used for the foreground subtraction, incomplete removal of Galactic synchrotron emission may leave a residual that masks the EoR signal.  One way to avoid the Faraday rotation problem is to measure total intensity, described by the Stokes I parameter, using a instrument that rejects the linearly polarized Stokes components Q and U.  However, all real instruments have various sources of leakage between these components. To design these instruments and plan future EoR observations one needs information on the polarized sky brightness in the EoR band near 150 MHz.  Recent studies at 350 MHz \\citep{2006AN....327..487D} report that the polarized component has much more structure on arc minute angular scales than the total intensity.  Polarized structure amounting to several Kelvin has been reported.  Scaling using a synchrotron spectral index of 2.6 \\citep{2008arXiv0802.1525D}, the polarized structure could be tens of Kelvin in the EoR band.  This means that instrumental leakage from Q/U to I could severely limit the ability to remove Galactic emission.  This paper places new constraints on polarized sky brightness in the EoR range of frequencies near 150 MHz. We find the level of polarized sky brightness is well below that expected by extrapolation from higher frequencies.  We discuss possible mechanisms which would lead to this strong depolarization, and suggest future techniques to differentiate between them. ",
        "conclusions": "The GMRT EoR data has resulted in the strongest upper limits to polarized emission at 150 MHz to date.  The observed polarization is an order of magnitude smaller than had been expected from higher frequency observations.  We note that comparison was not made on the same fields. We estimated the expected beam and depth polarization effects, but these are model dependent and poorly constrained by the data. Within the uncertainties, either mechanism could account for the lack of observed polarization.  If beam depolarization were the cause, it would imply a significant small scale magnetic field with structure at $l>10,000$.  Others have reported depth depolarization at similar RM$\\lambda^2$, so depth depolarization may play a role here as well. When compared to the sensitivities required for EoR observations, the upper bounds we report are still two orders of magnitude larger than the sensitivities required.  The measured polarization leakage at GMRT is less than 10\\%,  so a modest leakage correction, accurate to 10\\%, should make EoR observations feasible.  The low polarized emission we report is promising for the search for Reionization."
    },
    "0807/0807.3053_arXiv.txt": {
        "abstract": "{The chronology of bulge and disk formation is a major unsolved issue in galaxy formation, which impacts on our global understanding of the Hubble sequence.} {We analyse colours of the nuclear regions of intermediate redshift disk galaxies, with the aim of obtaining empirical information of relative ages of bulges and disks at $0.1 < z < 1.3$.} {We work with an apparent-diameter limited parent sample of 248 galaxies from the HST Groth Strip Survey. We apply a conservative criterion to identify bulges and potential precursors of present-day bulges based on nuclear surface brightness excess above the exponential profile of the outer parts and select a sample of 56 galaxies with measurable bulges. We measure bulge colours on wedge profiles opening on the semi-minor axis least affected by dust in the disk, and compare them to disk, and global galaxy colours.} {For $60\\%$ of galaxies with bulges, the rest-frame nuclear colour distribution shows a red sequence that is well fit  by passive evolution models of various ages, while the remainder $40\\%$ scatters towards bluer colours. In contrast, galaxies without central brightness excess show typical colours of star forming population and lack a red sequence. We also see that, as in the local Universe, most of the minor axis colour profiles are negative (bluer outward), and fairly gentle, indicating that nuclear colours are not distinctly different from disk colours. This is corroborated when comparing nuclear, global and disk colours: these show strong correlations, for any value of the central brightness prominence of the bulge. No major differences are found between the low and high inclination samples, both for the bulge and non-bulge samples.}  % {Comparison with synthetic models of red sequence bulge colours suggests that such red bulges have stopped forming stars at an epoch earlier than $\\sim 1$ Gyr before the observation. The correlation between nuclear and disk colours and the small colour gradients hints at an intertwined star formation history for bulges and disks: probably, most of our red bulges formed in a process in which truncation of star formation in the bulge did not destroy the disk.} ",
        "introduction": "\\label{sec:introduction}  Three reasons make bulge population ages useful for galaxy formation studies.  In all inside-out scenarios, the bulge regions should contain information on the earliest phases of star formation in galaxies.  Also, when compared to ages of elliptical galaxies, bulge ages should shed light on whether bulges and ellipticals share a formation history as proposed in CDM-based semi-analytical models \\citep{Baugh96,Kauffmann96bul}.  Finally, bulge population ages, when compared to ages of the disks they live in, should provide clues on whether bulges formed before disks, again as assumed in current CDM-based models, or after disks, as proposed in secular evolution models \\citep{Kormendy04}.  The current consensus emerging from local Universe studies points to bulges being uniformly old ($\\sim$10 Gyr) in early- to intermediate-type galaxies.  \\citet{Peletier99} show that optical-NIR colours are indistinguishable from those of ellipticals in Coma, suggesting similar ages of at most 2 Gyr younger than Coma ellipticals.  Independently, from stellar colour-magnitude diagrams rather than integrated colours, for the Milky Way bulge \\citet{Zoccali03} infer homogeneous stellar ages similar to those of globular clusters, or $\\ga$10 Gyr.  Bulges of later-type spirals have signs of harboring younger populations, and ongoing star formation; \\citep[e.g.][]{Peletier99,Carollo01}.  Studies of local galaxies have also shown that redder bulges live in redder disks \\citep{Balcells94}, i.e., colour-wise, bulges are more similar to their parent disks than to each other \\citep{Peletier96}. This suggests common evolution, although not necessarily that one forms from the other.  Because the bulge ages just mentioned do not appear to be much older than the oldest stars in nearby mature disks  (about 10~Gyr for the MW disk), the relative chronology of disk and bulge formation appears difficult to determine from the fossil record in nearby galaxies.  Studying galaxies at cosmologically-significant look-back times offers a means to date galaxian components seen at epochs closer to that of their formation, and may provide better time resolution to determine age differences between bulges and ellipticals, and between bulges and disks. This approach has been followed by several groups in recent years, with conflicting conclusions. \\citet{Ellis01} and \\citet{Menanteau01} find a prevalence ($30 - 50\\%$) of young field spheroids at $z \\sim 1$ in the two HDFs.  In contrast, \\citet{Koo05} emphasized that $85\\%$ of luminous ($M_B < -19$) field bulges at redshifts $z \\sim 0.8$ in the Groth strip are nearly as red ($U - B \\sim 0.50$) as local E/SO's, a result that suggests both an old metal-rich dominant population and a 'drizzling' contribution of star formation to the global colours.  \\citet{Koo05HDF} compared bulge colours for field and cluster galaxies, concluding that bulge colours are identical for field and cluster galaxies. \\citet{MacArthur07}, who isolated bulges for a sample of 137 spiral galaxies within the redshift range $0.1<z<1.2$ in the GOODS fields, find that bulge colour is related to mass. The most massive bulges are as old and red as massive spheroids, which would be consistent with a mass assembly at high redshift. Smaller bulges, in contrast, have quite diverse star formation histories, with significant star formation at $z < 1$.  Additional, indirect evidence for late formation of galaxy bulges was presented by \\citet{Hammer01}, who argued that some luminous, blue, compact sources at $0.5<z<1$ may be starburst phases leading to the formation of bulges of massive spirals.  While cosmic variance likely affects the diversity of results presented above, most likely their differences are strongly affected by the way bulges are selected, the way bulge colours are measured, and the way bulge age information is inferred from the measured colours by the various studies. \\citet{Ellis01} and \\citet{Menanteau01} work with 'field spheroidals' selected on the basis of concentration and asymmetry indices, as well as from visual classification, while \\citet{Koo05} define their samples of 'photobulges' on the basis of modeling surface brightness distributions as the sum of \\r14\\ and exponential components, using the \\GimTwoD\\ software \\citep{Simard02}.  Both methods allow for pure ellipticals to make it into their samples, which may be seen as a draw-back.  Moreover, \\citet{Ellis01} and \\cite{Menanteau01} may include objects with purely exponential surface brightness profiles and exclude bulges from star-forming spiral galaxies.  Colour measurements are carried out on each image pixel by \\citeauthor{Menanteau01}, while a global photo-bulge colour is derived by \\citeauthor{Koo05} from constrained bulge-disk fits on two bands.  Finally, \\citet{Menanteau01} infer ages from the internal dispersion of galaxy colours, while \\citet{Koo05} compare integrated rest-frame colours, and colour-magnitude diagrams, to those from local E-S0's and distant clusters.  The above studies are, probably, complementary in the estimation of population ages of the precursors of galaxy bulges. In this paper we analyse colours of a sample of bulge galaxies, which has been defined in a companion paper \\citep[hereafter Paper~I]{Dominguez08I}, with the aim of addressing similar questions to the mentioned studies. We aim to provide complementary clues to the formation of galaxy bulges, given our different approach.  Our methods are detailed in Paper~I.  Specifically, we work with a diameter-limited sample, and we define a prominence index $\\eta$ to identify a subsample of galaxies with measurable bulges. Section~\\ref{sec:data} sumarizes the approach we use to define the bulge sample and the method to measure colours. In Sect.~\\ref{sec:comparison} we compare our bulge sample with that of \\citet{Koo05}. In the remainder of the sections we discuss the redshift distribution of the bulge sample; the shape of the minor-axis colour profiles; rest-frame colours for bulges and non-bulge galaxies; and the comparison of bulge colours to global and disk colours.  A cosmology with  $\\Omega_M = 0.3$, $\\Omega_\\Lambda = 0.7$, $H_0 = 70$ \\kmsMpc\\ is assumed throughout.  Magnitudes are expressed in the Vega system. ",
        "conclusions": "\\label{sec:conclusions}  \\begin{enumerate}  \\item $60\\%$ of galaxies with high central brightness excess above the outer exponential ($\\eta > 1$) concentrate strongly at 'very red' colours, $(U-B) > 0.25$, defining a red sequence, which fits well passive evolution models of different ages. The remainder $40\\%$ span over a bluer and wider colour range.  \\item Galaxies without central brightness excess ($\\eta < 1$) show bluer colours and lack a red sequence. Galaxies show a colour bimodality that is related to the central brigthness excess, in the sense that more concentrated galaxies are redder.  \\item Nuclear and global colours are very similar to each other: redder bulges are embedded in redder galaxies. Disk colours also correlate with bulge colours. As a result, the colour difference between bulges and disks is smaller than that between bulges of different galaxies.  Colour profiles are smooth, and do not show a discontinuity at the bulge-disk boundary.  \\item The correlations between nuclear and disk regions are fulfilled for all types of galaxies, independently of the central brightness excess and the colour. The colour profiles are equally smooth in the samples with ($\\eta > 1$) and without ($\\eta > 1$) central concentration.  \\item A major merger origin for bulges, as proposed by others, needs to be able to account for the similarities in bulge and disk colours.  Such similarities lend support to models with coeval growth of bulges and disks, and to bulge formation models that do not destroy a pre-existing disk.  \\end{enumerate}"
    },
    "0807/0807.1326_arXiv.txt": {
        "abstract": "As an alternative explanation of the dimming of distant supernovae it has recently been advocated that we live in a special place in the Universe near the centre of a large void described by a Lema\\^itre-Tolman-Bondi (LTB) metric. The Universe is no longer homogeneous and isotropic and the apparent late time acceleration is actually a consequence of spatial gradients in the metric. If we did not live close to the centre of the void, we would have observed a Cosmic Microwave Background (CMB) dipole much larger than that allowed by observations. Hence, until now it has been argued, for the model to be consistent with observations, that by coincidence we happen to live very close to the centre of the void or we are moving towards it. However, even if we are at the centre of the void, we can observe distant galaxy clusters, which are off-centre. In their frame of reference there should be a large CMB dipole, which manifests itself observationally for us as a kinematic Sunyaev-Zeldovich (kSZ) effect. kSZ observations give far stronger constraints on the LTB model compared to other observational probes such as Type Ia Supernovae, the CMB, and baryon acoustic oscillations. We show that current observations of only 9 clusters with large error bars already rule out LTB models with void sizes greater than $\\sim 1.5$ Gpc and a significant underdensity, and that near future kSZ surveys like the Atacama Cosmology Telescope (ACT), South Pole Telescope (SPT), APEX telescope, or the Planck satellite will be able to strongly rule out or confirm LTB models with giga parsec sized voids. On the other hand, if the LTB model is confirmed by observations, a kSZ survey gives a unique possibility of directly reconstructing the expansion rate and underdensity profile of the void. ",
        "introduction": "Distant supernovae appear dimmer than expected in a matter-dominated homogeneous and isotropic FRW universe. The currently favoured explanation of this dimming is the late time acceleration of the universe due to a mysterious energy component that acts like a repulsive force. The nature of the so-called Dark Energy responsible for the apparent acceleration is completely unknown.  Observations seem to suggest that it is similar to Einstein's cosmological constant, but there is inconclusive evidence. There has been a tremendous effort in the last few years to try to pin down deviations from a cosmological constant, e.g.~with deep galaxy catalogues like 2dFGRS~\\cite{2dFGRS} and SDSS~\\cite{SDSS}, and extensive supernovae surveys like ESSENCE~\\cite{ESSENCE}, SNLS~\\cite{SNLS}, and SDSS-SN~\\cite{SDSS-SN}, and many more are planned for the near future e.g.~DES~\\cite{DES}, PAU~\\cite{PAU,Benitez2008}, BOSS~\\cite{BOSS} and JDEM~\\cite{JDEM}.  In the meantime, our realisation that the universe around us is far from homogeneous and isotropic has triggered the study of alternatives to this mysterious energy.  Since the end of the nineties it has been suggested by various groups~\\cite{Hellaby:1998,Celerier:1999hp,Tomita:2000jj,Moffat:2005yx, Garfinkle:2006sb,Enqvist:2007vb,Mattsson:2007tj,Wiltshire:2007zj,GBH:2008} that an isotropic but inhomogeneous Lema\\^itre-Tolman-Bondi universe could also induce an apparent dimming of the light of distant supernovae, in this case due to local spatial gradients in the expansion rate and matter density, rather than due to late acceleration. It is just a matter of interpretation which mechanism is responsible for the dimming of the light we receive from those supernovae. Certainly the homogeneous and isotropic FRW model is more appealing from a philosophical point of view, but so was the static universe and we had to abandon it when the recession of galaxies was discovered at the beginning of last century.  There is nothing wrong or inconsistent with the possibility that we live close to the centre of a giga parsec scale void. Such a void may indeed have been observed as the CMB cold spot~\\cite{Cruz:2006sv,Cruz:2006fy,Cruz:2008sb} and smaller voids have been seen in the local galaxy distribution \\cite{Frith:2003,Granett:2008}. The size and depth of the distant observed voids, i.e.~$r_0 \\sim 2$ Gpc and $\\Omega_M \\sim 0.2$ within a flat Einstein-de Sitter universe, seems to be consistent with that in which we may happen to live~\\cite{Tully:2007tp}, and could account for the supernovae dimming, together with the observed baryon acoustic oscillations and CMB acoustic peaks, the age of the universe, local rate of expansion, etc.~\\cite{GBH:2008}.  Moreover, according to the theory of eternal inflation~\\cite{Linde:1993xx}, rare fluctuations at the Planck boundary may be responsible for the non-perturbative amplification of local inhomogeneities in the metric, which would look like local voids in the matter distribution~\\cite{Linde:1994gy}. In the eternal inflation approach one assumes to be a typical observer whose local patch comes directly from a rare fluctuation within an inflationary domain at the Planck scale. Moreover, since it is a rare fluctuation it should be highly spherically symmetric, and the theory predicts that we should live close to the centre of such a void~\\cite{Linde:1994gy}. The size and depth of those voids depends on the theory of inflation on very large scales and may be a probe (perhaps the only probe) of the global structure of the universe.  In fact, observations suggest that if there is such a large void, we should live close to the centre, otherwise our anisotropic position in the void would be seen as a large dipole in the CMB.  Of course, we do observe a dipole, but it is normally assumed to be due to the combined gravitational pull of the Virgo cluster, and the Shapley super cluster. There is always the possibility that we live off-centre and we are moving towards the centre of the void, so that the two effects are partially cancelled, giving rise to the observed dipole. However, such a coincidence could not happen for all galaxies in the void and, in general, clusters that are off-centred should see, in their frame of reference, a large CMB dipole. Such a dipole would manifest itself observationally for us as an apparent kinematic Sunyaev-Zeldovich effect for the given cluster.  It is the purpose of this paper to study the very strong constraints that present observations of the kSZ effect already put on the LTB void models, and predict how near future observations from kSZ surveys like ACT or SPT will strongly rule out (or confirm) LTB models with giga parsec sized voids.  In section~2 we describe the general LTB void models, giving the corresponding Einstein-Friedmann equations, as well as parameterisations of their solutions. In a subsection we describe the GBH constrained model, where we assume the Big Bang is homogenous, and thus the model depends on a single function, the inhomogeneous matter ratio $\\Omega_M(r)$. In section~3 we study the induced dipole for off-centred clusters and compute the size of the analogue velocity of those clusters depending on the parameters of the void model. In section~4 we analyse present observations and give constraints on the model from current observations. In section~5 we then explore the prospects that future experiments like ACT, SPT and Planck will provide for strongly constraining or even ruling out LTB models of the universe. Finally, in section~6 we give our conclusions. ",
        "conclusions": "As an alternative explanation of the dimming of distant supernovae it has recently been advocated that we live in a special place in the Universe near the centre of a large void. The universe is no longer homogeneous and isotropic and the apparent late time acceleration is actually a consequence of spatial gradients in the metric. If we did not live close to the centre of the void, we would have observed a CMB dipole much larger than that allowed by observations. Hence, until now it has been argued, for the model to be consistent with observations, that by coincidence we happen to live very close to the centre of the void. However, even if we are at the centre of the void, we can observe distant galaxy clusters, which are off-centre. In their frame of reference there should be a large CMB dipole, which manifests itself observationally for us as a kinematic Sunyaev-Zeldovich effect.  In this paper we have studied the induced dipole for off-centred clusters due to the different trajectories of photons from the last scattering surface, and computed the size of the corresponding apparent velocity of those clusters with respect to the rest frame of the CMB LSS, depending on the parameters of the void model. We then analysed the present observations of the kSZ effect in a handful of clusters and gave very strong constraints on the size of the void in LTB models. In fact, for our specific constrained-GBH model the bounds are in conflict with other constraints from supernovae, baryon acoustic oscillations and CMB, at least at the 3-$\\sigma$ level, see Fig.~\\ref{fig:otherdata}, and therefore we conclude that the constrained-GBH void model is practically ruled out if the current interpretation of kSZ observations are correct.  At present, kSZ data leave small voids ($r_0 \\lesssim 800$ Mpc) unconstrained, independently of the inner density contrast ($\\Omega_{\\rm in}$) used, simply because the radius is so small that the void does not impact the cluster with the lowest redshift in the sample, see Table~1. In order to put limits on small voids with large density contrasts, or sudden transitions, we would need to use local large scale structure data that can be checked for homogeneity, or alternatively low redshift supernovae.  Current kSZ observations are limited in numbers, and are still not precise enough to make a single positive detection of the peculiar velocity of a cluster at the 2-$\\sigma$ level, see Fig.~\\ref{fig:obs}.  Hence, one cannot rule out that they are plagued by large systematic and/or random errors. This is about to change in the coming years, and we have made predictions of how systematic near-future observations from kSZ surveys like ACT or SPT could strongly rule out all LTB models with giga parsec sized voids. In the case that the LTB models were confirmed, the average apparent velocity profile as a function of distance would give a direct handle on the density and expansion rate of the void. This unique relationship makes by far kSZ observations the most powerful data for constraining LTB models.  If the sensitivity of experiments like ACT, SPT and Planck are as planned, and the present systematic errors are under control, it is expected that large voids will definitely be ruled out at many sigma. If these experiments yield 10 (100) well observed clusters we could reasonably expect to rule out voids of 800 (500) Mpc radius at the 3-$\\sigma$ level.  We hope to come back to this severe challenge to LTB models and redo the analysis with the new data in the near future."
    },
    "0807/0807.1921_arXiv.txt": {
        "abstract": "The Guitar Nebula is an H$\\alpha$ nebula produced by the interaction of the relativistic wind of a very fast pulsar, PSR B2224+65, with the interstellar medium.  It consists of a ram-pressure confined bow shock near its head and a series of semi-circular bubbles further behind, the two largest of which form the body of the Guitar.  We present a scenario in which this peculiar morphology is due to instabilities in the back flow from the pulsar bow shock. From simulations, these back flows appear similar to jets and their kinetic energy is a large fraction of the total energy in the pulsar's relativistic wind.  We suggest that, like jets, these flows become unstable some distance down-stream, leading to rapid dissipation of the kinetic energy into heat, and the formation of an expanding bubble.  We show that in this scenario the sizes, velocities, and surface brightnesses of the bubbles depend mostly on observables, and that they match roughly what is seen for the Guitar.  Similar instabilities may account for features seen in other bow shocks. ",
        "introduction": "\\label{sec:intro}  Many pulsars travel at high speed, and the collision between their relativistic winds and the interstellar medium leads to the formation of bow shocks.  These shocks are observed most readily at X-ray and optical wavelengths: the shocked relativistic wind will emit mostly synchrotron radiation, while the shocked interstellar medium will emit -- if it is partially neutral -- copious H$\\alpha$ emission (for a review, see \\citealt{gaens06}).  Arguably the most spectacular bow shock is the Guitar Nebula, made by one of the fastest pulsars known, PSR~B2224+65 \\citep*{cordrm93}. This H$\\alpha$ nebula has, as the name implies, a guitar-like shape, with a bright head, a faint neck, and a body consisting of two larger bubbles (see \\Fref{guitar}).  \\cite{cordrm93} suggest this morphology might reflect variations in either the pulsar energy injection rate or the interstellar medium density.  In this {\\em Letter}, we investigate whether instead the peculiar morphology could be due to instabilities in the jet-like flow of pulsar effluvium away from the bow shock.  Fast back flows are a natural consequence of bow shocks: the pulsar wind is greatly heated at the shock, which, for the usual case where cooling is slow, leads to a high pressure that drives a flow in the only direction available, to the back.  From simulations (e.g., \\citealt{buccadz05}), the flows seem similar to jets, being well-collimated and fast, and seem to carry most of the pulsar wind energy.  Jet-like flows are indeed seen in X-ray observations, which also show that only a small fraction of the energy is radiated (for a review, \\citealt{kargp08}).  So far, the simulations have not extended far to the back, but if the back flow is similar to a jet, one might expect it to become unstable further downstream.  From simulations of jets (e.g., \\citealt{bodo+98}), this instability would lead to mixing with the ambient medium and rapid dissipation of the kinetic energy into heat. The simulations have not followed what happens beyond this initial mixing, but it seems plausible that the material would expand rapidly and drive a bubble.  If so, it might initially expand faster than the pulsar motion, and gain further energy from the jet.  With time, however, it will slow down, and once the pulsar has moved sufficiently far ahead, the jet will become so long that it becomes unstable before reaching the bubble, and a new bubble will be formed.  We suggest the body of the guitar is made up of two such bubbles, while another one has just started to form near the head.  In \\Sref{model}, we describe our model in more detail, and in \\Sref{guitar} we compare it with the properties of the Guitar Nebula, finding qualitative agreement.  In \\Sref{discussion}, we discuss implications as well as ways in which our model could be tested. ",
        "conclusions": "We found that we could roughly reproduce the Guitar Nebula assuming the jet-like back flow from the pulsar bow shock becomes unstable and dissipates rapidly, causing expanding bubbles.  If this were to happen generally, one might expect other sources with jets or bow shocks to show Guitar-like bubbles, yet none appear to be known.  For jet sources, this may not be surprising: many jets are denser than the medium they move through, and hence more stable, and disruptions that do occur may be difficult to distinguish from, e.g., changes in jet orientation.  For other bow shocks, the absence of bubbles may partly be a selection effect: most have much larger stand-off radii than the Guitar, and hence any bubbles would be at correspondingly larger distances, where they might be missed, especially as they would be fainter than the bow shock (or even invisible if the expansion velocity became too low or if radiative effects became important; both perhaps relevant especially for stellar wind bow shocks).  The one possible exception is PSR B0740$-$28, which has a H$\\alpha$ bow shock with a relatively small stand-off radius of $\\theta_0=1\\farcs0$ as well as ``shoulders'' further behind \\citep{jonesg02}.  If related to an instability, one infers $15\\lesssim\\lambda/\\theta_0\\lesssim60$, of the same order as we see for the Guitar.  It would be interesting to obtain deeper images further behind the bow shock.  In some pulsar bow shocks, the shocked pulsar wind is observed directly, by its synchrotron emission (for an overview, \\citealt{kargp08}).  For many, including the Guitar \\citep{huib07}, emission is seen only close to the pulsar, likely at the pulsar wind termination shock.  Some, however, have much longer tails.  The longest belongs to the ``Mouse,'' associated with PSR J1747$-$2958. This nebula, with $\\theta_0\\approx\\!0\\farcs75$, shows a bulbous structure $\\sim\\!1\\farcm5$ behind the pulsar (the Mouse's body), but also a smooth, straight tail of $12\\arcmin$, without a clear end \\citep{gaen+04}.  Scaling with the stand-off distances, one might identify the Mouse's body with the equivalent of the Guitar's head bubble.  The long tail has a size equivalent to the bottom bubble, but, apart from changes in polarisation, shows little structure \\citep{yuseb87}.  This would seem inconsistent with any bubbles being formed, and thus is puzzling in the context of our model.  For two other pulsar bow shocks with long tails, the observations match expectations better.  For PSR J1509$-$5850, with $\\theta_0\\approx0\\farcs5$, the X-ray tail extends for $\\gtrsim\\!5\\farcm6$ and shows clear structure, with a change in brightness at $1\\farcm3$, a kink at~$3\\arcmin$, and a bright radio spot coincident with its end point \\citep{huib07b,karg+08}.  Comparing with the large bubbles in the Guitar, the typical length scale of $\\sim\\!1\\farcm5$ for the knots and kinks is about a factor 3 larger, roughly consistent with the ratio of the stand-off distances.  For PSR B1929+10, with $\\theta_0\\simeq2\\farcs3$, the tail extends up to $10\\arcmin$ and again shows substantial structure, with brightenings at $\\sim\\!2\\arcmin$ and $\\sim\\!5\\arcmin$, the latter coincident with a radio feature \\citep{beck+06,misapg07}.  Again scaling with the stand-off radii, the $5\\arcmin$ feature could be similar to the head bubble in the Guitar.  Overall, we conclude that our model of instabilities in a bow shock back flow roughly reproduces observations of the Guitar Nebula, without the need to appeal to variations in the density of the ambient medium, nor to energy sources beyond what is expected to be carried by the back flow.  It also seems consistent with what is seen in other pulsar bow shocks.  The model could be tested further both with observations and simulations.  Observationally, one test would be to measure the expansion velocities in the Guitar bubbles, either by determining proper motions, or by spectroscopy (from the broad component of the H$\\alpha$ profile, as done for non-radiative shocks in supernova remnants; \\citealt{raym91}).  Given the observed H$\\alpha$ surface brightness, this would allow one to estimate the ambient density, which should be similar to that at the location of the bow shock in our model, but substantially lower if the bubbles reflect density variations \\citep{cordrm93,chatc04}.  Simulations of bow shocks that extend to larger scales might show whether instabilities in fact lead to bubbles or rather to more continuous structure, or whether perhaps the process is sufficiently stochastic that both can occur (possibly leading to a shape like the Guitar's neck).  If bubbles form, the simulations might also shed light on details of the morphology, such as the closed appearance at the back of the head and bottom bubbles."
    },
    "0807/0807.2325_arXiv.txt": {
        "abstract": "In this work, exact solutions of static and spherically symmetric space-times are analyzed in $f(R)$ modified theories of gravity coupled to nonlinear electrodynamics. Firstly, we restrict the metric fields to one degree of freedom, considering the specific case of $g_{tt}\\,g_{rr}=-1$. Using the dual $P$ formalism of nonlinear electrodynamics an exact general solution is deduced in terms of the structural function $H_P$. In particular, specific exact solutions to the gravitational field equations are found, confirming previous results and new pure electric field solutions are found. Secondly, motivated by the existence of regular electric fields at the center, and allowing for the case of $g_{tt}\\,g_{rr}\\neq -1$, new specific solutions are found. Finally, we outline alternative approaches by considering the specific case of constant curvature, followed by the analysis of a specific form for $f(R)$. ",
        "introduction": "A central theme in cosmology is the perplexing fact that the Universe is undergoing an accelerated expansion \\cite{expansion}. Several candidates, responsible for this expansion, have been proposed in the literature, in particular, dark energy models and modified gravity. Amongst the modified theories of gravity, models generalizing the Einstein-Hilbert action have been proposed, where a nonlinear function of the curvature scalar, $f(R)$, is introduced in the action. These modified theories of gravity seem to provide a natural gravitational alternative to dark energy, and in addition to allow for a unification of the early-time inflation \\cite{Starobinsky:1980te} and late-time cosmic speed-up \\cite{Carroll:2003wy,Fay:2007gg}. These models seem to explain the four cosmological phases \\cite{Nojiri:2006be}. They are also very useful in high energy physics, in explaining the hierarchy problem and the unification of GUTs with gravity \\cite{Cognola:2006eg}. The possibility that the galactic dynamics of massive test particles may be understood without the need for dark matter was also considered in the framework of $f(R)$ gravity models \\cite{darkmatter}. One may also generalize the action by considering an explicit coupling between an arbitrary function of the scalar curvature, $R$, and the Lagrangian density of matter \\cite{Nojiri:2004bi}. Note that these couplings imply the violation of the equivalence principle \\cite{Olmo:2006zu}, which is highly constrained by solar system tests.  A fundamental issue extensively addressed in the literature is the viability of the proposed $f(R)$ models \\cite{viablemodels,Hu:2007nk,Sokolowski:2007pk}. In this context, it has been argued that most $f(R)$ models proposed so far in the metric formalism violate weak field solar system constraints \\cite{solartests}, although viable models do exist \\cite{Hu:2007nk,solartests2,Sawicki:2007tf,Amendola:2007nt}. The issue of stability \\cite{Faraoni:2006sy} also plays an important role for the viability of cosmological solutions \\cite{Nojiri:2006ri,Sokolowski:2007pk,Amendola:2007nt,Boehmer:2007tr, deSitter}. In the context of cosmological structure formation observations \\cite{structureform}, it has been argued that the inclusion of inhomogeneities is necessary to distinguish between dark energy models and modified theories of gravity, and therefore, the evolution of density perturbations and the study of perturbation theory in $f(R)$ gravity is of considerable importance \\cite{Boehmer:2007tr,Tsuj,Uddin:2007gj,Bazeia:2007jj}.  A great deal of attention has also been paid to the issue of static and spherically symmetric solutions of the gravitational field equations in $f(R)$ gravity \\cite{Clifton:2005aj,SSSsol,Multamaki:2006zb}. Solutions in the presence of a perfect fluid were also analyzed \\cite{Multamaki:2006ym}, where it was shown that the pressure and energy density profiles do not uniquely determine $f(R)$. In addition to this, it was found that matching the exterior Schwarzschild-de Sitter metric to the interior metric leads to additional constraints that severely limit the allowed fluid configurations. An interesting approach in searching for exact spherically symmetric solutions in $f(R)$ theories of gravity was explored in \\cite{Capozziello:2007wc}, via the Noether Symmetry Approach, and a general analytic procedure was developed to deal with the Newtonian limit of $f(R)$ gravity in \\cite{Capozziello:2007ms}. Analytical and numerical solutions of the gravitational field equations for stellar configurations in $f(R)$ gravity theories were also presented \\cite{Kainulainen:2007bt,Henttunen:2007bz,Multamaki:2007jk}, and the generalized Tolman-Oppenheimer-Volkov equations for these theories were derived \\cite{Kainulainen:2007bt}.  In the context of $f(R)$ modified theories of gravity, it was recently shown that power-law inflation and late-time cosmic accelerated expansion can be explained by a modified $f(R)$-Maxwell theory \\cite{Bamba:2008ja}, due to breaking the conformal invariance of the electromagnetic field through a non-minimal gravitational coupling. It is interesting to note that such a coupling may generate large-scale magnetic fields. Motivated by these ideas, we consider in this work $f(R)$ gravity coupled to nonlinear electrodynamics, and endeavor to search for exact solutions in a static and spherically symmetric set-up. In contrast to a non-minimal gravitational coupling, here conformal invariance is not broken.  In the context of nonlinear electrodynamics, a specific model was proposed by Born and Infeld in 1934 \\cite{BI} founded on a \\emph{principle of finiteness}, namely, that a satisfactory theory should avoid physical quantities to become infinite. The Born-Infeld model was inspired mainly to remedy the fact that the standard picture of a point particle possesses an infinite self-energy, and consisted on placing an upper limit on the electric field strength and considering a finite electron radius. Later, Pleba\\'{n}ski explored and presented other examples of nonlinear electrodynamic Lagrangians \\cite{Pleb}, and showed that the Born-Infeld theory satisfies physically acceptable requirements. Furthermore, nonlinear electrodynamics have recently been revived, mainly because these theories appear as effective theories at different levels of string/M-theory, in particular, in D$p-$branes and supersymmetric extensions, and non-Abelian generalizations (see Ref.~\\cite{Witten} for a review).  Much interest in nonlinear electrodynamic theories has also been aroused in applications to cosmological models \\cite{cosmoNLED}, in particular, in explaining the inflationary epoch and the late-time accelerated expansion of the universe \\cite{Novello}. It is interesting to note that the first {\\it exact} regular black hole solution in general relativity was found within nonlinear electrodynamics \\cite{Garcia,Garcia2}, where the source is a nonlinear electrodynamic field satisfying the weak energy condition, and recovering the Maxwell theory in the weak field limit. In fact, general relativistic static and spherically symmetric space-times coupled to nonlinear electrodynamics have been extensively analyzed in the literature: regular magnetic black holes and monopoles \\cite{Bronnikov1}; regular electrically charged structures, possessing a regular de Sitter center \\cite{Dymnikova}; traversable wormholes \\cite{Arellano1} and gravastar solutions \\cite{Arellano3}.  Thus, as mentioned above, motivated by recent work on a non-minimal Maxwell-$f(R)$ gravity model \\cite{Bamba:2008ja}, in this paper $f(R)$ modified theories of gravity coupled to nonlinear electrodynamics are explored, in the context of static and spherically symmetric space-times. This paper is outlined in the following manner: In section \\ref{sec2:action}, the action of $f(R)$ gravity coupled to nonlinear electrodynamics is introduced, and the respective gravitational field equations and electromagnetic equations are presented. In section \\ref{sec3:simple}, we restrict the metric fields to one degree of freedom, by considering the specific case of $g_{tt}=-g_{rr}^{-1}$, and using the dual $P$ formalism of nonlinear electrodynamics, we present exact solutions in terms of the structural function $H_P$. Subsequently, in section \\ref{sec4:2dof} we investigate the situation where the two metric fields are related via a power law in $r$, introducing additional parameters, and derive new specific solutions. In section \\ref{sec5:alternative}, we present alternative methods of finding exact solutions, first by considering the specific case of constant curvature, then by choosing a form for the $f(R)$, before we conclude in section \\ref{sec6:conclusion}. ",
        "conclusions": "\\label{sec6:conclusion}  The issue of exact static and spherically symmetric solutions in $f(R)$ modified theories of gravity is an important theme, mainly due to the analysis of weak field solar system constraints, and the generalization of exact general relativistic solutions to $f(R)$ gravity. In this work we have analyzed exact solutions of static and spherically symmetric space-times in $f(R)$ modified theories of gravity coupled to nonlinear electrodynamics. Firstly, the metric fields were restricted to one degree of freedom, by considering the specific case of $g_{tt}=-g_{rr}^{-1}$. Using the dual $P$ formalism of nonlinear electrodynamics, an exact general solution was found in terms of the structural function $H_P$. In particular, exact solutions to the gravitational field equations were found, confirming previous results and new pure electric field solutions were deduced. Secondly, by allowing two degrees of freedom for the metric fields, and motivated by the existence of regular electric fields at the center, new solutions were found. Finally, we have also briefly considered alternative approaches by analyzing the specific case of constant curvature and secondly, by considering a specific form for $f(R)$."
    },
    "0807/0807.0266_arXiv.txt": {
        "abstract": "We present a time-dependent radiative model for the atmosphere of the transiting planets that take into account the eccentricity of their orbit. We investigate the temporal temperature and flux variations due to the planet-star distance variability. We will also discuss observational aspects with {\\it Spitzer} measurements. ",
        "introduction": "{\\underline{\\it Why taking into account the eccentricity}}? The major theoretical models of atmospheric structure of exoplanets consider the planet only into two conditions: \\begin{itemize} \\item{} Either the flux received by the planet coming from the star is averaged over an hemisphere of the planet in a static case ; \\item{} or some sort of redistribution day to night of the incoming flux is included. \\end{itemize} In none of this cases the fact that the flux actually varies with time is taken into account. Even in the case of the less eccentric planet, WASP-10 b, with a low eccentricity of 0.057 (\\cite[Christian et al.  2008]{Christian08}) this represents a variation $\\sim$~25\\% in the received flux between periastron and apoastron   We focused our study to two planets: HD~17156b, the most eccentric transiting planet, and even if HD~80606b is not transiting, we choose this planet with the largest eccentricity to show the extreme case. Tab.~1 summarize the parameters of these two systems.    \\begin{table} \\begin{center} \\caption{Planetary and stellar parameters} \\label{tab1} {\\scriptsize \\begin{tabular}{l|l||c||c} \\hline &Parameters&HD 17156b &HD80606~b\\\\ \\hline Orbit% & $e$   & 0.67  & 0.93 \\\\ & $a$ [AU]  & 0.159 & 0.432 \\\\ & $d_{\\rm min}$ [AU]  & 0.052 &0.030 \\\\ & $d_{\\rm max}$ [AU]  &0.266  & 0.834\\\\ & $P$ [days]  &21.22 & 111.78 \\\\ & $\\omega$ [$^\\circ$]& 121 &300\\\\ \\hline Planet % & Radius % [\\rj] & 0.96 & 1.1$^{\\rm a}$\\\\   & Mass % [\\mj] & 3.11 & 5$^{\\rm b}$ \\\\ & $P_{\\rm spin}$ [days] &    3.8 &1.93\\\\ \\hline Star % & Type  & G0V   &  G5\\\\ & Radius [\\rsol]&1.47 &1.05  \\\\ &[Fe/H] & 0.24  & 0.43\\\\ \\hline  Refs.&        &  \\cite{Gillon07}& \\cite{LL08}\\\\ \\hline \\end{tabular}\\label{tab:param}   } \\end{center} \\vspace{1mm} \\scriptsize{ {\\it Notes:}\\\\ $^{\\rm a}$ This planet is not transiting. The value of the radius has been predicted by (4) ; \\\\ $^{\\rm b}$We adopted this value of the mass (inclination unknown). }  \\end{table} ",
        "conclusions": "We presented a still under development time-dependent radiative model for the atmosphere of the transiting planets. This model takes into account the eccentricity of their orbit and in its next version will be able to completely map the planet, including the rotation effects on the insolation.   The preliminary results shows already that an eccentric planet in under heating conditions is very different that if it were sitting at its semi major axis, even for a low eccentricity even if several steps need to be completed before reaching a fully self-consistent model."
    },
    "0807/0807.2780_arXiv.txt": {
        "abstract": "We report the discovery that substructures/subhaloes of a galaxy-size halo tend to fall in together in groups  in cosmological simulations, something that may explain the oddity of the MW satellite distribution.  The original clustering at the time of infall is still discernible in the angular momenta of the subhaloes even for events which took place up to eight Gyrs ago, $z \\sim 1$. This phenomenon appears to be rather common since at least $1/3$ of the present-day subhaloes have fallen in groups in our simulations. Hence, this may well explain the Lynden-Bell \\& Lynden-Bell ghostly streams. We have also found that the probability of building up a flattened distribution similar to the MW satellites is as high as $\\sim 80\\%$ if the MW satellites were from only one group and $\\sim 20\\%$ when five groups are involved.  Therefore, we conclude that the `peculiar' distribution of satellites around the MW can be expected with the CDM structure formation theory.  This non-random assignment of satellites to subhaloes implies an environmental dependence on whether these low-mass objects are able to form stars, possibly related to the nature of reionization in the early Universe. ",
        "introduction": "The discrepancy in numbers between the substructures/subhaloes resolved in a galaxy-size cold dark matter (CDM) halo and the satellites around the Milky Way (MW) has been a long standing issue for the `concordance' CDM structure formation theory. It implies a non-trivial mapping between the luminous satellites and the dark matter subhaloes at the (sub) galactic scales \\cite[(Zentner \\etal\\, 2005;]{zentner05} \\cite[Libeskind \\etal\\, 2007)]{libeskind07} through the astrophysics processes with baryons or the theory needs a major modification at a fundamental level (see e.g. \\cite[Kamionkowski \\& Liddle 2000)]{kl00}.  In the past ten years new attention has been drawn to the dynamical properties of the MW satellites. Starting with \\cite{lyndenbell95}, the existence of ghostly streams of satellites (dwarf galaxies and globular clusters) was postulated. These objects would share similar energies and angular momenta producing a strong alignment along great circles on the sky. Recently, the anisotropic distribution of satellites around the MW has been argued to be a problem for the CDM theory \\cite[(Kroupa, Theis \\& Boily, 2005;]{kroupa05} \\cite[Metz, Kroupa \\& Jerjen, 2007)]{metz07}.  The MW has approximately 20 satellites forming a disk-like structure while the simulated dark matter subhaloes usually distribute almost isotropically.  Here we report our findings of subhaloes falling in groups in dark matter simulations and its application on explaining the oddities of the dynamical properties of MW satellites, namely, the Lynden-Bell \\& Lynden-Bell ghostly steams and the great MW satellites disk.  Researches in the past showed that clusters of galaxies are built of galaxies coming in groups \\cite[(Knebe \\etal\\, 2004)]{knebe04}, but it was not clear whether a similar picture also applies at the (sub)galactic scale. We refer readers to a more detailed  description of our analyses and discussions on the group infall and its link to the environment in \\cite{lihelmi08}. ",
        "conclusions": "We have revisited the issue of the peculiar distribution and properties of the MW satellites and their link to the dark matter subhaloes. In particular, we have focused on the infall of substructures on to a Milky-Way like dark matter halo in a $\\Lambda$CDM cosmogony utilising a series of high-resolution dark-matter simulations.  We have found evidence of group infall on to the MW-like halo, which may explain the ghostly streams proposed by \\cite{lyndenbell95}.  We have also explored how this planar configuration may be obtained as a result of the infall of satellites in groups. The observed correlation in the angular momentum orientation of subhaloes naturally gives rise to disk-like configurations. For example, we find that if all subhaloes are accreted from just one group, a disk-like distribution is essentially unavoidable ($\\sim 80\\%$ probability), while for accretion from just two groups, the likelihood of obtaining a distribution as planar as observed is 40\\%.  Therefore the disky configuration of satellites is consistent with CDM if most satellites have their origin in a few groups.  Note that in our studies, we do not need to invoke the baryon-related physics to account for the dynamical properties of the MW satellites.  Thus both the `ghostly streams' and the `planar configuration' are manifestations of the same phenomenon: the hierarchical growth of structure down to the smallest galactic scales.  One of the possible implications of the reality of the ghostly streams is that its member galaxies formed and evolved in a similar environment before falling into the MW potential. This would have implications on the (oldest) stellar populations of these objects, such as for example, sharing a common metallicity floor \\cite[(Helmi \\etal\\, 2006)]{helmi06}. On the other hand, this implies that there should be groups that have failed to host any luminous satellites. This would hint at a strong dependence on environment on the ability of a subhalo to retain gas \\cite[(Scannapieco \\etal\\, 2001)]{scannapieco01}, or be shielded from re-ionization by nearby sources \\cite[(Mashchenko, Carignan \\& Bouchard 2004;]{mcb04} \\cite[Weinmann \\etal\\, 2007)]{weinmann07}.  Recent proper motion measurements of the Large and Small Magellanic clouds by \\cite{smc-mu}, as well as the simulations by \\cite{bc05} suggest that these systems may have become bound to each other only recently.  This would be fairly plausible in the context of our results.  The Clouds may well have been part of a recently accreted group and it may not even be necessary for them to ever have been a binary system."
    },
    "0807/0807.0971_arXiv.txt": {
        "abstract": "{ We find a new analytical solution for the chemical evolution equations, taking into account the delayed contribution of all low and intermediate mass stars (LIMS) as one representative star that enriches the interstellar medium. This solution is built only for star formation rate proportional to the gas mass in a closed box model. We obtain increasing C/O and N/O ratios with increasing O/H, behavior impossible to match with the Instantaneous Recycling Approximation (IRA). Our results, obtained by two analytical equations, are very similar to those found by numerical models that consider the lifetimes of each star. This delayed model reproduces successfully the evolution of C/O$-$O/H and $Y-O$ relations in the solar vicinity. This analytical approximation is a useful tool to study the chemical evolution of elements produced by LIMS when a galactic chemical evolutionary code is not available. }  \\resumen{ Encontramos una nueva soluci\\'on anal\\'{\\i}tica para las ecuaciones de evoluci\\'on qu\\'{\\i}mica tomando en cuenta la contribuci\\'on retrasada de todas las estrellas de $m < 8$ \\msun (LIMS) como una estrella representativa que enriquece al medio interestelar. Esta soluci\\'on es construida para tasa de formaci\\'on estelar proporcional a la masa de gas en un modelo de caja cerrada. Obtenemos incrementos en C/O y N/O cuando O/H aumenta, comportamiento imposible de igualar con IRA. Nuestros resultados, obtenidos por dos ecuaciones anal\\'{\\i}ticas, son muy similares a aquellos encontrados por modelos num\\'ericos que consideran el tiempo de vida de cada estrella. Este modelo retrasado reproduce la evoluci\\'on de C/O-O/H y $Y-O$ en la vecindad solar. Esta aproximaci\\'on anal\\'{\\i}tica es una herramienta util para estudiar la evoluci\\'on de elementos producidos por las LIMS cuando no se dispone de un c\\'odigo de evoluci\\'on qu\\'{\\i}mica. }  \\addkeyword{galaxies: abundances} \\addkeyword{galaxies: evolution} \\addkeyword{ISM: abundances}  \\begin{document} ",
        "introduction": "\\label{sec:intro}  Chemical evolution models are used to describe the temporal variation of the gas mass and the abundances of the different chemical elements that are present in the gas. Their importance relies on the fact that the chemical  history of the studied object (i.e. interstellar medium  in  galaxies and intergalactic medium) can be inferred. Moreover it is possible to get chemical information about stellar population properties and characteristics of the galaxies that we observe in the low redshift Universe.  The chemical evolution equations, shown by Tinsley (1974) and corrected later by Maeder (1992), are relatively complex and can be solved through numerical models. There are some analytical approximations that have the advantage of predicting the general behavior of chemical elements in a quick and easy way but some precision may be lost.  One of the most well known analytical approximations is the Instantaneous Recycling Approximation (IRA) (Talbot \\& Arnett 1971) where the star lifetimes are negligible compared with the age of the galaxies. This approximation has been widely used because it simplifies the solution to the chemical evolution equations, however, massive star (MS) lifetimes are on the order of $10^{6} - 10^{7}$ years while the lifetimes for the low and intermediate mass stars (LIMS) are on the order of $10^{8} - 10^{10}$ years comparable to the lifetime of a galaxy. Hence, IRA provides only a very rough approximation for elements produced by LIMS.   Serrano \\& Peimbert (1983) proposed an analytic approximation related to the delays in chemical enrichment in N/O-O/H relation assuming N and O yields increase with $Z$. They present closed and open models (with and without gas flows, respectively) that takes into account the delay on nitrogen production due to LIMS concluding that nitrogen must be an element mainly secondary.    Later on, Pagel (1989) presented another approximation for  open models introducing an arbitrary time delay in order to study the chemical evolution of element produced by SNIa and LIMS, such as Fe and Ba (through s-process). This time delay term makes the stars release the processed material to the interstellar medium (ISM) at a single time after the star formation. After this single time delay, the contribution of all type of star is instantaneous, like IRA.  The objective of this work is to present an alternative analytical solution to the chemical evolution equations that considers the LIMS lifetimes as only one group with  delay times during the whole evolution. Those delay times are different for each chemical element and are computed based on the characteristics of the stellar yields and the stellar population. The delayed contribution to the chemical enrichment of elements produced by this type of star at different times should give results with a precision intermediate between the results obtained by using IRA and the ones obtained by numerical codes. As an application, this work aims to reproduce the C/O \\textit{vs} O/H and $Y(O)$ histories indicated by the stars and HII regions at the solar vicinity.  This approximation is a simple tool for theoretical and observational astronomers who need to include chemical aspects in their computations or to interpret observational data when they do not have access to a numerical code of chemical evolution of galaxies. Our equation for the mass abundance of element $i$, $X_i(t)$, would replace the popular equation of $Z(t)$ obtained assuming IRA in the closed box regimen, erroneously applied for elements produced by LIMS (e.g. He, C, N). ",
        "conclusions": "\\begin{itemize}  \\item{We have found analytical equations for chemical evolution in the case of a closed box model and SFR proportional to the gas mass where the delayed enrichment by LIMS is represented by a single type of star. }  \\item{The delay of LIMS with respect to the galactic enrichment for He, C and N produces an increase on C/O and N/O with O/H and in Y with O in good agreement with the results obtained by numerical models. With IRA, the C/O and N/O values are constant with O/H in disagreement with observed data and with  model results that take into account all star lifetimes.}  \\item{For $\\mu = 0.1$ and  $Z_{pop}=0.02$, $Y(O)$, C/O, and N/O values show artificial secondary raises  due to the O dilution produced by LIMS and not because of the increase of He, C and N produced by the LIMS.}  \\item{The delayed approximation was probed successfully in the solar vicinity reproducing the main trends of C/O-O/H relation shown by dwarf stars in agreement with results obtained with a numerical model that considers all star lifetimes. That relation cannot be reproduced at all by the instantaneous recycling approximation.}  \\item{ The analytical equation (eq. 5) obtained by our approximation is a useful tool to know the chemical evolution of those elements produced by LIMS when no galactic chemical evolutionary code is available. }    \\end{itemize}   {\\it Part of this work was submitted in the Physics Undergraduate Program at the Universidad Nacional Aut\\'onoma de M\\'exico.}"
    },
    "0807/0807.2145_arXiv.txt": {
        "abstract": "We present a source and lens reconstruction for the optical Einstein ring gravitational lens system RXS J1131-1231. We resolve detail in the source, which is the host galaxy of a $z=0.658$ quasar, down to a resolution of 0.045 arc seconds (this is the size of the smallest conclusively resolved structures, rather than the pixel scale), using a Bayesian technique with a realistic model for the prior information. The source reconstruction reveals a substantial amount of complex structure in the host galaxy, which is $\\sim$ 8 kpc in extent and contains several bright compact substructures, with the quasar source residing in one of these bright substructures. Additionally, we recover the mass distribution of the lensing galaxy, assuming a simply-parameterised model, using information from both the quasar images and the extended images. This allows a direct comparison of the amount of information about the lens that is provided by the quasar images in comparison to the extended images. In this system, we find that the extended images provide significantly more information about the lens than the quasar images alone, especially if we do not include prior constraints on the central position of the lens. ",
        "introduction": "Gravitational lensing can be used as a powerful astrophysical tool for probing the density profiles of galaxies, and is one of the few ways in which dark matter can be detected \\citep[e.g.][]{2005MNRAS.363.1136K}. In addition, it often magnifies source objects by one to two orders of magnitude. This allows us to use the intervening gravitational lens as a kind of natural telescope, magnifying the source so that we can observe more detail than we would have been able to without the lens. This extra magnification provided by lensing has been very beneficial to studies of star formation and galaxy morphology at high redshift. Regions of the galaxy size and luminosity distribution that are inaccessible in unlensed observations are made (more) visible by lensing \\citep[e.g.][]{2000ApJ...528...96P, wayth, 2006ApJ...651....8B, 2007ApJ...671.1196M, 2008arXiv0804.4002D}. The properties of the lens galaxies (typically elliptical galaxies) can also be inferred from their lensing effect \\citep[e.g.][]{2006ApJ...649..599K, 2008arXiv0806.1056T}. Of course, gravitational lensing distorts the image of the source, as well as magnifying it. Thus, techniques have been developed that aim to infer the mass profile of the lens galaxy and the surface brightness profile of the source, given observed images \\citep[e.g.][]{2003ApJ...590..673W, 2006ApJ...637..608B}.  The aim of this paper is to carry out this process with the recently discovered gravitationally lensed quasar/host galaxy system RXS J1131-1231 \\citep{2003A&A...406L..43S}. This system consists of a quadruply imaged quasar at redshift $z=0.658$ lensed by a galaxy at $z=0.295$. At the time of its discovery, it was the closest known lensed quasar, with some evidence for an extended optical Einstein ring - the image of the quasar host galaxy. Initial simple modelling suggested that the quasar source was magnified by a factor of $\\sim$ 50. Thus, subsequent observations with the ACS aboard the Hubble Space Telescope \\citep[][hereafter C06]{2006A&A...451..865C} allow the recovery of the  morphology of the quasar's host galaxy down to a resolution of about 0.01 arc seconds - at least in principle, for the parts of the source nearest the caustic. Indeed, C06 presented a wide array of results based on HST observations (at 555nm and 814nm with ACS, and 1600nm with NICMOS), including a detailed reconstruction of the extended source.  The source reconstruction method used by C06 is based on lensing the image back to a pixellated grid in the source plane, setting the source surface brightnesses to equal the image surface brightness, and using a decision rule (in this case, the median) to decide on the value of a source pixel whenever two or more image pixels disagree about the value of the same source pixel. If the point spread function (PSF) is small or the image has been deconvolved (in C06, the deconvolution step was neglected for the purposes of the extended source reconstruction) and the lens model is correct, this method can expect to be quite accurate. However, in principle, the uncertainty in the lens model parameters and the deconvolution step should always be taken into account. In this paper, we focus our attention on the 555nm ACS image (the drizzled images, as reduced by C06, were provided to us), and the process by which we reconstruct the original, unlensed source from it. Any differences between our reconstruction and the C06 one can be attributed to the following advantages of our approach: PSF deconvolution, averaging over the lens parameter uncertainties, simultaneous fitting of all parameters, and the prior information that Bayesian methods are capable of taking into account: in the case of our model, that is the knowledge that the majority of pixels in an astrophysical sources should be dark \\citep{2006ApJ...637..608B}. The 555nm image is also of particular interest because its rest wavelength (335nm) probes regions of recent star formation in a galaxy with an AGN.  In the case of the Einstein Ring 0047-2808 \\citep{2006ApJ...651....8B}, our method was able to resolve structure down to scales of $\\sim$ 0.01 arcsec, a factor of five smaller than that obtainable in an unlensed observation with the Hubble Space Telescope and about double the resolution obtained by \\citet{2005ApJ...623...31D} using adaptive pixellation and least squares {\\it applied to exactly the same data}. This was possible because we used a prior distribution over possible sources that is more realistic as a model of our knowledge of an unknown astrophysical source, that is, we took into account the fact that it should be a positive structure against a dark background, a fact many methods (such as least squares and some popular regularisation formulas) implicitly ignore \\citep{2006ApJ...637..608B}. These differences between methods are likely to be most significant when data are sparse or noisy, and all methods tend to agree as the data quality increases and we approach the regime where the observed image uniquely determines the correct reconstruction. ",
        "conclusions": "In this paper, we have presented a detailed gravitational lens reconstruction of the optical extended source in the Einstein Ring RXS J1131-1231. The source is a medium sized galaxy ($\\sim$ 8 kpc in visible extent) with several compact bright emission regions. The substructures we found are in general agreement with those found by C06 in terms of their position, but we have shown that they are brighter and more compact than was previously determined. In addition, our reconstruction provides a clearer view of the substructures, including near the central regions of the source. The quasar resides in a bright emission region with an extent of about $\\sim$ 0.15 arcseconds or 1 kpc. It should be noted that the wavelength of the observations in the rest frame is 335 nm, so this reconstruction traces regions of recent star formation in the source galaxy.  We have also directly compared point images vs extended images with regard to how well each is able to constrain the lens model. We found that there is a significant gain to be made in taking into account all of the information from the extended images. It has been suggested that this is not true in general \\citep{2007arXiv0710.3159F}, although it really depends on the resolution and number of extended images, which in this case is high. Certainly, in using both, there is nothing to lose but CPU cycles. This system has the potential to become one of the most well-constrained gravitational lenses, with multiple images of the extended ring, quasar image positions and flux ratios in multiple bands, and time delay measurements available \\citep{timedelay, keeton}. Hence, it should be possible to carry out a detailed kiloarametric study of its mass profile to shed some light on the dark matter halo of the lens galaxy.  This paper was based on a single image of this system, the 555nm ACS image. Other HST images at different wavelengths (814nm, 1.6$\\mu$m) are available (C06) and can further constrain the lens model. Simultaneous multi-wavelength reconstructions are now becoming routine \\citep[e.g.][]{2007ApJ...671.1196M}. However, all of the structures in these images are in the same locations, and so a multi-wavelength reconstruction would not produce significantly different conclusions to those reached here. C06 note that in the near infrared image, the compact bright images are less pronounced compared to the diffuse background, which is what would be expected if the substructures are regions dense in hot young stars.  This study has relied on a number of common assumptions that future research will seek to relax. Extending lens reconstruction techniques to incorporate kiloparametric models of the source and the lens simultaneously is an ambitious task, but some steps are already being taken in that direction \\citep{2008arXiv0804.2827S}. Flexible lens modelling plus information from time delay measurements and other sources would be extremely valuable for studies of galaxy dark matter haloes. Also, explicit modelling of dust absorption by the lens galaxy is proving to be an important ingredient in the inversion of Einstein Rings and would be an essential part of future work on this system."
    },
    "0807/0807.0189_arXiv.txt": {
        "abstract": "We have observed a region of nebulosity first identified as starlight scattered by interstellar dust by Sandage (1976) using the GALEX ultraviolet imaging telescope. Apart from airglow and zodiacal emission, we have found a diffuse UV background of between 500 and 800 \\phunit\\ in both the \\galex\\ FUV (1350 -- 1750 \\AA) and NUV (1750 -- 2850 \\AA) bands. Of this emission, up to 250 \\phunit\\ is due to \\htwo\\ fluorescent emission in the FUV band. The remainder is consistent with scattering from interstellar dust with forward scattering grains of albedo about 0.4. These are the highest spatial resolution observations of the diffuse UV background to date and show an intrinsic scatter beyond that expected from instrumental noise alone. Further modeling is required to understand the nature of this scatter and its implications for the ISM. ",
        "introduction": "Ever since the first observations of diffuse ultraviolet radiation by \\citet{Hay69} and \\citet{Lillie}, there has been an effort to understand its distribution and its origin. Unfortunately, because of the difficulty of the observations and the faintness of the background, many of the early observations were conspicuous more by their disagreements than by the light they shed on the topic. The state of the observations and theories before 1990 have been reviewed by \\citet{Bowyer91} and \\citet{RCH91}.  There has been significant progress in more recent years, particularly in the far ultraviolet ($\\lambda < $ 1200 \\AA) where \\citet{JM99} and \\citet{JM04} have used spectroscopic data from the \\voyager\\ and \\fuse\\ ({\\it Far Ultraviolet Spectroscopic Explorer}) spacecraft, respectively, to trace the radiation field over many different locations in the sky. There have also been a number of observations at longer wavelengths, most recently by the SPEAR instrument \\citep[][and references therein]{Ryu}, but no systematic study of the UV background. The {\\it Galaxy Evolution Explorer} (\\galex) offers us the opportunity to extend coverage of the diffuse background to a significant fraction of the sky with a sensitivity of better than 100 \\phunit. In this work, we will report on one such observation: that of the nebulosity observed near M82 by \\citet{Sandage76}, as a template for our further work with a much larger data set.  This cloud is at a high Galactic latitude (38\\degr) with few nearby stars and was identified by \\citet{Sandage76} as a canonical high latitude dusty cloud illuminated by the Galactic interstellar radiation field (ISRF). Our expectation was that we would be able to differentiate between starlight scattered from the Galactic cloud and extragalactic light which would be shadowed by the foreground cloud. The observations presented here are the first to probe the diffuse UV background at a spatial resolution comparable to other surveys of dust emission, notably the IR. ",
        "conclusions": "We have obtained the highest spatial resolution images of the diffuse UV background to date with an effective spatial resolution of about 2\\arcmin. These observations are in a region of moderate optical depth with $\\tau\\ > 1$ where  \\citet{Sandage76} observed, and correctly identified, starlight scattered from interstellar dust. We have obtained \\galex\\ observations in both the FUV (1350 -- 1750 \\AA) and NUV (1750 -- 2850 \\AA) bands. After subtraction of the foreground airglow and zodiacal light, we were left with about 500 - 800 \\phunit\\ in both the FUV and NUV bands. The FUV/NUV ratio increased with increasing FUV emission suggesting the presence of \\htwo\\ fluorescent emission in the FUV ranging up to an integrated emission of about 250 \\phunit. Our values are consistent with those of \\citet{MHB90} in their scan of this region; however, we observe a spatial variability that was not possible with their observations.  When we link our data with the observations of G251.2+73.3 by \\citet{Haikala}, we find that the scattered UV light increases linearly with the IR emission for low optical depths but saturates at optical depths near unity. This is as expected given that the thermal IR emission is from the entire volume of the cloud because of the low optical depth in the IR. We have used the same models as in our earlier studies of the diffuse background and have found that a dust scattered component with $a = 0.4; g = 0.7$ is consistent with the data; ie., the dust is strongly forward scattering with a moderate albedo. However, the data show an intrinsic scatter much greater than can be attributed to photon noise alone which must reflect structure at a spatial scale of at least 2\\arcmin\\, possibly due to variations in the ISRF and in the  distribution of the interstellar dust.  We are now extending our analysis to a much larger body of \\galex\\ observations, both our own and archival data. Such an investigation will help resolve some of the uncertainties in this work such as the contribution of airglow and the zodiacal light. Previous observations of the scattering by interstellar dust were on much larger spatial scales and indicated a local origin to much of the background; ie., both nearby hot stars and interstellar dust were required. \\galex\\ observations will allow us to probe the diffuse background at much higher spatial resolutions and thus to investigate the small scale structure of the ISM."
    },
    "0807/0807.0376_arXiv.txt": {
        "abstract": "The GAMA survey aims to deliver 250,000 optical spectra (3--7\\AA~resolution) over 250 sq.\\ degrees to spectroscopic limits of $r_{AB} <19.8$ and $K_{AB}<17.0$ mag. Complementary imaging will be provided by GALEX, VST, UKIRT, VISTA, HERSCHEL and ASKAP to comparable flux levels leading to a definitive multi-wavelength galaxy database. The data will be used to study all aspects of cosmic structures on 1kpc to 1Mpc scales spanning all environments and out to a redshift limit of $z \\approx 0.4$. Key science drivers include the measurement of: the halo mass function via group velocity dispersions; the stellar, HI, and baryonic mass functions; galaxy component mass-size relations; the recent merger and star-formation rates by mass, types and environment. Detailed modeling of the spectra, broad SEDs, and spatial distributions should provide individual star formation histories, ages, bulge-disc decompositions and stellar bulge, stellar disc, dust disc, neutral HI gas and total dynamical masses for a significant subset of the sample ($\\sim 100$k) spanning both the giant and dwarf galaxy populations. The survey commenced March 2008 with 50k spectra obtained in 21 clear nights using the Anglo Australian Observatory's new multi-fibre-fed bench-mounted dual-beam spectroscopic system (AA$\\Omega$). ",
        "introduction": "Galaxy And Mass Assembly (GAMA) is a major expansion of the Millennium Galaxy Catalogue (MGC) survey (Liske et al 2003; Allen et al 2005; Driver et al 2005) and a natural extension of the extremely productive nearby ``Legacy'' surveys (e.g., SDSS, 2MASS, HIPASS etc).  In comparison to the superb SDSS survey GAMA will only sample 250 sq degrees of sky but will extend to significantly fainter spectroscopic limits (12$\\times$ the redshift density of SDSS main, 5$\\times$ stripe 82), to higher spatial ($0.6''$ FWHM) and spectral (3---7\\AA) resolutions, as well as moving to a far broader wavelength coverage (UV to Radio). GAMA has come about by parallel technological developments leading to a suite of new facilities whose survey sensitivities, resolutions, and capabilities are reasonably well matched. Until now the study of galaxies has generally been restricted to either large samples of limited wavelength data or multi-wavelength studies of small (and often biased) samples. However galaxy systems are extremely complex and diverse, exhibiting strong environmental and mass dependencies and containing distinct but interlinked components (AGN, nucleus, bulge, pseudo-bulge, bar, disc etc) and constituents (SMBH, plasma, stars, gas, dust etc). It then follows that a clear understanding of galaxy formation and evolution may only come about via the construction of a comprehensive survey which simultaneously samples all of these facets.  The GAMA team aims to provide this data. In addition to the provision of a generic galaxy database, the GAMA project also includes a number of more focussed science goals, in particular:  \\begin{description} \\item{1. } Measurement of the Halo Mass Function via virialised group velocity dispersions to directly test the {\\it numerical} prediction from CDM (and WDM) simulations. \\item{2. } Measurement of the dynamic, baryonic, HI and stellar mass functions to LMC masses versus redshift, environment, type, and component (as well as higher order relations, e.g., mass-spin [$M-\\lambda$]). \\item{3. } Measurement of the recent merger rates and star formation rates versus type, mass and environment over a 3---4 Gyr baseline. \\end{description} ",
        "conclusions": ""
    },
    "0807/0807.0006_arXiv.txt": {
        "abstract": "The orbital parameters of extra-solar planets have a significant impact on the probability that the planet will transit the host star. This was recently demonstrated by the transit detection of HD 17156b whose favourable eccentricity and argument of periastron dramatically increased its transit likelihood. We present a study which provides a quantitative analysis of how these two orbital parameters effect the geometric transit probability as a function of period. Further, we apply these results to known radial velocity planets and show that there are unexpectedly high transit probabilities for planets at relatively long periods. For a photometric monitoring campaign which aims to determine if the planet indeed transits, we calculate the significance of a null result and the subsequent constraints that may be applied to orbital parameters. ",
        "introduction": "There have been at least five cases in which planetary transits were detected through photometric follow-up of planets already known via their radial velocity (RV) discoveries. The case of HD 17156b (\\cite{bar07a}) is of particular interest since it is a 21.2 day period planet which happens to have a large eccentricity ($e = 0.67$) and an argument of periastron which places the periapsis of its orbit in the direction toward the observer and close to parallel to the line of sight, resulting in an increased transit probability.  Recent work by \\cite{bar07b} and \\cite{bur08} showed that higher eccentricities of planetary orbits will increase their transit probabilities and, consequently, expected yield for transit surveys. We demonstrate the combined effect of the eccentricity, $e$, and argument of periastron, $\\omega$, on transit probability. As shown by \\cite{kan07a}, the place in a planetary orbit where it is possible for a transit to occur (where the planet passes the star-observer plane perpendicular to the planetary orbit) is when $\\omega + f = \\pi / 2$. The probability of such a transit occurring, $P_t$, is given by \\begin{equation} P_t = \\frac{(R_p + R_\\star)(1 + e \\cos (\\pi/2 - \\omega))}{a (1 - e^2)}, \\label{transit_prob} \\end{equation} where $R_p$ and $R_\\star$ are the radii of the planet and star respectively, and $a$ is the semi-major axis.  The orbital configuration, especially with regards to the values of $e$ and $\\omega$, plays a major role in determining the likelihood of a planet transiting the parent star. ",
        "conclusions": ""
    },
    "0807/0807.0230_arXiv.txt": {
        "abstract": "We analyze the physical conditions of the outflow seen in QSO~2359--1241 (NVSS J235953--124148), based on high resolution spectroscopic VLT observations. This object was previously studied using Keck/HIRES data. The main improvement over the HIRES results is our ability to accurately determine the number density of the outflow.  For the major absorption component, level population from five different \\ion{Fe}{2} excited level yields $\\vy{n}{H}=10^{4.4}$ cm$^{-3}$ with less than 20\\% scatter. We find that the \\ion{Fe}{2} absorption arises from a region with roughly constant conditions and temperature greater than 9000~K, before the ionization front where temperature and electron density drop. Further, we model the observed spectra and investigate the effects of varying gas metalicities and the spectral energy distribution of the incident ionizing radiation field. The accurately measured column densities allow us to determine the ionization parameter ($\\log U_H \\approx -2.4$) and total column density of the outflow ($\\log N_H(\\rm{cm}^{-2}) \\approx 20.6$). Combined with the number density finding, these are stepping stones towards determining the mass flux and kinetic luminosity of the outflow, and therefore its importance to AGN feedback processes. ",
        "introduction": "In recent years, the potential impact of quasar outflows on their environment has become widely recognized  (e.g., Silk \\& Rees 1998, King 2003, Cattaneo et al.\\ 2005, Hopkins et al.\\ 2006).  Outflows are detected as absorption troughs in quasar spectra that are blueshifted with respect to the systemic redshift of their emission line counterparts. The absorption troughs are mainly associated with UV resonance lines of various ionic species (e.g., \\ion{Mg}{2}~$\\lambda\\lambda$2796.35,2803.53, \\ion{Al}{3}~$\\lambda\\lambda$1854.72,1862.79, \\ion{C}{4}~$\\lambda\\lambda$1548.20,1550.77, \\ion{Si}{4}~$\\lambda\\lambda$1393.75,1402.77, \\ion{N}{5}~$\\lambda\\lambda$1238.82,1242.80).  Some quasar outflows show absorption troughs from excited and metastable states. The ratio of the population level between the excited or metastable states and the ground state is sensitive to the number density and temperature of the plasma (Wampler, Chugai \\& Petitjean 1995; de~Kool et~al.\\ 2001). Therefore, accurate measurements of the column densities associated with both excited or metastable states and the ground state of a given ion can yield the gas number density of the outflow. In addition, these measurements and similar ones of troughs from other ions and elements allow us to determine the ionization equilibrium and total column density in the outflow (Arav et~al.\\ 2001; Arav et~al.\\ 2007).  Accurate column densities for the outflow's troughs are difficult to determine since the outflow does not cover the emission source homogeneously (Barlow 1997; Telfer et~al.\\ 1998; Arav 1997; Arav et~al.\\ 2003). Over the past several years, we have developed techniques for extracting reliable column densities for such situations (Arav et~al.\\ 1999a; Arav et~al.\\ 1999b; de~Kool et~al.\\ 2001, 2002a,b; Arav et~al.\\ 2002; Scott et al.\\ (2004) Gabel et al.\\ (2005); Arav et~al.\\ 2005). These efforts culminated with the analysis of spectroscopic VLT observations of the outflow seen in QSO~2359--1241 (Arav et~al.\\ 2008; hereafter Paper~I). These data contain absorption troughs from five resonance \\ion{Fe}{2} lines, as well as those from several other metal species and metastable excited state \\ion{He}{1}, with a resolution of $\\sim$7 km~s$^{-1}$ and signal-to-noise ratio per resolution element of order 100.  QSO~2359--1241 (NVSS J235953--124148; $E = 15.8$) is an intrinsically reddened ($A_V \\approx 0.5$), luminous ($M_B = -28.7$), radio-moderate, optically polarized ($\\sim 5\\%$), low-ionization broad absorption line quasar, at relatively low redshift $z \\approx 0.868$. See Brotherton et~al.\\ (2001) for further details. Brotherton et~al.\\ (2005) describes its X-ray spectrum. An initial investigation of the physical properties of the outflow in this object using $HST$ FOC and especially Keck HIRES spectra is described in Arav et~al.\\ (2001).  The VLT spectral data set of QSO~2359--1241 is described in detail in Paper~I. Its unprecedented high-quality allowed us to test a variety of absorber distribution models needed to derive reliable ionic column densities of the outflow (see Paper~I). In the present paper we report these column densities and use them to determine the physical conditions within the main component of the outflow ({\\bf e}, see Paper~I): the ionization equilibrium, total column density and number density of the absorbing material. To do so we use the photoionization code Cloudy (Ferland et al.\\ 1998) as well as a separate \\ion{Fe}{2} ion model (Bautista \\& Pradhan 1998).  The plan of the paper is as follows: In Section~2 we describe the column density measurements. In Section~3 we determine the physical conditions in the main component of the outflow. Finally, in Section~4 we summarize and discuss our results and provide a simple estimate of the outflow's distance from the central continuum source.  \\clearpage \\begin{deluxetable}{lrlrrrr} \\tabletypesize{\\scriptsize} \\tablecaption{Absorption lines identified in the VLT spectrum of QSO~2359--1241} \\tablewidth{0pt} \\tablehead{ \\colhead{ $\\lambda$}&\\colhead{$\\log(gf)$\\tablenotemark{a}}& \\colhead{Ion}&\\colhead{E$_{low}$(cm$^{-1}$)}& \\colhead{$g_{low}$}&\\colhead{E$_{up}$(cm$^{-1}$)}& \\colhead{$g_{up}$}} \\startdata 2764.62 &-1.95  &\\ion{He}{1}*  &159856 & 3& 196027 & 9\\cr 2829.92 &-1.74  &\\ion{He}{1}*  &159856 & 3& 195193 & 9\\cr 2945.98 &-1.58  &\\ion{He}{1}*  &159856 & 3& 193801 & 9\\cr 3188.69 &-1.16  &\\ion{He}{1}*  &159856 & 3& 191217 & 9\\cr 3889.80 &-0.72  &\\ion{He}{1}*  &159856 & 3& 185565 & 9\\cr 2852.97 & 0.270 &\\ion{Mg}{1}   &     0 & 1&  35051 & 3\\cr 2796.36 & 0.100 &\\ion{Mg}{2}   &     0 & 2&  35761 & 4\\cr 2803.54 &-0.210 &\\ion{Mg}{2}   &     0 & 2&  35669 & 2\\cr 1854.72 & 0.060 &\\ion{Al}{3}  &     0 & 2&  53917 & 4\\cr 1862.79 &-0.240 &\\ion{Al}{3}  &     0 & 2&  53683 & 2\\cr 1808.01 &-2.100 &\\ion{Si}{2}   &     0  &2  &55309  &4\\cr 1816.93 &-1.840 &\\ion{Si}{2}m* &   287  &4  &55325  &6\\cr 3934.83 & 0.134 &\\ion{Ca}{2}   &     0  &2  &25414  &4\\cr 3969.65 &-0.166 &\\ion{Ca}{2}   &     0  &2  &25192  &2\\cr 2576.87 & 0.433 &\\ion{Mn}{2}   &     0  &7  &38807  &9\\cr 2594.49 & 0.270 &\\ion{Mn}{2}   &     0  &7  &38543  &7\\cr 2606.46 & 0.140 &\\ion{Mn}{2}   &     0  &7  &38366  &5\\cr 2344.2139&  0.057& \\ion{Fe}{2}  &      0& 10&  42658&  8\\cr 2374.4612& -0.504& \\ion{Fe}{2}  &      0& 10&  42115& 10\\cr 2382.7652&  0.505& \\ion{Fe}{2}  &      0& 10&  41968& 12\\cr 2586.6500& -0.161& \\ion{Fe}{2}  &      0& 10&  38660&  8\\cr 2600.1729&  0.378& \\ion{Fe}{2}  &      0& 10&  38459& 10\\cr 2333.5156& -0.206& \\ion{Fe}{2}*&    385&  8&  43239&  6\\cr 2365.5518& -0.402& \\ion{Fe}{2}*&    385&  8&  42658&  8\\cr 2389.3582& -0.180& \\ion{Fe}{2}*&    385&  8&  42237&  8\\cr 2396.3559&  0.362& \\ion{Fe}{2}*&    385&  8&  42115& 10\\cr 2599.1465& -0.063& \\ion{Fe}{2}*&    385&  8&  38859&  6\\cr 2612.6542&  0.004& \\ion{Fe}{2}*&    385&  8&  38660&  8\\cr 2626.4511& -0.452& \\ion{Fe}{2}*&    385&  8&  38459& 10\\cr 2328.1112& -0.684& \\ion{Fe}{2}*&    668&  6&  43621&  4\\cr 2349.0223& -0.269& \\ion{Fe}{2}*&    668&  6&  43239&  6\\cr 2381.4887& -0.693& \\ion{Fe}{2}*&    668&  6&  42658&  8\\cr 2399.9728& -0.148& \\ion{Fe}{2}*&    668&  6&  42335&  6\\cr 2405.6186&  0.152& \\ion{Fe}{2}*&    668&  6&  42237&  8\\cr 2607.8664& -0.150& \\ion{Fe}{2}*&    668&  6&  39013&  4\\cr 2618.3991& -0.519& \\ion{Fe}{2}*&    668&  6&  38859&  6\\cr \\tablebreak 2632.1081& -0.287& \\ion{Fe}{2}*&    668&  6&  38660&  8\\cr 2338.7248& -0.445& \\ion{Fe}{2}*&    863&  4&  43621&  4\\cr 2359.8278& -0.566& \\ion{Fe}{2}*&    863&  4&  43239&  6\\cr 2405.1638& -0.983& \\ion{Fe}{2}*&    863&  4&  42440&  2\\cr 2407.3942& -0.228& \\ion{Fe}{2}*&    863&  4&  42401&  4\\cr 2411.2433& -0.076& \\ion{Fe}{2}*&    863&  4&  42335&  6\\cr 2614.6051& -0.365& \\ion{Fe}{2}*&    863&  4&  39109&  2\\cr 2631.8321& -0.281& \\ion{Fe}{2}*&    863&  4&  38859&  6\\cr 2345.0011& -0.514& \\ion{Fe}{2}*&    977&  2&  43621&  4\\cr 2411.8023& -0.377& \\ion{Fe}{2}*&    977&  2&  42440&  2\\cr 2414.0450& -0.455& \\ion{Fe}{2}*&    977&  2&  42401&  4\\cr 2622.4518& -0.951& \\ion{Fe}{2}*&    977&  2&  39109&  2\\cr 2629.0777& -0.461& \\ion{Fe}{2}*&    977&  2&  39013&  4\\cr 2332.00  & -0.720& \\ion{Fe}{2}*&   1873& 10&  44754&  8\\cr 2348.81  & -0.470& \\ion{Fe}{2}*&   1873& 10&  44447&  8\\cr 2360.70  & -0.700& \\ion{Fe}{2}*&   1873& 10&  44233& 10\\cr 2563.30  & -0.050& \\ion{Fe}{2}*&   7955&  8&  46967&  6\\cr 2715.22  & -0.440& \\ion{Fe}{2}*&   7955&  8&  44785&  6\\cr 2740.36  &  0.240& \\ion{Fe}{2}*&   7955&  8&  44447&  8\\cr 2756.56  &  0.380& \\ion{Fe}{2}*&   7955&  8&  44233& 10\\cr 2166.19  &0.230 &\\ion{Ni}{2}*  & 8394 &10  &54557 &10\\cr 2217.14  &0.480 &\\ion{Ni}{2}*  & 8394 &10  &53496 &12 \\cr 2223.61 &-0.140 &\\ion{Ni}{2}*  & 8394 &10  &53365 &10\\cr 2316.72  &0.268 &\\ion{Ni}{2}*  & 8394 &10  &51558 & 8\\cr \\enddata \\tablenotetext{a}{$gf$-values were taken from Kurucz (1995) for all but the \\ion{Fe}{2} lines with wavelengths given more than two decimal figures. $gf$-values for these come from Morton (2003).} \\end{deluxetable} \\clearpage ",
        "conclusions": "First, we compare the results of this investigation to the analysis of the Keck/HIRES observations of QSO~2359--1214 by Arav et al.\\ (2001; hereafter HIRES paper). Using the same MF87 SED, the HIRES analysis found $\\log N_H = 20.2$ and $\\log U_H = -2.7$, compared to $\\log N_H = 20.6$ and $\\log U_H = -2.4$ for the VLT analysis. These factors of $\\sim2$ differences are mainly attributed to using apparent optical depth methods to extract the \\ion{Fe}{2} and \\ion{He}{1}* column densities. As pointed out in the HIRES paper, the data was not of high enough signal-to-noise ratio to permit more sophisticated analyses.  Even so, the HIRES paper already showed that the outflow is not shielded by a hydrogen ionization front, a result confirmed by the VLT analysis.  The important leap in diagnostic power for the VLT data came from the ability to accurately measure the population levels of the excited \\ion{Fe}{2} levels, allowing us to pin point the number density of the outflow to $\\log(n_e)=10^{4.4}$ cm$^{-3}$ to better than 20\\% accuracy. This is both qualitatively and quantitatively a great improvement over the lower limit of $\\log(n_e)=10^{5}$ cm$^{-3}$ available from the HIRES data. This result is not accidental. The main reason we invested 6.5 hours of VLT observation on this outflow was precisely to yield a data set from which an accurate $n_e$ could be extracted. This determination of $n_e$ will allow us to determine the distance of the outflow from the central source and thus measure its mass flux and kinetic luminosity. This demonstrates the importance of taking high quality spectra of such outflows.  Other important results arising from the measured populations of the \\ion{Fe}{2} levels are the determination of a lower limit to the temperature of the \\ion{Fe}{2} region and the realization that the absorption spectrum forms before the hydrogen ionization front, beyond which the temperature and ionization drop sharply. The temperature determination was crucial in constraining a whole family of SEDs and gas metalicities that would yield very different temperatures in the \\ion{Fe}{2} region, and consequently allowing for a more secure determination of the total gas column density and ionization parameter. That the \\ion{Fe}{2} absorption occurs in a region of nearly constant conditions before the hydrogen ionization front is a key to being able to model the absorption spectrum.  Photoionization modeling allowed us to reproduce quite well the observed \\ion{Fe}{2} and \\ion{He}{1}* column densities in the main, {\\bf e}, component of the quasar outflow absorption spectrum. Reiterating, we found $\\log N_H \\approx 20.6$ and $\\log U_H \\approx -2.4$. The dominant error bars to these values come from the uncertainties in the assumed SED and gas metalicities and come to $\\sim$~0.3 dex.  Given the above gas density, ionization parameter, an estimate to an {\\em unobscured} incident bolometric luminosity of $\\sim4.7 \\times 10^{47}$ ergs~s$^{-1}$ based on the intrinsic reddening correction in Brotherton et~al.\\ (2001, 2005), and a standard cosmology ($H_o = 70$ km~s$^{-1}$~Mpc$^{-1}$, $\\Omega_\\Lambda = 0.70$, $\\Omega_m = 0.30$), we estimate a distance of component {\\bf e} of the outflow from the central continuum source of $\\sim3$ kpc. In a future paper we will similarly determine the physical conditions in the weaker, lower velocity outflow components {\\bf a--d}, as well as the distances and estimates of the kinetic luminosities for all components in the outflow, important to AGN feedback scenarios of galaxy evolution."
    },
    "0807/0807.1000_arXiv.txt": {
        "abstract": "Searching for transit timing variations in the known transiting exoplanet systems can reveal the presence of other bodies in the system. Here we report such searches for two transiting exoplanet systems, TrES-1 and WASP-2. Their new transits were observed with the 4.2m William Herschel Telescope located on La Palma, Spain. In a continuing programme, three consecutive transits were observed for TrES-1, and one for WASP-2 during September 2007. We used the Markov Chain Monte Carlo simulations to derive transit times and their uncertainties. The resulting transit times are consistent with the most recent ephemerides and no conclusive proof of additional bodies in either system was found. ",
        "introduction": "Transiting planets provide a wealth of information about exoplanetary systems. Short-term variations in the mid-eclipse times of the transits may reveal the presence of moons, trojans or other planets (\\cite[Holman \\& Murray 2005]{HolmanMurray2005}, \\cite[Agol \\etal\\ 2005]{Agol_etal05}, \\cite[Ford \\& Holman 2007]{FordHolman07}), whereas long-term variations could result from orbital precession (\\cite[Miralda--Escud\\'{e} 2002]{Miralda02}). This provides further constraints on theories of planetary system formation and evolution, as well on theories of planetary atmospheres and their composition. ",
        "conclusions": ""
    },
    "0807/0807.3005_arXiv.txt": {
        "abstract": "Pre-main sequence (PMS) stars are known to produce powerful X-ray flares which resemble magnetic reconnection solar flares scaled by factors up to $10^4$. However, numerous puzzles are present including the structure of X-ray emitting coronae and magnetospheres, effects of protoplanetary disks, and effects of stellar rotation. To investigate these issues in detail, we examine 216 of the brightest flares from 161 PMS stars observed in the Chandra Orion Ultradeep Project (COUP). These constitute the largest homogeneous dataset of PMS, or indeed stellar flares at any stellar age, ever acquired. Our effort is based on a new flare spectral analysis technique that avoids nonlinear parametric modeling. It can be applied to much weaker flares and is more sensitive than standard methods. We provide a catalog with $>30$ derived flare properties and an electronic atlas for this unique collection of stellar X-ray flares. The current study (Paper I) examines the flare morphologies, and provides general comparison of COUP flare characteristics with those of other active X-ray stars and the Sun. Paper II will concentrate on relationships between flare behavior, protoplanetary disks,  and other stellar properties.   Several results are obtained. First, the COUP flares studied here are among the most powerful, longest, and hottest stellar X-ray flares ever studied. Peak luminosities are in the range $31< \\log L_{X,pk}< 33$ erg~s$^{-1}$; rise (decay) timescales range from 1~hour to 1~day (few hours to 1.5 days); many peak temperatures exceed 100~MK. The scale of their inferred associated coronal structures is $0.5-10$~$R_{\\star}$. Second, no significant statistical differences in peak flare luminosity or temperature distributions are found among different morphological flare classes, suggesting a common underlying mechanism for all flares. Third, comparison with the general solar-scaling laws indicates that COUP flares may not fit adequately proposed power-temperature and duration-temperature solar-stellar fits. Fourth, COUP super-hot flares are found to be brighter but shorter than cooler COUP flares. Fifth, the majority of bright COUP flares can be viewed as enhanced analogs of the rare solar ``long-duration events''. ",
        "introduction": "}   All solar-type stars exhibit their highest levels of magnetic activity during their pre-main sequence (PMS) phase \\citep{Feigelson07}. This includes `superflares' with peak luminosities $\\log L_x \\ga 32$ erg s$^{-1}$ in the $0.5-8$ keV band, $10^4$ more powerful than the strongest flares seen in the contemporary Sun \\citep[e.g.][]{Tsuboi98, Grosso04, Favata05}. PMS stars thus join RS~CVn binary systems \\citep[e.g.][]{Osten07} as laboratories to study the physics of the most powerful magnetic reconnection events.  PMS stars are more distant and fainter than the closer RS~CVn systems, but hundreds of flaring PMS stars can be simultaneously studied due to their concentration in rich clusters.   The magnetic field structure of PMS stars, and thus the nature of their reconnection and flaring, may (or may not) qualitatively differ from other stars due to the presence of a protoplanetary disk during the early PMS stages. The intense high energy radiation from these PMS reconnection events may affect the physical and chemical properties of the surrounding circumstellar environment and play an important role in the formation of planets \\citep{Glassgold05,Feigelson07}.  A consensus has emerged during the past decade that PMS accretion is funneled by magnetic field lines linking the disk inner edge to the stellar surface \\citep[e.g.][]{Hartmann98, Shu00}.  However, while early theory assumed a dipolar field morphology, recent studies point to a complex multipolar field structure similar to the Sun's \\citep{Jardine06, Donati07, Long08}.   It is also unclear whether the X-ray flares occur primarily in large loops with both footprints anchored on the stellar surface, or in loops linking the stellar photosphere with the inner rim of the circumstellar disk \\citep{Isobe03, Favata05}. The first case may suffer instability due to centrifugal force \\citep{Jardine99} while the second case may load the loop with cool accreting material so that X-rays may not be produced \\citep{Preibisch05}.   The 13-day nearly continuous observation of $\\sim 1408$ PMS stars in the Orion Nebula, the Chandra Orion Ultradeep Project \\citep[COUP;][]{Getman05}, enables both studies of individual flare properties and statistical studies of flaring from Orion stars \\citep{Wolk05, Flaccomio05, Stassun06, Caramazza07, Colombo07}. COUP also provided a unique opportunity to study relatively rare superflares and long-duration flares. \\citet{Favata05} have analyzed the strongest 32 flares in the COUP dataset using a long-standing method of time resolved spectroscopy (TRS) modeled as cooling plasma loops. They concluded that at least 1/3 of these are produced by magnetic reconnection in very long coronal $5-20$~R$_{\\star}$ structures. Such structures were predicted in magnetospheric accretion models \\citep[e.g][]{Shu97} but not clearly identified before COUP. \\citet{Favata05} recognized that their sample was too small to quantitatively probe the relationship between long coronal flaring structure and disks or accretion.  The aim of the current study is to extend the flare sample of \\citet{Favata05} utilizing a more sensitive technique of flare analysis, the ``method of adaptively smoothed median energy'' (MASME) introduced by \\citet{Getman06}.  We combine this method with the astrophysical cooling loop models of \\citet{Reale97} to trace the evolution of the flare plasma in temperature-density diagrams and derive flaring loop sizes. The method allowed us to examine $216$ of the brightest flares from $161$ brightest COUP PMS stars. These constitute the largest homogeneous sample of powerful stellar flares ever acquired in the X-ray band. In \\citet{Getman08} (Paper~II), we use these results to study in detail the relationships between PMS X-ray flares, stellar properties, protoplanetary disks, and accretion.  Our flare analysis and the derived flare properties and classifications are presented in \\S \\ref{analysis}. Properties of the stars themselves are also provided. Global properties of our flares are considered in \\S \\ref{results} and compared to published studies of older stars. ",
        "conclusions": "}   We analyze 216 bright X-ray flares from the Chandra Orion Ultradeep Project which provides the longest nearly-continuous observation of a rich PMS stellar cluster in the X-ray band. Our effort is based on a new spectral analysis technique (MASME) that avoids nonlinear parametric modeling and is more sensitive than standard methods.  Flare loop parameters are derived from the well-established flare plasma model of \\citet{Reale97}.   We thus emerge with the largest dataset of PMS flares, or indeed stellar flares at any stellar age, with peak luminosities in the range  $31< \\log L_{X,pk}< 33$ erg~s$^{-1}$, several orders of magnitude more powerful than any solar flare. For each flare we provide a catalog of $>30$ derived flare properties including inferred sizes of associated coronal loops and flare morphological classes. We give an electronic atlas with flare lightcurves, temporal evolution plots of X-ray median energy, plasma temperature, emission measure, and derived temperature density diagram. This collection of empirical and model-dependent information can serve as a valuable testbed for stellar flare models.   The powerful COUP flares studied here have rise timescales ranging from 1~hour to 1~day and decay timescales ranging from a few hours to 1.5 days.  An important empirical result is that peak plasma temperatures are often 100~MK, in some cases $> 200$~MK. These temperatures are derived from a robust calibration of median energies; traditional time-resolve spectroscopy often does not have the time resolution to detect this brief super-hot phase. No significant differences in peak flare luminosity or temperature distributions are found among the wide range of morphological flare classes: typical fast-rise exponential-decays, step decays, double peaks, and slow-rise flat-top.  This suggests that all flare types arise from similar underlying magnetic reconnection mechanisms and geometries.   Comparison of the COUP flare properties with the general solar-stellar scaling laws of \\citet{Aschwanden08} presents surprising results. Our flares do not follow the solar-stellar trend between plasma peak emission measure and temperature, $EM_{pk} \\propto T_{obs,pk}^{4.7}$. The trend between flare duration and peak temperature is also absent. Super-hot COUP flares are found to be brighter but shorter in duration than cooler COUP flares. This is further developed in Paper II.   Compared to non-PMS systems, the inferred sizes of COUP flaring structures are remarkably large, ranging widely from $L=0.5$ to $10$~R$_{\\star}$. These large flaring structures must be associated with large-scale stellar magnetic fields. Rare long decay solar events associated with the largest known X-ray emitting structures are possible solar analogs to these COUP flares.   Our flare sample provides a valuable laboratory for the study of the physics  and astronomy of magnetic reconnection events.   This study (Paper I) examines flare morphologies, and provides general comparison of COUP flare characteristics with those of other active X-ray stars and the Sun. Paper II concentrates on relationships between flare behavior, protoplanetary disks (both passive and accreting), and other stellar properties including rotational periods and Keplerian corotation radii. Paper II further investigates super-hot COUP flares and magnetic field strength on Orion T-Tauri stars."
    },
    "0807/0807.4139_arXiv.txt": {
        "abstract": "We present the results of our study of astrometric stability of 200-in Hale (Mt. Palomar) and 10-m Keck II (Mauna Kea) telescopes, both with Adaptive Optics (AO) facilities. A group of nearby visual binaries and multiples was observed in near infrared, relative separations and position angles measured. We have also checked the influence of some systematic effects (e.g. atmospherical refraction, varying plate scale factor) on result and precision of astrometric measurements. We conclude that in visual binaries astrometrical observations it is possible to achieve much better precision than 1 miliarcsecond [$mas$], which in many cases allows detection of the astrometrical signal produced by planetary-mass object. ",
        "introduction": "Astrometry is thought to be the most promising method of exoplanets detection in the future. On the contrary to radial velocities measurements (RV), astrometry is almost independent on stellar physics, e.g. phenomena like the activity or pulsations, rather than the distance to the object. Todays interferometers are able to achieve precision at the level of mili- or microarcseconds, which is good enough to attempt exoplanet research.  Nevertheless, the same level is possible to achieve in small fields by CCD imaging with adaptive optics (AO) systems. To do that, one must subtract and correct some systematical effects in order to obtain a gaussian distribution of the measurements. In such case, the precision improves like $N^{-0.5}$, where $N$ is a number of single measurements. The goal of our studies was to check if a gaussian statistic can really be achieved with two top-class telescopes with AO systems: the 200-in Hale telescope + PHARO camera (Mt. Palomar) and 10-m Keck II telescope + NIRC2 (Mauna Kea), both working in infrared (IR). With these facilities we were imaging 17 objects, majority of which were M-type dwarf binaries, located closer than 20 $pc$ from the Sun. ",
        "conclusions": "Todays telescopes with AO systems allows us to perform astrometric measurements with precision well bellow 1 $mas$. This refers to a single-epoch observations, as well as to a long-term stability. For many cases such a precision means an ability to detect exoplanets around nearby stars.  \\begin{figure} \\includegraphics[width=0.95\\columnwidth]{Helminiak_fig1.eps} \\caption{Example of Allan variance vs. lag for uncorrected (dash-dotted) and corrected (solid line) measurements of separation components $x$ (left) and $y$. A small decrease is normal for finite series of measurements. Dotted line shows a --1-behavior for an ideal, infinitely long white-noise series.} \\end{figure}  \\begin{table} \\caption{Smallest $\\sigma_{\\rho}$ and corresponding detection limits for researched stars.} \\begin{tabular}{c|c|c|cc|cc|c} \\hline Star & Lowest & Dist. & Mass A & Limit for A & Mass B & Limit for B& Tel. \\\\ (GJ No.)& $\\sigma\\, [mas]$ & $[pc]$ & $[M_{\\odot}]$ & $[AU \\cdot M_J]$ & $[M_{\\odot}]$ & $[AU \\cdot M_J]$ &\\\\ \\hline 195  & 0.12  & 13.89 & 0.53  &  1.38 &  0.19 &  0.50 & Hale \\\\ 352  & 1.11  & 10.53 & 0.44  &  8.04 &  0.41 &  7.49 & Hale \\\\ 458  & 0.28  & 15.32 & 0.40  &  2.68 &  0.37 &  2.48 & Hale \\\\ 507  & 0.33  & 13.16 & 0.46  &  3.12 &  0.37 &  2.51 & Hale \\\\ 569B & 0.11  &  9.81 & 0.071 & 0.116 & 0.054 & 0.088 & Keck \\\\ 661  & 0.038 &  6.32 & 0.379 &  0.16 &  0.34 &  0.15 & Hale \\\\ 767  & 0.09  & 13.35 & 0.44  &  0.83 &  0.40 &  0.75 & Hale \\\\ 860  & 0.048 &  4.01 & 0.34  &  0.10 &  0.27 &  0.09 & Hale \\\\ 873  & 0.57  &  5.05 & 0.36  &  1.62 & unknown & unknown & Hale \\\\ 9071 & 0.20  & 13.89 & 0.53  &  2.22 &  0.49 &  2.05 & Hale \\\\ \\hline \\end{tabular} \\end{table}"
    },
    "0807/0807.1693_arXiv.txt": {
        "abstract": "Large volume cosmological simulations succeed in reproducing the large-scale structure of the Universe. However, they lack resolution and may not take into account all relevant physical processes to test if the detail properties of galaxies can be explained by the CDM paradigm. On the other hand, galaxy-scale simulations could resolve this in a robust way but do not usually include a realistic cosmological context.  To study galaxy evolution in cosmological context, we use a new  method that consists in coupling cosmological simulations and galactic scale simulations. For this, we record merger and gas accretion histories from cosmological simulations and re-simulate at very high resolution the evolution of baryons and dark matter within the virial radius of a target galaxy. This allows us for example to better take into account gas evolution and associated star formation, to finely study the internal evolution of galaxies and their disks in a realistic cosmological context.  We aim at obtaining a statistical view on galaxy evolution from z $\\simeq$ 2 to 0, and we present here the first results of the study: we mainly stress the importance of taking into account gas accretion along filaments to understand galaxy evolution. ",
        "introduction": "The morphology of galaxies in the Local Universe is well constrained by observations, but is still largely unexplained. Indeed, large volume cosmological simulations fail to reproduce realistic galaxies. For instance, the disks formed are often too concentrated : it is the ``angular momentum problem'', well known since the early work of \\cite{Navarro1991}.  It is still unclear whether this is an intrinsic problem of the $\\Lambda$CDM paradigm or if something (i.e. resolution, physical processes...) is missing in these simulations.  Another puzzle is the question of disk survival till z=0 (\\cite{Koda2007}). For instance, \\cite{Kautsch2006} study a large sample of edge-on spiral galaxies in the SDSS and find that a significant fraction of them (i.e. roughly one third) are bulgeless or ``superthin''. This is still unexplained by cosmological models. Indeed, $\\Lambda$CDM predicts that galaxy interactions are frequent (see e.g. the recent work by \\cite{Stewart2007}). More exactly, major mergers, that are well known to destroy disks to form ellipticals (\\cite{Barnes1991}) are rather rare, but minor mergers are much more common. These minor mergers can thicken disks, and if frequent enough could even form elliptical galaxies (\\cite{Bournaud2007}). The problem is then to find whether $\\Lambda$CDM predicts too many mergers, or if the satellites have properties and orbital parameters such that they have little influence on the galactic disks. Also, gas accretion along filaments could fuel a thin disk and counteract the effect of mergers (\\cite{Dekel2005}, \\cite{Keres2005}, \\cite{Ocvirk2008}).  To study the properties of galaxies at low and high redshift, it thus seems necessary to take the full cosmological context into account. Large scale cosmological simulations could of course achieve this goal and give a statistical view on galaxies at each redshift, but for now they mainly lack resolution at the galactic scale. On the contrary, small volume cosmological simulations like the one performed by \\cite{Naab2007} can resolve galactic scales in detail but are so time-consuming that obtaining a statistical sample is for now a challenge.  A first method to solve these problem is to use semi-analytical models, i.e. extracting merger trees from cosmological simulations and using different recipes to infer physical properties of galaxies (\\cite{Somerville2001}, \\cite{Hatton2003}, \\cite{Khochfar2005}). The drawback is that approximations are necessary.  Another possibility has been explored by \\cite{Kazantzidis2007}, \\cite{Read2007} and \\cite{Villalobos2008} : they extract merger histories from cosmological simulations and  re-simulate these histories at higher resolution. Nevertheless, they perform collisionless simulations with no gas component, neither in the main galaxy, nor in satellites, nor in filaments.  We here present a new approach where we re-simulate at high resolution a history given by a cosmological simulation, using self consistent realistic galaxies (the main galaxy and the satellites have a gas disk, a stellar disk and a dark matter halo), and we also take into account gas accretion from cosmic filaments. Our goal is to obtain a statistical sample of merger and accretion histories in a $\\Lambda$CDM context to simulate the resulting galaxies and to compare our results to observations at various redshifts.  After a description of the technique used, we will present our first results and emphasize the importance of gas accretion along filaments to understand galaxy evolution. ",
        "conclusions": "In order to study galaxy evolution in cosmological context, we have successfully developed a technique that allows us to perform high resolution simulations taking into account realistic merger and gas accretion histories.  The first two simulations shown here do not allow us to draw any general conclusion on galaxy evolution in a $\\Lambda$CDM context. Nevertheless, we can already confirm that even low mass satellites can thicken disks and that ellipticals from both through repeated minor mergers and major mergers. We also emphasize that gas accretion from filaments can allow to rebuild a thin disk in a galaxy, which proves the absolute necessity to take this accretion into account to understand galaxy evolution."
    },
    "0807/0807.0211.txt": {
        "abstract": "\\noindent We investigate scenarios in which a charged, long-lived scalar particle decouples from the primordial plasma in the Early Universe. We compute the number density at time of freeze-out considering both the cases of abelian and non-abelian interactions and including the effect of Sommerfeld enhancement at low initial velocity. We also discuss as extreme case the maximal cross section that fulfils the unitarity bound. We then compare these number densities to the exotic nuclei searches for stable relics and to the BBN bounds on unstable relics and draw conclusions for the cases of a stau or stop NLSP in supersymmetric models with a gravitino or axino LSP. ",
        "introduction": "%=======================================================================  The early Universe may have been populated by many exotic particles that, especially if charged, should have easily been in thermal equilibrium. No charged relic seems to have survived to the present day. In fact there are very strong upper bounds on the density of electromagnetically % charged and colored and/or colour charged particles with masses below 10--100 TeV from extensive searches for exotic nuclei \\cite{ch-relics}. The standard lore is therefore that only neutral relics may have survived until today.  However, it is possible that some unstable but very long-lived charged particle froze-out from thermal equilibrium and decayed much later to a neutral one. A typical example of this kind in supersymmetric models with R-parity conservation is the next-to-lightest supersymmetric particle~(NLSP) if the LSP and Cold Dark Matter is very weakly interacting like the axino~\\cite{axinolsp,crs02,axinolsp2} or the gravitino~\\cite{gravitinolsp, neutralNLSP}. Recently, such candidates have attracted a lot of attention, and indeed the signal of a charged metastable NLSP at colliders would be spectacular~\\cite{stauatcolliders,Fairbairn:2006gg}.  In general, strong bounds on the number density of any metastable relic with lifetime of about 1\\,s or longer are provided by Big Bang Nucleosynthesis (BBN)~\\cite{BBNrev}. They come from two classes of processes: on one hand injection of very energetic photons or hadrons from decays during or after BBN adds an additional non-thermal component to the plasma and can modify the abundances of the light elements~\\cite{neutBBN}; on the other hand, if the relic particle is electromagnetically charged, bound states with nuclei may arise that strongly enhance some of the nuclear rates and allow for catalysed production of e.g.\\ $^6Li$~\\cite{CBBN}. The bounds of the first type are very tight for lifetimes of the order of $10^4$\\,s and exclude, for instance, a neutralino NLSP with a gravitino LSP in the CMSSM~\\cite{neutralNLSP}. An electrically charged NLSP like the $\\tilde\\tau$ can instead escape the first class of constraints in part of the parameter space, but it is excluded for long lifetimes by bound state effects~\\cite{stauNLSP}. In the axino LSP case, the NLSP has a shorter lifetime; the BBN bounds are hence much weaker and both, neutralino and stau, NLSP are still allowed \\cite{axinolsp}.  In this paper, we investigate the most general case of a scalar charged thermal relic. We compute the number density and compare it to the bounds on exotic nuclei for stable particles and the BBN constraints for unstable ones. Similar studies have been carried out model-independently many years ago~\\cite{wolfram, Nardi:1990ku, unitarity} for stable relics and we will update and improve these computations.\\footnote{Recently the case of general EW charged relics as DM was also considered in full detail~\\cite{Cirelli07}.} % We mostly consider the role of the gauge interaction for two main reasons: i) the annihilation into gauge bosons is often the dominant channel for a charged particle and ii) it depends only on very few parameters, just the mass of the particle and its charge or representation. It is also enhanced by the Sommerfeld effect~\\cite{Sommerfeld},  analogous to heavy quark production at threshold,  which has previously been considered for dark matter annihilations in \\cite{gluino-splitsusy, Hisano:2006nn, Cirelli07, Freitas07, DMsommerfeld} and recently also in the context of leptogenesis in \\cite{Strumia:2008cf}. We discuss this Sommerfeld enhancement for the general abelian and non-abelian cases. Moreover, we compare the cross sections with the unitarity bound and update the unitarity limit on the mass of a stable relic.  Our main goal is to determine if it is at all possible to evade {\\it completely} either the exotic nuclei bounds or the BBN ones and how strongly the particle has to interact in this case. We then apply our findings to the Minimal Supersymmetric Standard Model and discuss in more detail the cases of the stau and stop NLSP.  The paper is organised as follows. In Section~2, we briefly review the computation of the number density from thermal freeze-out. The formulae for the annihilation cross section of a charged particle into gauge bosons are given in Section~3. Here we discuss abelian and non-abelian cases, the Sommerfeld enhancement and the unitarity cross section. Moreover, we compare the thermal averages with the first order in velocity expansion. The resulting relic density is discussed in Section~4. In Section~5, we review the constraints on stable and unstable relics. These are then applied in Section~6 to the concrete examples of relic staus and stops. Section~7 finally contains our conclusions. Details on the computation of the annihilation cross section and the case of massive gauge bosons are given in the Appendices A and B.  %======================================================================= ",
        "conclusions": "%======================================================================  We have studied the number density of a charged relic by computing the annihilation cross section into gauge bosons, including the Sommerfeld enhancement. We have found that the Sommerfeld factor increases the thermally averaged annihilation cross section by 20-50\\% and reduces the final yield even by a factor 2 or 3 for the $SU(3)$ case. Moreover the result is very sensitive on how the higher orders are resummed. Nevertheless the number density surviving the annihilation is still large and BBN constraints are relevant for most relics. They can be avoided completely only for very large $N$ for particles in the fundamental representation of $SU(N)$ ($N> 100 $ for $ m_X \\leq 10 $ TeV) or for cross sections nearly fulfilling the unitarity bound. For the cases of SM gauge groups, the allowed regions only correspond to very  light relic masses, where the number density is low enough, or to sufficiently heavy relic masses so that the decay takes place in the first stages of BBN. The latter allowed region depends strongly on the relic decay channel, and, in case of a gravitino LSP with conserved R-parity, also on the gravitino mass. Let us mention here that if R-parity is just marginally broken, the NLSP can decay with shorter lifetime through R-parity violating channels and the BBN constraints can be easily evaded for any NLSP while keeping the gravitino LSP as Dark Matter~\\cite{R-parity}.  More specifically, for the stau NLSP the light mass window has nearly completely been excluded by direct searches at LEP, even if the annihilation cross-section is maximal $ \\sim 4 \\sigma(\\tilde\\tau\\tilde\\tau^* \\rightarrow \\gamma\\gamma) $, unless the gravitino is lighter than a few tens of GeV, while the large mass region is unfortunately out of reach at the LHC for gravitino masses $m_{3/2} > 100 \\mbox{GeV} $. The detection of a quasi-stable stau at the LHC would then point to a scenario with relatively light gravitino mass, R-parity breaking or an axino LSP and could probably exclude the gravity mediated supersymmetry breaking scenario. In that case the determination of the stau lifetime and its decays will become crucial in distinguishing the different LSPs \\cite{stauatcolliders,bchrs05}.  The stop case is much less constrained thanks to the stronger annihilation cross-section, even if in this case the decay always produces mainly hadrons. We have practically no constraints if the LSP is an axino and even for a gravitino LSP, we can allow for relatively light stops up to approximately 700 GeV ($1\\;$ TeV for lifetimes below $10^{7}\\;$ s), if the annihilation cross section reaches the unitarity one after the QCD phase transition. The window between the present Tevatron bound around 250 GeV and 1 TeV should be surely completely covered by the LHC, the signature being a quasi-stable heavy fermionic meson. The detection of such a state would call for a non-minimal SUSY breaking sector with a coloured NLSP and a very weakly interacting LSP. In this case again only the analysis of the stop decays would allow to distinguish between the lightest states.   %======================================================================"
    },
    "0807/0807.4233_arXiv.txt": {
        "abstract": "The majority of Globular Clusters show chemical inhomogeneities in the composition of their stars, apparently due to a second stellar generation in which the forming gas is enriched by hot-CNO cycled material processed in stars belonging to a first stellar generation. Clearly this evidence prompts questions on the modalities of formation of Globular Clusters. An important preliminary input to any model for the formation of multiple generations is to determine which is today the relative number fraction of ``normal\" and anomalous stars in each cluster.  As it is very difficult to gather very large spectroscopic samples of Globular Cluster stars to achieve this result with good statistical significance, we propose to use the horizontal branch. We assume that, whichever the progenitors of the second generation, the anomalies also include enhanced helium abundance. In fact, helium variations have been recently recognized to be able to explain several puzzling peculiarities (gaps, RR Lyr periods and period distribution, ratio of blue to red stars, blue tails) in horizontal branches.  We summarize previous results and extend the analysis in order to infer the percentage in number of the first and second generation in as many clusters as possible. We show that, with few exceptions, approximately 50\\% or more of the stars belong to the second generation. In other cases, in which at first sight one would think of a simple stellar population, we give arguments and suggest that the stars might all belong to the second generation. We provide in Appendix a detailed discussion and new fits of the optical and UV data of NGC~2808, the classic example of a multiple helium populations cluster, consistently including a reproduction of the main sequence splittings and an examination of the problem of ``blue hook\" stars. We also show a detailed fit of the totally blue HB of M~13, one among the clusters that are possibly fully made up by second generation stars. We conclude that the formation of the second generation is a crucial event in the life of globular clusters. The problem of the initial mass function required to achieve the observed high fraction of second generation stars can be solved only if the initial cluster was much more massive than the present one and most of the first generation low mass stars have been preferentially lost. As shown by D'Ercole et al. by modelling the formation and dynamical evolution of the second generation, the mass loss due to the explosions of the type II supernovae of the first generation may be the process responsible for triggering the expansion of the cluster, the stripping of its outer layers and the loss of most of the first generation low-mass stars. ",
        "introduction": "\\label{sec:intro}  The observations of Globular Cluster (GC) stars are still to be interpreted in a fully consistent frame.  Nevertheless, a general consensus is emerging on the fact that most GCs can not be considered any longer ``simple stellar populations\" (SSP), and that ``self--enrichment\" is a common feature among GCs. This consensus follows from the well known ``chemical anomalies'', already noted in the seventies (such as the variations found in C and N abundances, the Na--O and Mg--Al anticorrelations). Recently observed to be present at the turnoff (TO) and among the subgiants \\citep[e.g.][]{gratton2001,briley2002, briley2004, cohen2005}, they must be attributed to some process of ``self--enrichment\" occurring at the first stages of the cluster life, as the same authors quoted above suggest. There was a first epoch of star formation that gave origin to the ``normal\" (first generation, hereinafter FG) stars, with CNO and other abundances similar to Population II field stars of the same metallicity. Afterwards, there must have been some other epoch of star formation (second generation, hereinafter SG), including material heavily processed through the CNO cycle.  This material either comes entirely from the stars belonging to the first stellar generation, or it is a mixture of processed gas and pristine matter of the initial star forming cloud. We can derive this conclusion as a consequence of the fact that there is no appreciable difference in the abundance of elements such as Ca and the heavier ones between ``normal\" and chemically anomalous stars belonging to the same GC. Needless to say, this statement {\\it does not} hold for $\\omega$ Cen, which must indeed be considered a small galaxy and not a typical GC. In the following, we will only examine ``normal clusters\", those which do not show signs of metal enrichment due to supernova ejecta.  The homogeneity in the heavy elements is an important fact that tells us, e.g., that it is highly improbable that the chemical anomalies are due to mixing of stars born in two different clouds, as there is no reason why the two clouds should have a unique metallicity. In addition, the clusters showing chemical anomalies have a large variety in metallicities, making the suggestion of mixing of two different clouds even more improbable.  The matter must have been processed through the hot CNO cycle, and not, or only marginally, through the helium burning phases, since the sum of CNO elements is the same in the ``normal\" and in the anomalous stars \\citep[e.g.][]{smith1996,ivans1999,cohenmelendez2005}. \\cite{carretta2005} find that actually the CNO is somewhat --but not much-- larger in the SG stars of some GCs. Therefore, the progenitors may be either massive asymptotic giant branch (AGB) stars \\citep[e.g.][]{ventura2001, ventura2002}\\footnote{If the \\cite{carretta2005} CNO data really indicate that a limited number of third dredge up \\citep[e.g.][]{ibenrenzini1983} episodes plays a (small) role in the nuclear processing of the matter giving origin to the SG, the massive AGB progenitors are possibly favoured.} or fast rotating massive stars \\citep{decressin2007}.  In both cases, models show that the ejected material must be enriched in helium with respect to the pristine one. The higher helium content has been recognized to have a strong effect on horizontal branch (HB) morphology, possibly helping to explain some features (gaps, hot blue tails, second parameter) which, until now, have defied explanation \\citep{dantona2002}. Along these lines, a variety of problems has been examined: the extreme peculiarity of the HB morphology in the massive cluster NGC~2808, \\citep{dantonacaloi2004}; the second parameter effect in M~13 and M~3 \\citep{caloi-dantona2005}; the peculiar features in the RR Lyr variables and HB of NGC~6441 and NGC~6388 \\citep{caloidantona2007}. The presence of strongly enhanced helium in peculiar HB stars has been confirmed, for NGC 2808 and NGC 6441, by spectroscopic observations \\citep{moeler2004,busso2007}.  Beside this spectroscopic evidence, an unexpected feature has recently appeared from photometric data: the splitting of the main sequence in NGC 2808.  After first indications from a wider than expected colour distribution \\citep{dantona2005}, recent HST observations by \\cite{piotto2007} leave no doubt that there are at least three different populations in this cluster. This came after the first discovery of a peculiar blue main sequence in $\\omega$~Cen \\citep{bedin2004}, interpreted again in terms of a very high helium content \\citep{norris2004, piotto2005}. The above mentioned cases can be considered as ``extreme'' ones, in the sense that no explanation had been attempted for them before the hypothesis of helium-enriched populations. Less critical situations, such as the HB bimodality in NGC 1851 and 6229, had been tentatively explained in terms of a unimodal mass distribution with a large mass dispersion (0.055 -- 0.10 \\msun, Catelan et al. 1998). But even such a rather artificial assumption could not help in the case of NGC 2808, which, as hinted before, finds a possible solution only in terms of varying helium in multiple stellar generations. Therefore, we consider appropriate to apply to the less peculiar cases the solution found plausible for the most peculiar ones. In fact, notice that split main sequences and strong bimodalities are only the tip of the iceberg of the self--enrichment phenomenon. In most clusters the higher helium abundances remain confined below Y$\\sim 0.30$, and the presence of such stars would not be put in evidence either from main sequence observations \\citep{dantona2002,salaris2006}, or from a naif interpretation of stellar counts on the HB, as we shall discuss in Sect. \\ref{ngc6441}.  If we wish to shed light on the entire process of formation of GCs we must have a rough idea of the total number of SG stars. We will then analyze the HB in terms of populations differing in Y, using one or more of the following peculiar features: \\begin{enumerate} \\item bimodal HBs and HB gaps; \\item presence of blue HB stars and very long period RR Lyr's in high Z clusters; \\item peaked number vs. period distribution of RR Lyr's; \\item blue--HB clusters. \\end{enumerate} In this paper we show the results of such an interpretation for several GCs. We start from a reanalyis of the NGC~2808 data, taking into account the results by \\cite{piotto2007} for the main sequence, and the ultraviolet HST data by \\cite{castellani2006}; we summarize the results already published and discuss briefly the other clusters. The table of the derived FG and SG percentages is the basis to discuss the clusters' dynamical evolution required to produce the high fraction of stars presently belonging to the SG. ",
        "conclusions": "In all GCs examined in this work, a large fraction of the stellar population takes origin from secondary star formation episodes. Notice that we have examined only a fraction of the clusters with HB morphology or RR Lyr period distributions similar to those described here, so that we can suggest that the results of this work probably hold for a larger population of Galactic GCs. While the most massive clusters have extreme helium enhancements, also moderately massive clusters show a considerable degree of helium variation.  We reached our goals by examining in detail GCs that have unexplained features in their HBs, and extending the results to clusters with similar features. The HB morphology is one of the important features: clusters having a bimodal or multimodal HB are most easily interpreted by the coexistence of multiple generations with different helium content. Also a unique SG, whose stars had different degrees of mixing with pristine matter and thus ended up with different helium contents, is a possible solution. In any case, the extreme helium rich populations (in \\ocen\\ and NGC~2808 at least) are neatly separated from the other MS stars so that they should have a well defined independent origin.  Further, we used the period distribution of RR Lyrs in several clusters in order to reject the hypothesis of a unique Y value with a relatively large spread of mass loss on the RGB, that has been the standard way of interpreting the whole HB colour extension, but is inconsistent with most period distributions.  We re--examined in detail the HB distribution in NGC~2808, and obtained a helium distribution consistent with the main sequence recent data by \\cite{piotto2007}. Also the UV data of this cluster find a good interpretation in terms of population with varying helium content, if we make the further hypothesis of deep mixing to understand the location of the blue hook stars. We also show that simulation of M~13 UV data is well explained with populations having different helium.  After this analysis, then, we must face the problem that the SG formation is not a peculiarity of a few very massive clusters, but must be the normal way in which a GC is formed in our Galaxy. It is almost obvious, and has often been discussed in the literature, that the ejecta of a unique first stellar generation with a normal initial mass function (IMF) can not produce enough mass to give origin to such a large fraction of second generation stars (see, e.g., the case made by Bekki and Norris (2004), for the blue main sequence of $\\omega$Cen). The only solution to the IMF problem is that {\\it the starting initial mass from which the first generation is born was much larger than today first generation remnant mass} (at least a factor 10 to 20 larger), so that the processed ejecta of the first generation provide enough mass  to build up the second one. There are two possible ways of producing this result: the first possibility is that all these GCs formed within a dwarf galaxy environment \\citep{bekki2006,bekki2007}. There, GCs may be formed by mixing of pristine gas with the winds of the very numerous massive AGB stars evolving in the field of the dwarf galaxy, and later on the dwarf galaxy is dynamically destroyed. A second possibility has been recently suggested by \\cite{dercole2008}. They  show that the SG stars are preferentially born in the inner core of a FG cluster, where a cooling flow collects the gas lost by the FG stars. The massive stars that explode as SN~II were preferentially concentrated in the cluster core. After the mass loss due to the supernovae type II explosions, the cluster expands, and begins losing the stars ---mainly of FG--- going out of the tidal radius. Thus the cluster may be destroyed, unless the gas lost by the most massive AGB stars begins collecting in the core and forms the SG, that initially does not take part in the cluster expansion. The study of the cluster dynamical evolution, followed by means of N-body simulations, shows that the cluster preferentially loses FG stars; these simulations show that high SG/FG number ratio can be achieved and SG-dominated clusters may survive."
    },
    "0807/0807.1766_arXiv.txt": {
        "abstract": "We present the first multi-color view of the scattered light disk of the Herbig Ae star HD 163296, based on coronagraphic observations from the \\textit{Hubble Space Telescope} Advanced Camera for Surveys (ACS). Radial profile fits of the surface brightness along the disk's semi-major axis indicates that the disk is not continuously flared, and extends to $\\sim$540 AU.  The disk's color (V-I)=1.1 at a radial distance of 3$\\farcs$5 is redder than the observed stellar color (V-I)=0.15.  This red disk color might be indicative of either an evolution in the grain size distribution (i.e. grain growth) and/or composition, both of which would be consistent with the observed non-flared geometry of the outer disk.  We also identify a single ansa morphological structure in our F435W ACS data, which is absent from earlier epoch F606W and F814W ACS data, but corresponds to one of the two ansa observed in archival HST STIS coronagraphic data.  Following transformation to similar band-passes, we find that the scattered light disk of HD 163296 is 1 mag arcsec$^{-2}$ fainter at 3$\\farcs$5 in the STIS data than in the ACS data.  Moreover, variations are seen in (i) the visibility of the ansa(e) structures, in (ii) the relative surface brightness of the ansa(e) structures, and in (iii) the (known) intrinsic polarization of the system.  These results indicate that the scattered light from the HD 163296 disk is variable.  We speculate that the inner disk wall, which Sitko et al. suggests has a variable scale height as diagnosed by near-IR SED variability, induces variable self-shadowing of the outer disk.  We further speculate that the observed surface brightness variability of the ansa(e) structures may indicate that the inner disk wall is azimuthally asymmetric. ",
        "introduction": "\\label{intro}  Investigations of the nature and evolution of dust grains in proto-planetary systems are motivated in part by our desire to understand the birth and evolution of planetary bodies, which originate from these systems.  It is well established that Herbig Ae/Be stars \\citep{her60} are intermediate mass pre-main-sequence stars, analogous to the more familiar low mass T Tauri stars, which contain copious amounts of circumstellar gas and dust \\citep{wat98}.  A wealth of observational evidence \\citep{man97,oud99,man00, vin02,eis04} suggests that the geometry of this gas and dust takes the form of a circumstellar disk. By contrast, our understanding of fundamental properties such as grain composition and size distribution, as well as the time evolution of these parameters, are much less understood.  Analysis of ISO spectra led \\citet{mee01} to develop a 2-part evolutionary classification of Herbig systems; Group I sources were characterized as being slightly younger systems with flared outer disks, while Group II sources were characterized as slightly more evolved systems which have flatter disks owing to grain growth and settling.  While \\citet{mee01} invoked dust settling to explain the differences in flaring between Group I and II sources, it has also been suggested that the outer regions of some Herbig disks may experience self-shadowing, owing to a ``puffed-up'' inner circumstellar disk rim \\citep{nat01,dul01,dul04,ise05}. Recent models which include inflated inner disk rims have proven to be successful in explaining the near-IR and interferometric properties of several Herbig Ae stars \\citep{ise06}.  Conclusive confirmation of the geometry of the outer disk regions of Herbig Ae/Be stars inferred by SED-based studies has generally not occurred.  HD 163296 is a young (4 Myr; \\citealt{van00}), nearby (122 pc; \\citealt{van98}) Herbig Ae star.  While the star is not deeply embedded in a natal star formation dust cloud \\citep{the85}, it still displays clear signs of active accretion via the presence of jets and Herbig Haro knots \\citep{gra00,dev00,was06}.  Hence, HD 163296 appears to be in a transition phase between optically thick, extremely young pre-main-sequence stars and the much more optically thin, near zero-age-main-sequence debris disk type stars \\citep{ard04,kal05,kal06,gol06}.  Indirect techniques \\citep{bjo95} and resolved imaging \\citep{man97,gra00,ise07} have confirmed the presence of a circumstellar disk associated with HD 163296.  HST STIS white-light coronagraphic observations spatially resolved the outer regions of this disk, and detected evidence of disk structure including an inner annulus of reduced scattering and a bright outer ring or ansae of material \\citep{gra00}.  The inner region of HD 163296's circumstellar disk, as diagnosed by near-IR (0.8 - 5 $\\mu$m) SED monitoring, exhibits evidence of variability, possibly owing to changes in the inner disk wall \\citep{sit06}.  Attempts to link processes which dominate the inner regions of Herbig Ae disks to morphological features observed in the outer disk regions, i.e. determining the relative roles of self-shadowing, dust settling, and disk flaring \\citep{dul04}, requires a wealth of observational data which diagnose the inner and outer disk regions.  In this paper, we present a multi-epoch, multi-color view of the HD 163296 scattered light disk, providing a significant improvement in the documented behavior of the outer regions of HD 163296's circumstellar disk.  We describe the observational details and data reduction procedures applied to our ACS data in Section \\ref{obs-red}.  In Section \\ref{acs-charact}, we document the basic features of the HD 163296 scattered light disk in our ACS data, and compare these results to earlier epoch scattered light imaging in Section \\ref{stis}.  The possible mechanisms which might explain the observed scattered light variability are discussed in Section \\ref{origin}.  We provide a summary of our main results in Section \\ref{summary}. ",
        "conclusions": "\\label{summary}  We have presented the first multi-color, multi-epoch analysis of resolved optical scattered light imaging of the Herbig Ae star HD 163296.  To summarize the observational properties of these data: \\begin{enumerate}  \\item We spatially resolved the HD 163296 scattered light disk over radial distances of 2$\\farcs$9-4$\\farcs$4 (350-540 AU) in the HST/ACS F435W, F606W, and F814W filters.  \\item Radial profiles of the surface brightness along the semi-major axis of the disk follow a $\\sim$r$^{-4}$ power law behavior, which is indicative of a non-continuously flared disk.  \\item The (V-I) color of the disk at 3$\\farcs$5, $\\sim$1.1, is significantly redder than the stellar (V-I) color, 0.15.  This red disk color appears to be spatially uniform at the SNR of our data, and is consistent with that expected for a disk which has experienced grain growth and is at least partially self-shadowed.  \\item A single ansa structure is present in the SE disk quadrant of the 2004 epoch ACS data (F435W filter), while no structure was observed in the 2003 epoch ACS data (F606W and F814W filters), despite the fact that the limiting detection magnitude of these latter data was sufficient to probe structure at the surface brightness observed in 2004.  This morphological feature is spatially coincident with the fainter ansa reported in 1998 STIS observations \\citep{gra00} of HD 163296; the brighter ansa seen in the NW disk quadrant of the STIS data are absent from our ACS epoch data.  \\item The scattered light disk was observed to be significantly brighter in each of three filters of ACS observations from 2003-2004, as compared to white-light STIS observations in 1998.  Combining the multi-color ACS data to crudely approximate the STIS bandpass suggests that the ACS epoch data are $\\sim$1 mag arcsec$^{-2}$ brighter than the STIS data.  Along with the variability of the visibility and surface brightness of ansa structure(s) in the disk, these results \\textbf{\\textit{suggest the scattered light disk of HD 163296 is variable}}.  \\end{enumerate}  We expect that the characterization of the basic behavior of HD 163296's scattered light disk that we have provided in this study will serve as an important foundation to future efforts to model the multi-wavelength behavior of the system.  Although speculative, we suggest that one plausible explanation for the origin of the observed variability is: \\begin{enumerate}  \\item The scale height of the inner disk wall is believed to inflate and deflate on time-scales of less than a few years, based on IR SED monitoring \\citep{sit06}.  The variable scale height of the inner disk wall could induce variable self-shadowing of the outer disk, hence produce the observed overall variability of the scattered light disk.  \\item We suggest that the ansa structure(s), which appear to be discernible only during periods of enhanced self-shadowing, represent a localized region of scatterers which are at a scale height (at least marginally) above the projected shadow.  As the features are not discernible during periods of more complete illumination of the outer disk, we believe that a localized enhancement in the scale height of the disk is more likely to produce the observed phenomenon than a localized density enhancement of the scatterers (i.e. a clump).  \\item The relative surface brightness variability of the NW versus SE ansa structures during epochs in which they are visible suggests that the inner disk wall might be azimuthally asymmetric, hence produce an azimuthally asymmetric shadow on the outer disk.  Based on the short ($\\leq$1 week, Sitko 2007 personal communication) orbital time-scale of material located at the inner wall, we suggest that a series of resolved scattered light images of HD 163296, obtained during an epoch of enhanced self-shadowing, would provide a test of this particular suggested phenomenon.  \\end{enumerate}"
    },
    "0807/0807.3555_arXiv.txt": {
        "abstract": "Near-infrared photometric observations of the Type IIn SN 2005ip in NGC 2906 reveal large fluxes ($>$1.3 mJy) in the $K_s$-band over more than 900 days.  While warm dust can explain the late-time $K_s$-band emission of SN 2005ip, the nature of the dust heating source is ambiguous.  Shock heating of pre-existing dust by post-shocked gas is unlikely because the forward shock is moving too slowly to have traversed the expected dust-free cavity by the time observations first reveal the $K_s$ emission.  While an infrared light echo model correctly predicts a near-infrared luminosity plateau, heating dust to the observed temperatures of $\\sim$1400-1600 K at a relatively large distance from the supernova ($\\ga 10^{18}$~cm) requires an extraordinarily high early supernova luminosity ($\\sim1\\times10^{11}$~\\lsolar).  The evidence instead favors condensing dust in the cool, dense shell between the forward and reverse shocks.  Both the initial dust temperature and the evolutionary trend towards lower temperatures are consistent with this scenario.  We infer that radiation from the circumstellar interaction heats the dust.  While this paper includes no spectroscopic confirmation, the photometry is comparable to other SNe that do show spectroscopic evidence for dust formation.  Observations of dust formation in SNe are sparse, so these results provide a rare opportunity to consider SNe Type IIn as dust sources. ",
        "introduction": "\\label{sec_intro}  For nearly 40 years, core-collapse SN events have been considered as possible sources of dust in the universe \\citep{cernuschi67,hoyle70}.  More recent studies have proposed that the core-collapse supernovae may be the primary sources of dust in the early universe \\citep{todini01, nozawa03, dwek07}.  Several models \\citep{todini01,nozawa03,nozawa08} succeed in producing the large amounts of dust observed at high redshifts, and recent data present the first evidence for a supernova origin for dust in an object at $z>6$ \\citep{maiolino04}.  Nonetheless, direct observational evidence for dust formation in supernovae remains sparse, even in the local universe \\citep[][and references therein]{meikle07}.  Any newly formed dust would produce a late-time near-infrared excess in comparison to the blackbody optical spectrum.  \\citet{merrill80} was the first to detect such an infrared excess from a supernova (SN 1979C).  Since then, only a handful of near-infrared excesses associated with core-collapse events have been observed \\citep[e.g.][and references therein]{dwek92,gerardy02,pozzo04,meikle06,smith08a}.  Late-time near-infrared emission, however, is not unique to dust formation.  Several possible mechanisms may give rise to such an excess due to the complex nature of the surrounding environment of most core-collapse supernovae (see Figure \\ref{f1}).  (1) As the ejecta expand and cool, dust condenses and is radiatively heated.  (2) The hot gas behind the forward shock collisionally heats pre-existing circumstellar gas and dust.  (3)  The SN peak luminosity heats pre-existing dust and produces an ``IR echo.''  (4) Dust condenses in either the forward or reverse shocks and is heated by radiative shock emission.  The first scenario requires dust to condense from the ejecta as it expands and cools.  The physical conditions in an expanding SN support this possibility.  In particular, large abundances of the necessary elements, cooling of the ejecta via expansion, and dynamical instabilities can result in dust formation \\citep{lucy91,dwek92a,meikle93,roche93}.  In this scenario the primary energy source is the radioactive power released in the expanding gas, which sets the luminosity of the dust.  It is possible, however, for other energy sources to exist, such as circumstellar interactions.  The second scenario describes pre-existing circumstellar dust heated directly by shock interactions.  The scenario is more likely to occur at later times.  Assuming a spherically symmetric circumstellar distribution of dust,  \\citet{dwek83b} shows that a SN peak-luminosity of $\\sim10^{10}~$\\lsolar~creates a dust free cavity, via photo-evaporation, with a radius of $\\sim3\\times10^{17}$~cm (0.1 light years).  For typical expansion velocities ($\\sim5000$~\\kms), this cavity delays the onset of emission by a couple of years.  In addition, dust sputtering may limit the grain radiation efficiency so that only a small portion of the total shock power is converted into radiation \\citep{draine81}.  \\begin{figure*}[t] \\plotone{f1.eps} \\caption{An illustration of the environment surrounding core-collapse supernovae similar to SN 2005ip (adapted from \\citet{smith08a}).  Dust may form in a number of locations.   Following the initial supernova explosion, the SN ejecta expand and cool sufficiently for dust formation to occur.   A shocked region forms in between a forward shock in the circumstellar gas and a reverse shock in the supernova ejecta.  If the gas is able to cool efficiently, dust condensation is likely to occur in the cool, dense shell behind the shock.  Radiative cooling is more likely at the reverse shock because of the higher density, lower shock velocity, and the possibility of heavy element enrichment.  Cooling at the forward shock is less likely, but the presence of clumps in the circumstellar medium, as indicated by the relatively narrow lines observed in SNe IIn, may allow radiative shocks.} \\label{f1} \\end{figure*}  In the third scenario, the energy from the SN peak-luminosity once again creates a dust free cavity.  The remaining shell of dust warms to high temperatures but does not vaporize.  Due to light travel time effects, the thermal radiation from the dust grains reaches the observer over an extended period of time \\citep{dwek83b}.  The infrared luminosity plateau occurs on year long time scales, which are set by the light travel time across the inner edge of the cavity.  An IR echo's spectrum and temporal evolution are useful for estimating the mass-loss history of the progenitor star, and to investigate the geometry and composition of the circumstellar structure \\citep{dwek83b,dwek85,emmering88}.  In the final scenario, the ejecta collide with the pre-existing circumstellar medium, which creates a forward shock.  The circumstellar interaction decelerates the blast wave, thereby simultaneously creating a reverse shock.  The deceleration converts some of the kinetic energy of the shock into X-rays and visible light.  Dust grains are likely to condense in the cool, dense shell that forms behind the radiative shocks as they undergo a thermal instability \\citep{pozzo04}.  Typically, the post-shock conditions surrounding reverse shocks are more conducive to dust condensation as the chemically rich, dense ejecta pass through the shock front and cool.  The decelerated blast wave associated with the reverse shock is more likely to cool in sufficiently dense circumstellar environments.  Type IIn events typically have the densest environments of any SNe, as they are defined by the ``narrow'' H lines that originate from clumps in the pre-existing dense H-rich shell that is moving at a relatively slow velocity \\citep{schlegel90}.  Therefore, one might suppose dust formation in the cool, dense shell behind the reverse shock to be associated with Type IIn events.  Unfortunately, Type IIn events are rare, consisting of only $\\sim$2-3\\% of all core-collapse SNe \\citep{galyam07}, and few well-studied events exist.  This paper presents an analysis of the observed late-time near-infrared emission in the Type IIn SN 2005ip from days \\firstday-\\lastday~post-discovery.  In \\S \\ref{sec_observations}, we present the observations, data reduction techniques, and photometry.  In \\S \\ref{sec_analysis}, we discuss the different possible dust heating mechanisms and compare the SN to SNe with similar characteristics.  The observations suggest that either an IR echo or dust formation in the post-shocked gas explain the long lived IR excess, but that dust formation is more likely.  \\S \\ref{sec_conclusion} summarizes these results. ",
        "conclusions": "\\label{sec_conclusion}  This paper presented near-infrared observations of SN 2005ip for the first \\lastday~days following detection.  A large $K_s$-band luminosity persists even as the supernova's $J$-band luminosity falls.  Among a variety of potential mechanisms, dust condensation in the cool, dense shell downstream from the reverse shock is the likely source of the observed infrared emission.  We are able to rule out other mechanisms for late-time near-infrared emission.  While dust condensation in the supernova ejecta is a likely possibility, the duration of the observed near-infrared light curve is inconsistent with the quickly declining radioactive heating source.  Shock/mechanical heating of pre-existing dust grains from prior mass loss is not possible because the near infrared excess appears quickly ($<$100 days), providing insufficient time for the expanding ejecta to cross the dust-free cavity formed by the supernova.  An IR echo successfully explains the fairly uniform luminosity over the period of observation.  The duration of the infrared excess, however, implies a cavity size that requires a peak supernova luminosity that is much larger than observed for SN 2005ip and is comparable to the most luminous supernova ever observed.  On the other hand, there is mounting evidence to support dust condensation in the shock interaction region.  SN 2005ip shares many properties with SNe 1998S and 2006jc, for which spectra reveal direct evidence of dust condensation in the reverse shock \\citep{pozzo04,smith08b}.  While we currently have no such spectra for SN 2005ip, the observed infrared luminosities and temperatures are consistent with dust forming in shocked, cooled gas.  The observed X-rays indicate the infrared emission is most likely associated with the shock interaction region.  Furthermore, the initial dust temperature of $\\sim1400-1600$~K is consistent with the dust condensation temperature of carbon-rich grains, which we expect to find in the ejecta encountering the reverse shock.  The temperature evolution is also somewhat consistent with the decline predicted by an expanding cool, dense shell and diluted radiation field.  Unlike the IR echo, this scenario is physically plausible and a large initial luminosity is not required.  Of course, near-infrared photometry is only a small part of a much larger picture.  Infrared spectra can better constrain the dust temperature, infrared flux, and emission/absorption efficiency than can be obtained with near-infrared photometry alone.  Mid-infrared spectra also hold the prospects of revealing the mineralogy of the heated dust, distinguishing between carbon-rich (graphite) and oxygen-rich (silicate) grains.  Based on the dust composition and mass loss history, the nature of the progenitor system may be deduced.  Optical spectra are also of interest in that they can potentially reveal more evidence for dust formation and heating via circumstellar interaction \\citep[e.g.][]{fransson02}  The overall similarity between SN 1995N and SN 2005ip suggests that SN 1995N be considered in the context of the dust condensation model, as opposed to the echo model favored by \\citet{gerardy02}.  Indeed, \\citet{pastorello05} show that the late-time optical and infrared luminosities of SN 1995N are comparable, which suggests a correlation between the ejecta/circumstellar medium interaction and the warm dust.  The IR echo model also predicts a optical light echo from the scattered light, as well as an infrared echo produced by the absorption of the supernova peak luminosity \\citep{chevalier86}.  Yet detailed spectroscopic observations of SN 1995N suggest the late visible emission is due to circumstellar interaction and not an echo \\citep{fransson02}.  These spectra reveal evidence for absorption in the diminution of the red sides of emission lines compared to the blue sides, as would be expected with dust formation.  Furthermore, the authors present evidence for an increased helium abundance, CNO-burning products, and Ly$\\alpha$-pumped fluorescence of Fe II lines, all of which suggest a complex interaction between the ejecta and a pre-existing, dense circumstellar medium.  The dust temperature of SN 1995N on day 730, $700-800$~K, is slightly smaller than for SN 2005ip.  While this difference may be not be statistically significant, it is possible that SN 1995N had either a significantly higher expansion velocity or a different grain composition.  Our discussion of SN 1995N and the data for SN 2005ip indicate that Type IIn SNe may be sources of dust in our universe.  This association would explain the consistent late-time infrared emission observed in several Type IIn events \\citep[e.g.][]{pastorello02,gerardy02,pozzo04}.  The rarity of Type IIn events, which represent only $\\sim$2-3\\% of all core-collapse SNe \\citep{galyam07}, makes collecting complete light curves difficult.  Nonetheless, an interesting study would include follow-up observations of all Type IIn events over the past several years to determine which events are still bright in the infrared.  We also reiterate the call by \\citet{gerardy02} for simultaneous UV and X-Ray observations to better constrain $\\tau$ and the amount of shock radiation being reprocessed by the dust grains."
    },
    "0807/0807.4143_arXiv.txt": {
        "abstract": "\\noindent We investigate the possibility that a (light) hidden sector extra photon receives its mass via spontaneous symmetry breaking of a hidden sector Higgs boson, the so-called \\emph{hidden-Higgs}. The hidden-photon can mix with the ordinary photon via a gauge kinetic mixing term. The hidden-Higgs can couple to the Standard Model Higgs via a renormalizable quartic term -- sometimes called the \\emph{Higgs Portal}. We discuss the implications of this light hidden-Higgs in the context of laser polarization and light-shining-through-the-wall experiments as well as cosmological, astrophysical, and non-Newtonian force measurements. For hidden-photons receiving their mass from a hidden-Higgs we find in the small mass regime significantly stronger bounds than the bounds on massive hidden sector photons alone. ",
        "introduction": "Many extensions of the Standard Model contain so-called hidden-sectors which interact only very weakly with the {known particles from the visible sector}. Due to their feeble interactions particles in these hidden sectors can be easily missed in conventional collider experiments. Therefore, the bounds on their masses are often very weak and even masses in the sub-eV regime are possible. Such small masses, however, open the possibility that these particles may be detectable in low energy high precision experiments. Moreover, they could leave observable footprints in astrophysics and cosmology. This could therefore open a new window into particle physics which could give us crucial complementary information about the underlying laws of nature.  One interesting class of hidden sector particles is additional U(1) gauge bosons, {\\it i.e.}~hidden-photons. For example many models arising from string compactifications contain extra U(1) gauge particles under which the Standard Model particles are uncharged. Accordingly the only renormalizable interaction of the hidden-photon with the Standard Model is via mixing of the hidden-photon with the ordinary electromagnetic photon~\\cite{Okun:1982xi,Holdom:1985ag,Foot:1991kb}. Current constraints on this mixing are shown in Fig.~\\ref{bounds1}.  As can be seen from Fig.~\\ref{bounds1} the bounds depend crucially on the mass of the extra photon. In particular for very small masses the bounds become very weak. In this note we want to investigate if knowledge of the mechanism which generates a mass for the hidden-photon can improve the bounds. In principle a mass for the hidden-photon can be generated either via a Higgs mechanism or via a St\\\"uckelberg mechanism \\cite{Stueckelberg:1938}. Here, we focus mainly on the case of the Higgs mechanism.  \\begin{figure}[tb] \\begin{center} \\includegraphics[width=0.7\\linewidth]{HP.eps} \\end{center} \\vspace{-2ex} \\caption[...]{\\small Current bounds on hidden-sector photons from analyzing the magnetic fields of Jupiter and Earth~\\cite{Goldhaber:1971mr}, Coulomb law tests \\cite{Williams:1971ms,Bartlett:1988yy} (gold), electroweak precision data \\cite{Feldman:2007wj} (lightgray), searches of solar hidden-photons with the CAST experiment (purple) \\cite{Popov:1991,Popov:1999,Andriamonje:2007ew,Redondo:2008aa} and light-shining-through-walls (LSW) experiments \\cite{Cameron:1993mr,Robilliard:2007bq,Chou:2007zz,Ahlers:2007rd,Ahlers:2007qf} {(grey)} as well as CMB measurements of the effective number of neutrinos and the blackbody nature of the spectrum (black) \\cite{Mangano:2006ur,Ichikawa:2006vm,Komatsu:2008hk,Jaeckel:2008fi}. Improvements of the solar bounds can be achieved using the SuperKamiokande detector or upgrading the CAST experiment \\cite{Gninenko:2008pz}. The region $m_{\\gamma^{\\prime}}\\lesssim {\\rm few}\\,\\,{\\rm meV}$ could be tested by an experiment using microwave cavities \\cite{Jaeckel:2007ch,Penny} or experiments searching for magnetic fields leaking through a superconducting shielding~\\cite{Jaeckel:2008sz}. \\label{bounds1}} \\end{figure}  The crucial difference between the Higgs and the St\\\"uckelberg mechanism is that the gauge boson acquires a mass from the expectation value of a physical boson. As we will see in the following this additional boson will open new avenues of detection. Moreover, the additional physical boson also allows for a new possible renormalizable coupling to the Standard Model. The hidden sector Higgs can mix with the Standard Model Higgs via a quartic term~\\cite{Foot:1991bp}, sometimes called the Higgs-Portal~\\cite{portalref}.  The paper is organized as follows. In the next Sect.~\\ref{gauge} we will present the essentials of the hidden-photon hidden-Higgs system including a gauge kinetic mixing term with the photon. This will already lead us to our first main conclusion. In processes where the momentum transfer is greater than the mass of the hidden-photon the hidden-Higgs behaves essentially like a minicharged particle and corresponding (strong) astrophysical and cosmological bounds apply. Then in Sect.~\\ref{magnetic} we discuss the effects of a strong magnetic field as relevant for laser polarization and light-shining-through-walls experiments. Again we find that the bounds improve significantly. Moreover we suggest possible ways to experimentally distinguish between the Higgs and St\\\"uckelberg mechanisms. In Sect.~\\ref{Portal} we then include effects of electroweak symmetry breaking and a possible mixing of the hidden-Higgs with the Standard Model Higgs via a Higgs-Portal term. Bounds from fifth-force experiments provide interesting constraints on the Higgs-Portal term which are independent of the size of the kinetic mixing. Finally, in Sect.~\\ref{conclusions} we summarize and conclude. ",
        "conclusions": "\\begin{figure}[tb] \\begin{center} \\includegraphics[width=0.7\\linewidth]{HP_MCP.eps} \\end{center} \\vspace{-2ex} \\caption[...]{\\small Bounds on the kinetic mixing parameter for massive hidden sector photons (cf. the caption of Fig.~\\ref{bounds1}). Regions labeled in italic are the the bounds that apply if the mass arises from a Higgs mechanism, and the Higgs boson appears as a minicharged particle. We have bounds from a SLAC beamdump experiment, invisible orthopositronium decays, light-shining-through-walls experiments (LSW), big bang nucleosynthesis (BBN), and energy loss considerations in supernovae~(SN1986a), white dwarfs and red giants (see~\\cite{Davidson:1993sj,Badertscher:2006fm,Davidson:2000hf,Ahlers:2007qf}). Notice that we have assumed $g_Xq_X=e$ (so $|\\chi|=|q_\\theta|$) and $m_{\\gamma'}\\simeq m_h$. For $g_X q_X \\ll e$, typically $m_h\\gg m_{\\gamma'}$ and the new bounds move upwards and to the left.\\label{newbounds}} \\end{figure}  Extra `hidden' U(1) gauge bosons appear in many extensions of the Standard Model. The bounds on these hidden sector photons depend crucially on their mass. This mass can arise either via a St\\\"uckelberg mechanism or from a Higgs mechanism. In this paper we have investigated if one can use knowledge about the mechanism that generates the mass to improve the bounds. In particular, we have focused on the case of the Higgs mechanism. The crucial point in the case of the Higgs mechanism is that it provides an extra degree of freedom which leads to additional experimental and observational constraints. Indeed, at large momentum transfer as, {\\it e.g.}, in the interior of stars, a light hidden-Higgs behaves as a minicharged particle. A similar behavior is found inside strong magnetic fields. This can be used to translate bounds on minicharged particles into bounds on massive hidden sector photons. As can be seen from Fig.~\\ref{newbounds} this leads to a dramatic strengthening of the (astrophysical as well as laboratory) bounds for small masses.  The hidden-Higgs field $\\theta$ also provides new potential couplings to the Standard Model. In particular, it allows for a renormalizable interaction with the Standard Model Higgs $\\phi$ via a so-called Higgs-Portal term $\\kappa|\\theta|^2|\\phi|^2$. This coupling leads to fifth-force type couplings which can be used to obtain bounds on this coupling which are independent of the size of the kinetic mixing between photon and hidden-photon (cf.~Fig.~\\ref{nonnewton}). This opens the `Higgs-Portal' to the physics of light hidden sectors."
    },
    "0807/0807.3588_arXiv.txt": {
        "abstract": "We show that a consistent fit to observed secondary eclipse data for several strongly irradiated transiting planets demands a temperature inversion (stratosphere) at altitude. Such a thermal inversion significantly influences the planet/star contrast ratios at the secondary eclipse, their wavelength dependences, and, importantly, the day-night flux contrast during a planetary orbit. The presence of the thermal inversion/stratosphere seems to roughly correlate with the stellar flux at the planet. Such temperature inversions might caused by an upper-atmosphere absorber whose exact nature is still uncertain. ",
        "introduction": "Theoretical modeling of atmospheres of extrasolar giant planets (EGP) is a young field, about one decade old, and as youngsters usually are, it is very active, restless, and sometimes quite unpredictable. It had recently undergone a transition from a purely care-free stage (no observations were available) to a more difficult stage where there already are some observations to be fit by the theory.  In this paper, we briefly describe our recent efforts in this area. We will try to convince the reader that despite its young age and many shortcomings, the theory is actually doing quite well. ",
        "conclusions": "We have found (\\cite{BHBKC08}; \\cite{BBH08}) that a consistent fit to all data at secondary eclipse for several strongly irradiated transiting planets (HD 209458b, HD 149026b, and possibly HD 189733b), and very likely a non-transiting planet $\\upsilon$ And b, requires that their atmospheres have temperature inversions -- stratospheres -- at altitude. Such a thermal inversion affects: (i) planet/star contrast ratios at the secondary eclipse; (ii) their wavelength dependences; and (iii) day-night flux contrast during a planetary orbit. Moreover, the presence of the thermal inversion/stratosphere seems to roughly correlate with the total irradiated flux.  Temperature inversion is caused either by TiO/VO, as first suggested by \\cite{HBS03}, or by another, as yet unidentified, opacity sources. These may be tholins, polyacetylenes, or various non-equilibrium compounds. We invoke such extra absorbers, because a cold-trap effect can operate to deplete the upper atmosphere of TiO/VO. However, one may speculate that with ongoing mass loss and/or rotational shear instabilities the atmosphere may be partially replenished in TiO/VO. Therefore, while TiO/VO might be responsible for the formation of thermal inversions in the strongly irradiated planets, the exact nature of the absorber must be viewed as very uncertain."
    },
    "0807/0807.4682_arXiv.txt": {
        "abstract": "We consider the Friedmann-Robertson-Walker cosmologies of theories of gravity that generalise the Einstein-Hilbert action by replacing the Ricci scalar, $R$, with some function, $f(R)$.  The general asymptotic behaviour of these cosmologies is found, at both early and late times, and the effects of adding higher and lower powers of $R$ to the Einstein-Hilbert action is investigated.  The assumption that the highest powers of $R$ should dominate the Universe's early history, and that the lowest powers should dominate its future is found to be inaccurate.  The behaviour of the general solution is complicated, and while it can be the case that single powers of $R$ dominate the dynamics at late times, it can be either the higher or lower powers that do so.  It is also shown that it is often the lowest powers of $R$ that dominate at early times, when approach to a bounce or a Tolman solution are generic possibilities.  Various examples are considered, and both vacuum and perfect fluid solutions investigated. ",
        "introduction": "We study here the dynamics of Friedmann-Robertson-Walker (FRW) universes in $f(R)$ theories of gravity.  These theories are derived from generalisations of the usual Einstein-Hilbert Lagrangian of General Relativity (GR), such that \\begin{equation} \\label{gravL} \\mathcal{L} = f(R), \\end{equation} and have been considered extensively in the literature (see e.g. \\cite{early1,early2,early3,early4}).  Specification of the function $f(R)$ defines the theory, and GR can be seen to be the special case $f =R$.  Such theories have drawn considerable interest as they found success in early attempts to create a perturbatively re-normalisable quantum field theory of gravity \\cite{qgrav}, as well as turning up more recently in the effective actions of string theory \\cite{string1,string2}.  In cosmology these theories have been used extensively in attempts to explain the late-time accelerating expansion of the Universe \\cite{accel1,accel2}, cosmological inflation \\cite{infl1,infl2,infl3} and the nature of the initial singularity \\cite{singul1,singul2,singul3}. For a recent review see \\cite{review}.  In considering generalised $f(R)$ theories of gravity it is often implicitly assumed that at late times in the evolution of the Universe it should be the lowest powers of $R$ that dominate the gravitational Lagrangian.  That is, at late times we should have $R \\rightarrow 0$, and the Universe should behave as if it were governed by a gravitational Lagrangian of the form \\begin{equation} \\label{low} \\mathcal{L}_0 = \\lim_{R \\to 0} f(R), \\end{equation} which is often presumed to correspond to the Einstein-Hilbert Lagrangian, although other limits have been considered in attempts to address the apparent late-time acceleration of the Universe \\cite{accel1,accel2}.  Conversely, the introduction of higher powers of $R$ into the gravitational Lagrangian has often been assumed to mean that at early times the Universe should behave as if governed by the vacuum dynamics of a Lagrangian \\begin{equation} \\label{high} \\mathcal{L}_{\\infty} = \\lim_{R \\to \\infty} f(R). \\end{equation} The picture is then one of a universe that starts off at high $R$, dominated by a Lagrangian of the form (\\ref{high}), and that subsequently expands until $R$ becomes small and the gravitational dynamics are dominated by a Lagrangian of the form (\\ref{low}). It is the purpose of this paper to determine the veracity of such assumptions.  This is achieved by studying the dynamical evolution of FRW universes, governed by theories with general $f(R)$.  The asymptotic behaviour of the general solution to the Friedmann equations is then investigated, and used to evaluate the extent to which the afore mentioned behaviour may be considered generic.  The general solutions of FRW cosmologies governed by $f(R)$ theories of gravity have been studied previously by a number of authors, in a number of different contexts.  Much of this work has made use of the dynamical systems approach, which has been used to study specific classes of $f(R)$ in isotropic cosmologies in \\cite{Power, FRW1,FRW2}, and anisotropic cosmologies in \\cite{Bianchi1,Bianchi2,Bianchi3}.  Exact analytic expressions have been found for the general FRW solutions of some $f(R)$ theories in \\cite{exact}, and the dynamical systems approach applied to general $f(R)$ has been considered in \\cite{gen}\\footnote{See \\cite{crit} for a criticism of this work.  The present study is free from the defects pointed out in \\cite{crit}.}. For studies of spherically symmetric and weak field solutions see \\cite{review,Power,weak,ppn}, and references therein. The approach used here is a generalisation of the analysis performed in \\cite{Power}, where theories of the form $f \\propto R^n$ were considered.  We find here that the late-time attractors of FRW cosmologies for general $f(R)$ have various different forms, and that the asymptote toward which the general solution is attracted depends upon the initial conditions.  Some of these solutions correspond to the lowest powers of $f(R)$ dominating at late-times, and others to the highest powers.  Expanding universes with powers of $R$ lower than $R^2$ dominating their dynamics generically appear to exhibit the former behaviour, while universes with powers of $R$ greater than $R^2$ dominating appear to generically exhibit the latter.  Expanding universes with higher powers of $R$ dominating their early evolution therefore appear unlikely to evolve to a state where the Einstein-Hilbert term dominates.  We also find that there usually exist multiple early-time attractors for the general solution.  These can take on different forms, but generically it appears that they either evolve as $a \\sim t^{\\frac{1}{2}}$, toward a big bang singularity in the past, or that they approach a point of inflexion, where the scale factor is constant. This is in good agreement with the analytic general solutions for $f \\propto R^n$ found in \\cite{exact}.  These results do not mean that a period of inflation cannot occur (as indeed it appears to if $f \\sim R^2$), but it does mean that the general solution does not generally start off inflating  (even if $f \\sim R^2$).  It also means that the picture of the highest powers of $R$ dominating the earliest stages of the Universe's evolution may not be an accurate one.  We begin in section 2 by giving the FRW field equations for $f(R)$ theories, together with some simple power-law particular solutions that will later appear as asymptotes of the general solution.  In section 3 we use a dynamical systems approach to determine the form of the general solution for vacuum cosmologies.  The phase space of the general solution is two dimensional, and the location and stability of critical points in this space is determined.  In section 4 we perform a similar analysis for the case of perfect fluid cosmologies.  The phase space of the general solution is now three dimensional, and the location and stability of critical points is again determined.  In section 5 we consider the effect of adding higher and lower powers of $R$ to the Einstein-Hilbert action. Section 6 provides a discussion of the results, and the appendix gives some special cases that are of particular interest. ",
        "conclusions": "We have considered in this paper the evolution of spatially flat FRW universes governed by $f(R)$ theories of gravity.  The Friedmann equations (\\ref{Friedmann1})-(\\ref{conservation}) were transformed into an autonomous system of first-order differential equations, and a dynamical systems analysis was performed.  The location and stability of all critical points in the phase space were found, for both vacuum and perfect fluid cosmologies, and for general $f(R)$.  It was shown that the simple power-law solutions, given by equations (\\ref{vacuum})-(\\ref{radiation}), often act as the early and late-time asymptotes of the general solution.  The general behaviour of $f(R)$ FRW cosmologies is complicated.  The phase space of solutions is often divided into sub-spaces by invariant manifolds, through which the trajectories describing the general solution cannot pass.  The asymptotic past of general solutions can contain points of inflexion, or big-bang singularities that can be approached in different ways.  Similarly, future behaviour can be seen to be able to asymptote toward matter dominated expansion, (\\ref{matter}), vacuum domination, (\\ref{vacuum}), or various other forms.  Whatsmore, the three dimensional phase space of solutions, in the presence of a perfect fluid, allows for the possible existence of strange attractors, and chaotic behaviour.  Nevertheless, despite the complicated behaviour exhibited by the general solutions, it is still possible to make statements about the effects of modifying the Einstein-Hilbert action to more general functions of $R$.  It can be said that theories that contain lower powers of $R$ often have a stable asymptote that corresponds to the expanding vacuum dominated solution, (\\ref{vacuum}), of that lowest power, and that therefore behave as if governed by a gravitational Lagrangian of the form (\\ref{low}) at late times.  However, if a theory contains any higher powers of $R$ then there is also often a stable asymptote that corresponds to the expanding vacuum dominated solution, (\\ref{vacuum}), of that highest power, and that therefore behaves as if governed by a Lagrangian of the form (\\ref{high}) at late times. Theories containing both higher and lower powers of $R$ can then asymptote, at late times, to regimes in which either the lowest {\\it or} highest powers or $R$ dominate.  In either case the consequent evolution is that of a gravitational Lagrangian dominated by a single power of $R$.  Which solutions asymptote to high $R$ domination, and which to low $R$ domination, depends on the initial conditions, and the form of $f(R)$.  Using illustrative examples we have shown that if a power of $R$ higher than $R^2$ dominates at some point then the generic behaviour of expanding solutions is to higher $R$.  If a power of $R$ lower than $R^2$ dominates, then the trend is to lower $R$.  A term $R^2$ in the Lagrangian is then a special case; if it dominates then the expanding attractor corresponds to exponential growth, and acts as a separatrix between the higher or lower powers of $R$ dominating the future dynamics of the Universe. Those solutions with low $R$ dominating at late-times expand eternally, while those with high $R$ dominating approach either a big-rip singularity, or exponential expansion.  Big-rips often occur if there exists a single power of $R$ that dominates at late times, and exponential expansion occurs if $B \\rightarrow 1$, as is the case for some infinite power series, such as $f \\sim \\exp \\{ R \\}$.  This late-time acceleration does not appear to be a good candidate for the apparent acceleration we observe, however, as it cannot follow from a period of Einstein-Hilbert domination in which $R \\rightarrow 0$.  The asymptotic past of the general solutions is similarly complicated.  A variety of behaviours seems possible, including big-bang singularities and bounces, where the scale factor reaches a non-zero minimum.  Big-bang singularities are often approached with the scale-factor behaving as in the Tolman solution, (\\ref{radiation}).  It is interesting to note that even trajectories which undergo an early period of inflation (such as those solutions in which a power of $R^2$ dominates at some point) do not generically follow such expansion indefinitely into the past, but rather have a big bang or bounce at some point in their past, prior to the onset of inflation.  While the behaviour found here is quite complicated, with numerous different asymptotes possible, it is certainly not the most general case that one may consider.  We have limited ourselves here to spatially flat FRW universes.  Relinquishing the criterion of spatial flatness would lead to more complicated behaviours still, and one may also consider inhomogeneous and/or anisotropic cosmologies, or cosmologies with multiple fluids.  It is not clear whether or not the behaviour identified above would hold in these more general cases or not.  What does seem clear, however, is that the simple picture of the highest powers of $R$ dominating at early times, and lower powers dominating at late times, is unlikely to be accurate.  \\appendix"
    },
    "0807/0807.4820_arXiv.txt": {
        "abstract": "Data points are placed in bins when a histogram is created, but there is always a decision to be made about the number or width of the bins.  This decision is often made arbitrarily or subjectively, but it need not be.  A jackknife or leave-one-out cross-validation likelihood is defined and employed as a scalar objective function for optimization of the locations and widths of the bins.  The objective is justified as being related to the histogram's usefulness for predicting future data.  The method works for data or histograms of any dimensionality. ",
        "introduction": "There are many situations in experimental science in which one is presented with a collection of discrete measurements $\\xx_j$ and one must bin those points into a set of finite-sized bins $i$, with centers $\\XX_i$ and full-widths $\\DD_i$, to create a histogram of numbers of points $N_i$, or the equivalent when the points have non-uniform weights $w_j$.  The problem of binning comes up, for example, when one needs to plot a data histogram, when one needs to perform least-square fitting of a probability distribution function, and when one wants to compute entropies or other measurements on the inferred data probability distribution function.  The choice of bin centers and widths often seems arbitrary.  However, there is a non-arbitrary choice, derived below, which emerges when the histogram is thought of as an estimate of the probability distribution function of whatever process generated the data.  If the binning is too coarse, the histogram does not give much information about the shape of the probability distribution function.  If the binning is too fine, bins become empty and the histogram becomes noisy, so it in some sense ``overfits'' the data.  The best binning lies in between these extremes and can be found simply and quickly by a ``jackknife'' or cross-validation method, that is, by excluding data subsamples and using the non-excluded data to predict the excluded data.  This is not the only data-based binning-choice approach\\footnote{\\raggedright see, for example, Knuth,~K.~H., ``Optimal data-based binning for histograms,'' arXiv:physics/0605197, and references cited therein.}, but it is simple and sensible.  In what follows, we are going to consider a data histogram, which we imagine as a set of bins $i$, with centers $\\XX_i$ and widths (or multi-dimensional volumes) $\\DD_i$.  Equivalently (and perhaps more usefully), the parameterization of the bins can be described by a set of edges $\\XX_{(i-1/2)}$ so the centers become $\\XX_i=\\left(\\XX_{(i-1/2)}+\\XX_{(i+1/2)}\\right)/2$ and the widths become $\\DD_i=\\left|\\XX_{(i+1/2)}-\\XX_{(i-1/2)}\\right|$.  These bins will get filled by a set of (possibly multi-dimensional) data points $\\xx_j$, leading to each bin $i$ containing a number of data points $N_i$.  We will also make reference to the binning function $i(\\xx)$ which, for a given data value $\\xx$, returns the bin $i$. ",
        "conclusions": "I have shown that when a histogram of data needs to be made, there \\emph{is} a non-arbitrary choice of binning.  Some qualitative observations follow. \\begin{itemize} \\item The optimal bin widths get smaller as the number of data points goes up or as the features in the (true) probability distribution function get narrower. \\item The results are more sensitive to the smoothing parameter $\\alpha$ when the number of empty or near-empty bins becomes significant. \\item The jackknife likelihood makes discontinuous jumps as the bin edges cross individual data points.  For this reason, the likelihood does not have well-defined derivatives.  Some care must be taken that the optimization technique does not depend on having a differentiable likelihood function. \\item There is nothing special about one-dimensional or two-dimensional distribution functions; this is easily generalized to $n$-dimensional distributions.  However, it takes a lot of data points to measure a distribution function in $n$ dimensions when $n$ is large; I understand that the required number of data points scales worse than $e^n$ [need ref]. \\item There is nothing special about equal-width binning; I simply chose this to make the optimization problem easily tractable and the results easily presentable. \\item This method makes no reference to the \\emph{errors} or \\emph{uncertainties} on the measurements $\\xx_j$.  Effectively, I have assumed that the errors are small relative to any real features in the probability distribution function.  In practice, it is rarely useful to have more than a few bins per the width of your error distribution, if all the points have similar uncertainties. \\item There is often an additional choice about what minimum and maximum data values to allow for histogramming.  This choice also ought to be made in a non-arbitrary fashion if there are data points that will be excluded by the choice. \\item Finally, there is nothing special about the ``tophat'' binning model used in the above examples.  Everything can be generalized to smoothly overlapping bins, in which points are assigned fractionally to multiple bins.  In general, smoother binning models make for more well-behaved derivatives of the jackknife likelihood and therefore more straightforward optimization.  This can also all be generalized to kernel-smoothing techniques for density estimation, which ought to be made the subject of a separate note. \\end{itemize}  \\paragraph"
    },
    "0807/0807.1634_arXiv.txt": {
        "abstract": "We study direct and indirect detection possibilities of neutralino dark matter produced non-thermally by e.g. the decay of long-lived particles, as is easily implemented in the case of anomaly or mirage mediation models. In this scenario, large self-annihilation cross sections are required to account for the present dark matter abundance, and it leads to significant enhancement of the gamma-ray signature from the Galactic Center and the positron flux from the dark matter annihilation. It is found that GLAST and PAMELA will find the signal or give tight constraints on such nonthermal production scenarios of neutralino dark matter. ",
        "introduction": "While there are lots of cosmological evidence of the dark matter in the universe \\cite{Jungman:1995df,Bertone:2004pz}, its detailed properties remain largely undetermined. Requirements for the dark matter particle are the following. (1) It reproduces the present matter density of the universe. In terms of the density parameter, $\\Omega_m h^2 \\sim 0.11$ must be satisfied where $h(\\sim 0.70)$ is the Hubble parameter in units of 100~km/s/Mpc \\cite{Komatsu:2008hk}. (2) It is electrically neutral. (3) It is cold, which means that its free-streaming length ($\\lambda_{\\rm FS}$) is not so long as to seed the structure formation satisfactory, and this requires $\\lambda_{\\rm FS} \\lesssim 1$~Mpc.  In fact many candidates of dark matter are proposed in the framework of physics beyond the standard model. In particular, supersymmetry (SUSY) provides interesting candidates. If $R$-parity is conserved, the lightest SUSY particle (LSP) becomes stable and contributes present matter density of the universe. Among SUSY particles, the gravitino and (lightest) neutralino are possible candidates of the LSP dark matter. From the viewpoint of detection possibility, the gravitino dark matter is undesirable because its interaction strength with ordinary matter is Planck-suppressed.\\footnote{ Recently it is pointed out that the detection of inflationary gravitational wave background can help the situation \\cite{Nakayama:2008ip}. } In the following our focus is the neutralino dark matter, which may have distinct signatures of direct and/or indirect detection.  Usually neutralinos are assumed to be produced thermally as in the following scenario \\cite{Jungman:1995df,Kolb:1990}. In the early universe with temperature $T\\gtrsim 1$~TeV, SUSY particles including neutralinos are thermalized and their number density is given by $\\sim T^3$. As the temperature decreases, their thermal abundance receives Boltzmann suppression factor and eventually they decouple from thermal bath at the freeze-out temperature $T_f\\sim m_{\\rm L}/20$ where $m_{\\rm L}$ denotes the LSP mass. After that the number density of the LSP per comoving volume remains constant until now and hence contributes as dark matter of the universe. The resultant abundance of the LSP is estimated as \\begin{equation} Y_{\\rm L}\\equiv \\frac{n_{\\rm L}}{s} \\sim \\frac{1}{T_f M_P \\langle \\sigma v \\rangle}, \\end{equation} where $\\langle \\sigma v \\rangle$ denotes the annihilation cross section of the LSP and $M_P$ is the reduced Planck scale $(=2.4\\times 10^{18}$~GeV).  However, such a thermal relic scenario does not always hold in realistic SUSY models. For example, there often exists Polonyi or moduli field in order to break SUSY and gives rise to correct order of gaugino masses. Those singlet scalar fields generally have long lifetime and decay after freeze-out of the LSP, yielding substantial amount of LSPs. Actually Polonyi/moduli dominate the universe before they decay, and hence reheat the universe again with very low reheating temperature of $O(1)$~MeV-$O(1)$~GeV depending on their masses \\cite{Moroi:1994rs,Kawasaki:1995cy}. In this case a large amount of LSPs are produced non-thermally by the Polonyi/modulus decay, and hence large annihilation cross section is needed to account for the present dark matter abundance. Therefore, taking into account nonthermal production mechanism may significantly change the properties of the LSP and its direct/indirect detection signatures.  Thus in this paper we study direct/indirect detection signatures of non-thermally produced neutralino dark matter. As for direct detection, there are some ongoing and planned projects devoted to detect scattering signals of the LSP with nucleons such as CDMS \\cite{CDMS} and XENON \\cite{XENON}. As for indirect detection, many possible ways are proposed. First, neutralinos accumulated in the Galactic Center annihilate each other and produce line and continuum gamma rays. Such gamma-ray signals can be searched by satellite experiments (GLAST \\cite{GLAST}) or ground-based Cerenkov telescope (HESS \\cite{HESS}, MAGIC \\cite{MAGIC}, CTA \\cite{CTA}).  Second, anti-matter such as positrons or anti-protons are produced by the annihilation of the neutralinos. Since these particles are diffused by galactic magnetic fields during their propagation to the Earth, we need to solve its propagation in a diffusion model to discuss their flux on the Earth \\cite{Longair}. Fortunately, the positron flux is less sensitive to the precise diffusion model since magnetic fields easily dissipates their energy through the propagation and positrons come only from near the Earth. These anti-matter signals can be detected PAMELA \\cite{PAMELA} and AMS-02 \\cite{AMS02}.  Third, neutralinos trapped in the Sun annihilate and produce high-energy neutrinos. Super Kamiokande \\cite{Desai:2004pq}, AMANDA \\cite{Andres:1999hm}, IceCube \\cite{Ahrens:2003ix} and planned KM3NeT \\cite{KM3NeT} experiments search high energy muon signals, which arise from high-energy neutrino interaction with Earth matter. We investigate characteristic signals of nonthermal neutralino dark matter on these experiments.  For the sake of concreteness, we stick to two SUSY breaking models : minimal anomaly-mediated SUSY breaking model \\cite{Randall:1998uk} and mirage-mediation model \\cite{Choi:2004sx}. The former model predicts the wino-like neutralino LSP in broad parameter regions. Since the wino-like neutralino with a mass of ${\\cal O}(100)$ GeV has too large annihilation cross sections, its thermal abundance becomes too small to account for the present dark matter abundance. Hence we need to consider some non-thermal production processes of the neutralino dark matter. The latter model contains a heavy modulus field and non-thermal production of the neutralino dark matter is expected naturally. As we will see later, the large annihilation cross section of the neutralino dark matter is a general feature of non-thermal production scenario, and hence our results are less sensitive to the model construction.  A similar subject was studied in Ref.~\\cite{Profumo:2004ty}, where it was pointed out that the large annihilation cross section of the neutralino yields enhancement of the anti-matter signals. We emphasize that such an enhancement is rather generic feature when considering the nonthermal production scenario of the dark matter, and its detection may be directly related to the early Universe cosmology, in particular, the existence of late-decaying particles and their decay temperature. Also we have performed more detailed parameter analyses both in the anomaly-mediated SUSY breaking and mirage-mediation models, including the gamma-ray signature as well as anti-matter searches.  This paper is organized as follows. In Sec.~\\ref{sec:NT} we review some non-thermal production mechanisms of the neutralino dark matter. The minimal anomaly-mediated SUSY breaking model and the mirage-mediation mode are taken as examples. In Sec.~\\ref{sec:detection} detection possibilities of nonthermal dark matter are discussed. These include direct detection using recoil of nuclei by the neutralino, gamma-ray flux from the neutralino annihilation at the Galactic Center, positron flux from the annihilation near the Earth, and high energy neutrino flux from the annihilation in the Sun. Sec.~\\ref{sec:conclusion} is devoted to our conclusions.  For calculating these direct/indirect detection rates, we have utilized DarkSUSY code \\cite{Gondolo:2004sc}. ",
        "conclusions": "\\label{sec:conclusion}  We have discussed direct and indirect detection signatures of neutralino dark matters produced non-thermally with a very low reheating temperature $T_d\\sim {\\cal O}(1) \\rm{MeV} -{\\cal O}(1)\\rm{GeV}$. In this scenario, the self-annihilation cross section of dark matter should be large enough to account for the present relic abundance, within the rage consistent with bounds provided in \\cite{Mack:2008wu}. In SUSY models, such a large annihilation cross section is naturally realized for the neutralino dark matter with significant wino or Higgsino components. In the case of bino-like neutralino, such a large annihilation cross section can be obtained by s-channel Higgs resonance. In both cases, the large annihilation cross section leads to the enhancement of gamma-ray signals and the positron flux from the dark matter annihilation, and it becomes promising to detect the non-thermally produced dark matter with $T_\\chi\\lsim 1{\\rm GeV}$ by these indirect detection experiments. In other words, the indirect detection experiments may give us clues to explore the history of the universe with the temperature up to $1~{\\rm GeV}$.  Obviously, the consideration with other observations is important and essential to make definitive conclusion about the non-thermal production scenario. For example, it is known that the large annihilation cross section of the LSP affects the Big-Bang-Nucleosynthesis (BBN), and non-thermally produced dark matter with $T_\\chi\\lsim {\\cal O}(100) {\\rm MeV}$ may be severely constrained by the observation of $^6$Li abundance \\cite{Jedamzik:2004ip}. And also the combination with collider experiments may be most important. The upcoming Large Hadron Collider (LHC) experiments are expected to discover new particles relevant to the EWSB in the standard model, and there may appear dark matter candidates. Once such a dark matter candidate is discovered, we may have insight on its production mechanism in the universe by comparing the theoretical calculation of the cross section with cosmological and astrophysical observations, such as the dark matter abundance and its direct and indirect detection signatures. As explained in this paper, large indirect detection signals are characteristic features for the non-thermally produced dark matter, and combined with LHC experiments we may probe the nature of dark matter by these experiments.  {\\it Note added:} While finalizing this manuscript, Ref.~\\cite{Grajek:2008jb} was submitted to the preprint server, which studied similar subject to our present work. While they focus on non-thermally produced wino and higgsino like dark matter, we have also studied bino-like one in the s-channel resonance region, and performed detailed parameter analyses in MAMSB and mMSB models. The main conclusion seems to be consistent with ours."
    },
    "0807/0807.3966_arXiv.txt": {
        "abstract": "Using the Spitzer Space Telescope, we have obtained 3.6--24\\,$\\mu$m photometry of 38 radio galaxies and 24 quasars from the 3CR catalog at redshift $1<z<2.5$. This 178~MHz-selected sample is unbiased with respect to orientation and therefore suited to study orientation-dependent effects in the most powerful active galactic nuclei (AGN).  Quasar and radio galaxy subsamples matched in isotropic radio luminosity are compared. The quasars all have similar spectral energy distributions (SEDs), nearly constant in $\\nu$\\,F$_{\\nu}$ through the rest 1.6--10~\\micron\\ range, consistent with a centrally heated dust distribution which outshines the host galaxy contribution. The radio galaxy SEDs show larger dispersion, but the mean radio galaxy SED declines from rest 1.6 to 3\\,$\\mu$m and then rises from 3 to 8\\,$\\mu$m. The radio galaxies are on average a factor 3--10 less luminous in this spectral range than the quasars.  These characteristics are  consistent with composite emission from a heavily reddened AGN plus starlight from the host galaxy. The mid-infrared colors and radio to mid-infrared spectral slopes of individual galaxies are also consistent with this picture.  Individual galaxies show different amounts of extinction   and host galaxy starlight, consistent with the orientation-dependent unified scheme. ",
        "introduction": "When exploring the general evolution of galaxy populations across cosmic times, a particular challenge is to distinguish between black hole and star-forming activity. Star formation and obscuring dust go hand in hand, and black-hole-driven active galactic nuclei (AGN) are also surrounded by dust mainly distributed in a disk/torus-like geometry (Antonucci 1993). There is evidence that AGN mainly power the near- and mid-IR emission (NIR, $\\sim$2\\,$\\mu$m; MIR, $\\sim$10\\,$\\mu$m) from hot nuclear dust, while starbursts contribute mainly to the far-infrared  (FIR, $\\sim$60\\,$\\mu$m) luminosity (e.g., Rowan-Robinson 1995, Vernet et al.\\ 2001, Schweitzer et al.\\ 2006). Using the MIR/FIR luminosity ratio as an indicator for the relative AGN and starburst contributions, numerous studies have found an increase of AGN/starburst activity with total luminosity and redshift, but the validity of this trend is still under discussion because of selection effects on the various samples. More seriously, an unfavorable AGN orientation could cause MIR obscuration (e.g. Pier \\& Krolik 1993), leading to a fundamental observational degeneracy: a low MIR/FIR luminosity ratio can be due to either a high star-forming contribution or to an AGN in which the hot dust is obscured. The spectral energy distribution (SED) of an obscured AGN may thus mimic that of a starburst-powered source. While this degeneracy has now been widely examined at low/intermediate luminosity and redshift ($z<1$), it has still to be explored for the most luminous sources at high redshift ($z>1$). In order to assess galaxy and AGN evolution in the universe, we therefore need to understand this AGN/starburst degeneracy for a population of luminous high-redshift sources. A crucial step towards this is to study the orientation dependence of the NIR and MIR emission of high-redshift AGN.  Orientation-dependent effects can only be tested and quantified with AGN samples having type~1 (unobscured) and type~2 (obscured) subsamples matched in isotropic emission. The clean AGN tracers --- optical, [\\ion{O}{3}]~$\\lambda$5007\\,\\AA, NIR, and X-ray ($\\la$10~keV) --- all fail to fulfill this requirement. The [\\ion{O}{2}]~$\\lambda$3727\\AA ~emission, while isotropic (Hes et al.\\ 1993), is probably dominated by extended starbursts and shocks (Best et al.\\ 2000) rather than by the AGN. Therefore, the only feasible way is low-frequency (meter-wavelength) radio selection because the integrated emission from the radio lobes is optically thin and essentially isotropic. This makes radio-loud AGN particularly attractive for studying orientation-dependent properties at other wavelengths and, after sorting out the influence of radio jets/lobes on the emission, for generalizing conclusions about orientation-dependent effects to the much larger population of radio-quiet AGN.  The brightest low-frequency-selected AGN sample is the 3CR compilation (Spinrad et al.\\ 1985). The powerful double-lobed radio galaxies (henceforth simply called radio galaxies) are supposed to be misaligned quasars (Barthel 1989). Based on {\\it IRAS} coadded scans and a few individual detections, Heckman et al.\\ (1992, 1994) already noted an average MIR/FIR difference between 3CR quasars and radio galaxies. More comprehensive MIR and FIR spectrophotometry from {\\it ISO} is in hand (as compiled by Siebenmorgen et al.\\ 2004 and by Haas et al.\\ 2004) as well as from {\\it Spitzer} (e.g., Shi et al.\\ 2005, Haas et al.\\ 2005, Ogle et al.\\ 2006, Cleary et al.\\ 2007), providing a basis to study the $z < 1$ 3CR objects. These sources are, however, a factor of five less radio-luminous on average than the most powerful radio sources seen at higher redshift, and the lower indicated accretion power may reflect different source physics. The higher-luminosity population can be sampled by the 3CR sources at  $1<z<2.5$, which have radio luminosities similar to those of the most powerful radio sources at even higher redshift ($2.5<z<6$). However, with the exception of a few targets (Siebenmorgen et al.\\ 2004, Seymour et al.\\ 2007), the high-$z$ 3CR sample has not been well observed in the rest frame NIR and MIR.  This paper is based on {\\it Spitzer} observations of 62 of the 64 high-$z$ 3CR sources. It focuses on the hot nuclear dust emission and its obscuration in the most luminous type-2 AGN. We use a $\\Lambda$CDM cosmology with $H_0 =71$~km/s/Mpc, $\\Omega_{{\\rm m}} = 0.27$ and $\\Omega_{\\Lambda} = 0.73$.   \\begin{figure*}[htb] \\begin{center} \\includegraphics[height=14.5cm, angle=90]{f1.eps} \\caption{Infrared versus radio luminosity of the 3CR sample at $z>1$ prior to normalization.  'x' symbols denote quasars; circles and squares denote radio galaxies. Superposed crosses indicate radio galaxies that show evidence of silicate absorption (\\S\\ref{sec_rg}). The vertical long-dashed lines mark the range of our luminosity-matched  quasar and radio galaxy subsamples. The dotted lines indicate $L_{8\\micron}/L_{\\rm 178\\,MHz}$ ratios of 1, 10, and 100. The radio galaxies are grouped into several SED classes in Fig.\\,\\ref{fig_nir_mir_cc} and \\S\\ref{sec_rg}. The color-coding and symbols are:  green circle (A), red circle (B), red square (C), blue square (D), blue circle (E). The two low-excitation radio galaxies 3C\\,68.1 and 3C\\,469.1 are labeled with their 3C numbers, as are sources outside the luminosity range of our analysis. } \\label{fig_ir_radio_lum} \\end{center} \\end{figure*} ",
        "conclusions": "The 3CR sample at $1<z<2.5$ represents the most luminous steep-spectrum quasars (type~1 AGN) and powerful double-lobed radio galaxies (type~2 AGN).  This sample is nearly unbiased by orientation.  We have defined subsamples of 19 quasars and 33 radio galaxies matched in isotropic rest 178\\,MHz luminosity and have obtained {\\it Spitzer} 3.6--24\\,$\\mu$m photometry.  The main results are:  \\begin{itemize} \\item[1)] Quasars all have  similar energy distributions in the rest frame  1.6--10\\,$\\mu$m range, and their ratio of MIR to radio luminosity is also nearly constant.  This is consistent with results seen previously in lower-redshift samples.  \\item[2)]   The rest frame 1.6--10\\,$\\mu$m SEDs of radio galaxies can  be explained as reddened quasars, consistent with orientation-dependent unification. Various amounts of extinction of the AGN emission combined with addition of host galaxy starlight can explain the diversity of radio galaxy SEDs.  \\item[3)] If the extinction is sufficiently large, there is no need to invoke a beamed synchrotron contribution to explain the MIR luminosity difference between quasars and radio galaxies. The actual amount of extinction has to be derived from additional observations.  \\item[4)] The above results hold also for splitting our sample in redshift and luminosity; within our sample we do not find any trends with redshift or luminosity.  \\item[5)] At rest frame 8\\,$\\mu$m, quasars are 3 times more luminous than radio galaxies. If this difference applies also to high-redshift, radio-quiet AGN, then MIR (24\\,$\\mu$m) surveys are strongly biased in favour of type-1 and against type-2 AGN. This will make it very difficult to resolve the AGN/starburst degeneracy with only broadband SEDs, and spectral line diagnosis will be required.  \\end{itemize}  While our near-mid-IR SEDs provide a fundamental set of high luminosity AGN templates, we expect to derive more stringent conclusions from proper two-dimensional radiative transfer modelling in combination with  {\\it Spitzer} MIR spectra, {\\it Chandra} X-ray observations, and {\\it Herschel} far-IR/sub-mm data."
    },
    "0807/0807.4541_arXiv.txt": {
        "abstract": "Transiting planets provide a unique opportunity to study the atmospheres of extra-solar planets. Radiative hydrodynamical models of the atmosphere provide a crucial link between the physical characteristics of the atmosphere and the observed properties. Here I present results from 3D simulations which solve the full Navier-Stokes equations coupled to a flux-limited diffusion treatment of radiation transfer for planets with 1, 3, and 7 day periods. Variations in opacity amongst models leads to a variation in the temperature differential across the planet, while atmospheric dynamics becomes much more variable at longer orbital periods. I also present 3D radiative simulations illustrating the importance of distinguishing between optical and infrared opacities. ",
        "introduction": "The abundance of close-in planets coupled with the wealth of information that can be deduced from the fraction that transit their host stars is now allowing us to probe our understanding of the physics of atmospheres through radiative hydrodynamical modelling. A number of approaches to this problem exist within the community. The physics of radiation and hydrodynamics are both extremely important, and there are several approaches currently utilized. Below I separate the approaches to both the dynamical and radiation aspects. Although this is not an exhaustive list, I believe it encapsulates the most widely utilized approaches in the field.\\\\ \\\\ {\\bf Dynamical Methods}\\\\ \\indent Equivalent Barotropic Equations (2D)\\\\ \\indent Shallow Water Equations (3D) \\\\ \\indent Full Navier-Stokes Equations (3D)\\\\ {\\bf Radiative Transfer Methods}\\\\ \\indent Relaxation Methods \\\\ \\indent Flux-limited Radiative Diffusion \\\\ \\indent Coupled Radiation/Thermal Energy Equations \\\\ \\indent Full Wavelength Dependent Radiative Transfer \\\\ \\\\ Ideally, we would solve the full 3D Navier-Stokes coupled with wavelength dependent radiative transfer. However, such a task is well beyond our computational resources, so groups working in this area select a method from one or both lists with which to tackle the problem. Although selecting methods with many built-in assumptions allows for high resolution studies, one runs the risk of missing fundamental physics crucial for a full understanding of the atmospheres. On the flip side, as with all numerical studies, the more exact your method, the more difficult and longer it takes to simulate. For a comparison of methods please see \\cite[Dobbs-Dixon \\& Lin (2008)]{dobbsdixon2008} or \\cite[Showman et al (2008)]{showman2008}.  In the following proceedings I will describe results from simulations utilizing the full 3D Navier-Stokes equations (described in Section (2)). In Section (3), I couple these with the flux-limited radiative transfer in simulations which solve a single equation for the energy of the radiation and gas components. These simulations exhibit a range of behavior for different opacities and orbital periods. I show that the local value of opacity is crucial in determining the efficiency of heat re-distribution to the night-side. In Section (4) I present results of simulations again solving the full Navier-Stokes equations, but now solving separate energy equations for the radiation and gas components. Decoupling these equations allows for the introduction of multiple opacities corresponding to both incident stellar and local radiation. As I show, this distinction is crucial in reproducing the observed upper atmosphere temperature inversions. I conclude in Section (5) with a discussion of the results. ",
        "conclusions": "In this proceeding I have reviewed results from full three-dimensional radiative hydrodynamical simulations of the atmospheres of close-in planets. Simulations for a range of atmospheric opacities and semi-major axis were performed. The results indicate that both stellar proximity and opacity are both crucial factors in determining the behavior of planetary atmospheres. Increasing planetary opacities yields less efficient advection and colder night-side temperatures. Furthermore, planets situated at 1 and 3 day orbital periods exhibit a fairly stable banded jet structure, with jet width increasing with rotation period. On the other-hand, dynamics of planets with 7 day orbital periods can be much more variable, with wandering jets and the formation of vortices on the night-side.  The decoupling of the thermal and radiative energies, as well as the introduction of multiple opacities relevant for impinging optical stellar photons vrs local re-radiated infrared photons allows us to explore in greater detail temperature structure near the top of the atmosphere. Radiative solutions presented in Section (4) produce temperature inversions similar to those observed and demonstrate the feasibility of simultaneously studying these radiative features and the resulting hydrodynamics. The existence of an inversion depends critically on the ratio of the optical to infrared opacities and the behavior of this ratio with radius.  Given the importance of both optical and infrared opacities it is fruitful to consider how we expect atmospheric properties to vary as we change these two quantities independently. \\cite[Fortney et al. (2008)]{fortney2008} suggested a classification planets based on the presence or absence of TiVO in the upper atmosphere. Although the presence of TiVO in the upper atmosphere is controversial, it would provide the necessary absorption opacity to cause an inversion. Fortney's classifications of pM and pL classes correspond to those with and without detected inversions. Here extend this classification in terms of the individual absorption ($\\kapa$) and Planck ($\\kapp$)opacities explored in this proceeding.  As explained in Section (3), the night-side temperature is governed by the efficiency of advection, which in turn is determined by the cooling timescale where the incident radiation is deposited. The cooling timescale is directly related to the local Planck opacity $\\kapp$; large $\\kapp$ will result in a large day-night temperature differential. Such a distinction is observable from current and ongoing phase monitoring programs. On the other hand, the absorption opacity for the incident stellar light $\\kapa$ will play an important role in the presence or absence of a temperature inversion. As described in Section (4) the ratio $\\gamma=( \\frac{\\kapa}{\\kapp} ) ^{1/4}$ will determine the temperature structure in the upper atmosphere. If the radial structure of $\\gamma e^{-\\tau_{\\star}}$ is such that it peaks in the upper atmosphere there will be an inversion. In these terms, the presence of TiVO (a pL planet) would lead to large $\\kapa$, but an inversion would occur only if $\\kapp$ does not also increase. In Figure (\\ref{fig:diagram}) I have unified these two concepts onto a single diagram. A measurement of both the magnitude of day-night temperature differential and a determination of the presence or absence of an inversion will uniquely tell you both the optical and infrared opacities and place the planet in one of the quadrants. Such a determination will ultimately tell us much about the formation and evolution of these planets.  \\begin{figure} \\centering \\resizebox{8.0cm}{!}{\\includegraphics{./figs/diagram.ps}} \\caption[]{The gross behavior of irradiated atmospheres as you change the absorption and local Planck opacities. Increasing optical absorption opacity ($\\kapa$) is given in the horizontal axis, while the vertical axis gives the local opacity to re-radiated photons ($\\kapp$). The diagonal line demarcates the division between planets with an inversion ($\\gamma>1$) and those without ($\\gamma<1$), where $\\gamma^4=\\kapa/\\kapp$. Phase measurements coupled to inversion determinations will uniquely place you in a given quadrant.} \\label{fig:diagram} \\end{figure}"
    },
    "0807/0807.4294_arXiv.txt": {
        "abstract": "TeV $\\gamma$ rays from distant astrophysical sources are attenuated due to electron-positron pair creation by interacting with ultraviolet/optical to infrared photons which fill the universe and are collectively known as the extra-galactic background light (EBL).  We model the $\\sim$0.1--10~eV starlight component of the EBL derived from expressions for the stellar initial mass function, star formation history of the universe, and wavelength-dependent absorption of a large sample of galaxies in the local universe.  These models are simultaneously fitted to the EBL data as well as to the data on the stellar luminosity density in our local universe.  We find that the models with modified Salpeter A initial mass function together with Cole et al. (2001) or Hopkins \\& Beacom (2006) star formation history best represent available data.  Since no dust emission is included, our calculated EBL models can be interpreted as the lower limits in the $\\sim$0.1--1~eV range.  We present simple analytic fits to the best-fit EBL model evolving with redshift.  We then proceed to calculate $\\gamma$-ray opacities, and absorption of $\\sim$10--300~GeV $\\gamma$-rays coming from different redshifts.  We discuss implications of our results for the Fermi Gamma Ray Space Telescope and ground-based Air Cherenkov Telescopes. ",
        "introduction": "Stars are the dominant sources of electromagnetic radiation in the universe after the cosmic microwave background (see, e.g., Fukugita \\& Peebles 2004).  They emit radiation longward from ultraviolet to infrared wavelengths.  However, photons with wavelength $\\lesssim 2~\\mu$m are highly absorbed by the dust in the host galaxies and only a fraction of the radiation emitted by the stars escape to the inter-galactic medium and form a diffuse background or EBL (see, e.g., Baldry \\& Glazebrook 2003; Driver et al. 2008).  The dust in the host galaxies, heated by the starlight, also radiate in the infrared wavelengths and contribute to the EBL density at $\\sim 10^{12}$ Hz. It is the direct starlight component, $\\lesssim 2~\\mu$m or $\\gtrsim 0.1$~eV, that affects the propagation of $\\lesssim 5$~TeV $\\gamma$-rays from distant sources.  Indeed, the very soft spectral energy distribution ($dN/dE \\propto E^{-\\Gamma}$) with $\\Gamma \\gtrsim 3$ observed from several TeV blazars at high redshift ($z\\gtrsim 0.1$) such as PKS 2155-304 (Aharonian et al. 2005), H 2356-309 (Aharonian et al. 2006a); 1ES 1218+304 (Albert et al. 2006); 1ES 1101-232 (Aharonian et al. 2006b); 0347-121 (Aharonian et al. 2007), 1ES 1011+496 (Albert et al. 2007) and 3C 279 (Albert et al. 2008), and their cutoff at $\\gtrsim 1$~TeV are hints that high energy $\\gamma$ rays from these sources are absorbed by the EBL UV/optical photons (Persic \\& de Angelis 2008).  Lower energy ($<$TeV) $\\gamma$-rays from high redshift sources such as gamma-ray bursts (GRBs) and blazars can also probe the EBL starlight component.  Calculation of the opacity of the universe to $\\gamma$-rays by $\\gamma\\gamma \\to e^+e^-$ process dates back to Nishikov (1961), followed by Gould \\& Shr\\'eder (1966) and Fazio \\& Stecker (1970). More recently Malkan \\& Stecker (1998, 2001); Primack et al. (1999); Kneiske, Mannheim \\& Hartmann (2002); Kneiske et al. (2004); Primack, Bullock and Somerville (2005); Stecker, Malkan \\& Scully (2006) calculated EBL models adopting either a phenomenological approach or Monte Carlo galaxy formation code.  These models trace the general trend of the data, which may be fitted with a combination of two or more modified blackbody spectra for its two distinct peaks at the infrared and optical wavebands (Dermer 2007).  Significant uncertainty in data and large dispersion among models led to an indirect method to constrain the EBL, namely by estimating change in spectral slope from distant TeV blazars due to $\\gamma\\gamma$ absorption (Stecker \\& de Jager 1993; Stanev \\& Franceschini 1998; Mazin \\& Raue 2007). However, such a method generally does not include possible absorption at the source (see, e.g., Reimer 2007) and presumes a source spectrum.  In this paper, we build models of the EBL starlight component ($\\sim$0.1--10~eV) directly from the stellar thermal surface radiation.  Emission from an individual star during its main-sequence lifetime is well approximated as a blackbody with a mass-dependent temperature.  The post-main-sequence lifetime of a star is very short compared to its main-sequence lifetime and their contribution at the UV-optical wave bands is not significant.  They can, however, contribute significantly to longer wavelengths due to their increased luminosity in the post-main-sequence phase (Finke et al., in preparation).  Only a small fraction of the stars with mass $\\gtrsim 8M_\\odot$ produce supernovae and even a smaller fraction produce GRBs. Emission from these sources dominate the diffuse MeV background (Watanabe et al. 1999; Ruiz-Lapuente, Cass\\'e \\& Vangioni-Flam 2001; see, however, Strigari et al. 2005; Inoue, Totani \\& Ueda 2008). Emission from quasars and AGNs, on the other hand, dominate the diffuse X-ray background (Mushotzky et al. 2000).  Contributions from all these sources add only a small fraction to the total cosmic electromagnetic energy density in the $\\sim$ 0.1--10~eV range, and we also ignore that.  The estimated lifetimes of individual stars depend on their masses and the assumed cosmology, which is the standard $\\Lambda$CDM with ($h, ~\\Omega_m, ~\\Omega_\\Lambda$) = (0.7,~0.3,~0.7), and the Hubble constant $H_0 = 70 h_{0.7}$~km~s$^{-1}$~Mpc$^{-1}$. Summing over contributions from stars of all masses formed in the history of the universe then gives us the diffuse emission or EBL.  A sum over contributions from individual stars radiating at a given redshift corresponds to the luminosity density or the stellar energy emissivity of the universe at that redshift.  The initial mass function (IMF), which is the distribution of stars by mass, and the star formation rate (SFR), which is the mass that forms stars per unit comoving volume per unit time, are two uncertain but related parameters in our calculation.  We form classes of models by choosing different combinations of these parameters and compare, in the UV-optical band, the luminosity density data of the local universe found from the surveys of nearby galaxies.  The same models are then compared with EBL data.  Note that there are no adjustable free parameters in our calculation once we choose a particular model.  Finally we use one of our best-fit models to calculate the $e^\\pm$ pair production opacities in the $\\sim$10--300~GeV energy range at different redshifts.  These results are applicable to high-energy emission from distant sources such as GRBs and blazars detected by the currently operating Fermi Gamma Ray Space Telescope and Air Cherenkov Telescopes such as HESS, MAGIC and VERITAS.  In Sec. \\ref{sec:formalism} we outline the formalism of our method and introduce different models in Sec. \\ref{sec:SFR_models} upon which we base our EBL calculation.  We report our results in Sec. \\ref{sec:results}, compare with results from previous authors as well as calculate EBL evolution with redshift.  We discuss implications of our results for $\\gamma$-ray astronomy in Sec. \\ref{sec:implications}, calculating $\\gamma\\gamma \\to e^+e^-$ absorption opacity ($\\tau_{\\gamma\\gamma}$) for high energy $\\gamma$ rays from sources at different redshifts.  Conclusions of the work for EBL and $\\gamma$ ray absorption are given in Sec. \\ref{sec:conclusions}. ",
        "conclusions": "\\label{sec:conclusions}  We have derived a class of well-defined models for the spectral energy density of the EBL.  The models with modified Salpeter A initial mass function, a single power-law mass-luminosity relation and Cole et al. (2001) or Hopkins \\& Beacom (2006) star formation history reasonably fit the EBL UV-optical data and the luminosity density in our local universe.  Our models are based on the underlying assumption that the bulk of the EBL radiation between $\\sim 0.1$ -- 10 eV is due to stellar radiations absorbed by dust, which can be determined from recent analyses of galaxies by Driver et al.\\ (2008).  This approach differs from models by Stecker, Malkan \\& Scully (2006) based on luminosity evolution of galactic spectral energy distributions, which is limited to the accuracy of the available observational data used in the survey. Without need of a population synthesis code (e.g., models by Primack, Bullock \\& Somerville 2005) or fits to the results of such a code (e.g., models by Kneiske, Mannheim \\& Hartmann 2002), our model is based on well-studied results from stellar astronomy.  The sources of uncertainties in our models are the (i) main-sequence age of the star and stellar luminosity, discussed earlier, (ii) star formation history and initial mass function, and (iii) dust absorption.  We have already discussed point (ii) in some details using five SFR+IMF models.  Note that in all those models, the IMF was assumed to be independent of redshift.  In principle the normalization of the IMF or even its shape may depend on $z$.  Nevertheless, a universal IMF fits SFR data reasonably well.  The evolution of the dust absorption model with redshift, which we have not taken into account, is of potentially greater concern.  At high redshift it is more reasonable to assume that the dust absorption (escape) fraction would be higher (lower), so that the stellar contribution to the EBL at high $z$ would be less than if using a constant absorption fraction.  A more detailed examination of these issues are under further study (Finke et al., in perperation).  We have provided an analytic fit to our best-fit EBL model and its evolution with redshift.  This result can be used to calculate, as we have done in this work, opacity of the universe to $\\sim$10--300~GeV $\\gamma$-rays relevant for the high energy data from the {\\em Fermi Gamma Ray Space Telescope} and Air Cherenkov Telescopes, and estimating unknown quantities such as the spectrum and energy at production of the distant GRBs and blazars such as 3C 279."
    },
    "0807/0807.3820_arXiv.txt": {
        "abstract": "A detailed analysis of new and existing photometric, spectroscopic and spatial distribution data of the eccentric binary V731~Cep was performed. Spectroscopic orbital elements of the system were obtained by means of cross-correlation technique. According to the solution of radial velocities with {\\em UBVR$_c$} and {\\em I$_c$} light curves, V731~Cep consists of two main-sequence stars with masses M$_{1}$=2.577 (0.098) M$_{\\odot}$, M$_{2}$=2.017 (0.084) M$_{\\odot}$, radii R$_{1}$=1.823 (0.030) R$_{\\odot}$, R$_{2}$=1.717 (0.025) R$_{\\odot}$, and temperatures T$_{eff1}$=10700 (200) K, T$_{eff2}$=9265 (220) K separated from each other by \\textit{a}=23.27 (0.29) R$_{\\odot}$ in an orbit with inclination of 88$^{\\circ}$.70 (0.03). Analysis of the O--C residuals yielded a rather long apsidal motion period of $U$$=$10000(2500) yr compared to the observational history of the system. The relativistic contribution to the observed rates of apsidal motion for V731~Cep is significant (76 per cent). The combination of the absolute dimensions and the apsidal motion properties of the system yielded consistent observed internal structure parameter (log$\\bar{k}_{2,obs}$ = $-$2.36) compared to the theory (log$\\bar{k}_{2,theo}$ = $-$2.32). Evolutionary investigation of the binary by two methods (Bayesian and evolutionary tracks) shows that the system is $t=$133(26) Myr old and has a metallicity of $[M/H]=-0.04(0.02)$ dex. The similarities in the spatial distribution and evolutionary properties of V731~Cep with the nearby ($\\rho\\sim$3$^{\\circ}$.9) open cluster NGC 7762 suggests that V731~Cep could have been evaporated from NGC~7762. ",
        "introduction": "Study of eclipsing binaries is still the most effective way of determining the absolute parameters of stars, especially from the spectroscopic and photometric analysis of detached double-lined eclipsing binaries, masses and radii of the components can be obtained with a precision of $\\sim$1 per cent (e.g. Southworth et al. 2005; Bak\\i\\c{s} et al. 2008), the limit precision for stellar evolutionary tests (Andersen 1991). Among the stars, those in the upper main sequence band are particularly useful in empirical tests of the convection formulae used in various evolution codes, while systems containing unevolved stars are useful in testing opacity and metallicity effects in near-ZAMS (Zero Age Main Sequence) models. V731~Cep, the binary discussed in the present paper, is of the latter type.  The variability of V731~Cep (GSC~4288~0168; brightness at maximum V$\\sim$10.5 mag; orbital period $P$$\\sim$6.06 d) was discovered by Bak\\i\\c{s} et al. (2003) (hereafter B03). A brief history of V731~Cep was given by Bak\\i\\c{s} et al. (2007) (hereafter B07) who presented photometric light curves in {\\em BVR$_c$} and {\\em I$_c$} bands and limited spectroscopic observations. Within the scope of a project to study close eclipsing binaries of SB2 type, we were able to obtain new times of minima as well as new spectroscopic data for V731~Cep.  The system has an eccentric orbit ($e=0.0165$), which makes it an important astrophysical tool for the investigation of internal structures of its components, if the rotation period of the apsides is precisely determined. The apsidal motion period together with system geometry allow the computation of the observed internal structure constant (ISC) to be compared with theoretical internal structure computations. The primary motivation for the present paper is to obtain the parameters of the close binary system V731~Cep, to discuss the evolutionary status of the system using the system parameters with the latest evolutionary models, and to investigate the apsidal motion of the orbit for estimation of the internal structures of the component stars. ",
        "conclusions": "\\subsection{Absolute Dimensions and Distance of the System}  Combination of spectroscopic orbital elements (Table~6) with light curve elements (Table~7) yields the absolute dimensions of the system, which are presented in Table~8. The adopted temperature T$_{eff1}$$=$10700 K and mass $M_{1}$$=$2.577$M_{\\odot}$ of the primary component correspond to the spectral type of a normal B8.5-type main sequence star, while the adopted temperature T$_{eff2}$$=$9265 K and mass $M_{2}$$=$2.017$M_{\\odot}$ of the secondary component are in good agreement with the spectral type of a normal A1.5-type main sequence star (i.e. Strai\\v{z}ys \\& Kuriliene 1981). However, the radii of the components, $R_{1}$$=$1.823$R_{\\odot}$ and $R_{2}$$=$1.717$R_{\\odot}$, are in better agreement with the same spectral type stars but at closer locations to ZAMS, suggesting a young age of the system, as determined in \\S6.2.  The synchronization time-scale for the component stars of the V731~Cep system is, following Zahn (1977), in the order of 15 Myr, which is smaller than the age of 120 Myr estimated from the isochrones (see \\S6.2). To compare the observed rotational velocities with the synchronization velocities listed in Table~8, we have modelled Si II doublets (6347.091 \\AA, 6371.359 \\AA) with model atmosphere grids using ATLAS9 and SYNTHE codes under Linux (Kurucz 1993). The modeling yielded equatorial rotational velocities of V$_{rot1}$$=$19(3) km s$^{-1}$ and V$_{rot2}$$=$18(3) km s$^{-1}$ for the primary and secondary components, respectively. Although errors in the observed rotational velocities are in the order of 15 per cent, the synchronization velocities seem to be slightly below the observational measurements. The asynchronization of the components with the orbit should be confirmed by analyzing new spectra with a higher S/N ratio and more absorption lines in a larger spectral range.  \\begin{table*} \\small \\caption{Close binary stellar parameters of V731~Cep. Errors of parameters are given in parenthesis.} \\label{table8} \\begin{tabular}{lccc}\\hline\\hline Parameter                          & Symbol  & Primary & Secondary  \\\\ \\hline Mass (M$_\\odot$)                   & \\emph{M}       & 2.577(0.098) & 2.017(0.084)    \\\\ Radius (R$_\\odot$)                 & \\emph{R}       & 1.823(0.030) & 1.717(0.025)    \\\\ Separation (R$_\\odot$)             & \\emph{a}     & \\multicolumn{2}{c}{23.27(0.29)} \\\\ Surface gravity (cgs)              & log $g$        &4.304(0.011)& 4.273(0.011)   \\\\ Integrated visual magnitude (mag)  &  \\emph{V}      &\\multicolumn{2}{c}{10.54(0.01)}\\\\ Integrated colour index (mag)       & $B-V$          &\\multicolumn{2}{c}{0.09(0.01)}\\\\ Colour excess (mag)                 &$E(B-V)$&\\multicolumn{2}{c}{0.13(0.03)}\\\\ Visual absorption (mag)            & $A_{V}$ &\\multicolumn{2}{c}{0.40}\\\\ Intrinsic colour index (mag)        & $(B-V)_{0}$&\\multicolumn{2}{c}{-0.04(0.02)}\\\\ Component intrinsic colour index (mag) & $(B-V)$ & -0.073(0.020) & 0.016(0.020)  \\\\ Temperature (K)                    & $T_{eff}$ & 10700(200) & 9265(220) \\\\ Spectral type                      & Sp        & B8.5 V   & A1.5 V        \\\\ Luminosity (L$_\\odot$)             & log \\emph{L}& 1.618(0.035) & 1.292(0.043)\\\\ Computed synchronization velocities (km s$^{-1}$)& V$_{synch}$ & 15.6(0.2) & 14.3(0.2) \\\\ Observed rotational velocities (km s$^{-1}$) & V$_{rot}$ & 19(3) & 18(3) \\\\ Bolometric magnitude (mag)         &$M_{bol}$& 0.705(0.088) & 1.519(0.108)    \\\\ Velocity amplitudes (km s$^{-1}$)  &$K_{1,2}$& 85.18(1.72) & 108.84(1.73)    \\\\ Absolute visual magnitude (mag)    &$M_{v}$  & 1.104(0.054) & 1.631(0.070)    \\\\ Bolometric correction (mag)        &\\emph{BC}& -0.399      & -0.112      \\\\ Distance (pc)                      &\\emph{d} & \\multicolumn{2}{c}{809(30)} \\\\ Systemic velocity (km s$^{-1}$)    &$V_{\\gamma}$ & \\multicolumn{2}{c}{0.62(0.94)} \\\\ Parallax (mas)                     &$\\pi$        & \\multicolumn{2}{c}{1.236(0.044)*} \\\\ Proper motion (mas yr$^{-1}$)      &$\\mu_\\alpha cos\\delta$, $\\mu_\\delta$ & \\multicolumn{2}{c}{-1.5 (2.6), -3.1 (2.5)**} \\\\ Space velocities (km s$^{-1}$)     & $U, V, W$  & \\multicolumn{2}{c}{7.59(9.03), 4.71(4.31), -9.75(9.62)}\\\\ \\hline {\\em * In this study.} \\\\ {\\em ** NOMAD Catalog (Zacharias 2005).} \\end{tabular} \\end{table*}   Using the brightness of the system listed in Table~1 together with the light contributions of the components listed in Table~7, the intrinsic magnitudes of the components were calculated and are presented in Table~9. During the derivation of the intrinsic magnitudes, the interstellar extinction in {\\em B} and {\\em V} bands were adopted from the $Q$-method while the following relations of Fiorucci \\& Munari (2003) were used for the determination of extinction in $R_c$ and $I_c$.  \\begin{eqnarray} \\nonumber (R_{c})_0 = R_c - 2.494 \\times E(B-V), \\\\ (I_{c})_0 = I_c - 1.753\\times E(B-V). \\end{eqnarray}  \\noindent where ($R{_c})_0$ and ($I{_c})_0$ are the de-reddened magnitudes.   \\begin{table} \\begin{center} \\caption{De-reddened magnitudes of stars in V731~Cep system.} \\label{table9} \\begin{tabular}{lccccc}\\hline\\hline & {\\em B}     & {\\em V}     & {\\em R$_c$} & {\\em I$_c$} & Err.\\\\ \\hline {\\em Primary}   & 10.59 & 10.66 & 10.67 & 10.74 & 0.02 \\\\ {\\em Secondary} & 11.21 & 11.19 & 11.17 & 11.20 & 0.02 \\\\ \\hline \\end{tabular} \\end{center} \\end{table}   The de-reddened visual apparent magnitude and optical absolute magnitude presented in Table~8 allowed us to derive a distance of 809(30) pc to the system. To compare the distance of V731~Cep system using a different method, a luminosity-colour relation (Bilir et al. 2008) which has been formed for detached binary systems with main-sequence components was used in this study. The near-infrared magnitudes of the system were taken from the {\\em 2MASS} Point Sources Catalogue of Cutri et al. (2003) and are shown in Table~1. For de-reddenig near-infrared magnitude and colours of the system, the following formulae (Bilir, G\\\"uver, Aslan 2006; Ak et al. 2007; Bilir et al. 2008) were used:  \\begin{eqnarray} \\nonumber J_{o}=J-0.884\\times E(B-V), \\\\ (J-H)_{o}=(J-H)-0.322\\times E(B-V), \\\\ \\nonumber (H-K_{s})_{o}=(H-K_{s})-0.187\\times E(B-V). \\end{eqnarray}  All the colours and magnitude with subscript ``0'' show the de-reddened ones. The colour excess $E(B-V)=0.13$  was estimated in a direction to V731~Cep by using {\\em Q}-method (see \\S5.1). The near-infrared absolute magnitude of V731~Cep system was estimated by the luminosity-colour relation, $M_{J}=5.228(J-H)_{o}+6.185(H-K_{s})_{0}+0.608$, of Bilir et al. (2008) and the distance of the system calculated as $733(50)$ pc by using the photometric parallaxes method. The photometric distance of $809(30)$ pc given in Table~8 is consistent with the $733(50)$ pc distance estimated by luminosity-colour relation obtained for detached binary systems.    \\subsection{Internal Structure}  Binary systems with apsidal motion allow us to determine the ISC, which is an important parameter of stellar evolution models. The observed apsidal motion period of $\\textit{U}$ = 10000(2500) yr, corresponding to a total rate of $\\dot{\\omega}$ = 0.00060(0.00015) $^\\circ$ cycle$^{-1}$ was obtained in \\S6. The relativistic contribution to the apsidal motion in case of V731~Cep is substantial $\\dot{\\omega}_{rel}$ = 0.00045 $^\\circ$ cycle$^{-1}$, or about 75 per cent of the total observed rate (Gim\\'enez 1985). After correcting for this effect, an average ISC was derived to be log $\\bar{k}_{2,obs}$ = $-$2.36 under the assumption that the component stars rotate pseudosynchronously. This value is in very good agreement with the theoretical prediction of log $\\bar{k}_{2,theo}$ = $-$2.34 according to new evolutionary models of Claret (2004) with the standard chemical composition of (X,Z) = (0.70, 0.02). It should, however, be noted that the present apsidal motion solution is still tentative due to relatively short observational history of V731 Cep compared to the apsidal motion period. Therefore, accurate eclipse timings are strongly needed in a decade or more in order to say more definite on the apsidal motion parameters and the related ISC.   \\subsection{Evolutionary Stage and Age of the System}  We investigated the evolutionary status of the system by means of the Bayesian method and constructing an H-R diagram for the component masses in log T$_{eff}$-log L plane. We used a slightly modified version of the Bayesian estimation method idealized by J{\\o}rgensen \\& Lindegren (2005), which is designed to avoid statistical biases and to take error estimates of all observed quantities into consideration. Estimation of age and metal abundance of the components was made by using the web interface\\footnote{http://stev.oapd.inaf.it/$\\sim$lgirardi/cgi-bin/param} based on the Bayesian method of da Silva et al. (2006). Including their errors, the effective temperatures, visual brightness, metal abundance of components and distance to the system were the parameters used in the web interface to obtain the best matching model parameters (i.e. surface gravity (log $g$), radii ($R_{1,2}$), masses ($M_{1,2}$) and age of components ($t_{1,2}$)) to Padova isochrones by the Bayesian method. Since the metal abundance of the components was initially not known, a range of metal abundance (i.e. $-$0.10 $<$ $[M/H]$ $<$ +0.10 dex) was selected and the values in this range were used with 0.02 dex steps. The output model parameters for each component were compared by means of $\\chi^2$ test with the absolute dimensions of the components listed in Table~8. Consequently, the minimum $\\chi^2$ yielded simultaneously the metal abundance and age of the components as $[M/H]_1=-0.06(0.02)$ dex and $t_{1}=116(15)$ Myr for the primary and $[M/H]_2=-0.02(0.02)$ dex and $t_{2}=150(15)$ Myr for the secondary component. From these values, we adopted the mean metal abundance and mean age of the system to be $[M/H]=-0.04(0.02)$ dex and $t=133(26)$ Myr, respectively.  Interpretation of the evolutionary status of V731~Cep requires also the construction of H-R diagrams using the latest theoretical evolutionary models. In Fig.~8, the components of V731~Cep are shown in the log T$_{eff}$-log $L$ plane together with Yonsei-Yale (Y2) evolutionary tracks (Yi, Demarque, Kim et al. 2001; hereafter YDK) for different masses. Evolutionary tracks for the exact masses of the components were also computed using the code provided by YDK. Among the tracks computed for the exact masses with their errors, those with $[M/H]$$=$-0.024 dex ($Z$$=$0.0172) metallicity models match the locations of the components within the error limits in the log T$_{eff}$-log $L$ plane. We also computed a set of isochrones using a metallicity of $[M/H]$$=$ $-$0.024 dex in Y2 models. In Fig.~8, two isochrones ($t=$100 Myr and $t=$120 Myr) are plotted for comparison. It was found that 120 Myr age is the best fitting isochrone to the locations of both components. Although, the most reliable method of metal abundance determination is the atmosphere modeling of spectral lines, in the present study, it seems fair to conclude that the two methods (Bayesian methods and evolutionary tracks) used for estimation of the metallicity and age of the system are in excellent agreement within the error limits. The metallicity we found for the components should be confirmed with the atmosphere modeling of metallic absorption lines in spectra taken in wider wavelength range.   \\begin{figure} \\centering \\resizebox{90mm}{!} {\\includegraphics[]{fig8.eps}} \\caption{Evolutionary tracks for individual component masses and isochrone curves best matching the location of the components in log T$_{eff}$ - log $L$ plane. The primary and secondary stars are shown with filled and empty circles, respectively.} \\label{fig8} \\end{figure}    \\subsection{Possible Origin of V731~Cep}  One of the formation regions of early-type stars is open clusters. To find a possible formation region for V731~Cep, nearby young open clusters were investigated. Among others, NGC~7762 was found to be the closest open cluster and the most similar in chemical composition to V731~Cep. Chincarini (1966) estimated the distance of NGC~7762 to be 750 pc and the age 266 Myr. In a more recent study relating on the cluster, Patat \\& Carraro (1995) proposed a similar distance of 800 pc to NGC~7762. Nevertheless, Patat \\& Carraro (1995) estimated an older age for the cluster at 1.8 Gyr and less metallicity compared to the Sun. The age (133 Myr) we adopted for V731~Cep system seems to agree more with the age (266 Myr) that Chincarini (1966) calculated for NGC~7762. In addition to this, the distance and metal abundance estimation of NGC~7762 by Patat \\& Carraro (1995) are in agreement with the distance and metal abundance of V731~Cep found in this study. The most reliable evidence that V731~Cep is evaporated from NGC~7762 can be obtained from the spatial distribution and the metallicity of V731~Cep.  In calculation of the kinematical properties of V731~Cep, the systemic velocity, distance and proper motion components listed in Table~8 were used in the algorithm given by Johnson \\& Soderbloom (1987). The computed space velocity components with their errors are given in Table~8. The total space velocity of 14 km s$^{-1}$ for V731~Cep is in agreement with the space velocities of young stars. Although the space velocity of the cluster could not be computed due to its unavailable RV data, the distance of 55 pc between V731~Cep and NGC~7762, and similar age and metallicity distribution of V731~Cep with NGC~7762, suggest that V731~Cep could be evaporated from NGC~7762. However, it is necessary to carry out precise photometric and spectroscopic observations of the member stars of NGC~7762 in order to form more definite conclusions on the history of V731~Cep in relation with NGC~7762. \\\\ \\\\  \\textbf{Acknowledgements} \\\\ The authors would like to thank to Dr. Antonio Cabrera-Lavers for his help in extracting extinction data from RC stars and to the referee, Prof. E. F. Guinan, for his useful comments that improved the readability of this paper. This study is part of a project funded by \\c{C}OMU Scientific Research Foundation under project code BAP2008/37. The research of MW was supported by the Research Program MSM0021620860 of the Ministry of Education of Czech Republic. Participation of MZ in this study was endorsed by the grant GA CR 205/06/0217 of the Czech Science Foundation. \\\\"
    },
    "0807/0807.4631_arXiv.txt": {
        "abstract": "We present the magnetic landscape of the polar region of the Sun that is unprecedented in terms of high spatial resolution, large field of view, and polarimetric precision. These observations were carried out with the Solar Optical Telescope aboard \\emph{Hinode}. Using a Milne-Eddington inversion, we found many vertically-oriented magnetic flux tubes with field strength as strong as 1 kG that are scattered in latitude between $70\\,^{\\circ}$ $\\sim$ $90\\,^{\\circ}$. They all have the same polarity, consistent with the global polarity of the polar region. The field vectors were observed to diverge from the center of the flux elements, consistent with a view of magnetic fields that expand and fan out with height. The polar region is also covered with ubiquitous horizontal fields. The polar regions are the source of the fast solar wind channelled along unipolar coronal magnetic fields whose photospheric source is evidently rooted in the strong field, vertical patches of flux.  We conjecture that vertical flux tubes with large expansion around the photosphere-corona boundary serve as efficient chimneys for Alfv\\'en waves that accelerate the solar wind. ",
        "introduction": "The Sun's polar magnetic fields are thought to be the direct manifestation of the global poloidal fields in the interior, which serve as seed fields for the global dynamo that produces the toroidal fields responsible in active regions and sunspots. The polar regions are also the source of fast solar winds. Although the polar regions are of crucial importance to the dynamo process and acceleration of the fast solar winds, its magnetic properties are poorly known. Magnetic field measurements in the solar polar regions have long been a challenge: variable seeing combined with the strong intensity gradient and the foreshortening effect at the solar limb greatly increases the systematic noise in ground-based magnetographs. Nevertheless, pioneering observations have been carried out for the polar regions \\citep{tw90, lvz94, o04a, of04, hkm97, br07}. These observations typically have provided only the measurements of the line-of-sight magnetic component. Full Stokes polarimetry has also been carried, but as with most of the ground-based observations described above, the spatial resolution of those measurements was limited by seeing \\citep{bw96}. Another limitation of the past polar observations is that they have been restricted to individual polar faculae within a small field of view, and have not provided us with a global magnetic landscape of the polar region except for \\emph{GONG/SOLIS} \\citep{jh07}.  We investigated the properties of photospheric magnetic field in polar regions using the Solar Optical Telescope, SOT \\citep{st07, ys07, ki07, ts07}. SOT is a diffraction-limited (0\".2\\,-\\,0\".3) Gregorian telescope with filter-graph and spectro-polarimeter aboard the satellite \\emph{Hinode} \\citep{tk07}. These observations are unprecedented in terms of their very high spatial resolution, wide field of view, and high polarimetric sensitivity and accuracy in measurements of vector magnetic fields. ",
        "conclusions": "We discovered that the poloidal field near the pole has a form of unipolar flux tubes scattered in the polar region rather than a weak extended field. If the polar field with the same total magnetic flux $\\Phi \\sim BfS$ is uniformly distributed ($S$ is the total magnetic area), the estimated effective field strength would be about 10G as described above. The total magnetic energy is then proportional to $B^{2}fS = B\\Phi$. Thus, the surface poloidal magnetic energy is approximately 90 times larger than the case for the uniform magnetic field, if we take $B \\sim 900$G, corresponding to the peak of the energy PDF in Figure 5 (b). The equi-partition field strength $B_{e}$ is the field strength where magnetic energy is equal to kinetic energy of surface granular motion: $B_{e}=\\sqrt{4 \\pi \\rho v^{2}}$. The typical equipartition field strength $B_{e}$ is about 400 Gauss for granules with a velocity of $v = 2 \\times10^{5}$ cm s$^{-1}$ and $\\rho$ given in section 4.2. The magnetic field strength for the majority of patches is larger than the equi-partition field strength.  The observed unipolar strong flux tubes scattered in the polar region are considered to represent seed poloidal fields for toroidal fields \\citep{wns89a, wns89b}. Magnetic flux is transported to the polar regions with meridional flows and supergranular diffusion in the flux-transport dynamo model \\citep{dc99}. Since magnetic field has a form of such isolated flux tubes with super equi-partition strength instead of diffuse weak mean-field assumed in the flux transport dynamo \\citep{dc99}, flux transport on the sun would be done via an aerodynamic (drag) force against magnetic tension force, and may be more difficult than the case for the mean field case assumed in the models.  If the flux tubes seen on the surface of the Sun are maintained inside the Sun, this would affect a known difficulty in $\\Omega$-mechanism \\citep{wme56} to generate intense toroidal field: smaller amplification factor is needed to generate the same toroidal field from the poloidal field with intrinsic field strength of 1kG than from the averaged 10G field, and thus may be achievable within a solar cycle. We, however, recognize that there would remain a serious energetic problem, if the toroidal field strength indeed reaches 100kG \\citep{sc96, mr06}.  Total flux of vertical magnetic field at the polar region estimated here is at most $7.2 \\times 10^{21}$ Mx at the solar minimum, while various measurements on the total magnetic flux of single active region indicate $\\sim 10^{22}$Mx \\citep{lc07, jc07, mt08}. Thus, the measured total polar flux barely corresponds to that of single active region. The total toroidal flux would increase with time during the winding-up process by differential rotation, and the concept of the $\\Omega$-mechanism would be viable with the observations presented here.  The transient horizontal magnetic field discovered in the polar region appears to have properties similar to those found in quiet Sun and in active regions \\citep{bwl07, rc07, os07, ri07, ri08}. In particular, PDFs of magnetic field strength for the polar region (Figure 5 (a)), quiet Sun, and active regions \\citep{ri08} are remarkably similar, suggesting a common local dynamo process \\citep{fc99} taking place all over the Sun.  The X-ray telescope and EUV imaging spectrometer aboard \\emph{Hinode} observed remarkable activity in the polar regions in a form of micro flares and jets \\citep{sa07, ci07}. The lateral spreading of the vertical flux tubes to large area may be located well above the formation height of the two Fe lines, since there is no clear positional correlation between the horizontal fields and the vertical fields as seen in Figure 4. These X-ray jets could be due to magnetic reconnection at the lateral magnetic contacts with the horizontal fields and/or transient emergence of separate bipolar field lines \\citep{ssh98, sh92}.  In conclusion, the magnetic landscape of the polar region is characterized by vertical kG patches with super equi-partition field strength, a coherency in polarity, lifetime with time scale of $5\\,-\\,15$ hours, and the ubiquitous weaker transient horizontal fields. The life time of the magnetic concentrations in the quiet Sun is estimated to be 2500 s for $2.5 \\times 10^{18}$Mx, and 40 ks for $10 \\times 10^{18}$Mx with \\emph{MDI} \\citep{hage99}. It is important to clarify similarities and differences between the polar region and the quiet Sun with \\emph{Hinode}. This will be discussed in our subsequent paper."
    },
    "0807/0807.0224_arXiv.txt": {
        "abstract": "{}{Theory of random processes provides an attractive mathematical tool to describe the fluctuating signal from accreting sources, such as active galactic nuclei and Galactic black holes observed in X-rays. These objects exhibit featureless variability on different timescales, probably originating from an accretion disc.}{We  study the basic features of the power spectra in terms of  a general framework, which permits semi-analytical determination of the power spectral density (PSD) of the resulting light curve. We consider the expected signal generated by an ensemble of spots randomly created on the accretion disc surface. Spot generation is governed by  Poisson or by Hawkes processes. The latter one represents an avalanche mechanism and seems to be suggested by the observed form of the power spectrum. We  include general relativity effects shaping the signal on its propagation  to a distant observer.}{We analyse the PSD of a spotted disc light curve and show the accuracy of our semi-analytical approach by comparing the obtained PSD with the results of Monte Carlo simulations. The asymptotic slopes of PSD are $0$ at low frequencies and they drop to $-2$ at high frequencies, usually with a  single frequency break. More complex two-peak solutions also occur. The amplitude of the peaks and their frequency difference depend on the inherent timescales of the model, i.e., the intrinsic lifetime of the spots and the typical duration of avalanches.}{At intermediate frequencies, the intrinsic PSD is influenced by the individual light curve profile as well as by the type of the underlying process. However, even in cases when two Lorentzians seem to dominate the PSD, it does not necessarily imply that two single oscillation mechanisms operate simultaneously. Instead, it may well be the manifestation of the avalanche mechanism. The main advantage of our approach is an insight in the model functioning and the fast evaluation of the PSD.} ",
        "introduction": "It is widely accepted that massive black holes with accretion discs reside in cores of active galactic nuclei, where most activity originates and X-rays are produced (e.g., \\citeauthor{1990agn..book.....B} \\citeyear{1990agn..book.....B}). The observed light curves, $f\\,\\equiv\\,f(t)$, show irregular, featureless fluctuations with a very complex behaviour, practically at every studied frequency \\citep{2006ASPC..360.....G}. Variability has been traditionally analysed by the Fourier method \\citep{1992scma.book.....F}. Remarkably, a number of similarities appear between the properties of massive black holes in galactic nuclei and those in X-ray binaries, suggesting that some kind of universal rescaling operates according to central masses of  these systems \\citep{1998Natur.392..673M}. This concerns also the X-ray power spectra (e.g., \\citeauthor{2003ApJ...593...96M} \\citeyear{2003ApJ...593...96M}; \\citeauthor{2006Natur.444..730M} \\citeyear{2006Natur.444..730M}).  Light curves can be characterised by an appropriate estimator of the source variability which, in the mathematical sense, is a functional: $f\\rightarrow\\mathcal{S}\\left[f\\right]$. We accept the idea that the signal resulting from a spotted accretion disc is intrinsically stochastic, likely originating from turbulence. Hence, $\\mathcal{S}\\left[f\\right]$ is a random value. It can be a number (for example, the `rms' characteristic), function of a single variable (for example, the power-spectral density -- PSD) or of many variables (e.g., poly-spectra, rms--flux relation, etc). The appropriate choice depends on the type of information we seek and the quality of data available. The PSD is a traditional and widely utilised method to examine variable signals, and the AGN light curves are no exception. A typical signal can be represented by a broad band PSD with the tendency towards flattening at low frequency \\citep{1987Natur.325..694L,1987Natur.325..696M,1993ApJ...414L..85L,1993ARA&A..31..717M,2002MNRAS.332..231U}.  There is an ongoing debate about the overall shape of the PSD and the occurrence of the break frequency or, possibly, two break frequencies at which the slope of the PSD can change \\citep{1999ApJ...510..874N,2003ApJ...593...96M}. In the case of a widely-studied Seyfert galaxy, MCG--6-30-15, \\cite{2005MNRAS.359.1469M} have closely examined the slope of PSD, namely its bending, with RXTE and XMM-Newton data. It is worth noticing that the accurate fits to the X-ray sources seem  to exhibit a multi-Lorentzian structure rather than a simple power law.  The same is true for the  best studied example, the Seyfert~1 galaxy Akn~564 \\citep{2007MNRAS.382..985M}.  It was proposed \\citep{1991A&A...245..454A,1991A&A...246...21Z,1992AIPC..254..251W} that hot spots contribute to the variability of the AGN variability,  or that they could even be the dominant process shaping the variability pattern. These spots should occur on the disc surface following its intermittent irradiation by localised coronal flares \\citep{1979ApJ...229..318G,2001MNRAS.328..958M,2004A&A...420....1C}. Here, the ``spots'' represent a somewhat generic designation for non-axisymmetric features evolving on the disc surface in connection with flares. They share the bulk orbital motion with the underlying disc. The observed signal is thus modulated by relativistic effects as photons propagate towards a distant observer.  Various schemes have been discussed in which the fluctuations of the disc emissivity at different points of space and time are mutually interconnected in some way. In particular, the avalanche model \\citep{1999MNRAS.306L..31P,2002MNRAS.333..800Z,2005MNRAS.359..308Z} seems to be physically substantiated within the framework of magnetically-triggered flares and spots. It is also a promising model capable to reproduce, for example, the broken power-law PSD profiles. Notice, however, that other promising ideas were proposed (e.g., \\citeauthor{1994PASJ...46...97M} \\citeyear{1994PASJ...46...97M}; \\citeauthor{1997MNRAS.292..679L} \\citeyear{1997MNRAS.292..679L}), provoking the question of whether a common mathematical basis could reflect the entire range of models and provide us with general constraints, independent of (largely unknown) model details.  We add to this model by applying the method of random point processes \\citep{1965cox}. Interestingly enough, a rather formal approach can provide useful  analytical formulae defining the basic form of the expected power  spectrum. Apart from this practical aspect, we suggest that the concept discussed here offers much better insight into various influences that shape the expected form the power spectrum. These are very attractive features especially with respect to avalanche models, which may have different flavours, typically with a vast parameter space, thus proving very demanding to examine in a systematic manner.  Even more important is that the adopted formalism provides a very general tool and allows for a broader perspective on different mechanisms of variability \\citep{2007arXiv0711.4772P}. We develop the idea in a systematic way and give the explicit correspondence between our approach and some of the above-mentioned and widely-known scenarios \\citep{1991A&A...245..454A,1999MNRAS.306L..31P}. This description provides semi-analytical solutions, convenient to search through a broad parameter space. Our results can help to identify how the intrinsic properties of individual flares and the relativistic effects influence the overall PSD. In particular, we can identify those situations in which a doubly-broken power law occurs.  We consider stochastic processes (e.g., \\citeauthor{1971aitp.book.....F} \\citeyear{1971aitp.book.....F}; \\citeauthor{1994hsmp.book.....G} \\citeyear{1994hsmp.book.....G}) in the framework appropriate for modelling the accretion disc variability. In particular, in Sect.~\\ref{models} we consider a simple model of a spotted accretion disc constrained by the following three assumptions about the creation and evolution of spots: (i)~each spot is described by its time and place of birth ($t_j$, $r_j$ and $\\phi_j$) in the plain of the disc; (ii)~every new occurrence starts instantaneously; afterwards, the emissivity decays gradually to zero (the total radiated energy is of course finite); and (iii)~the intrinsic emissivity is fully determined by a finite set of parameters which form a vector, $\\mbox{\\boldmath $\\xi$}_j$, defining the light curve profile. Later on, we will consider the modulation of the intrinsic emission by the orbital motion and relativistic lensing. The disc itself has a passive role in our considerations; we will treat it as a geometrically thin, optically thick layer lying in the equatorial plane.  Because of the apparently random nature of the variability, we adopt a stochastic model in which the creation of spots is governed by a random process. The assumption that spots are mutually statistically independent seems to be a reasonable (first) approximation, however, we find that we do need to introduce some kind of relationship between them. This connection is discussed in Sect.~\\ref{relationship}. The statistical dependence among spots can be introduced in several ways. In Sect.~\\ref{results}, we explore in detail the specific models of interrelated spots using the formalism of Hawkes-type processes. Conclusions are summarised in Sect.~\\ref{conclusions}. Finally, in the Appendix we provide some mathematical prerequisites, which the reader may find useful to understand the general background of the paper, and we also summarise the mathematical notation. ",
        "conclusions": "\\label{conclusions} We adopted the viewpoint that the variability pattern is determined by the interplay among the bulk orbital motion, relativistic effects, and the intrinsic changes of the inner accretion disc. We concentrated our attention solely on the PSD characteristics. The spots have a certain kind of memory in our model.  We gave several examples in which the PSD changes the slope and certain break frequencies. The frequency of the break depends on the interplay of model properties, i.e., the intrinsic form of the spot light curves, which determine the individual contributions to the total signal together with the avalanche mechanism. {\\em The location of spots on the disc and the inclination of the source define the importance of relativistic effects.}  In some cases, a double break occurs and the overall PSD profile is then approximated by a broken power law. This is a promising feature in view of applications to real sources with accreting black holes. The broken power-law profile either resembles a combination of the Lorentzians or, in some cases, an intermediate power law develops and connects the two peaks across the middle frequencies. The change of the PSD slope is clearly visible and well-defined in some cases, though under most circumstances it appears rather blurry. The low-frequency limit of the PSD slope is a constant; the-high frequency behaviour depends mainly on the shape of the spot emission profile, including the general relativity effects. In our calculations the emissivity was decaying exponentially and the slope of the PSD was equal to $-2$ at high frequencies. {\\em In between those two limits the intrinsic PSD is influenced by both the individual light curve profile and the underlying process.}  It is interesting to notice that the doubly-broken power law occurs only for certain assumptions about the intrinsic light curves of the individual spots or avalanches -- their onset and the decay; in other cases the break frequencies are not well defined, or the broken power-law PSD is not preferred at all. {\\em We stress that if two Lorentzian seem to dominate the PSD (i.e., two peaks show up), it still does not necessarily mean that two single oscillation mechanisms operate simultaneously. Instead, it may well be the manifestation of the avalanche mechanism.}  We employed a general statistical approach to the variability of a black hole accretion disc with orbiting spots that continuously arise and decay. The origin and evolution of spots were described by Poissonian and Hawkes' processes, the latter one representing a category of avalanche models. We derived analytical formulae for the PSD, Eqs.\\ (\\ref{PoissPSD}) and (\\ref{HawkPSD}), and we discussed their limitations and accuracy. {\\it The main advantage of the analytical form is the insight into the properties and the fast evaluation that captures the main trends of the PSD shape.}  It is worth noting that the PSD does not maintain all information about the light curves that can be studied by Fourier methods \\citep{1992ApJ...391..518V}. Extensions have been discussed and compared with real data \\citep{1993ApJ...402..432K,1997MNRAS.288...12K,1999ApJ...510..874N,2008arXiv0802.0391V}, but this would go beyond the scope of our present work.  Our approach allows us to investigate the influence of the assumed mechanism, which describes the creation of parent spots and of subsequent cascades of daughter spots. In particular, we can discuss the PSD slope at different frequency ranges and locate the break frequencies depending on the model parameters. The relationship between the mathematical nature of the process and the PSD of the resulting signal is an interesting consequence of this investigation, as it provides a way to grasp and rigorously constrain the physical models of the source. Therefore we believe that the method that we described is very helpful for identifying the underlying mechanisms that shape the PSD in black hole accreting sources."
    },
    "0807/0807.2017_arXiv.txt": {
        "abstract": "We study second-order cosmological perturbations in scalar-tensor models of dark energy that satisfy local gravity constraints, including $f(R)$ gravity. We derive equations for matter fluctuations under a sub-horizon approximation and clarify conditions under which first-order perturbations in the scalar field can be neglected relative to second-order matter and velocity perturbations. We also compute the skewness of the matter density distribution and find that the difference from the $\\Lambda$CDM model is only less than a few percent even if the growth rate of first-order perturbations is significantly different from that in the $\\Lambda$CDM model. This shows that the skewness provides a model-independent test for the picture of gravitational instability from Gaussian initial perturbations including scalar-tensor modified gravity models. ",
        "introduction": "The constantly accumulating observational data \\cite{Perl} continue to confirm that the Universe has entered the phase of an accelerated expansion after the matter-dominated epoch. To reveal the origin of dark energy (DE) responsible for this late-time acceleration is one of the most serious stumbling block in modern cosmology \\cite{review,CST}. The first step toward understanding the nature of DE is to find a signature whether it originates from some modification of gravity or it comes from some exotic matter with negative pressure. If gravity is modified from Einstein's General Relativity, this leaves a number of interesting experimental and observational signatures that can be tested. Especially local gravity experiments generally place tight bounds for the parameter space of modified gravity models.  So far many modified gravity DE models have been proposed-- ranging from $f(R)$ gravity \\cite{fR} ($R$ is a Ricci scalar), scalar-tensor theory \\cite{stensor,Boi} to braneworld scenarios \\cite{brane}. The $f(R)$ gravity is presumably the simplest generalization to the $\\Lambda$-Cold Dark Matter ($\\Lambda$CDM) model ($f(R)=R-\\Lambda$). Nevertheless it is generally not easy to construct viable $f(R)$ models that satisfy all stability, experimental and observational constraints while at the same time showing appreciable deviations from the $\\Lambda$CDM model. In order to avoid that a scalar degree of freedom (scalaron) as well as a graviton becomes ghosts or tachyons we require the conditions $f_{,RR}>0$ and $f_{,R}>0$ \\cite{Star07}. These conditions are also needed for the stability of density perturbations \\cite{staper}. For the existence of a matter-dominated epoch followed by a late-time acceleration, the models need to be close to the $\\Lambda$CDM model ($m \\equiv Rf_{,RR}/f_{,R} \\approx +0$) in the region $R \\gg R_0$ ($R_0$ is the present cosmological Ricci scalar) \\cite{AGPT}. Moreover the mass of the scalaron field in the region $R \\gg R_0$ is sufficiently heavy for the compatibility with local gravity experiments \\cite{Nava,CT,TUT,Brax}. Finally, for the presence of a stable de Sitter fixed point at $r \\equiv -Rf_{,R}/f=-2$, we require that $0 \\le m(r=-2) \\le 1$ \\cite{Faraoni,AGPT}. The models proposed by Hu and Sawicki \\cite{Hu07} and Starobinsky \\cite{Star07} satisfy all these requirements. They take the asymptotic form, $f(R) \\simeq R-\\mu R_c [1-(R/R_c)^{-2n}]$ ($\\mu>0, R_c>0, n>0$), in the region $R \\gg R_c$ ($R_c$ is roughly the same order as $R_0$). See Refs.~\\cite{Li,AT,Appleby,Tsuji08,NO07} for other viable $f(R)$ models.  The main reason why viable $f(R)$ models are so restrictive is that the strength of a coupling $Q$ between dark energy and non-relativistic matter (such as dark matter) is large in the Einstein frame ($Q=-1/\\sqrt{6}$) \\cite{APT}. In the region of high-density where local gravity experiments are carried out, the scalaron field $\\phi$ needs to be almost frozen \\cite{Nava,CT} with a large mass through a chameleon mechanism \\cite{KW} to avoid that the field mediates a long ranged fifth force. Cosmologically this means that the field does not approach a kinematically driven $\\phi$ matter-dominated era ($``\\phi$MDE'' \\cite{coupled}) in which the evolution of scale factor is non-standard ($a \\propto t^{1/2}$ \\cite{APT}). The deviation from the $\\Lambda$CDM model becomes important as the mass of the scalaron gets smaller so that the field begins to evolve slowly along its potential. In other words the effect of modified gravity manifests itself from the late-time matter era to the accelerated epoch \\cite{Star07,Hu07}. This leaves a number of interesting observational signatures for the equation of state of DE \\cite{AT,Tsuji08}, matter power spectra \\cite{staper,Star07,Tsuji08} and convergence spectra in weak lensing \\cite{TT,Schmidt}.  One can generalize the analysis in $f(R)$ gravity to the theories that have arbitrary constant couplings $Q$ \\cite{TUMTY}. In fact this is equivalent to Brans-Dicke theory \\cite{BD} with a scalar-field potential $V(\\phi)$. By designing the potential so that the field mass is sufficiently heavy in the region of high density, it is possible to satisfy both local gravity and cosmological constraints even when $|Q|$ is of the order of unity \\cite{TUMTY}. The representative potential of this type is given by $V(\\phi)=V_0[1-C(1-e^{-2Q\\phi})^p]$ $(V_0>0, C>0, 0<p<1)$, which covers the $f(R)$ models of Hu and Sawicki \\cite{Hu07} and Starobinsky \\cite{Star07} as special cases. Especially when $|Q|$ is of the order of unity, these models lead to the large growth of matter density perturbations ($\\delta \\propto t^{(\\sqrt{25+48Q^2}-1)/6}$) at a late epoch of the matter era compared to the standard growth ($\\delta \\propto t^{2/3}$) at an early epoch. This gives rise to a significant change of the spectral index of the matter power spectrum relative to that in the $\\Lambda$CDM model \\cite{Star07,Tsuji08}. Moreover it was recently shown that the convergence power spectrum in weak lensing observations is subject to a large modification by the non-standard evolution of matter perturbations \\cite{TT}.  In this paper we shall study another test of modified gravity DE models mentioned above by evaluating a normalized skewness, $S_3=\\langle \\delta^3 \\rangle/\\langle \\delta^2 \\rangle^2$, of matter perturbations. The skewness provides a good test for the picture of gravitational instability from Gaussian initial conditions \\cite{Ber}. If large-scale structure grows via gravitational instability from Gaussian initial perturbations, the skewness in a Universe dominated by a pressureless matter is known to be $S_3=34/7$ in General Relativity \\cite{Peebles}. Even when cosmological constant is present at late times, the skewness depends weakly on the expansion history of the Universe (less than a few percent) \\cite{Marc,Benabed,Kofman}. This situation hardly changes in open/closed universes \\cite{curved} and Dvali-Gabadadze-Poratti braneworld models \\cite{Mul}. One can see some difference for the models that are significantly different from Einstein gravity--such as Cardassian cosmologies \\cite{Freese,Mul}, modified gravity models that respect Birkhoff's law \\cite{Lue}. In the context of dark energy coupled with dark matter, it was shown in Ref.~\\cite{Luca04} that the skewness can be a probe of the violation of equivalence principle between dark matter and (uncoupled) baryons.  In Brans-Dicke theory with cosmological constant $\\Lambda$ the skewness has been calculated in Ref.~\\cite{BD2} under the condition that the Brans-Dicke field is massless. In this case the evolution of scale factor during the matter-dominated epoch is given by $a(t) \\propto t^{(2\\omega_{\\rm BD}+2)/ (3\\omega_{\\rm BD}+4)}$ \\cite{BD2}, where $t$ is a cosmic time and $\\omega_{\\rm BD}$ is a Brans-Dicke parameter. If the field is massless, the Brans-Dicke parameter is constrained to be $\\omega_{\\rm BD}>40000$ \\cite{Hoyle} from solar-system experiments. This shows that the evolution of the scale factor in the matter era is very close to the standard one: $a(t) \\propto t^{2/3}$. We note that an effective gravitational ``constant'' that appears as a coefficient of matter density perturbations is also subject to change in Brans-Dicke theory. However it was found that the skewness in such a case is given by $S_3=(34\\omega_{\\rm BD}+56)/ (7\\omega_{\\rm BD}+12)$ \\cite{BD2} during the matter era, which is very close to the standard one ($S_3=34/7$) under the condition $\\omega_{\\rm BD}>40000$.  The $f(R)$ gravity corresponds to theory with the Brans-Dicke parameter $\\omega_{\\rm BD}=0$ \\cite{Chiba}. Even in this situation, if the scalaron field has a potential whose mass is sufficiently large in the region of high density, the $f(R)$ models can pass local gravity constraints as in the models proposed in Refs.~\\cite{Hu07,Star07}. In such cases, compared to Brans-Dicke theory with a massless field, it is expected that the skewness may show significant deviations from that in General Relativity. Since the evolution of scale factor and matter perturbations is different from that in the massless case, we can not employ the result of the skewness presented above.  In this paper we study second-order perturbations and the skewness for Brans-Dicke theory in the presence of a potential $V(\\phi)$. This is equivalent to the scalar-field action given in Eq.~(\\ref{action}) by identifying the coupling $Q$ with the Brans-Dicke parameter $\\omega_{\\rm BD}$ via the relation $1/(2Q^2)=3+2\\omega_{\\rm BD}$. In the massless case the solar-system constraint, $\\omega_{\\rm BD}>40000$, gives the bound $|Q| \\lesssim 10^{-3}$, but it is difficult to find some deviations from General Relativity in such a situation. Our interest is the case in which the coupling $Q$ is of the order of $0.1 \\lesssim |Q| \\lesssim 1$ with the field potential that has a sufficiently large mass in the high-density region. This analysis includes viable $f(R)$ models \\cite{Hu07,Star07} recently proposed in the literature. We would like to investigate how much extent the skewness differs from that in the $\\Lambda$CDM model. We also derive conditions under which the contribution coming from first-order field perturbations can be neglected relative to second-order matter and velocity perturbations by starting from fully relativistic second-order perturbation equations. ",
        "conclusions": "\\label{conclude}  In this paper we have studied the evolution of second-order matter density perturbations in a class of modified gravity models that satisfy local gravity constraints. We have considered the scalar-tensor action (\\ref{action}), which is equivalent to Brans-Dicke action (\\ref{action0}) with the correspondence $1/(2Q^2)=3+2\\omega_{\\rm BD}$. In the presence of a field potential it is possible to satisfy local gravity constraints (LGC) even when $|Q|$ is of the order of unity. In fact the potential (\\ref{potential}) is designed to have a large mass in the region of high density for the consistency with LGC. This covers the models proposed by Hu and Sawicki \\cite{Hu07} and Starobinsky \\cite{Star07} in the context of $f(R)$ gravity ($Q=-1/\\sqrt{6}$).  Starting from second-order relativistic equations of cosmological perturbations, we have derived the equation (\\ref{basic3}) of matter density fluctuations approximately. In so doing we employed the approximation that first-order perturbations in the scalar field $\\phi$ is neglected relative to second-order matter and velocity perturbations. This is valid under the conditions (\\ref{QH}) and (\\ref{delcon}), both of which can be naturally satisfied for the values of $Q$ we are interested in ($0.1 \\lesssim |Q| \\lesssim 1$). Compared to the $\\Lambda$CDM model, the effective gravitational constant $G_{\\rm eff}$ is subject to change at the late epoch of the matter era. This leads to the larger growth rate of first-order matter perturbations ($\\delta_k^{(1)} \\propto t^{(\\sqrt{25+48Q^2}-1)/6}$) compared to the standard case ($\\delta_k^{(1)} \\propto t^{2/3}$).  The skewness of matter distributions is determined by the second-order growth factor $E_a$ relative to the squared of the first-order growth factor $D$. In the ``scalar-tensor regime'' where the effective gravitational constant is given by $G_{\\rm eff} \\simeq (1+2Q^2)/8\\pi F$, we have derived the analytic expression (\\ref{anaes}) of the skewness in the matter-dominated epoch. In the ``General Relativistic regime'' where $G_{\\rm eff} \\simeq 1/8\\pi F \\simeq G$, we have reproduced the standard value $S_3=34/7$ in the Einstein-de Sitter Universe. In modified gravity models with $|Q| \\lesssim 1$, the analytic value (\\ref{anaes}) of the skewness in the asymptotic regime of the matter era is different from the value $34/7$ only less than a few percent. Even if we take into account the evolution of perturbations during the accelerated phase, the difference of the skewness relative to the $\\Lambda$CDM model remains to be small. The above result comes from the fact that the ratio of the second-order growth rate relative to the first-order one has a weak dependence on the coupling $Q$.  When $|Q|={\\cal O} (1)$ the growth rate of first-order matter perturbations is significantly different from that in the $\\Lambda$CDM model. This gives rise to large modifications to the matter power spectrum as well as to the convergence spectrum in weak lensing, while the skewness is hardly distinguishable from that in $\\Lambda$CDM model. This fact can be useful to discriminate large coupling scalar-tensor models among many other dark energy models from future high-precision observations."
    },
    "0807/0807.0012_arXiv.txt": {
        "abstract": "We investigate the GeV emission from gamma-ray bursts (GRBs), using the results from the Energetic Gamma Ray Experimental Telescope (EGRET), and in view of the {\\it Gamma-ray Large Area Space Telescope (GLAST)}. Assuming that the conventional prompt and afterglow photons originate from synchrotron radiation, we compare an accompanying inverse-Compton component with EGRET measurements and upper limits on GeV fluence, taking Klein-Nishina feedback into account. We find that EGRET constraints are consistent with the theoretical framework of the synchrotron self-Compton model for both prompt and afterglow phases, and discuss constraints on microphysical parameters in both phases. Based on the inverse-Compton model and using EGRET results, we predict that {\\it GLAST} would detect GRBs with GeV photons at a rate $\\gtrsim 20$ yr$^{-1}$ from each of the prompt and afterglow phases. This rate applies to the high-energy tail of the prompt synchrotron emission and to the inverse-Compton component of the afterglow. Theory predicts that in a large fraction of the cases where synchrotron GeV prompt emission would be detected by {\\it GLAST}, inverse-Compton photons should be detected as well at high energies ($\\gtrsim 10$ GeV). Therefore {\\it GLAST} will enable a more precise test of the high-energy emission mechanism. Finally, we show that the contribution of GRBs to the flux of the extragalactic gamma-ray background measured with EGRET is at least 0.01\\% and likely around 0.1\\%. ",
        "introduction": "\\label{sec:Introduction}  Cosmological gamma-ray bursts (GRBs) have released a tremendous amount of energy in the past and present Universe. Their emission covers very wide range of frequencies: a highly variable prompt phase radiates $\\sim$100 keV gamma rays, while a subsequent afterglow radiates radio to X-ray photons. It is likely that the bulk of these photons are emitted by gyration of relativistic electrons in magnetic fields---e.g., synchrotron radiation. The relativistic electrons are accelerated in either internal dissipation (for prompt emission) or external shocks (for afterglows). For reviews, see, \\citet{Piran2005,Meszaros2006,Nakar2007}.  GeV photons were detected as well from several GRBs by the Energetic Gamma Ray Experimental Telescope (EGRET) on board the {\\it Compton Gamma Ray Observatory (CGRO)} \\citep{Schneid1992,Sommer1994,Hurley1994,Schneid1995,Gonzalez2003}. The data are still not sufficient for us to firmly infer emission mechanisms of these GeV gamma rays, but the most promising mechanism is synchrotron self-Compton (SSC) scattering \\citep*[e.g.,][]{Meszaros1994,Waxman1997,Wei1998,Chiang1999, Panaitescu2000,Zhang2001,Sari2001,Guetta2003}. This is because the relevant emission parameters such as the energy fraction of the GRB jets going to electrons ($\\epsilon_e$) and magnetic fields ($\\epsilon_B$) are relatively well measured from the afterglow spectra as well as light curves; the typical values are $\\epsilon_e = 0.1$ and $\\epsilon_B = 0.01$ \\citep[e.g.,][]{Panaitescu2001,Yost2003}. In the prompt emission, $\\epsilon_e$ is similar or even higher, as evident from the high efficiency of this phase, while $\\epsilon_B$ is not well constrained.  Thus, there should be a significant inverse-Compton (IC) component accompanying the synchrotron radiation in both the afterglow and prompt emission. The luminosities of the synchrotron and IC are expected to be comparable as IC-to-synchrotron luminosity ratio is roughly given by $(\\epsilon_e / \\epsilon_B)^{1/2}$, according to theory \\citep[e.g.,][]{Sari2001}.  In this paper, we explore the GeV gamma-ray emission of GRBs in the context of SSC mechanism.\\footnote{Our analysis and conclusions are applicable also if the MeV and/or radio-X-ray afterglow emission mechanism is not synchrotron but another type of emission from relativistic electrons that gyrate in a magnetic field, such as jitter radiation \\citep{Medvedev00}.}  Besides the several GRBs detected by EGRET, there are many others for which upper bounds on the fluence were obtained \\citep{Gonzalez2005}. These $\\sim$100 GRBs should also be compared with the predictions of SSC model, because the fluence upper limits in the EGRET energy band are comparable to the fluence of prompt emission collected by Burst And Transient Source Experiment (BATSE) instrument onboard {\\it CGRO}. As the experimental bound is already strong, while theoretical models of SSC process predict a large fluence for the EGRET energy range, we derive meaningful constraints from EGRET data analysis on the physics of the high-energy emission mechanisms of GRBs. This approach is different from (and therefore complementary with) that in previous studies \\citep*[e.g.,][and references therein]{Dermer00, Asano2007, Ioka2007, Gou2007, Fan2008, Murase2008, Panaitescu2008}, where the prediction of gamma-ray flux relies only on theoretical models and sub-GeV observations. We instead use EGRET data in order to infer the GeV emission and constrain the theoretical models.  We use our results to predict the expected number of GRBs that would be detected by the {\\it Gamma-ray Large Area Space Telescope (GLAST)}. The {\\it GLAST} satellite is equipped with the Large Area Telescope (LAT), which is an upgraded version of EGRET. Since revealing the high-energy emission mechanisms of GRBs are one of the important objectives of {\\it GLAST}, our prediction should give a useful guideline. Finally, we apply our results to estimate the contribution of GRBs to the diffuse extragalactic gamma-ray background (EGB), which was also measured by EGRET \\citep[][see, however, \\citealt*{Keshet2004b} for a subtle issue of Galactic foreground subtraction]{Sreekumar1998,Strong2004}.  This paper is organized as follows. In \\S~\\ref{sec:IC}, we summarize the predictions of SSC model for the prompt (\\S~\\ref{sub:prompt}) and afterglow (\\S~\\ref{sub:afterglow}) phases. Section~\\ref{sec:Constraint on high-energy emission with EGRET} is devoted for analysis of the GRB fluence data by EGRET, from which distributions of fluence in the GeV band are derived. We then use these distributions to argue prospects for GRB detection with {\\it GLAST} in \\S~\\ref{sec:GLAST}, and implications for EGB from GRB emissions in \\S~\\ref{sec:EGB}. In \\S~\\ref{sec:conclusions}, we give a summary of the present paper. ",
        "conclusions": "\\label{sec:conclusions}  The {\\it GLAST} satellite would enable us to test high-energy emission mechanisms of GRBs. If this emission will be found to be consistent with SSC then its observations would constrain physical parameters such as $\\epsilon_{e}/\\epsilon_B$ ratio and the bulk Lorentz factor of the jet, $\\Gamma_{b}$. The EGRET instrument on board {\\it CGRO}, while less sensitive than the {\\it GLAST}-LAT detector, identified several BATSE GRBs with GeV photons. In addition, stringent upper limits for $\\sim$100 GRBs were put on fluences in the GeV band by analyzing the EGRET data \\citep{Gonzalez2005}.  In this paper, we further extended this EGRET result, comparing with the SSC emission model. Following theoretical models of SSC, we assumed that there is a linear correlation between fluences in BATSE and EGRET energy bands, and that the proportionality coefficient $\\eta$ follows a log-normal distribution. We found that the predictions from the SSC model using canonical parameter values is fully consistent with EGRET fluence measurements and upper limits for both the prompt and afterglow phases. During the course of showing this result, we properly took the Klein-Nishina feedback effect into account in the theoretical calculation. The best-fit value of the coefficient was $\\log \\eta \\simeq -1.5$ for both the prompt and afterglow emissions, and it is already stringent enough to test the SSC scenario. The limits for the prompt emission phase are for the synchrotron radiation, and thus if we consider the IC component as well, the value of $\\eta$ could be larger by up to one order of magnitude.  The obtained $\\eta$ distribution, together with the BATSE fluence distribution, gives the expected fluence distribution in the GeV band, which is shown in Figure~\\ref{fig:dndf_egret}. As the {\\it GLAST}-LAT detector covers EGRET energy band, we can predict the detectable number of GRBs with {\\it GLAST} from the distribution of $F_{\\rm EGRET}$, given the {\\it GLAST}-LAT sensitivity. Our conservative estimate using the five-photon criterion is that about $\\sim$20 GRBs among those detected with GBM would be detected with {\\it GLAST}-LAT each year. This number could be even larger if we use fewer-photon criteria. The fluence distribution can also be used to estimate the GRB contribution to the EGB intensity. We found that the contribution would be at least $\\sim$0.01\\% but is likely to be as large as $\\sim$0.1\\%."
    },
    "0807/0807.4340_arXiv.txt": {
        "abstract": "We estimate the temporal change of magnetic flux perpendicular to the solar surface in a decaying active region by using a time series of the spatial distribution of vector magnetic fields in the photosphere. The vector magnetic fields are derived from full spectropolarimetric measurements with the Solar Optical Telescope aboard \\textit{Hinode}. We compare a magnetic flux loss rate to a flux transport rate in a decaying sunspot and its surrounding moat region. The amount of magnetic flux that decreases in the sunspot and moat region is very similar to magnetic flux transported to the outer boundary of the moat region. The flux loss rates [$(dF/dt)_{loss}$] of magnetic elements with positive and negative polarities are balanced each other around the outer boundary of the moat region. These results suggest that most of the magnetic flux in the sunspot is transported to the outer boundary of the moat region as moving magnetic features, and then removed from the photosphere by flux cancellation around the outer boundary of the moat region. ",
        "introduction": "How and where is magnetic flux in sunspots removed from the photosphere? Sunspot umbrae are sometimes split by formation of a light bridge, which is a bright lane crossing the umbra \\citep{Bray1964}. Small magnetic features with a typical size less than 2$\\arcsec$ called moving magnetic features \\citep[MMFs;][]{Harvey1973} are generally observed in the moat region that surrounds a decaying, mature sunspot. It has been reported that MMFs appear not only in the decaying phase of sunspots but also in the growing phase \\citep[e.g.][]{Wang1991}. The MMFs mostly appear around the outer boundary of the sunspot, moving almost radially outward during their lifetime ranging from a few minutes to 10 hr \\citep{Harvey1973, Zhang2003, Hagenarr2005}. The formation of the light bridges and MMFs is closely related to the fragmentation and disintegration of the sunspot magnetic flux. Indeed, it has been observed that the net flux carried away from the sunspot by MMFs is larger than the flux decrease in the sunspot \\citep{Martinez2002, Kubo2007a}. This indicates that MMFs can be responsible for the flux loss of the sunspot.   The mutual apparent loss of magnetic flux is often observed in the line-of-sight magnetograms when one magnetic polarity element collides with another polarity magnetic element in the photosphere. This apparent flux loss is called ``magnetic flux cancellation'' as a descriptive term. It is observed that moving magnetic features often collide with apparently static opposite polarity magnetic features around the outer boundary of the moat region \\citep{Martin1985, Yurchyshyn2001, Chae2004, Bellot2005} and widely believed that understanding the flux cancellation process around the moat boundary is the key to understand the dissipation of the sunspot flux from the photosphere. Three models have been proposed by \\citet{Zwaan1987} to describe the flux cancellation: (1) retraction of magnetic fields that connect an emerged bipole, (2) submergence of $\\Omega$-loop formed by magnetic reconnection between the canceling two bipoles above the photosphere, and (3) emergence of U-loop due to reconnection below the photosphere. As expected in these models, horizontal magnetic fields are formed between canceling magnetic elements \\citep{Wang1993, Chae2004, Kubo2007b}. However, whether upward or downward motions are observed in the cancellation sites depends on the events and the positions in the cancellation sites \\citep{Harvey1999, Kubo2007b}. Therefore, the nature of the physical process driving magnetic flux cancellation is still an open issue.   This study attempts to address a basic question how much magnetic flux is carried away from the sunspot to the outer boundary of the moat region and is subsequently removed from the photosphere. Because it has been difficult to measure the magnetic field vector under stable seeing conditions for the period longer than a typical lifetime of MMFs, the flux loss rate of the sunspot, flux transport rate due to MMFs, and flux cancellation rate have been independently estimated by using different data sets. A time series of spectropolarimetric measurements with the Solar Optical Telescope \\citep[SOT;][]{Tsuneta2008} aboard the \\textit{Hinode} satellite \\citep{Kosugi2007} allows us, for the first time, to estimate an accurate flux change without any effects of atmospheric seeing. Moreover, the high spatial resolution observations with the SOT decrease the likelihood of spurious magnetic cancellation events, i.e., those for which magnetic elements with opposite polarities are located entirely within a resolution element and will dramatically increase the reliability of the results presented. ",
        "conclusions": "The averaged flux change and averaged flux transport in Figure~\\ref{flux_summary} are described in units of Mx s$^{-1}$, so that we can compare these values directly. The observed flux change [$(dF/dt)_{Obs}$] in each region and the observed flux transport [$(F_v)_{Obs}$] at its boundaries would have a relationship: \\begin{equation} (\\frac{dF}{dt})_{Obs} = (\\frac{dF}{dt})_{Emerge} - (\\frac{dF}{dt})_{Loss} \\pm (F_v)_{Obs}.\\label{eq_flux_relation} \\end{equation} The increase of magnetic flux due to flux emergence [$(dF/dt)_{Emerge}$] can be neglected in this case, because we select the period without any significant flux emergence. Thus, we can estimate an actual flux loss rate [$(dF/dt)_{Loss}$] in each region from the observations.  The total of flux decrease rates in the sunspot and unipolar regions ($dF/dt$ = -3.2 - 4.8 = -8.0$\\times 10^{15}$ Mx s$^{-1}$) is almost equal to the flux transport rate at the outer boundary of the unipolar region for the positive polarity ($F_v$ = 7.4$\\times 10^{15}$ Mx s$^{-1}$). This means that most of magnetic flux that disappeared in the sunspot and unipolar regions is carried away to the mixed polarity region. Note that we do not trace each magnetic element, and all of the MMFs that separated from the sunspot may not reach at the outer boundary of the unipolar region. However, we make Figure~\\ref{flux_summary} from the observations for about 12 hr, which is twice as long as the period that magnetic elements with the average horizontal speed of 0.5 km s$^{-1}$ need to move through the unipolar region of a width of 14$\\arcsec$. The increase of the positive flux in the mixed polarity region supports the migration of positive flux into the mixed polarity region. Both the increase of the positive flux in the mixed polarity region and the flux transport for the positive polarity elements at the outer boundary of the mixed polarity region are smaller than the positive flux transported from the unipolar region. Therefore, the magnetic flux that is carried away from the sunspot (and moat region) mostly disappears in the mixed polarity region, especially near the outer boundary of the moat region.  One issue is that the positive flux carried away from the sunspot region ($F_v$ = 7.8$\\times 10^{15}$ Mx s$^{-1}$) is bigger than decrease of the positive flux in the sunspot region ($dF/dt$ = -3.2$\\times 10^{15}$ Mx s$^{-1}$). This tendency was also reported in the previous work with a lower (about $1\\arcsec$) spatial resolution \\citep{Kubo2007a}. As a result of no flux emergence in the sunspot region, the flux transport rate should be less than the flux decrease rate in the sunspot region. That is to say that the flux transport rate is overestimated at the outer boundary of the sunspot region. In the calculation of horizontal velocities, we use the apodization window with 1$\\arcsec$, which is lower than the spatial resolution of the magnetic field maps. Such a lower spatial resolution of the horizontal velocity maps probably causes the overestimation of the flux transport rate at the outer boundary of the sunspot region. Fuzzy, small magnetic elements with a short lifetime have been observed around the outer boundary of decaying sunspots \\citep{Zhang2007, Kubo2008}. These fuzzy magnetic elements have higher outward motion and smaller magnetic flux than those of usual MMFs. In the estimation of flux transport at the outer boundary of the sunspot, magnetic flux is mostly represented by MMFs with small horizontal velocity and large magnetic flux, but its horizontal velocity is represented by the fuzzy magnetic elements. We believe that these fuzzy magnetic elements correspond to a fluctuation of field strength or a fluctuation of inclination of penumbral magnetic fields, and thus do not contribute to the flux loss of the sunspot. Further investigation using spectropolarimetric measurements with a higher cadence will be necessary to know the nature of such fuzzy magnetic elements and their impact on the presented calculations.  The magnetic flux of negative polarity decreases in the mixed polarity region, although the negative flux converges from the inner and outer boundaries of the mixed polarity region. Considering that the negative flux moves into the mixed polarity region with the average rate of 1.6$\\times 10^{15}$ Mx s$^{-1}$, the actual flux loss rate [$(dF/dt)_{Loss}$] in the mixed polarity region may be as large as 3.9$\\times 10^{15}$ Mx s$^{-1}$ from Equation~(\\ref{eq_flux_relation}). This flux loss rate of negative polarity is balanced by the actual flux loss rate of the positive polarity (3.9$\\times 10^{15}$ Mx s$^{-1}$), which is a difference between the flux transport rate ($F_v$ = 7.4 - 0.8 = 6.6$\\times 10^{15}$ Mx s$^{-1}$) into the mixed polarity region and the flux increase rate ($dF/dt$ = 2.7$\\times 10^{15}$ Mx s$^{-1}$) there. The flux loss rates with both polarities in the mixed polarity region are consistent with the cancellation rates in active regions \\citep{Chae2000,Chae2004,Kubo2007b}. Furthermore, most of the magnetic elements with negative polarity are located in contact with the positive elements. These results suggest that magnetic flux cancellation at the outer boundary of the moat region is essential for the removal of the sunspot magnetic flux from the photosphere."
    },
    "0807/0807.1565_arXiv.txt": {
        "abstract": "Using three newly identified galaxy clusters at $z\\sim1$ (photometric redshift) we measure the evolution of the galaxies within clusters from high redshift to the present day by studying the growth of the red cluster sequence.  The clusters are located in the Spitzer Infrared Array Camera (IRAC) Dark Field, an extremely deep mid-infrared survey near the north ecliptic pole with photometry in 18 total bands from X-ray through far-IR. Two of the candidate clusters are additionally detected as extended emission in matching Chandra data in the survey area allowing us to measure their masses to be $M_{500}= 6.2 \\pm 1.0 \\times 10^{13}$ and $3.6 \\pm 1.1 \\times 10^{13}$ \\msun.  For all three clusters we create a composite color magnitude diagram in rest-frame B-K using our deep HST and Spitzer imaging.  By comparing the fraction of low luminosity member galaxies on the composite red sequence with the corresponding population in local clusters at $z=0.1$ taken from the COSMOS survey, we examine the effect of a galaxy's mass on its evolution.  We find a deficit of faint galaxies on the red sequence in our $z\\sim1$ clusters which implies that more massive galaxies have evolved in clusters faster than less massive galaxies, and that the less massive galaxies are still forming stars in clusters such that they have not yet settled onto the red sequence. ",
        "introduction": "The redshift range from $z=1$ to the present day is a particularly dynamic epoch in the history of groups and clusters as evidenced by the evolution of the morphology-density relation and increasing fraction of blue galaxies with increasing redshift \\citep{capak2007, butcher1984}. Interestingly, cluster ellipticals at $z\\sim1$ already have a narrow distribution of red colors \\citep[the red cluster sequence (RCS);][]{blakeslee2003,vandokkum2001}.  There is some debate about the mechanism by which these cluster galaxies arrive onto the red sequence.  It is difficult to distinguish whether these red ellipticals all formed their stars and did their merging at $z>3$, then stopped forming stars when they entered the cluster environment \\citep{ford2004}; or if they are the product of the merging of gas-poor systems which do not produce star formation \\citep{vandokkum2005}.  We investigate whether the red population is still in the process of forming at $z=1$ or if indeed assembly has already finished at higher redshift by studying the presence of the faint end of the RCS at $z=1$ and comparing it to the present epoch.  We measure the ratio of faint to bright RCS galaxies in a sample of three $z\\sim1$ clusters from the IRAC Dark Field (described below).  These clusters have the benefit of extremely deep $3.6\\,\\micron$ data which allows us to study the faint end of the luminosity function at rest frame near-IR, which traces the peak of the spectral energy distribution in galaxies.  A confirmed deficit of faint galaxies on the RCS would imply that more massive galaxies have evolved in clusters faster than less massive galaxies. A constant fraction of faint red galaxies between $z=1$ and the present would require a formation mechanism where galaxies of all masses have already joined the red sequence at redshifts higher than one.  There is evidence that the faint end of the RCS is not completely in place by z=0.8 \\citep[][and references therein]{delucia2007,koyama2007}, although at least some clusters at these redshifts appear to have complete RCSs to M*+3.5 \\citep{andreon2006}.  These questions are ideally addressed with deep infrared surveys of clusters at high redshift.  In the last four years deep and wide area surveys in the mid-infrared using the {\\it Spitzer Space Telescope} have substantially opened a new window on galaxy and star formation at $0<z<3$.  {\\it Spitzer} now routinely produces imaging of large fields to higher resolution and fainter depths than previously possible.  IRAC, the mid-infrared camera on-board Spitzer, takes images at 3.6, 4.5, 5.8, and 8.0\\,\\micron \\citep{fazio2004}.  The shorter wavelengths provide a direct measurement of the stellar content of galaxies at redshifts as high as three.  The longer wavelength channels sample emission from polycyclic aromatic hydrocarbons (PAHs) in low redshift galaxies, as well as direct thermal emission from hot dust.  The deepest such survey is the dark current calibration field for the mid-infrared camera, commonly known as the ``IRAC Dark Field''.  These deep mid-IR data are supplemented by additional 14 band photometry including Palomar u',g',r',i', {\\it HST/ACS} F814W, MMT $z^\\prime$, Palomar J, H,\\& K , Akari 11 \\& $15\\,\\micron$, Spitzer MIPS 24 \\& $70\\,\\micron$, and Chandra ACIS-I imaging.  This extensive, multiwavelength dataset gives us the distinctly unique opportunity to study the development of the red sequence in galaxy clusters at redshift one.  The second paper on galaxy clusters in this series will discuss the role of star formation in the evolution of clusters at $z=1$ by examining the {\\it Spitzer} 24\\,\\micron \\ data in conjunction with the morphological information from the HST ACS dataset.   This paper is structured in the following manner. In \\S\\ref{observations} \\& \\S\\ref{photz} we discuss the data and derived photometric redshift determination.  Details of the cluster search and cluster properties are presented in \\S\\ref{search}.  In \\S\\ref{cmds} \\& \\S\\ref{results} we present the color magnitude diagrams of the candidate clusters and results of the faint-to bright ratios of red sequence galaxies.  In \\S\\ref{discuss} we discuss the implications for the evolution of cluster galaxies.  Throughout this paper we use $H_0=70$km/s/Mpc, $\\Omega_M$ = 0.3, $\\Omega_\\Lambda$ = 0.7.  With this cosmology, the luminosity distance at z=1 is 6607 Mpc, but the angular diameter distance is a factor of $(1+z)^2$ less, or 1652 Mpc.  All photometry is quoted in the AB magnitude system. ",
        "conclusions": "\\label{discuss}   We find that there are fewer faint red cluster galaxies at high redshift than low redshift in comparison to the number of bright red galaxies.  There are two possible explanations for this.  There could be overall less faint cluster galaxies at z=1 at all colors than the present epoch, or the faint end of the red sequence is not yet in place, and instead those galaxies which will fill the faint end at z=0, are still forming stars at higher redshifts.  The first explanation, that there are overall, at all colors, fewer faint galaxies at z=1 than at z=0, goes directly against the theory of hierarchical formation.  We expect that clusters at higher redshifts will have more faint galaxies than today, and that over time the faint galaxies will merge into brighter, more massive galaxies.  This will have the effect of clusters at higher redshifts having steeper faint end slopes of the luminosity function than today's flatter slopes, or exactly the opposite trend of what would be suggested if overall there were less faint galaxies at $z=1$.  While we do not think this is a likely solution, it is interesting that in a simulation of cluster luminosity functions, \\citet{khochfar2007} find a slight trend that the faint end slope of clusters does steepen from redshift 1 to 0, whereas overall they find the hierarchical formation trend of flattening slope from z=6 to 0.  This is probably due to noise in their calculation.  It would be nice to know how the overall faint end slope of clusters is evolving from z=1 to 0, however we do not have a large enough sample to attack this problem.  It is more likely that the deficit of faint red galaxies at high z is due to the fact that those faint galaxies are still forming stars, and so have bluer colors at those redshifts.  This implies that a) bright, massive galaxies have already shut off their star formation in clusters by z=1, and b) faint, less massive galaxies are still forming stars in clusters at z=1.  We know that clusters exhibit red sequences at higher z, so it is not new that massive galaxies have already undergone some process, excited by the cluster environment, which keeps them from forming stars, and lands them on the red sequence.  This is further evidence for the popular theory of 'downsizing', in which the more massive galaxies evolve first.  However, all galaxies in clusters do not follow the same evolutionary processes, instead evolution from the blue cloud to the red sequence seems to be mass dependent.  Whatever process stops star formation in clusters appears has not yet happened at z=1 to the less massive galaxies. Those galaxies are still being allowed to form enough stars to stay off of the red sequence.  The red population of clusters is not yet fully in place by z=1.   We do not compare the specific values of faint-to-bright ratios of this work with those in the literature as different wavelengths, areas, and definitions of faint and bright are used.  However the trend for a deficit of faint galaxies at high redshifts is in agreement with the works of \\citet{delucia2007, gilbank2008, stott2007} and \\citet{koyama2007}, but in contrast to \\citet{andreon2007}.  The cause of the difference is unclear, but we note that we use a much redder, wider filter set at high redshifts (rest-frame B-K) than \\citet{andreon2007} (rest-frame u-g).  We also compare with the same rest-frame color at low redshift without applying a k-correction.   Future work in this field requires a large enough sample of confirmed clusters at high-z with consistent observations well below $M^*$ to be able to split the sample on cluster properties.  There are hints that the faint-to-bright ratios are cluster mass or richness dependent, but the literature shows contradictory trends for these effects likely due to the small samples used to date \\citep{gilbank2008, delucia2007}. With this evidence for ongoing star formation in clusters at z=1, in paper two of this series we will examine the star formation rates and morphologies of the cluster galaxies using deep SPITZER MIPS $24\\,\\micron$ and HST ACS F814W data. Additionally as the community continues to build a larger sample of high redshift clusters we will be able to study their properties, in particular their suitability for dark energy number count surveys \\citep{wang2004}."
    },
    "0807/0807.3560_arXiv.txt": {
        "abstract": "We recently reported the discovery of a regular corrugation pattern in the \\HI\\ disk of the isolated, edge-on spiral galaxy IC~2233. Here we present measurements of the vertical structure of this galaxy at several additional wavelengths, ranging from the far ultraviolet to the far infrared. We find that undular patterns with amplitude $\\lsim 5''$ ($\\lsim$250~pc) are visible in a variety of \\popI\\ tracers in IC~2233, including the young-to-intermediate age stars, the \\HII\\ regions, and the dust. However, the vertical excursions become less pronounced in the older stellar populations traced by the mid-infrared light. This suggests that the process leading to the vertical displacements may be linked with the regulation of star formation in the galaxy. We have also identified a relationship between the locations of the density corrugations and small-amplitude ($\\lsim$5~\\kms) velocity undulations in the \\HI\\ rotation curve. We are able to exclude several possible mechanisms for the origin of the observed corrugations, including tidal interaction from a companion, Parker instabilities, or a galactic bore. Global gravitational instabilities appear to be the most likely explanation, although local perturbations may also be important. ",
        "introduction": "It is well known that the disks of late-type galaxies are not completely flat. Warping of the outer gas disk into an ``integral sign'' shape is a nearly ubiquitous feature of spiral galaxies, even those that appear to be quite isolated (Garc\\'\\i a-Ruiz et al. 2002). A significant fraction of galaxies also have warped stellar disks (S\\'anchez-Saavedra et al. 1990; Reshetnikov \\& Combes 1998). For several decades it has been recognized that the \\HI\\ disk and other \\popI\\ components of the Milky Way exhibit an additional type of vertical structure in the form of systematic displacements from the mean plane known as ``corrugations'' (e.g., Gum et al. 1960; Quiroga 1974,1977; Lockman 1977; Sanders et al. 1984). These corrugations are present along both radial and azimuthal directions and have scales ranging from $\\sim$50 to 350~pc in amplitude and from one to several kiloparsecs in wavelength (Spicker \\& Feitzinger 1986).  Despite predictions that corrugations may be a common phenomenon in disk galaxies, systematic searches for these features have been limited, and consequently, little has been known about their true frequency of occurrence in external galaxies. To date, evidence of corrugations has been reported for only a handful of cases outside the Milky Way: in the emission nebulae of M31 (Arp 1964), in the stellar (optical) light of three edge-on, late-type galaxies (Florido et al. 1991,1992), and in the velocity structure of the face-on spiral NGC~5427 (Alfaro et al. 2001).  The origin of disk corrugations is a matter of longstanding debate given the small number of galaxies in which they have so far been detected (see the reviews by Alfaro \\& Efremov 1996 and Alfaro 2003). Suggestions have included: gravitational instabilities (e.g., Nelson 1976); tidal interactions (Edelsohn \\& Elmegreen 1997), collisions of high-velocity clouds with the disk (Franco et al. 1988; Santill\\'an et al. 1999); interaction of spiral waves with the gaseous disk (Alfaro et al. 2001); and the undular mode of the Parker instability (Franco et al. 2002). Depending on their origin, the presence or absence of corrugations can provide important clues into the role of magnetic fields on galaxy disk structure (e.g., Spicker \\& Feitzinger 1986; Santill\\'an et al. 2000; Franco et al. 2002), the formation of molecular clouds (e.g., Nelson \\& Matsuda 1980), the degree of self-gravity of galaxy disks (e.g., Revaz \\& Pfenniger 2004), or the timescales of recent interactions (e.g., Edelsohn \\& Elmegreen 1997). Moreover, current evidence suggests that whatever the physical process(es) responsible for corrugations, they are closely linked to the mechanisms responsible for the formation of the dense gas condensations and therefore with the regulation of star formation (e.g., Nelson \\& Matsuda 1980; Alfaro et al. 1992; Alfaro \\& Efremov 1996; Alfaro 2003).  In a recent paper (Matthews \\& Uson 2008; hereafter MU08), we presented  optical and \\HI\\ imaging observations of the nearby, edge-on spiral galaxy IC~2233. There we reported the discovery of a corrugated pattern in the \\HI\\ disk of this galaxy---i.e., we found that the \\HI\\ layer of IC~2233 exhibits a remarkably regular pattern of positive and negative vertical displacements from the kinematically-defined midplane. This pattern has a wavelength of $\\sim150''$ (7~kpc) and an amplitude that increases with distance from the center of the galaxy, reaching a maximum of $\\sim3''$ ($\\sim$150~pc).\\footnote[3]{Throughout this paper we assume a distance to IC~2233 of 10~Mpc (see MU08).} [For comparison, the FWHM thickness of the \\HI\\ layer ranges from $\\sim$\\as{10}{4} near the disk center to $\\sim$\\as{20}{3} toward the outer disk (MU08)]. Subsequently, we have explored the vertical structure of IC~2233 using a variety of additional tracers of the stellar and interstellar components of its disk. We have also examined the relationship between the peaks of the vertical \\HI\\ excursions with both \\HI\\ density enhancements along the disk and with velocity perturbations in the disk rotation curve. Here we report the discovery of undular patterns in several additional \\popI\\ tracers in IC~2233. In contrast, we find the vertical displacements are minimal in the old stars. We also report the discovery of a link between the locations of the peak excursions from the midplane in \\HI\\ with features in the \\HI\\ rotation curve. In \\S~\\ref{discussion} we  discuss the implications of these findings for understanding the origin of the disk corrugations seen in IC~2233 and other galaxies. ",
        "conclusions": "} As noted above, a number of mechanisms have been proposed to account for the observed corrugations in disk galaxies. We now discuss the suitability of several of these mechanisms for explaining the observed corrugations in IC~2233.  \\subsection{Tidal Interactions\\protect\\label{tides}} Numerical modelling by Edelsohn \\& Elmegreen  (1997) demonstrated that tidal interactions with a low-mass companion are one means of producing a corrugated disk structure. In this model, the companion generates a spiral-shaped wave  that propagates inwards from the edge of the disk, resulting in perpendicular motions. Although the simulations of Edelsohn \\& Elmegreen were performed for a  Milky Way-like galaxy and a perturber roughly 8\\% as massive, the results exhibit qualitative similarities to our observations of IC~2233. However, MU08 found that IC~2233 appears to be an extremely isolated galaxy, exhibiting no obvious signs of a recent interaction or minor merger. This casts significant doubt upon tidal effects as being the primary driver of the vertical disk structure of IC~2233. In addition, since tidally induced corrugations are excited from the outer regions of the disk (see Edelsohn \\& Elmegreen 1997), this does not provide a natural explanation for the ``null'' seen at the kinematic center of IC~2233 (see \\S~\\ref{data} and Figure~\\ref{fig:corplots}). Interestingly, at least one of the galaxies found by Florido et al. (1991) to have disk corrugations (NGC~4244), also appears to be quite isolated (Dahlem et al. 2005). A search for corrugations in additional isolated systems, coupled with deep searches for low-mass companions, could provide better statistical constraints on the likely role of companions in exciting corrugated patterns.  While an intrusion from a companion galaxy seems unlikely in the case of IC~2233, the influence of impacts from very low-mass perturbers ($\\lsim10^{6}~M_{\\odot}$), such as high-velocity clouds, cannot be entirely ruled out. Based on their energetics, high-velocity cloud impacts were the explanation favored by Alfaro et al. (1991) to account for vertical displacements of the midplane of the Milky Way and their associated star formation activity (see also Franco et al. 1988).  Moreover, IC~2233 is known to have a vertically extended \\HI\\ component that is kinematically distinct from the main disk component (MU08), and whose presence suggests that  the circulation of \\HI\\ material between the plane and higher galactocentric latitudes is actively occurring (e.g., Fraternali \\& Binney 2008). While it would be difficult to explain how a {\\it global} corrugation structure as regular as that seen in the \\HI\\ disk of IC~2233 could be explained by a series of cloud impacts, such  events may help to account for {\\it local} vertical displacements and offsets between certain disk components. For example, such an event offers one possible explanation for the differences in the vertical structure of various disk tracers that is observed over a portion of the northern half of IC~2233 as well as the perturbation in the \\HA\\ rotation curve observed near this position (see \\S~\\ref{results}).  \\subsection{Galactic Bores} The results presented in Figure~\\ref{fig:velres} seem to exclude another class of model for explaining the vertical meanderings of the IC~2233 disk. In particular, our findings appear to be inconsistent with a scenario where corrugations result from gas being hydraulically ``lifted'' off the plane by a spiral arm as a result of interaction of the spiral density wave with a thick, magnetized, gas disk (Martos \\& Cox 1998). In this so-called ``galactic bore'' model, the lifting effect is expected to show an even symmetry about midplane  (see Alfaro et al. 2001), resulting in some thickening of the plane and no net velocity excursions in an edge-on galaxy.  \\subsection{Parker Instabilities} In contrast to the galactic bore picture, the undular mode of the Parker (1966) instability in a thick, magnetized gas disk predicts a velocity pattern similar to that seen in Figure~\\ref{fig:velres} (see Franco et al. 2002). However, an additional prediction of this model is that density enhancements should be present at the locations of maximum deviation from the mean plane (e.g., Hanawa 1995). Indeed, such a trend has been observed in the corrugated Carina-Sagittarius arm of the Galaxy (Franco et al. 2002). While IC~2233 does show a series of \\HI\\ density enhancements along its disk (MU08), we find no clear correlation between the locations of these clumps and the maxima in the corrugated pattern in the \\HI\\ disk. Because corrugations resulting from the undular mode of the Parker instability occur along the azimuthal direction, line-of-sight integration effects may make it more difficult to match the location of density enhancements and vertical displacements in an edge-on galaxy. However, for the case of IC~2233, the Parker instability  model encounters an additional challenge in that magnetic fields in small spiral galaxies with rotational velocities comparable to IC~2233 are typically quite weak (e.g., Chy{\\. z}y et al. 2007).  \\subsection{Spontaneous Gravitational Instabilities} The observation that the corrugated structure of IC~2233 is more pronounced in the younger stellar components than in the old stars would seem to argue against a purely gravitational origin for the corrugations, since both gas and stars will be influenced by gravity. However, gas is in general more dynamically responsive than stars owing to its dissipative nature and its lower velocity dispersion (e.g., Kim \\& Ostriker 2007). In addition, gaseous distortions may be amplified by self-gravity and/or magnetic effects (e.g, Edelsohn \\& Elmegreen 1997). For these reasons, even if the corrugations have a purely gravitational origin, it is not entirely unexpected to find the largest amplitude corrugations in the gaseous components  of the galaxy. In this picture, the youngest stars (which still reside near their birth locations) would be expected to exhibit a corrugated structure, while dynamical evolution and heating would slowly erode these signatures in the older stars.  To determine whether an origin for the corrugations in IC~2233 via gaseous instabilities is feasible, it is valuable to compare the ages of the stellar populations participating in the corrugations with relevant dynamical timescales in the disk. For example, using equation 6-85 from Binney and Tremaine (1987), we can estimate the oscillation period, $P$, for stars in the disk of IC~2233. Assuming a  mass equal to the dynamical mass of the galaxy ($M_{d}=1.55\\times10^{10}~M_{\\odot}$), a disk radius of 9.2~kpc, and a maximum rotational velocity $V_{\\rm max}=85$~\\kms\\ (MU08), then at intermediate galactocentric radii (near $r$=4.5~kpc), we find $P\\sim5\\times10^{6}$~years. This is comparable to the lifetime of an \\HII\\ region, implying that the hot young stars giving rise to these ionized regions must be born near their present off-planar locations. If the corrugations arise initially in the gas, then the presence of corrugations in the NUV and FUV light (corresponding to stars with ages $<10^{8}$ years) implies that the midplane deviations need to persist for only a few oscillation periods to be visible in these bands. The mix of stars dominating the $R$-band light in IC~2233 is more uncertain, although the very blue color of this galaxy ($B-R$=0.67; MU08) suggests that this band may too be weighted toward moderately young stars, with ages no more than a few oscillation periods.  As discussed by Fridman \\& Polyachenko (1984), vertical displacements in a galactic disk may arise  from two types of internal gravitational mechanisms: spiral density waves and vertical bending oscillations (see also Fridman et al. 1998). Variations of both types of mechanisms could be linked with disk corrugations, even in galaxies with negligible magnetic fields. For example, Nelson (1976) suggested that corrugations could arise from the perturbations of a spiral potential on a differentially rotating gas layer (see also Nelson \\& Matsuda 1980), while Masset \\& Tagger (1997) postulated that their origin could be the result of a non-linear coupling between a spiral density wave and two so-called ``warp waves''. In the model of Masset \\& Tagger, a transmitted wave results in an integral sign warp, while a reflected wave produces corrugations. Several arguments suggest that IC~2233 is likely to exhibit at least rudimentary spiral structure (see MU08), but unfortunately its edge-on orientation makes it impossible to discern the details of this spiral pattern and thus to examine any possible relationship to the observed corrugations.  Vertical bending oscillations in galaxy disks have been explored by various authors as a possible origin for galaxy warps (e.g., Sellwood 1996; Revaz \\& Pfenniger 2004), but to our knowledge, relatively little attention has been paid to these oscillations as a possible origin of disk corrugations (although see Sparke 1995; Fridman et al. 1998; Griv \\& Chiueh 1998). While the sophisticated numerical simulations needed to explore this phenomenon are beyond the scope of the present paper, we point to the results of some existing studies that suggest that such an exploration might be fruitful. For example, in certain galaxy disk models, Sellwood (1996) found higher order bending modes whose amplitude and wavelength are reminiscent of the corrugation pattern of IC~2233 (e.g., his Figure~1d), although it remains to be explored how these results would translate to a disk containing both gas and stars. More recently, Revaz \\& Pfenniger (2004) used multi-component disk models to show that bending instabilities can be long-lived (i.e., persisting through several rotations) provided that disks have sufficient self-gravity. In their models it is assumed that some portion of the galaxy's dark matter resides in the disk rather than in a halo. Revaz \\& Pfenniger explored only the low-order ($m=0$ and $m=1$) bending modes relevant to producing galaxy warps; however, a search for higher-order modes, as well as an extension of their models to include the dissipational behavior of the gas would be of interest."
    },
    "0807/0807.1109_arXiv.txt": {
        "abstract": "We use a series of high-resolution N-body simulations of a `Milky-Way' halo, coupled to semi-analytic techniques, to study the formation of our own Galaxy and of its stellar halo. Our model Milky Way galaxy is a relatively young system whose physical properties are in quite good agreement with observational determinations. In our model, the stellar halo is mainly formed from a few massive satellites accreted early on during the galaxy's lifetime. The stars in the halo do not exhibit any metallicity gradient, but higher metallicity stars are more centrally concentrated than stars with lower abundances.  This is due to the fact that the most massive satellites contributing to the stellar halo are also more metal rich, and dynamical friction drags them closer to the inner regions of the host halo. ",
        "introduction": "Our own galaxy - the Milky Way - is a fairly large spiral galaxy consisting of four main stellar components: (1) the thin disk, that contains most of the stars with a wide range of ages and on high angular momentum orbits; (2) the thick disk, that contains about 10-20 per cent of the mass in the thin disk and whose stars are on average older and have lower metallicity than those in the thin disk; (3) the bulge, which contains old and metal rich stars on low angular momentum orbits; and (4) the stellar halo which contains only a few per cent of the total stellar mass and whose stars are old and metal poor and reside on low angular momentum orbits.  While the Milky Way is only one Galaxy, it is the one that we can study in unique detail. Over the past years, accurate measurements of ages, metallicities and kinematics have been collected for a large number of individual stars, and much larger datasets will become available in the next future thanks to a number of ongoing and planned astrometric, photometric and spectroscopic surveys. This wealth of detailed and high-quality observational data provides an important benchmark for current theories of galaxy formation and evolution.  In the following, we outline the main results of a recent study of the formation of the Milky Way and of its stellar halo in the context of a hybrid cosmological approach which combines high-resolution simulations of a `Milky Way' halo with semi-analytic methods. We refer to \\cite[De Lucia \\& Helmi (2008)]{DeLucia_Helmi_2008} for a more detailed description of our method and of our results. ",
        "conclusions": "We have combined high-resolution resimulations of a `Milky Way' halo with semi-analytic techniques to study the formation of our own Galaxy and of its stellar halo. The galaxy formation model used in our study has been used in a number of previous studies and has been shown to provide a reasonable agreement with a large number of observational data both in the local Universe and at higher redshifts (De Lucia \\& Helmi 2008 and references therein). Our study demonstrates that the same model is able to reproduce quite well the observed physical properties of our own Galaxy. The agreement is not perfect: our model Galaxy contains about twice the gas observed in the Milky Way, and the model bulge is slightly less massive and substantially less metal rich than the Galactic bulge. A detailed comparison between model results and observational measurements of metallicity distributions is complicated by the use of an instantaneous recycling approximation which is not appropriate for the iron-peak elements, mainly produced by supernovae Type Ia. Relaxing of this approximation in future work, will allow us to carry out a more detailed comparison with observed chemical compositions, and to establish similarities and differences between present-day satellites and the building blocks of the stellar halo.  Our model stellar halo is made up of very old stars (older than $\\sim 11$~Gyr) with low metallicity, although higher than what is known for the stellar halo of our Galaxy. Most of the stars in the halo are contributed by a few relatively massive satellites accreted early on during the galaxy's lifetime. The building blocks of the stellar halo lie on a well defined mass-metallicity relation. Since the most massive galaxies are dragged closer to the inner regions of the halo by dynamical friction, this produces a stronger concentration of more metal rich stars, in qualitative agreement with recent observational measurements. The numerical resolution of our simulations is too low for the study of spatially and kinematically coherent structures in model stellar halo. Higher resolution simulations are needed for this kind of study."
    },
    "0807/0807.3335.txt": {
        "abstract": "{In a previous work, three bright mid-infrared/radio sources were discovered in the nuclear region of startburst/AGN galaxy NGC1365.} {The present study aims at confirming that these sources are indeed young and massive ``embedded'' clusters, and at deriving their physical parameters, such as extinction, age and mass.}  {Using ISAAC and VISIR at the VLT we obtained maps and low resolution spectra in the near- and mid-infrared. The resulting datasets are first interpreted by comparing the observations with images and spectra of the close-by young cluster R136 in the Large Magellanic Cloud and then by using model predictions for both the nebular emission lines and the spectral energy distribution of the sources.}  {We produce maps of the region containing the three sources in the R, J, Ks, L' bands and at 12.8\\micro~and perform their accurate relative positioning. We also provide spectra in the ranges 1.8-2.4\\micro, 3.3-4.0\\micro, 8.1-9.3\\micro~and 10.4-13.2\\micro. The spectral energy distribution of the three sources rises with wavelength.  Emission lines from ionised hydrogen and molecular hydrogen are detected, as well as PAH emission.  Conspicuous [NeII]12.8\\micro~line is also present, while neither the [ArIII]\\,8.9\\micro~nor the [SIV]\\,10.4\\micro~lines are detected.  This provides a stringent constraint on the age of the sources: we argue that they are relatively evolved young clusters (6-8\\,Myr). Owing to their ionising photon emission rates and ages, they must be extremely massive clusters (of the order of $10^7$\\,\\msol). Their mid-infrared spectral energy distribution suggests the presence of two components: (1) an optically thin component, with a continuum comparable to that of R136, and (2) an optically thick component which might be related to subsequent or on-going episodes of star formation. We anticipate that these sources are good candidates for evolving according to a bi-modal hydrodynamical regime, in which matter is trapped at the centre of a compact and massive cluster and generates further star formation.}  {} ",
        "introduction": "\\label{Introduction} Starburst regions in close-by galaxies were first resolved in a population of star clusters in the early nineties, thanks to the high angular resolution of the Hubble Space Telescope. The first galaxy-target, NGC1275 \\citep{Holtzman92}, showed a population of young and massive compact clusters.  Soon after, similar objects were encountered in a wide variety of environments, such as dwarf galaxies \\citep{O'Connell94,Hunter94,O'Connell95,Leitherer96,Gorjian96}, interacting galaxies \\citep{Whitmore93,Conti94,Shaya94,Whitmore95,Meurer95} and circumnuclear star-forming rings \\citep{Benedict93,Barth95,Bower95,Maoz96}.  Their possible parental link with classical and well-studied globular clusters was proposed and, in subsequent studies, these young massive clusters (abbreviated YMCs) became increasingly referred to as objects likely to evolve into globular clusters after a few Gyrs. Their masses are greater than $10^5$\\,\\msol, with radii smaller than 5\\,pc and ages below 100\\,Myr. In the literature, the term super star cluster (SSC) is often associated with YMCs which are bright in the visible, hence suffer little extinction. We prefer to use the generic term YMC in all cases, and specify, whenever needed, whether the YMC is still dust embedded (embedded YMC) or naked (UV-bright YMC).  The youngest YMCs discovered so far may be younger than 1\\,Myr, and show up as heavily dust-embedded HII regions. Examples of such extragalactic embedded YMCs known to date are still scarce: some were found and discussed in the Antennae galaxies NGC4038/39 \\citep{Mirabel98,Gilbert00}, in Henize 2-10 \\citep{Kobulnicky99}, in NGC5253 \\citep{Gorjian01}, in SBS 0335-052 \\citep{Plante02}, in IIZw40 \\citep{Beck02}, in NGC1365 and NGC1808 \\citep{Galliano05a,Galliano08a}, in NGC7582 \\citep{Wold06}.  As proposed in \\citet{Johnson04}, parallel evolutionary sequences can be imagined for the formation of massive stars and for the formation of massive clusters. Both types of objects start as HII regions deeply embedded in a dust cocoon. In the case of a massive star, this stage is identified as ultra compact HII region (UCHII region) and in the case of a YMC as ultra dense HII region (UDHII region) \\citep{Kobulnicky99}. They are both inconspicuous in the visible and near-infrared (NIR), while bright in the mid-infrared (MIR) and far-infrared (FIR). They are also intense sources of thermal radio continuum as well as of line emission from ionised gas. The embedding material eventually dissipates off and they become detectable as UV-bright sources. The extent to which this parallel is sustainable remains an open question: similarities and differences along the two sequences are worth investigating and might bring clues about the conditions for star formation in galaxies, and in particular for the formation of massive star clusters and globular clusters.  We present in this paper the first step of a thorough analysis and modelling of embedded YMCs in nearby spiral galaxies. The three YMCs we are interested in here are located in the starburst circumnuclear region of the Seyfert2 galaxy NGC1365 (distance 18.6\\,Mpc, hence 1\\arcsec~corresponds to 90\\,pc). \\citet{Galliano05a} discovered these sources in the MIR and demonstrated that they coincide with bright thermal centimetre radio sources detected by \\citet{Sandqvist95}. They have also been recently detected in the CO molecule by \\citet{Sakamoto07}. These authors infer a mass of molecular material of the order of 10$^9$\\msol~in the central 2\\,kpc diameter region, which includes the three YMCs under discussion.  In \\citet{Galliano05a}, their ages were grossly estimated from their radio spectral indexes and found to be of a few (~3) Myrs. Accordingly, their masses were estimated to be of the order of 10$^6$\\,\\msol.  These clusters are located at the inner Linblad resonance (ILR) in their host galaxy, and also within around~1\\,kpc from its active galactic nucleus (AGN). Their environment has a slightly above solar metallicity, as expected from the observed metallicity gradient in the disc of NGC1365 discussed by e.g. \\citet{Dors05}.  The questions to be addressed in the current paper remain basic ones. What are the spectral characteristics of such YMCs? Are NIR and/or MIR data sufficient to get a trustable insight on the properties of embedded YMCs? Do we miss a substantial part of the phenomenon by observing only in the NIR? Which physical parameters can be safely derived for these objects, either in a direct way or through a comparison with models?  The paper is structured as follows: the first part is devoted to the presentation and discussion of the data, while subsequent sections deal with the physical interpretation of the sources. In Sec.~\\ref{Data}, we describe the acquisition and reduction of the dataset, which consists of NIR/MIR images and spectra, all collected at the ESO/VLT using the instruments ISAAC and VISIR. We first discuss the images and perform the relative registration of the maps at different wavelengths (Sec.~\\ref{Registration}). We then describe and discuss the spectra (Sec.~\\ref{Spectra}), which display intense nebular lines and NIR/MIR rising continua. In Sec.~\\ref{Measurements}, we discuss the uncertainties on the quantities measured from the images and from the spectra.  In Sec.~\\ref{R136}, we perform a comparison of the data for the three YMCs in NGC1365 with comparable data obtained for R136 in the LMC, one of the nearest known YMC.  To do so, we use WFI and IRAC images, an ISO SWS spectrum and a wide field ISO CVF spectrum of R136. In Sec.~\\ref{YMC basic parameters}, we derive basic parameters for the YMCs based upon their emission lines, via a comparison with predictions from a library of photo-ionisation models generated with the code \\texttt{CLOUDY}.  In Sec.~\\ref{SED}, we analyse the NIR/MIR infrared emission of the sources, both their spectral energy distribution (SED) and their line emission, performing a more complex modelling with the dusty stellar population evolution code \\texttt{GRASIL}.  Then, in Sec~\\ref{specific evolution}, we attempt to position our results in the light of the theoretical evolution of very massive clusters, considering the bi-modal hydrodynamic solution for re-inserted matter, as proposed by \\citet{Silich07}~and \\citet{Tenorio07}. Finally, the conclusions and perspectives of our work are highlighted in Sec.~\\ref{conclusion}.   %----------------------------------------------------------------------------------------------------------------------------------- ",
        "conclusions": ""
    },
    "0807/0807.1244.txt": {
        "abstract": "We create new U, B, V, R and I-band light curve templates of type Ia Supernovae (SNe Ia) and re-analyze 122 nearby (redshift $<$ 0.11) SNe Ia using a new ``Multi-band Stretch method,'' which is a revised Stretch method (cf. \\citealt{per97,gol01}) extended to five bands. We find (i) our I-band template can fit about 90\\% of SNe Ia I-band light curves, (ii) relationships between luminosity, colours and stretch factors, (iii) possible sub-groups of SNe Ia, and (iv) the ratio of total to selective extinction $R$ in other galaxies can be consistent with that in the Milky Way under the assumption that SNe Ia have diversity in their intrinsic colour. Based on these results, we discuss how to select subsets of SNe Ia to serve as good distance indicators for cosmology. We find two possibilities: one is to choose ``BV bluest'' ($-0.14 < (B-V)_{max} \\le -0.10$) objects and the other is to use only SNe Ia which occur in E or S0 galaxies. Within these subsets, we find the root mean square (r.m.s.) of peak B-band magnitudes is 0.17 mag (``BV bluest'' sample) and 0.12 mag (E or S0 sample). ",
        "introduction": "SNe Ia are some of the most important objects in observational cosmology because of their role as distance indicators. SNe Ia are bright enough to be observed at high redshifts, they have uniform peak luminosity (cf.\\citealt{bra92}), and they are expected to evolve less from low z to high z than other objects such as galaxies (cf. \\citealt{rie99b}). Many studies have used SN Ia as distance indicators to show the accelerating expansion of the universe (e.g \\citealt{rie98,per99,kno03,rie04,ast06,woo07}). All of these works are based on the empirical relationship between light curve shapes and peak luminosity (cf. \\citealt{phi93}). From the 1990s onward, as more low-redshift SNe Ia have been found, the relation has been revised (cf. \\citealt{ham96,rie99a,alt04,rei05}).  Recently, the number of well-measured SNe Ia has grown larger and larger, especially at intermediate and high redshifts. There are many supernova survey programs currently running; for example, the Lick Observatory and Tenagra Observatory Supernova Searches (LOTOSS\\footnote{http://astro.berkeley.edu/\\~{}bait/lotoss.html}), the Nearby Supernova Factory (SNfactory, \\citealt{ald02}), the SDSS-II Supernova Survey (SDSS SN Survey, \\citealt{fri08,sak08}), Equation of State: SupErNovae trace Cosmic Expansion (ESSENCE, \\citealt{woo07}), and the SuperNova Legacy Survey (SNLS, \\citealt{pri05}). Most of SNe Ia found in these surveys are observed in multiple passbands and over many epochs. These high quality samples have given us clues to understand the nature of SN Ia. For example, \\cite{ham96c} showed that brighter SNe Ia, like SN1991T, tend to occur in galaxies with active star formation, while fainter SNe Ia, such as SN1991bg, tend to occur in galaxies which lack star formation (see also \\citealt{how01,van05,sul06}). \\cite{ben05} discussed the intrinsic diversity of the spectroscopically normal SN Ia (called \"Branch normal\", \\citealt{bra93}). \\cite{qui07} pointed out there are two subgroups within Branch normal SN Ia. Unusual SNe Ia have also been found: SN2000cx and SN2002cx were reported as an extreme sample of spectroscopically peculiar SN Ia \\citep{li01,li03}.; \\cite{how06} reported SNLS-03D3bb was a Super-Chandrasekhar-mass supernova with an exceptionally high luminosity and low kinetic energy. H$\\alpha$ emission in the spectra of SN2002ic and SN2005gj at early phases indicates the existence of circumsteller material around the progenitor (\\citealt{den04,ald06} and so on).  Although SNe Ia seem to be good standard candles for cosmological studies, there are several issues we must understand in order to use them optimally. The main issues are (1) intrinsic photometric and spectroscopic diversity of SNe Ia and (2) extinction in host galaxies. Both issues severely limit the use of SNe Ia as a standard candles if we can not deal with them properly. Recent studies have revealed an intrinsic diversity of SNe Ia; for example, \\cite{man05} and \\cite{sul06} showed there are two types of SN Ia, called ``prompt'' and ``delayed'', which differ in the distribution of their stretch factors, and \\cite{ber06} demonstrated how spectroscopic diversity can affect significantly the intrinsic colours of SNe Ia and K-corrections. The latter issue is important if we want to use larger SN Ia samples. Correction of extinction by dust in the host galaxy is necessary because almost all SNe Ia may be affected by the dust. Previous works insisted that the ratio of total to selective extinction, $R_{B}$, in other galaxies is smaller than that of the Milky Way; see \\cite{alt04} (hereafter ALT04), \\cite{phi99}, and \\cite{kno03}.  In this paper, we discuss these issues using a sample of nearby SNe Ia. We re-analyze published data on nearby SNe Ia with an improved ``Multi-band Stretch method'' which is based on the usual ``stretch method'' (\\citealt{per97,gol01}). ",
        "conclusions": "In this work, we re-analyze U, B, V, R and I-band light curves of 122 previously-published nearby (redshift $<$ 0.11) SNe Ia, which are mainly from JHA06, \\cite{rie99a} and \\cite{ham96}. We devise the ``Multi-band Stretch method'', which is a revised Stretch method extended to fitting up to five bands. We create new U, B, V, R and I-band light curve templates from well-observed SNe Ia for our ``Multi-band Stretch method''. It is known that the I-band light curves of SNe Ia shows much variation. However, we found our I-band template fitted about 90\\% of SNe Ia I-band light curves very well.  Using the results of our light curve fitting, we examine the relationships between luminosity, colours at maximum light, and stretch factors. In \\S5, we show a possible sub-group of SNe Ia which have different colours from the main group of SN Ia; this sub-group tends to be found in E or S0 galaxies. In \\S6, we find the relationship between inverse B-band stretch factor and U, B, V, R and I-band luminosity, using colour-selected SNe Ia to avoid contamination from the outlying sub-group. We obtain a stretch - magnitude relationship for ``BV bluest'' SNe Ia without colour correction as $(2.28 \\pm 0.42) \\times s_{(B)}^{-1} - (21.49 \\pm 0.41)$, which is steeper than the slope of the corresponding stretch-magnitude relation in ALT04. The r.m.s. of the residuals from this relationship is 0.17 mag. If we attempt to reduce the dispersion in luminosity by using colour corrections, we find that for distant (redshift $>$ 0.02) SNe Ia, $M_B = 0.98 \\times s_{(B)}^{-1}- 2.28 \\times (B - V)_{max} - 19.95$ with an r.m.s. of 0.27 mag. The r.m.s. becomes 0.12 mag when we select only SNe Ia hosted by E or S0 galaxies. In \\S7, we statistically investigate the properties of dust in host galaxes, and find that the ratio of total to selective extinction $R$ can be consistent with that of dust in the Milky Way if we assume that SNe Ia have some diversity in their intrinsic colours.  Based on these results, in \\S8, we discuss how to select subsets of SNe Ia which may be good distance indicators for cosmology. We find two possibilities: one is to use only ``BV bluest'' SNe Ia with broad light curves, and the other is to use only SNe Ia inside E or S0 galaxies. To use these samples for precision cosmology, we need to study the properties of SNe Ia from the low-z to the high-z universe."
    },
    "0807/0807.2566_arXiv.txt": {
        "abstract": "{ Very long baseline interferometric observations of the radio galaxy 3C~120 show a systematic presence of gradients in Faraday rotation and degree of polarization across and along the jet. These are revealed by the passage of multiple superluminal components throughout the jet as they move out from the core in a sequence of 12 monthly polarimetric observations taken with the VLBA at 15, 22, and 43 GHz. The degree of polarization has an asymmetric profile in which the northern side of the jet is more highly polarized. The Faraday rotation measure is also stratified across the jet width, with larger values for the southern side. Superposed on this structure we find a localized region of high Faraday rotation measure ($\\sim$ 6000 rad m$^{-2}$) between approximately 3 and 4 mas from the core. This region of enhanced Faraday rotation may result from the interaction of the jet with the ambient medium, which may also explain the stratification in degree of polarization. The data are also consistent with a helical magnetic field in a two-fluid jet model, consisting of an inner emitting jet and a sheath of nonrelativistic electrons. ",
        "introduction": "Helical magnetic fields may play an important role in the dynamics and emission of relativistic jets in active galactic nuclei (AGN), specially in the formation and collimation processes \\citep{2002Sci...295.1688K,2008Natur.452..966M}. If jets are surrounded by a sheath of nonrelativistic electrons, it is possible to search for these helical magnetic fields by looking for Faraday rotation measure (RM) gradients across the jet \\citep*[e.g.,][]{Blandford:1993fk}. Such gradients have been observed across the jet in 3C~273 \\citep{2002PASJ...54L..39A,2008ApJ...675...79A,2005ApJ...626L..73Z,2005ApJ...633L..85A} and other sources \\citep[e.g.,][]{2004MNRAS.351L..89G}, however they do not seem to be a universal feature \\citep{2003ApJ...589..126Z}.  \\begin{figure*}[t!] \\center{\\includegraphics[scale=0.85]{Gomez_fig1.eps}} \\caption{\\footnotesize Map of the mean value of the rotation measure across epochs. Data with standard deviation larger than 1000 rad m$^{-2}$ were discarded. Bars indicate the mean value of the RM-corrected EVPAs, with all displayed pixels having a standard deviation smaller than 30$^{\\circ}$ (96\\% under 20$^{\\circ}$). Contours show the 22 GHz total intensity at epoch 2001.00 for reference. The thick lines indicate the direction of the slices shown in Fig.~\\ref{slices}.} \\label{rm_mean} \\end{figure*}  Thanks to its proximity ($z=0.033$), Very Long Baseline Array (VLBA) observations of the radio galaxy 3C~120 are capable of resolving the jet across its width, revealing a very rich structure in total and polarized flux \\citep{1998ApJ...499..221G,1999ApJ...521L..29G,2000Sci...289.2317G,2001ApJ...561L.161G,2008ApJ...681L..69G,2001ApJ...556..756W,2005AJ....130.1418J,2002Natur.417..625M,2007ApJ...665..232M}. In addition, evidence for the presence of a helical magnetic field has been found in 3C~120 by analyzing the motion and polarization of superluminal components \\citep{2001ApJ...561L.161G,2005ApJ...620..646H}. ",
        "conclusions": ""
    },
    "0807/0807.2799_arXiv.txt": {
        "abstract": "Observations of the intergalactic medium (IGM) suggest that quasars reionize HeII in the IGM at $z \\approx 3$.  We have run a set of $190$ and $430$ comoving Mpc simulations of HeII being reionized by quasars to develop an understanding of the nature of HeII reionization and its potential impact on observables.  We find that HeII reionization heats regions in the IGM by as much as $25,000 \\, \\Kelvin$ above the temperature that is expected otherwise, with the volume-averaged temperature increasing by $\\sim 12,000 \\, \\Kelvin$ and with large temperature fluctuations on $\\sim 50$ Mpc scales.  Much of this heating occurs far from quasars by photons with long mean free paths.  We find a temperature-density equation of state of $\\gamma -1 \\approx 0.3$ during HeII reionization, but with a wide dispersion in this relation having $\\sigma_{T} \\sim 10^4$ K.  HeII reionization by the observed population of quasars cannot produce an inverted relation ($\\gamma - 1 < 0$).  Our simulations are consistent with the observed evolution in the mean transmission of the HeII Ly$\\alpha$ forest.  We argue that the heat input from HeII reionization is unable to cause the observed depression at $z \\approx 3.2$ in the HI Ly$\\alpha$ forest opacity as has been suggested.  We investigate how uncertainties in the properties of QSOs and of HeII Lyman-limit systems influence our predictions. ",
        "introduction": "\\label{firstpage}  In the standard picture for the reionization history of the Universe, radiation from Population II stars ionized the intergalactic HI at $z>6$ as well as the HeI, converting the vast majority of the intergalactic helium to HeII.  However, these stars cannot ionize HeII, and at $z \\approx 3$ quasars, with their harder UV spectrum, doubly ionize the intergalactic helium.  To test this model, many observations are targeting high redshifts to probe hydrogen reionization (e.g., \\citealt{fan02, taniguchi05, bouwens07, kashikawa06, stark07, totani06}).  In this picture, HeII reionization occurs at redshifts for which there is substantially more data on the state of the intergalactic medium (IGM).   In fact, a number of observations suggest that HeII reionization happened at $z \\sim 3$.  Two measurements of the mean transmission in the HI Ly$\\alpha$ forest have noted an upward bump at $z \\approx 3.2$ \\citep{bernardi03, faucher07}, which \\citet{theuns02} interpreted as arising from a temperature increase of the IGM during HeII reionization (but see \\citet{faucher07} for alternative explanations). An increase in the average temperature of the IGM would also decrease the small-scale fluctuations in the HI Ly$\\alpha$ forest. \\citet{ricotti00} and \\citet{schaye00} measured the temperature from the widths of the narrowest lines in the HI Ly$\\alpha$ forest and claimed to have detected a sudden increase in the temperature of $\\Delta T \\sim 10^4$ K between $z = 3.5$ and $3$.  Photo-heating during HeII reionization is the only known process that could be responsible for such an effect (e.g., \\citealt{miralda94, abel99}).  However, a subsequent study by \\citet{mcdonald01b} using a similar method and \\citet{zaldarriaga01} using the HI Ly$\\alpha$ forest power spectrum did not confirm this sudden increase in temperature, but rather found a constant temperature at mean density of $T_0 \\approx 17,000$ K for $2 < z < 4$. Temperatures of $17,000\\; K$ are difficult to explain without HeII reionization occurring at $z \\sim 3$ \\citep{hui03}, and it is unclear whether as sudden an increase in temperature as \\citet{ricotti00} and \\citet{schaye00} find is even expected theoretically.  If a substantial fraction of the helium is in HeII ($\\gtrsim 1\\%$), this would produce a Gunn-Peterson absorption trough in the spectra of high-redshift quasars at wavelengths blueward of HeII Ly$\\alpha$. Observations of HeII Ly$\\alpha$ forest absorption at $2.8 < z < 3.3$ find $10$s of comoving Mpc regions with no detected transmission \\citep{jakobsen94, davidsen96, hogan97, reimers97, heap00}.  These troughs may signify the presence of diffuse intergalactic HeII. However, current data, which consist of only a few quasar sight-lines, do not rule out the intergalactic HeII being primarily ionized and in photo-ionization equilibrium with a weak background \\citep{giroux97, fardal98,heap00}. The Cosmic Origins Spectrograph, which NASA plans to install on the Hubble Space Telescope in 2009, will increase the quantity and quality of HeII Ly$\\alpha$ forest sight-lines.  As the intergalactic HeII becomes progressively more ionized, the extragalactic UV background will harden around the ionization energy of HeII at $54.4 \\; \\eV$.  This hardening will affect the ionization state of intergalactic metals.  \\citet{songaila98} observed a sharp evolution at $z \\approx 3$ in the column density ratios in SIV (Ionization Potential $= 45.1$ eV) to CIV ($64.5$ eV) absorbers. \\citet{boksenberg03} found evidence for a more gradual hardening of the background between $2 < z <4$ from the column density ratios of NV ($98 \\, {\\rm eV}$) to CIV.  Finally, by simultaneously fitting to multiple metal lines that originate from the same absorption systems, \\citet{agafonova05} and \\citet{agafonova07} inferred a background spectrum that is hardening at $z \\approx 3$ near $4$ Ry.     Measurements from the $z \\sim 3$ HI Ly$\\alpha$ forest ignore the effects of a patchy HeII reionization process.  For example, estimates of the photo-ionizing background and the IGM temperature from the forest assume a power-law temperature-density relation. Cosmological parameter studies from the HI Ly$\\alpha$ forest power spectrum make a similar assumption. Different regions can have vastly different HeII reionization histories, resulting in a more complicated distribution of temperatures and pressure-smoothing scales than is commonly adopted. Realistic simulations of HeII reionization will help quantify the level at which this process biases these measurements.\\footnote{Of note, \\citet{lai06} used simple models for HeII reionization to show that it has a surprisingly small effect on the Ly$\\alpha$ forest power spectrum on large scales, modifying it at a $ \\lesssim 5\\%$ level for wavevectors $k \\lesssim 5 \\, {\\rm comoving} \\; \\Mpc^{-1}$.}  The aim of this paper is to run realistic simulations of HeII reionization to understand the morphology of this process as well as its effect on observables. We concentrate primarily on its impact on quantities that are sensitive to the IGM temperature, but we also study the effect of HeII reionization on the transmission in the HeII Ly$\\alpha$ forest.  \\citet{sokasian02} and \\citet{paschos07} have performed the most realistic simulations of patchy HeII reionization to date.  There are several differences between our work and these earlier investigations.  Both of these studies employed volumes $\\leq 100^3$ comoving $\\Mpc^3$.  Here, we examine HeII reionization in $186^3$ and $429^3$ comoving Mpc$^3$ volumes, providing a more representative cosmic sample.  However, both \\citet{sokasian02} and \\citet{paschos07} simulated HeII reionization as a post-processing step on top of cosmological simulations that included gas dynamics.  Our study instead uses N-body simulations, which result in a less realistic model for the gas distribution, but afford a larger dynamic range.  Furthermore, \\citet{sokasian02} assumed sharp ionizing fronts and ignored the detailed temperature evolution. \\citet{paschos07} did not calculate the gas temperature self-consistently.  Our calculations capture the width of the ionization fronts and the temperature in a consistent manner. Finally, in contrast to previous studies, our work presents a large set of radiative transfer simulations in order to survey the parameter space.   In Section \\ref{code}, we describe the details of our code.  The models for the quasar sources are described in Section \\ref{sources}. Section \\ref{sims} presents the simulations.  Finally, Section \\ref{observations} addresses the implications HeII reionization has on observations of the HI and HeII Ly$\\alpha$ forests.  Throughout, we use a $\\Lambda$CDM cosmology with $n_s = 1$, $\\sigma_8 = 0.8$, $\\Omega_m = 0.27$, $\\Omega_{\\Lambda} = 0.73$, $\\Omega_b = 0.046$, and $h = 0.7$, consistent with the most recent WMAP results \\citep{komatsu08}.  All distances are in comoving units unless specified otherwise.  An overbar over a variable signifies a volume average, and $x_Y$ is the fraction of helium/hydrogen that is in ionization state $Y$. ",
        "conclusions": "We have run a set of simulations of HeII reionization to understand the structure of HeII reionization and its effect on several observables.  We find that for a late reionization of HeII by quasars: \\begin{itemize}  \\item The popular assumptions that the HeIII ionization fronts are sharp and that there is uniform heating within the front (and no heating outside of it) do not yield a realistic model for HeII reionization.  While the ionization fronts are still fairly localized, hard photons stream far from their sources and are absorbed ahead of the front.  These photons inject at large distances a significant fraction of the energy radiated above the HeII Lyman-limit.  \\item The average temperature at the mean density is increased by $\\approx 12,000 \\; {\\rm K}$ over the average temperature of the gas in the absence of HeII reionization for $\\bar{\\alpha}_{\\rm UV} = 1.6$. This temperature increase is consistent with estimates that assume that the IGM absorbs all photons with energies less than a few hundred eV during HeII reionization. Regions that are ionized last are ionized by the hardest radiation, reaching $T > 30,000 \\; {\\rm K}$ in our fiducial model. If the spectral index of the QSOs is different from our fiducial value of $\\bar{\\alpha}_{\\rm UV} = 1.6$, the average temperature can be be significantly altered.  \\item Poisson fluctuations in the number of QSOs rather than their spatial clustering shape the structure of the ionization and temperature fluctuations on $\\lesssim 50$ Mpc scales because rare $L \\sim L_*$ quasars ionize the HeII.  HeII reionization produces large temperature fluctuations on $50$ Mpc scales, and it results in the ionization and HeII photo-heating fluctuations being essentially uncorrelated with the density fluctuations on the Jeans scale, the scale most relevant for HI and HeII Ly$\\alpha$ forest statistics.  \\item Measurements of the $z \\sim 3$ forest suggest $T_0 \\approx 20,000\\, {\\rm K}$. Without invoking an exotic heating mechanism, a late HeII reionization epoch is required to produce this temperature. The amount of additional heat provided by HeII reionization in our simulations is enough to produce the inferred temperatures.   \\item To the extent that there is a $T$-$\\Delta_b$ relation, HeII reionization by quasars leads to a temperature-density equation of state of $\\gamma -1 \\approx 0.3$ for $0.1 \\lesssim \\bar{x}_{\\rm HeIII} \\lesssim 0.9$.  HeII reionization by QSOs cannot result in an inverted relation ($\\gamma - 1 < 0$) as has been claimed.  \\item Our simulations of a $z\\sim 3$ HeII reionization process produce a similar evolution in the HeII Ly$\\alpha$ mean transmission to what has been estimated.  Better observations of the mean transmission and its scatter will be able to rule out models presented here.  \\item In our simulations, the heating from HeII reionization by quasars is unable to produce any semblance of the observed depression in the $z \\approx 3.2$ HI Ly$\\alpha$ forest opacity.  The HeII reionization process is too extended in redshift ($\\Delta z > 1$) to be responsible for this narrow feature ($\\Delta z \\lesssim 0.4$).   \\end{itemize}  The morphology of HeII reionization is considerably different than that of hydrogen reionization (e.g., \\citealt{mcquinn07}).  During hydrogen reionization, current models find that the growth of ionized bubbles is more collective; hundreds or even thousands of galaxies within a bubble contribute to its growth, leading to the bubble structure tracing the distribution of galaxies and to tens of Mpc HII regions (e.g., \\citealt{zahn07}).  During HeII reionization, the growth is more stochastic.  Regions that happen to host a ``nearby'' quasar (within $\\sim 30$ Mpc) are ionized by that quasar. Because the QSO bubbles are so extended, the ionized structures are larger than during HI reionization.  Another significant difference stems from the m.f.p. of the ionizing photons to be absorbed in diffuse gas.  The spectrum from quasars is harder than that of stars -- our best guess for what ionizes the hydrogen -- the number density of helium is $70$ times smaller at $z= 3$ than hydrogen at $z= 6$, and the cross section is $4$ times smaller.  These three factors conspire to make the typical m.f.p. for a HeII ionizing photon megaparsecs rather than kiloparsecs, as it is for HI ionizing photons during hydrogen reionization.  We have seen that some HeII ionizing photons free stream far from their sources, partially ionizing and heating up these regions.  If stars reionize the hydrogen, ionization and heating occurs within HII bubbles.   A definitive identification of the redshifts of HeII reionization will place constraints on the sources that produce $> 4$ Ry photons and will aid studies of the HI Ly$\\alpha$ forest.  The data from previous observations, if analyzed properly, may be able to confirm whether quasars reionize HeII at $z \\approx 3$.  In addition, future observations of the HI and HeII Ly$\\alpha$ forest will soon be available with Sloan Digital Sky Survey III\\footnote{www.sdss3.org} and the Cosmic Origins Spectrograph on the Hubble Space Telescope. These telescopes will significantly increase the number of sight-lines in the HI and HeII Ly$\\alpha$ forests.  It is timely to make predictions for the effect of HeII reionization on these observations."
    },
    "0807/0807.1942_arXiv.txt": {
        "abstract": "The outer disks of galaxies present a unique laboratory for studying the process of disk formation. A considerable fraction of observed disks exhibit a break in their surface brightness profiles. The ubiquity of these features points to a crucial aspect of disk formation which must be explained. Recent theoretical work suggests that such breaks are related to significant amounts of radial migration. We discuss the current observational evidence which supports this picture. ",
        "introduction": "The majority of the thin disk forms out of gas quiescently cooling and collapsing inside the host dark matter halo following the last major merger \\citep{Brook:2004}. While the inner parts of the stellar thin disk are entangled with the pre-merger material, the outer parts evolve in relative solitude, modulo interactions with substructure components present within the host halo. The outer parts of galactic disks therefore provide us with a direct view of disk assembly in progress.  Since \\citet{van-der-Kruit:1979,van-der-Kruit:1987} it has been known that the surface brightness profiles of disk galaxies may not always follow a simple single-exponential law. In a recent work using data from the Sloan Digital Sky Survey (SDSS), \\citet{Pohlen:2006} showed that $\\sim60\\%$ of nearby disk galaxies have downward-bending surface brightness breaks, traditionally termed ``truncations''. Disk breaks have also been observed in the distant universe out to a redshift of z$\\sim$1 \\citep{Perez:2004,Trujillo:2005,Azzollini:2008a}, further implying that they are a generic feature of disk formation.  Several theories for the formation of downward-bending breaks have been suggested. The most common interpretations include angular momentum-limited collapse \\citep{van-der-Kruit:1987, van-den-Bosch:2001}, star formation threshold either due to a critical gas density \\citep{Kennicutt:1989} or a lack of a cool equilibrium ISM phase \\citep{elmegreen:1994,Schaye:2004}. Alternatively, breaks have also been attributed to angular momentum redistribution \\citep{Debattista:2006, Foyle:2008}. ",
        "conclusions": "We have argued that outer disk breaks are a phenomenon that is not only common in observed systems, but that its mechanism of formation may provide important insights into disk evolution as a whole. Substantial observational evidence taken in the context of recent theoretical models suggests that at least a fraction of outer parts of late-type galactic disks form due to substantial radial migration of fully-formed stellar material. Such evolution affects these extreme outer regions of galaxies, and profoundly impacts the properties of stellar populations in the entire disk."
    },
    "0807/0807.1490_arXiv.txt": {
        "abstract": "According to previous investigations the effect of diffusion in the stellar atmospheres and envelopes of subdwarf B (sdB) stars with luminosities $10 \\la L / L_{\\odot} \\la 100$ strongly depends on the presence of weak winds with mass loss rates $\\dot M \\la 10^{-12} M_{\\odot}/ \\rm yr$. These calculations with the mass loss rate as a free parameter have shown that it is hardly possible to reproduce the measured abundances of helium and metals simultaneously. A possible reason is the decoupling of metals, which preferably absorb the photon momentum, from hydrogen and helium in the wind region. In the present paper it will be investigated if ``chemically homogeneous\" winds, as assumed in previous investigations, with mass loss rates $\\dot M \\leq 10^{-12} M_{\\odot} / \\rm yr$ can exist. From an observational point of view the existence of weak winds in sdB stars is unclear. Only in the most luminous ones possible wind signatures have been detected. Therefore it will be investigated if according to the theory of radiatively driven winds the existence of weak winds is plausible. A stellar mass $M_{*} = 0.5 M_{\\odot}$ is assumed.  The results for effective temperatures $T_{\\rm eff} = 35000$, $30000$ and $25000 \\rm K$, metallicities $0.1 \\leq Z/Z_{\\odot} \\leq 1$ predict decreasing mass loss rates with increasing surface gravity. Dependent on the luminosity and metallicity the mass loss rates are between about $10^{-11} M_{\\odot} / \\rm yr$ and zero. If at all, chemically homogeneous winds can exist for the most luminous sdB stars only. For the other ones selective winds are expected which should lead to additional changes of the surface composition.  In sdB stars, hot white dwarfs and HgMn stars (which are chemically peculiar main sequence stars) the measured metal abundances are tendencially lower than the ones predicted from diffusion calculations which assume an equilibrium between gravitational settling and radiative levitation. Only for helium in almost all cases the measured abundances are larger than the predicted ones, but usually lower, below the solar value. This may be an indication that the abundance anomalies of metals are preferably due to the selective winds, whereas the helium deficiencies are due to gravitational settling, which for still unknown reasons is less effective than expected in an undisturbed stellar atmosphere. ",
        "introduction": "The abundance anomalies in subdwarf B (sdB) stars, white dwarfs and chemically peculiar main sequence stars are believed to be at least partially due to the effect of diffusion in the stellar atmosphere and envelope. Several attempts have been made to explain the abundances with the effect of gravitational settling which may be counteracted by radiative levitation. As the radiative force on an element decreases with increasing abundance due to saturation effects, the surface composition can be predicted from the equilibrium condition between the inward gravitational force and the outward radiative force. However, an agreement between predicted and measured abundances requires the absence of disturbing processes like mass loss or convective mixing. This may be the reason why in many cases the agreement is not satisfactory. A comparison between predicted and measured metal abundances for sdB stars \\citep[e.g.][]{t12_ber88,t12_ohl00,t12_chay06,t12_nat08} and hot white dwarfs \\citep[e.g.][]{t12_son02,t12_son05,t12_good05} shows that the measured abundances are tendencially lower than the predicted ones, although a few exceptions exist (e.g.\\ silicon in hot white dwarfs).  According to recent spectral analyses of sdB stars \\citep[e.g.][]{t12_gei08,t12_gei08b, t12_otol06,t12_edel06,t12_blan06}  especially the hotter ones with $T_{\\rm eff} > 30000 \\rm K$ in many cases show strong deficiencies of light metals like Al, Mg, O and Si by more than a factor of $100$ in comparison to the solar abundances, whereas enrichments of elements heavier than iron by a factor of $100$ are not unusual. Qualitatively, these abundance patterns show some similarity to those found in HgMn stars, which are a subgroup of the chemically peculiar main sequence stars reviewed by \\citet{t12_smith96}. The HgMn stars with $10000 \\rm K \\la T_{\\rm eff} \\leq 16000 \\rm K$, $\\log g \\approx 4.0$ are characterized by low rotational velocities and weak or non-detectable magnetic fields. Some light metals (e.g.\\ Al, N) tend to be deficient, whereas heavy metals (e.g.\\ Hg, Mn, Pt, Sr, Ga) may be enriched up to several orders of magnitude. Some of the most recent spectral analyses are from \\citet{t12_zav07} and \\citet{t12_adel06}. The measured abundances of iron group elements \\citep{t12_sea96,t12_jom99} as well as nitrogen \\citep{t12_roby99} tend to be lower than predicted from equilibrium diffusion calculations. Only for mercury abundances larger than predicted have been detected \\citep{t12_prof99}.  So in hot white dwarfs, sdB and HgMn stars a common tendency seems to be present, according to which the metal abundances are lower than expected from the equilibrium condition between gravitational settling and radiative levitation. In addition there is a large scatter of abundances from star to star. Even for stars with similar stellar parameters the abundances may be different. This points to some time-dependent process and not to an equilibrium state. In main sequence stars the chemically peculiar phenomenon is restricted to stars with low rotational velocities. In addition the presence of magnetic fields may be of importance, because magnetic fields may change the radiative acceleration or suppress convection \\citep{t12_tur03}. White dwarfs and sdB stars, however, are always more or less chemically peculiar. Up to now no correlation between abundance anomalies and magnetic field strengths has been found \\citep{t12_otol05}. In hot DA white dwarfs no outer convection zones should exist, because hydrogen is preferably ionized and helium is strongly deficient. In sdB stars a thin superficial convection zone with a mass depth of the order $10^{-12} M_{*}$ may be present only for helium abundances $\\rm He / \\rm H \\geq 0.01$ by number \\citep{t12_gro85}. So it seems to be unlikely that magnetic fields are of decisive importance for the explanation of the surface compositions.  In hot DAO white dwarfs helium is detectable, but in many cases it is deficient in comparison to the solar value \\citep[e.g.][]{t12_napi99}. The same is true for the majority of sdB's \\citep[e.g.][]{t12_edel03,t12_lisk05} and the helium deficient main sequence stars (e.g.\\ in HgMn stars helium usually is deficient). In contrast to the metal abundances, however, the abundances of helium are always larger than predicted from equilibrium diffusion calculations \\citep{t12_ven88,t12_mic89,t12_mic79}. These results may be explained with the effect of gravitational settling which, however, somehow must be disturbed. Possible reasons for this disturbance of the equilibrium may be the presence of turbulence \\citep{t12_vauc78} or mass loss. \\citet{t12_fon97} and \\citet{t12_ub98} predicted helium abundances as a function of time in the presence of weak winds. The results have shown that for mass loss rates of the order $10^{-13} M_{\\odot} / \\rm yr$ helium sinks much more slowly than in the case of an undisturbed stellar atmosphere. Within the lifetimes of sdB stars near the extended horizontal branch ($\\approx 10^{8} \\rm yr$) the helium abundance would gradually decrease from the solar value to $\\rm He / \\rm H \\approx 10^{-4}$ by number. This could explain why the helium abundances usually are in this range.  If this scenario with weak winds were the correct explanation for both the helium and the metal abundance anomalies, then it should be possible to find a mass loss rate for which all abundances can be explained simultaneously. The calculations of \\citet{t12_ub01} for the elements H, He, C, N and O have shown that this is hardly possible. According to these calculations for solar initial composition helium should always be more deficient than the metals. No mass loss rate exists which leads to deficiencies of C and O by more than a factor of $100$, whereas helium is deficient by a factor of ten only. This, however, is not an unusual composition in sdB stars \\citep[e.g.][]{t12_heb00}. Moreover in the presence of winds with mass loss rates of the order $10^{-13} M_{\\odot} / \\rm yr$ which are required to explain the helium abundances, the proposed pulsation mechanism of some sdB stars \\citep{t12_char97,t12_fon03} would become questionable. As mass loss tends to level out concentration gradients, the reservoir of iron (or other iron group elements) in the stellar envelope needed to explain the pulsations should be destroyed in time scales which are much shorter than the lifetimes of sdB stars. According to \\citet{t12_chay04} and \\citet{t12_fon06} for a mass loss rate of $\\dot M = 6\\times 10^{-15} M_{\\odot} / \\rm yr$ the reservoir would be destroyed on about $10^{7} \\rm yr$. For $\\dot M = 10^{-13} M_{\\odot} / \\rm yr$ the matter in mass depths $\\la 10^{-7} M_{*}$, where the reservoir is expected, would be blown away within one million years only.  Probably the most crucial assumption in these diffusion calculations with mass loss has been that the winds are ``chemically homogeneous\". If $\\dot M_{l}$ is the mass loss rate of an element $l$ and $\\zeta_{l}$ its mass fraction in the photosphere, then this assumption states that $\\dot M_{l} = \\zeta_{l} \\dot M$, where $\\dot M$ is the total mass loss rate.  Such a chemically homogeneous wind prevents (if $\\dot M \\ga 10^{-11} M_{\\odot} / \\rm yr$) or retards (if $\\dot M < 10^{-11} M_{\\odot} / \\rm yr$) gravitational settling. However, it does not directly change the surface composition. The opposite case would be a ``selective\" wind in which the mass loss rates of the individual elements are essentially independent of each other. A selective wind should lead to additional changes of the surface composition, which have not yet been taken into account in the calculations.  In radiatively driven winds of hot stars the photon momentum is absorbed preferably by the metals \\citep[see e.g.][]{t12_abb82,t12_vink01}, whereas the contribution of hydrogen and helium is small. Thus the metals are accelerated and move outwards. If the flow of metals is sufficiently large, then due to Coulomb collisions with the metals hydrogen and helium are accelerated as well. For this purpose, in dense winds a small velocity difference between the metals and hydrogen and helium is sufficient. Then it should be a good approximation that all elements have the same velocity and that the wind is chemically homogeneous. If, however, the flow of metals is small, then it may happen that the coupling of the various constituents due to collisions is not sufficiently effective and that hydrogen and helium are left behind. This scenario may lead to pure metallic winds such as investigated by \\citet{t12_bab95} for main sequence A stars.  Mass loss has been detected in subdwarf O stars (sdO) which are more luminous than sdB's \\citep{t12_ham81,t12_rau93}. Up to now there is no observational proof for the existence of winds in sdB stars. From a quantitative analysis of $\\rm H \\alpha$ line profiles of 40 sdB stars \\citep{t12_max01}, a comparison of synthetic NLTE $\\rm H \\alpha$ line profiles from static model atmospheres with the observations revealed perfect matches for almost all stars. Only in the four most luminous sdB's anomalous $\\rm H \\alpha$ lines with a small emission at the line center have been detected, which possibly are signatures of weak winds \\citep{t12_heb03}. For the case $T_{\\rm eff} = 36000 \\rm K$, $\\log g = 5.5$ and $\\log L/L_{\\odot} = 1.51$, from a spectral synthesis of $\\rm H \\alpha$ with his wind code \\citet{t12_vink04} found a similar behaviour of $\\rm H \\alpha$ if the existence of a weak wind with $\\dot M \\approx 10^{-11} \\rm M_{\\odot} / \\rm yr$ is assumed.  In Sect.~2 it will be investigated for a stellar mass $M_{*} = 0.5 M_{\\odot}$  if winds with mass loss rates $\\dot M \\la 10^{-12} M_{\\odot} / \\rm yr$ can be chemically homogeneous. The arguments are similar as in the investigations for more luminous stars e.g.\\ from \\citet{t12_ow2002}, \\citet{t12_spr92} and \\citet{t12_krt03}. All metals are lumped together into one mean metal which is accelerated due to the absorption of photon momentum. This is some simplification, in \\citet{t12_krt06} the individual elements are considered separately. Hydrogen, helium and the free electrons are denoted as ``passive plasma\" which can only be accelerated due to collisions with the outflowing metals. From the calculations e.g.\\ of \\citet{t12_abb82} and \\citet{t12_vink01} as well as from own calculations as described in Sect.~3 there is no indication that the radiative force on hydrogen and helium is not negligible. The mean radiative acceleration on the metals exceeds the one on hydrogen and helium by at least a factor of $100$.  In Sect.~3 the results of mass loss calculations for sdB stars according to the original theory of radiatively driven winds of \\citet{t12_ca1975} are presented and are compared with the predictions from the mass loss recipe of \\citet{t12_vink02}. For these wind models, which are obtained from a one component description of the wind, it is again checked if the metals may be coupled to hydrogen and helium. In Sect.~4 the consequences of the results for the surface composition of sdB stars are discussed. ",
        "conclusions": "In the $T_{\\rm eff} - \\log g$ diagram of Fig.~2 for $25000 \\mathrm{K} \\leq T_{\\rm eff} \\leq 40000 \\rm K$ and for various metallicities the lines are shown above which according to the mass loss calculations as described in Sect.~3 chemically homogeneous winds may exist. \\begin{figure} \\centering \\includegraphics[width=11cm,bb=132 75 625 425,clip=true]{klaus2.eps} \\caption{Lines in the $T_{\\rm eff} - \\log g$ diagram above which chemically homogeneous winds may exist for $Z/Z_{\\odot} = 0.1$, $1/3$ and $1$, respectively. Squares and circles represent the sdB stars analyzed by \\citet{t12_max01} and \\citet{t12_lisk05}, respectively. Filled symbols represent the sdB's with peculiar $\\rm H \\alpha$ line profiles, which may indicate the presence of a weak wind.} \\end{figure} It can be seen that for the majority of sdB stars this is not possible if the metallicity is solar or subsolar. They are below the line for $Z / Z_{\\odot} = 1$. The sdB stars introduced in the diagram have been analyzed by \\citet{t12_max01} and \\citet{t12_lisk05}. According to the assumptions of the present paper in those ones with peculiar $H \\alpha$ line profiles (represented by filled symbols) chemically homogeneous winds may indeed exist if the metal abundances are not too far below the solar value. The predicted mass loss rates agree with the ones from the mass loss recipe of \\citet{t12_vink02} and are of the order $10^{-10}$ to $10^{-11} M_{\\odot} / \\rm yr$. However, as explained in Sect.~3, this agreement does not exclude that the mass loss rates are overestimated. So this result is still questionable. The existence of winds may depend on the abundance of one element only, which preferably contributes to the radiative acceleration.  From the results it is clear that chemically homogeneous winds with mass loss rates $\\dot M \\la 10^{-12} M_{\\odot} / \\rm yr$ cannot exists. From arguments similar as in the paper of \\citet{t12_ow2002}, this result can be obtained without the calculation of a wind model. The mass loss rate of the metals alone must be at least of the order $10^{-14} M_{\\odot} / \\rm yr$. Otherwise hydrogen and helium cannot be accelerated throughout the wind. This result could be questioned only if a wind solution could be found for which the terminal velocity of the metals is clearly lower than the surface escape velocity. \\citet{t12_krt00} suggested that the terminal velocity could indeed be lower than $v_{\\rm esc}$ (by a factor of the order two only, however), if the wind switches to a ``shallow\" solution with an abrupt lowering of the velocity gradient. However, \\citet{t12_ow2002} and \\citet{t12_krt02} argued that these solutions are unstable.  If hydrogen and helium cannot be expelled from the star, then pure metallic winds may exist. As the outflowing metals in the stellar atmosphere not only have to overcome the gravitational force, but in addition the frictional force due to collisions with protons and helium particles, the metal abundances should be lower than predicted from equilibrium diffusion calculations (if concentration gradients are negligible). For these metals, for which the mass loss rate is sufficiently small, measured and predicted abundances should be in agreement. For metals with non-zero mass loss rate this scenario should lead to abundances varying with time as has been discussed by \\citet{t12_sea96, t12_sea99} for iron group elements in the envelopes of HgMn stars.  If both hydrogen and helium are in hydrostatic equilibrium, then measured helium abundances should be approximately in agreement with the ones predicted from equilibrium diffusion calculations. The fact that in the various types of helium deficient chemically peculiar stars the measured ones are larger, seems to indicate that gravitational settling in general is less effective than expected in an undisturbed stellar atmosphere. Moreover the existence of two distinct sequences of sdB stars which are characterized by an offset in the helium abundance \\citep{t12_edel03,t12_lisk05} can hardly be explained with one atmospheric effect alone. It may point to a dependence on the star's history. Several scenarios of single star and binary evolution of the sdB and the hotter sdO stars are under discussion \\citep{t12_str07}."
    },
    "0807/0807.1729_arXiv.txt": {
        "abstract": " ",
        "introduction": "Recent times  have seen cosmology enter a  precision era. Measurements of the cosmic microwave background (CMB) have lead to estimates of the standard  6 cosmological  parameters  (see, for  example,  the 5  year results  from WMAP~\\cite{Dunkley:2008ie}  and references  therein). Of particular interest  here are the estimates of  the amplitude, $A_{\\rm s}$,  and   spectral  index,  $n_{\\rm  s}$,  of   the  scalar  density fluctuations     which     are      measured     to     be     $n_{\\rm s}=0.963^{+0.014}_{-0.015}$    and    $A_{\\rm   s}=(2.41\\pm0.01)\\times 10^{-9}$.  It is presumed  that the  density fluctuations  are created during an  inflationary era through quantum  effects. Constraining the nature of this epoch on the basis of the observations is now the focus of  much  effort~\\cite{Peiris:2006sj,Kinney:2006qm,Martin:2006rs} with the ultimate aim of making contact with particle physics models around the Grand Unified Theory (GUT) scale~\\cite{Lyth:1998xn}.  In  addition to  these fluctuations  of quantum  origin, there  may be others  created   by  topological   defects  formed  naturally   as  a consequence   of   a  phase   transition   at   the   end  of   hybrid inflation~\\cite{Linde:1993cn}. Such fluctuations cannot be the primary source  of CMB  anisotropies, but  they could  provide  a sub-dominant component~\\cite{Wyman:2005tu},   which  can   have   some  interesting effects. In particular,  it has been shown that  if cosmic strings are formed,  then   $n_{\\rm  s}=1$  can   be  made  compatible   with  the data~\\cite{Battye:2006pk,Bevis:2007gh,Battye:2007si}.    This   is  of interest  since   many  of  the  models  of   inflation  motivated  by fundamental theories predict larger values for the spectral index than the WMAP best  fit value.  For example, SUSY  $F$- and $D$-term hybrid inflationary        models~\\cite{Copeland:1994vg,Dvali:1994ms,DTerminf} predict $n_{\\rm  s}>0.98$~\\cite{Panagiotakopoulos:1997qd}, while even larger values  may be required if one  considers gravitino constraints on the reheat  temperature or the bounds on  the cosmic string tension in  $D$-term  models.    Brane  inflation~\\cite{Dvali:1998pa}  in  its simplest form predicts $n_{\\rm s}>0.97$, but again, more sophisticated implementations    may   lead    to    a   scale-invariant    spectral index~\\cite{Haack:2008yb}.  Most work on inflation plus defect scenarios has focused on strings. However, the recent suggestion~\\cite{Cruz:2007pe,Cruz:2008sb} that the previously identified cold spot in the CMB~\\cite{Cruz:2006fy} could be due to cosmic textures~\\cite{Turok:1989ai,Turok:1990gw} motivates our consideration of hybrid inflation scenarios involving global defects. Such dynamically unstable defects produce a spectrum of anisotropies qualitatively similar to strings~\\cite{Bevis:2004wk} and are likely to have similar properties in respect of the degeneracy with $n_{\\rm s}$. We note, in addition, the other cold spots detected on smaller scale in the CMB which do not appear to be associated with the Sunyaev-Zel'dovich effect~\\cite{GenovaSantos:2008hf}.  Textures are most easily implemented in non-supersymmetric models, where they can emerge from a scalar potential of the form~\\cite{VilenkinShellard} \\begin{equation} \\label{V:texture} V(\\Phi)=\\frac14 \\lambda (\\Phi_i\\Phi_i-\\phi_0^2)^2\\,, \\end{equation} where $\\lambda$  is a dimensionless coupling constant  and $\\phi_0$ is the symmetry  breaking scale.  The  fields $\\Phi_i$ are  a fundamental representation  of  a global  ${\\rm  O}(4)$-symmetry.  This  potential induces spontaneous  symmetry breaking ${\\rm  O}(4)\\to{\\rm O}(3)$, for which  the   vacuum  manifold   is~$S^3$,  where  $S^3$   denotes  the 3-dimensional  sphere.  Since  $\\pi_3(S^3)=\\mathbbm{Z}$, there  is the possibility of non-trivial topological configurations called textures. However, they  cannot be stable  by Derrick's theorem and  hence, once formed at  the phase transition,  they unwind. This  unwinding process takes  place  in a  self-similar  fashion  such  that the  density  of textures scales relative to the  background. The latter has been shown to be  a more  general feature  of global field  models with  a broken $O(N)$ symmetry~\\cite{Turok:1991qq}.  The process of unwinding creates density               fluctuations               and               CMB anisotropies~\\cite{Turok:1990gw,Durrer:1990mk,Notzold:1990jt}. The CMB anisotropies result in a distribution  of hot and cold spots, which by the central  limit theorem, appear as approximately  Gaussian on small scales.   On larger  scales, they  may  be revealed  by more  detailed analysis.  By  taking into  account various observational  biases, the analysis of  Cruz {\\it  et al.}~\\cite{Cruz:2007pe} has  suggested that the large  cold spot identified  in the CMB  would result from  such a model with \\begin{equation} \\label{eta:obs} \\phi_0=(8.7^{+2.1}_{-3.0})\\times 10^{15}~{\\rm GeV}\\,. \\end{equation}  In this  paper, we study  supersymmetric (SUSY) scenarios  of $F$-term hybrid inflation  that realise the global  symmetry ${\\rm SU}(2)\\times {\\rm  U}(1)_X$.   We also  consider  the  possibility  that the  ${\\rm U}(1)_X$-symmetry is  gauged, giving rise to a  semi-local model. Even though  in the  context of  SUSY hybrid  inflation, models  with local symmetries are more commonly discussed, we note that the present paper is  not  the first  one  in  which the  use  of  global symmetries  is considered.  In one of the  seminal papers on SUSY $F$-term inflation, the use  of a spontaneously  broken global ${\\rm  U}(1)_X$-symmetry is suggested~\\cite{Copeland:1994vg},  and  the  possibilities  of  larger symmetries and textures  are mentioned. To the best  of our knowledge, however,  neither an explicit  analysis of  the particle  spectrum has been  performed  so far,  nor  have  the  consequences of  the  higher symmetry been investigated in  detail within the framework of $F$-term hybrid inflation.  An important  aspect of the  model-building is that  unless additional ${\\rm  SU}(2)$-invariant  lifting   terms  are  included,  the  vacuum manifold is  larger than ${\\rm  SU}(2)\\sim S^3$.  This is  because the holomorphic nature of the superpotential is consistently maintained by doubling  the number of  degrees of  freedom in  the symmetry-breaking fields, such  as the  field $\\Phi$ in~(\\ref{V:texture}).   We describe here in detail  how these additional degrees of  freedom in the vacuum manifold  can be  lifted  such that  it  is indeed  reduced to  $S^3$. Furthermore,  we include the  effect of  radiative corrections  on the predictions  for  the  initial   power  spectrum  created  during  the inflationary phase as first performed in Ref.~\\cite{Dvali:1994ms}. The breaking of a local instead of a global symmetry was first proposed in this work  and following this  reference most articles  on SUSY-hybrid inflation consider only local symmetries.  This may also be due to the fact that quantum gravitational effects can violate global symmetries. Since  no complete  model of  quantum gravity  exists yet,  we  do not consider this possibility  here and take ${\\rm SU}(2)$  to be an exact symmetry. Such an assumption  is not entirely unrealistic, since ${\\rm SU}(2)$ is the only group that is self-non-anomalous, independently of the representation of the chiral fermions in the theory.  The organisation of the paper  is as follows: in Section~2 we describe in detail the  SUSY hybrid texture model based  on the global symmetry ${\\rm SU(2)}\\times  {\\rm U(1)}_X$.  We also calculate  its scalar mass spectrum  and discuss  mechanisms for  giving masses  to  the massless moduli  fields   present  in  the   model.   Formal  aspects   of  the model-building have been relegated to the appendices.  In Section~3 we analyse  possible scenarios for  reheating after  inflation, including constraints from  possible overproduction  of gravitinos and  Big Bang Nucleosynthesis (BBN).  In addition, we  estimate the effect  that the texture  Goldstone  bosons may  have  on  stellar cooling.   Section~4 discusses a semi-local realisation of the SUSY hybrid model, where the aforementioned global  ${\\rm U(1)}_X$ symmetry is promoted  to a local one.   In Section~5  we study  the inflationary  dynamics of  the SUSY hybrid  model and analyse  the possibility  whether the  texture scale $\\phi_0$ given in~(\\ref{eta:obs})  can naturally be implemented within this  model, without  being in  conflict with  other  cosmological and astrophysical constraints discussed  in Section~3.  Finally, Section~6 summarises the main conclusions of our study. ",
        "conclusions": ""
    },
    "0807/0807.3206_arXiv.txt": {
        "abstract": "In this paper we use coherently integrated visibilities (see separate paper in these proceedings\\cite{jorgensen:2008}) to measure the properties of binary stars. We use only the phase of the complex visibility and not the amplitude. The reason for this is that amplitudes suffer from the calibration effect (the same for coherent and incoherent averages) and thus effectively provide lower accuracy measurements. We demonstrate that the baseline phase alone can be used to measure the separation, orientation and brightness ratio of a binary star, as a function of wavelength. ",
        "introduction": "\\noindent Binary stars are important for calibrating evolutionary stellar models. Because stellar models are sensitive to the parameters of the stars, it is important to obtain the highest possible accuracy of the stellar parameters. In this paper we will demonstrate how to extract brightness ratio and separation vector from measurements of the complex visibility phase only. This is significant because the visibility phase does not suffer from the calibration effects of visibility amplitudes, and therefore much higher precision can be obtained. Further, there are more complex visibility baseline phases than closure phases such that using baseline phases instead of closure phases yields more information. Finally, complex visibility phases generally have better SNR than closure phases. ",
        "conclusions": "\\noindent We have demonstrated that coherently integrated visibility phases can be used to obtain high-quality measurements of the fundamental parameters of binary stars. The measurements are accurate to the point where careful calibration of wavelengths and bandpasses is necessary. This calibration was not available for the data sets considered in this paper, but is routinely performed on more recent NPOI data sets."
    },
    "0807/0807.3801_arXiv.txt": {
        "abstract": " ",
        "introduction": "By  the end of twentieth century the overall picture of star formation process in dense interstellar clouds became fully formed. Contraction (collapse) of protostellar clumps is initiated by gravitation and external pressure; thermal pressure, rotation of the clumps and magnetic field counteract contraction. Thus, general scenario of star formation is defined by a complex interaction of these factors. However, many fundamental problems of star formation remain unsolved.    For instance, the mechanisms that stimulate the collapse of low-mass prestellar cores remain unclear. Even more enigmatic is formation of massive stars. Apparently, more complex processes, like competitive accretion and merger of protostellar fragments are involved in the latter process in addition to spherical and disc accretion.  \\begin{table*}[p!] \\caption{Atomic and molecular lines in the 10~mkm to 2~cm range which are used for the study of star formation regions.} \\begin{tabular}{|l|c|l|} \\hline Molecule & Lines & Radiation mechanism \\& Specific problem \\\\ \\hline H$_{2}$O & 22~GHz  & Maser emission, Studies of protostellar discs, astrometry \\\\ CH$_{3}$OH & 25~GHz & Maser and thermal emission, Astrometry, determination of \\\\ &        & physical conditions in protoplanetary \\\\ &        &discs and regions of massive star formation \\\\ NH$_{3}$ & 24~GHz & Thermal emission, Tracer of physical conditions \\\\ &        &in the dense gas, especially in the late stages \\\\ &        &of the evolution of prestellar cores \\\\ CS & 49~GHz  &  A tracer of physical conditions in the dense gas\\\\ & 98~GHz  & \\\\ & 147~GHz & \\\\ & 244~GHz & \\\\ HCN & 89~GHz & A tracer of physical conditions in the dense gas \\\\ & 266~GHz & \\\\ HCO$^{+}$ & 89~GHz & Tracer of physical conditions \\\\ & & in the dense gas, including ionization degree \\\\ & 268~GHz & \\\\ HNC & 91~GHz  &  A tracer of physical conditions in the dense gas \\\\ & 272~GHz & \\\\ N$_{2}$H$^{+}$ & 93~GHz & Tracer of physical conditions \\\\ &        & in the dense gas, especially in the late stages \\\\ &  & of the evolution of prestellar cores and \\\\ &  & in the regions of massive star formation \\\\ CO & 115~GHz & The main tracer of the presence of diffuse molecular gas\\\\ & 230~GHz &\\\\ H$_{2}$CO & 140~GHz &  Tracer of physical conditions in the dense gas\\\\ NO        & 150~GHz &  Tracer of physical conditions in the dense gas \\\\ & 250~GHz &\\\\ H$_{2}$D$^{+}$ & 372~GHz & Tracer of physical conditions \\\\ &         & in the dense gas, especially, of the \\\\ &         & kinematics of the central regions of prestellar cores \\\\ C & 492~GHz & Tracer of physical conditions in diffuse gas,  \\\\ &         & PDR-regions, ultracompact HII regions \\\\ & 809~GHz & \\\\ C$^{+}$ & 1.9 THz & A tracer of physical conditions in diffuse gas, \\\\ & &PDR-regions, ultracompact HII regions\\\\ Si$^{+}$ &  8.6 THz & A tracer of physical conditions in the protoplanetary discs\\\\ H$_{2}$  & 10.7 THz & A tracer of physical conditions in the protoplanetary discs\\\\ Fe$^{+}$ & 11.5 THz & A tracer of physical conditions in the protoplanetary discs\\\\ S & 12.0 THz & A tracer of physical conditions in the protoplanetary discs\\\\ Fe & 12.5 THz & A tracer of physical conditions in the protoplanetary discs\\\\[1mm] \\hline \\end{tabular} \\end{table*}  \\begin{table*}[t!] \\caption{Requirements to the angular resolution in observations of different stages of star formation}  \\begin{tabular}{l|c|c|c} \\hline Stage & Typical  & Typical  & Angular resolution\\\\ & distance & scale & (10 diagrams per object) \\\\ \\hline Prestellar cores  & 140 pc & 0.1 pc\\phantom{0} & $15^{\\prime\\prime}$\\phantom{0,} \\\\ Hot cores, UCHII$^*$ & 2--4 kpc and more & 0.1 pc\\phantom{0} & ${<}1^{\\prime\\prime}$\\phantom{00} \\\\ protostellar objects & 400 pc & 0.01 pc & \\phantom{0}$0.5^{\\prime\\prime}$ \\\\ Outer regions of the discs & 100 pc & 1000 a.u. & $1^{\\prime\\prime}$\\phantom{,} \\\\ Inner regions of the discs, & 100 pc & \\phantom{0}100 a.u. & \\phantom{0}$0.1^{\\prime\\prime}$ \\\\ brown dwarfs discs &&&\\\\ \\hline \\end{tabular}  \\end{table*}  \\vspace*{5pt} For the later evolutionary stages, when the planetary systems form, the number of unsolved problems is none the less. There are still no  unique solutions for the problems of the nature of angular momentum transfer in the protoplanetary discs, their physical and chemical structure, the role of mixing in the  formation of chemical and mineralogical composition of protoplanetary matter (including  protosolar nebula).  These uncertainties are related, first, to the low energetics of the transformation of gas into stars, especially in its initial stages, which makes impossible studies of this process in the visual waverange. For instance, in the prestellar cores the temperature does not exceed 10~K and it is only slightly higher at the periphery of protostellar objects, protoplanetary  and more evolved debris discs. Therefore, a significant fraction of their radiation is emitted in submillimeter and millimeter waveranges. However, these waveranges are very informative \"--- the spectral range from 100~mkm to 20~mm contains thousands of lines of many dozens of interstellar molecules (table~1) which, in the absence of observed emission of molecular hydrogen, are the only source of information on composition, temperature, and kinematics of molecular clouds and star formation regions. Second, prestellar and, especially, protostellar objects are rather compact, but at the same time their structure is complex and its study demands high angular resolution (table~2). Just for this reason there is world-wide growing  interest to construction of sensitive detectors for millimeter and submillimeter wavebands, including interferometers.  Of course, the main expectations are related to the submillimeter interferometric system ALMA (located in the Atacama desert in Chile) which will consist of several tens of 12-m antennas with the maximum distance between them of 12~km~[1]. A somewhat lower scale project works successfully already. It is submillimeter interferometer SMA~[2] in Hawaii (USA) which consists of eight 6-m antennas. However, capabilities of the best of existing and planned ground-based instruments for the studies in millimeter and submillimeter ranges are limited by disturbances created by the Earth atmosphere. In the submillimeter range there are only several transparent windows in the Earth atmosphere, with transmittance factor that does not exceed \\vspace*{5pt} 60$\\%$ even in the Earth regions with the best astroclimate \\footnote{http://www.eso.org/projects/alma/specifications/ FreqBands.html.}. To some extent this problem will be solved with the launch of the space-born infrared telescope ``Herschel'', but its spatial resolution will be limited by its relatively small 3.5-m dish. More, in the long wavelength band ``Herschel'' will be sensitive up to 670!mkm only (ALMA will be sensitive up to 1~cm). Therefore, this telescope Will be unable to observe both the coldest clouds and many astrophysically interesting molecular lines, including, for instance, the lines associated with low-level rotational transition of the CO molecule.  A prerequisite to substantial increase in our knowledge of star formation would be construction of an extra-atmospheric submillimeter telescope which would conjoin sensitivity and spectral resolution with high angular resolution (which would mean possibility of its usage in the regime of an interferometer). It is by no means unimportant that location of the telescope in the space would allow homogeneous observations both in northern and southern    celestial hemispheres. Space telescope ``Millimetron'' which is a part of the Russian Federal Space Programme may become such an instrument. In Russia, problems of star formation are studied in most of astronomical research centers (see, for instance, the collection of works edited by Wiebe and Kirsanova~[3]). For this reason, high demand for observations by submillimeter telescope is expected from Russian scientists working in this field of astrophysics. In the present paper we describe some of the problems for which space-born submillimeter observatory will give crucially important results and justify some of requirements to its parameters. ",
        "conclusions": "Currently, several projects of millimeter and submillimeter range telescopes are already implemented or developed.  The problems related to different stages of star formation take an important place in their programmes. However, the sensitivity of ground-based instruments is  naturally limited by atmospheric absorption. A submillimeter telescope in the space, even operating in the one-mirror would allow to carry out extremely important investigations both of close and distant  star formation regions (thanks to high sensitivity) and will allow to solve a lot of enigmas of star formation not only in the vicinity of the Sun but also in a more extended region of the Galaxy. To some extent this goal will be achieved by ``Herschel'' project, but the angular resolution of `Herschel'' will be not high enough to solve many fundamental problems of star formation theory. In this respect, capabilities of the ``Earth--Space'' interferometer will be unbeaten and will allow to obtain highly valuable information on the structure of the star formation regions and protostellar (young stellar) objects that is unavailable by other means."
    },
    "0807/0807.0175_arXiv.txt": {
        "abstract": "{The mechanism of formation of the \\ion{He}{i} 10830~\\AA\\ triplet in cool stars has been subject of debate for the last 30 years. A relation between the X-ray luminosity and the \\ion{He}{i} 10830~\\AA\\ flux was found in cool stars, but the dominant mechanism of formation in these stars (photoionization by coronal radiation followed by recombination and cascade, or collisional excitation in the chromosphere), has not yet been established.} {We use modern instrumentation (NOT/SOFIN) and a direct measurement of the EUV flux, which photoionizes \\ion{He}{i}, to investigate the formation mechanism of the line for the most active stars which are frequently  excluded from analysis.} {We have observed with an unprecedented resolution ($R\\sim 170,000$) the \\ion{He}{i} 10830~\\AA\\ triplet in a set of 15 stars that were also observed with the Extreme Ultraviolet Explorer (EUVE) in order to compare the line strengths with their EUV and X-ray fluxes.} {Active dwarf and subgiant stars do not exhibit a relation between the EUV flux and the equivalent width of the \\ion{He}{i} 10830~\\AA\\ line.  Giant stars however, show a positive correlation between the strength of the \\ion{He}{i} 10830~\\AA\\ absorption and the EUV and X-ray fluxes.  The strength of the \\ion{C}{iv} 1550~\\AA\\ emission does not correlate with coronal fluxes in this sample of 15 stars.} {Active dwarf stars may have high chromospheric densities thus allowing collisional excitation to dominate photoionization/recombination processes in forming the \\ion{He}{i} 10830~\\AA\\ line. Active giant stars possess lower gravities, and lower chromospheric densities than dwarfs, allowing for photoexcitation processes to become important.  Moreover, their  extended chromospheres allow for scattering of infrared continuum radiation, producing  strong absorption in \\ion{He}{i} and tracing wind dynamics.} ",
        "introduction": "Observations of the solar corona and chromosphere reveal that regions with copious X-ray emission (emitted by  bright points or active regions in the corona) also have enhanced \\ion{He}{i} 10830~\\AA\\ absorption that arises in the chromosphere (Fig.~\\ref{fig:solar}). Conversely, chromospheric regions located below solar coronal holes, where X-rays are diminished, show weakened \\ion{He}{i} absorption \\citep{zir75,she81,dup96}. The \\ion{He}{i} 10830~\\AA\\ line, actually a triplet (10829.081~\\AA, 10830.250~\\AA, and 10830.341~\\AA), is a transition between the lower, metastable level (2$^3$S) in the triplet series of \\ion{He}{i} and the 2$^3$P level (Fig.~\\ref{fig:levels}). The metastable 2$^3$S level can be populated only through collisional excitation from the ground level or through recombination and deexcitation from upper levels. Two mechanisms have been proposed   for the population of the lower 2$^3$S level. In the photoionization-recombination (PR) mechanism \\citep{gol39}, X-rays and EUV radiation ($\\lambda <$504~\\AA) from the corona photoionize the neutral helium from the ground state; photoionization is followed by recombination, and the electrons cascade to populate the lower levels of \\ion{He}{i}, especially the metastable 2$^3$S level. Scattering of the local infrared continuum produces an absorption line.  Models suggest \\citep{and97} the PR process is important in the Sun at temperatures $<$10,000~K. However, the opacity of the chromosphere can limit the efficiency of the photoionization mechanism. In an alternative mechanism, electron collisions from the ground and the  metastable 2$^3$S  level of \\ion{He}{i} dominate photoionization. In this case, a temperature $\\sim$ 20,000-30,000~K is required, and high densities in the chromosphere and transition region enhance the process. A combination of the two processes can exist as well. Arguments for each of the mechanisms can be found in \\citet{zir82,sim82,smi83,wol84,obr86, zar86,lan95,and95,and97,piet04}, and references therein.  If the PR mechanism dominates the formation of the 10830~\\AA\\ transition, a correlation is expected between the strength of the helium line and the radiation field at $\\lambda <$504~\\AA\\ located in the EUV and X-ray bands. Several studies  have compared the equivalent widths (EW) of the \\ion{He}{i}~10830~\\AA\\ line and the X-ray flux in late-type stars in order to establish a relation between those parameters. Theses studies show that stronger X-ray emission yields  stronger 10830~\\AA\\ absorption in both dwarfs (of spectral type F7 or later) and giants \\citep{zir82,obr86,zar86}. However the RS~CVn active binary systems were generally not included in these analyses. Many of the  previous observations made use of photographic plates to measure the 10830~\\AA\\ line, and correlated the line strengths with the X-ray fluxes observed by the  Einstein satellite (IPC and HRI instruments) which spanned the energy range  0.1~--~4 keV. In this paper we present measurements of the 10830~\\AA\\ line in very active cool stars and binaries taken with modern instrumentation and high resolution.  The use of high spectral resolution is essential in order to discern blends with telluric lines. Additionally, EUV fluxes are obtained as measured directly with the Extreme Ultraviolet Explorer (EUVE). Since the photoionization edge of \\ion{He}{i} is located at 504~\\AA\\ (0.02 keV), the EUV fluxes are expected to relate  closely  to the photoionization rate of helium.  The chromospheric line of \\ion{He}{i} at 10830~\\AA\\ can indicate bulk mass motions in the atmospheres of cool luminous stars.  The lower level of the transition is metastable, and a large population can build up in this level. If the gas itself is moving outward in a wind, the helium in the metastable level in the outflow scatters  photospheric infrared radiation, and the line profile can reveal the dynamics of the atmosphere. If there is a significant contribution by the photoionization-recombination process, then the level population is independent of the local thermodynamic conditions in the chromosphere.  In all cases, detailed modeling demonstrates \\citep[see][]{dup92} that in luminous stars, the Helium atom is formed further out in the atmosphere than other optical diagnostics of mass flow such as H$\\alpha$ and \\ion{Ca}{ii} H \\& K lines, making it an extremely sensitive probe of the acceleration region of a stellar wind. Even in a dwarf star, such as the Sun, models demonstrate that the 10830~\\AA\\ line is formed above the \\ion{Ca}{ii} K core and H$\\alpha$ \\citep{avr98} so that this Helium transition is more favorable for detecting regions where acceleration is likely to occur.   Most importantly, observations of the 10830~\\AA\\ line in both Sun and stars demonstrate that high outflow velocities are  observed in the line profiles of cool stars \\citep[see e.g.][]{obr86, dup92, dup96, edw03, dup05}. Given the high spectral resolution of our observations, the presence of winds in the stars of the sample can be explored. The cool giants are especially good targets.  In Sect. 2 we  describe the observations. Sect. 3 compares the \\ion{He}{i} line strengths to other parameters. Results are discussed in Sect. 4; conclusions can be found in Sect. 5.  \\begin{figure} \\centering \\hspace{10mm} \\includegraphics[width=0.45\\textwidth]{figure1.eps} \\caption{Spectra of different solar regions: coronal hole, quiet Sun, and an active region  \\citep[from][with spectra from an active region courtesy of M. Penn]{dup96}.} \\label{fig:solar} \\end{figure} ",
        "conclusions": "The \\ion{He}{i} $\\lambda$10830 line responds differently to the EUV radiation field between the dwarf  and giant stars in our sample. Active dwarf stars reach a 'saturated' equivalent width in $\\lambda$10830 in the presence of a strong radiation field.  This behavior appears consistent with model calculations in which high chromospheric densities allow collisional excitation to dominate photoionization/recombination processes in forming the line. Giant stars, with lower chromospheric densities than dwarfs, show increased \\ion{He}{i} absorption related to an increased EUV radiation field and the absorption is  strengthened by an extended expanding atmosphere scattering the line. The \\ion{He}{i} line in giant stars has a photoionization-recombination component that appears to dominate the line-forming process. Detailed radiative transfer calculations would be helpful to assess the contribution of collisions to line formation in the giant stars."
    },
    "0807/0807.4929_arXiv.txt": {
        "abstract": "By observing the transits of exoplanets, one may determine many fundamental system parameters. I review current techniques and results for the parameters that can be measured with the greatest precision, specifically, the transit times, the planetary mass and radius, and the projected spin-orbit angle. ",
        "introduction": "Henry Norris Russell~(1948) once delivered a lecture here in Cambridge entitled ``The royal road of eclipses,'' about the determination of accurate parameters for eclipsing binary stars, and the promise that such systems held for progress in stellar astrophysics. Given the rapid progress on display at this meeting, it is clear that exoplanetary science too has its royal road: the royal road of transits. Figure~1 illustrates the happy situation in which the planet's orbit is viewed nearly edge-on, and the planet undergoes transits and occultations.  \\begin{figure}[b] \\begin{center} \\includegraphics[width=29pc]{Pedagogical.eps} \\caption{Illustration of transits and occultations. During a transit, the planet blocks a fraction of the starlight.  Afterwards, the planet's brighter dayside comes into view and the total flux rises. The total flux drops again when the planet is occulted by the star.} \\label{fig1} \\end{center} \\end{figure}  \\begin{table} \\caption{Properties that have been measured, or that might be measured in the future, through precise observations of transiting planets.} \\label{tab1} \\begin{tabular}{lclc}\\hline {\\bf Property} & {\\bf Refs.} & ~~~~~~~{\\bf Property} & {\\bf Refs.} \\\\ \\hline Orbital period & 1,2 & ~~~~~~~Planet-planet interactions (short-term) & 19,20 \\\\ Orbital inclination & 1,2 & ~~~~~~~Planet-planet interactions (long-term) & 21,22 \\\\ Planetary mass & 1,2 & ~~~~~~~Mutual orbital inclinations & 20,23 \\\\ Planetary radius & 1,2 & ~~~~~~~Planetary rings & 24,25 \\\\ Stellar obliquity & 3,4 & ~~~~~~~Satellites & 9,24 \\\\ Orbital eccentricity & 5,6 & ~~~~~~~Relativistic precession & 26,27 \\\\ Stellar limb darkening & 7 & ~~~~~~~Parallax effects & 28, 29 \\\\ Star spots & 8,9 & ~~~~~~~Apsidal motion constant & 30 \\\\ Thermal emission & 5,10 & ~~~~~~~Stellar differential rotation & 31 \\\\ Absorption spectrum & 11,12 & ~~~~~~~Oblateness and obliquity & 32,33 \\\\ Albedo & 13,14 &  ~~~~~~~Variations in stellar radius & 34 \\\\ Phase function & 15 & ~~~~~~~Yarkovsky effect & 35 \\\\ Effective radiative time constant & 16 & ~~~~~~~Planetary wind speed & 36 \\\\ Trojan companions & 17,18 &  ~~~~~~~Artificial planet-sized objects & 37 \\\\ \\hline \\end{tabular}  {\\it Non-exhaustive list of references:} (1) Charbonneau et al.~(2000). (2) Henry et al.~(2000). (3) Queloz et al.~(2000). (4) Winn et al.~(2005). (5) Charbonneau et al.~(2005). (6) Bakos et al.~(2007). (7) Knutson et al.~(2007a). (8) Silva~(2003). (9) Pont et al.~(2007). (10) Deming et al.~(2005). (11) Charbonneau et al.~(2002). (12) Vidal-Madjar et al.~(2003). (13) Rowe et al.~(2006). (14) Winn et al.~(2008a). (15) Knutson et al.~(2007b). (16) Langton \\& Laughlin~(2008). (17) Ford \\& Gaudi~(2006), (18) Madhusudhan \\& Winn (2008). (19) Holman \\& Murray~(2005). (20) Agol et al.~(2005). (21) Miralda-Escud\\'e (2002). (22) Heyl \\& Gladman~(2007). (23) Fabrycky, D., these proceedings. (24) Brown et al.~(2001). (25) Barnes \\& Fortney~(2004). (26) P\\'al \\& Kocsis~(2008). (27) Jordan \\& Bakos (2008). (28) Scharf~(2007). (29) Rafikov~(2008). (30) Ragozzine \\& Wolf~(2008). (31) Gaudi \\& Winn~(2007). (32) Seager \\& Hui~(2002), (33) Barnes \\& Fortney~(2003). (34) Loeb~(2008). (35) Fabrycky~(2008). (36) Spiegel et al.~(2007). (37) Arnold~(2005). \\end{table}  Table~1 summarizes the information that has been obtained---or that is obtainable in principle---through precise observations of transits and occultations. This table is surely incomplete. Every few months, a new and creative application of transit observations is proposed. I was asked to discuss some of the measurements that can be made with the highest signal-to-noise ratio. In the best cases, we can measure orbital periods with 8 significant digits; transit times to within a fraction of a minute; the planetary mass and radius to within a few per cent; and the stellar obliquity (or at least its sky projection) to within a few degrees. ",
        "conclusions": ""
    },
    "0807/0807.1021_arXiv.txt": {
        "abstract": " ",
        "introduction": "This is the third paper of a series that studies a relativistic cosmology modelling the relativistic generalization of the single fractal Newtonian model advanced by Pietronero (1987; see also Coleman \\& Pietronero 1992), in which the galactic clustering problem is studied by assuming that the large-scale structure of the universe can be described as being a self-similar fractal system.\\footnote{ \\ The models investigated in this series of papers are in the realm of classical cosmology. No hypotheses concerning inflationary cosmology have been considered.}  In Ribeiro (1992a, hereafter paper I) I argued that the recent all sky redshift surveys (de Lapparent, Geller \\& Huchra 1986; Saunders et al. 1991) present observations consistent both with the old Charlier hypothesis of hierarchical clustering  and with fractals, where the latter is in essence a more precise conceptualization of the scaling idea implicit in the hierarchical clustering hypothesis. In paper I Pietronero's (1987) basic hypotheses were assumed in order to propose similar ones in a relativistic context, and I obtained observational relations compatible with fractals in Tolman's spacetime and devised a numerical strategy for finding particular Tolman solutions representing a fractal behaviour along the backward null cone. In Ribeiro (1992b, hereafter paper II) I studied analytically the Einstein-de Sitter model in the context of the theory developed in paper I. By treating the Einstein-de Sitter model as a special case of Tolman's spacetime, I found that it does not appear to remain homogeneous along the past null geodesic, has a volume (average) density which vanishes asymptotically and that it also shows no single fractal features along the backward null cone. The apparent inhomogeneity of the Einstein-de Sitter model is explained by the fact that densities measured along the geodesic go through different hypersurfaces of constant $t$, where each one has a different value for the proper density.  This paper continues the study of these cosmologies and presents relativistic fractal solutions obtained by following the numerical simulation strategy already devised in paper I. By {\\it fractal solutions} I mean solutions where the fractal system is smoothed-out and its average density follows the de Vaucouleurs density power-law. These solutions represent fractal behaviour along the backward null cone and they were obtained for all three types of Tolman dust models, namely, elliptic, parabolic and hyperbolic. By analysing these solutions we conclude that the only ones with features that may represent real astronomical observations are of hyperbolic type. As we are studying the Tolman spacetime as a region (or regions) possibly surrounded by a Friedmann universe (see paper I), if we adopt the fitting condition approach for interpreting the match between the two spacetimes (Ellis \\& Stoeger 1987; Ellis \\& Jaklitsch 1989) we then conclude that the Friedmann background required is also hyperbolic and we would be living in an open, ever expanding, universe.  This paper is organized as follows. In \\S 2 is shown a summary of the observational relations developed in papers I and II plus some minor extensions which will be necessary here, and in \\S 3 I discuss the initial conditions and the algorithm used to find Tolman numerical solutions along the past light cone. In \\S 4 I present the numerical solutions for the spatially homogeneous special cases and \\S 5 shows fractal solutions for all types of Tolman models. \\S 6 discusses these fractal solutions in terms of comparison with the spatially homogeneous cases, fitting condition, evolution of the most realistic fractal model in terms of real observations and relations with homothetic self-similarity. In this section I also express my criticism of criticisms of fractal cosmology. I finish in \\S 7 with the conclusions on this and the preceding papers. ",
        "conclusions": "\\subsection{A Metric for a Smoothed-out and Averaged Fractal}  In the previous section it was shown that Tolman fractal solutions do exist, and that the only ones compatible with observations are the hyperbolic type solutions obtained by means of the specializations given by equations (\\ref{t29}) and (\\ref{t30}).\\footnote{ \\ Notice however that fractal solutions of elliptic and parabolic types may, in principle, be obtained from other more complex specializations of the arbitrary functions than the ones considered in this paper, and these solutions could be compatible with observations.} As in equations (\\ref{t29}) the function $\\beta(r)$ is not constant, this means that the model has no simultaneous big bang. In other words, in a model of this sort the big bang singularity hypersurface occurred at different proper times in different locations, and the age of the universe is different for different observers at different radial coordinates. More specifically, as $\\beta(r)$ in equations (\\ref{t29}) is an increasing function, regions at smaller $r$ are younger than at bigger $r$, and the youngest region of the model is ``here'', at $r=0$.  An universe model where some regions are older than others is not as odd in terms of accepted ideas of galaxy formation as it might seem at first. Inasmuch as the observed universe is lumpy, in a spatially homogeneous Friedmann universe where $\\rho = \\rho (t)$ and the big bang singularity is simultaneous, there must be density fluctuations $\\delta \\rho / \\rho$ of some kind in order to form galaxies, and it is necessary to have some sort of metric perturbations for that to happen. So, at the era of galaxy formation, which may be defined as a hypersurface of constant time in order to agree with our intuition in an unperturbed metric, in the perturbed metric the overdensities $\\rho_{\\rm o}$ occur at time $t_{\\rm o}$ and the underdensities $\\rho_{\\rm u}$ occur at $t_{\\rm u}$, and $t_{\\rm o}$ must be different from $t_{\\rm u}$, otherwise there would not be any fluctuation at all. In other words, due to the density fluctuations $\\delta \\rho / \\rho$, in the perturbed metric a hypersurface of constant density no longer coincides with a hypersurface of constant time. Therefore, a deviation of the Friedmann metric from spatial homogeneity, even if it is small, is essential for lumpiness and, hence, some regions will inevitably have different local times than others. In the perturbed metric we may even define the era of galaxy formation as being a hypersurface of constant density. In other words, a non-simultaneous big bang seems inevitable in order to form galaxies in the standard scenario, even if those differences in local times are small.  Note that this discussion assumes that a hypersurface of simultaneity is defined by a specific value of the proper time, which is a logical thing to do in an unperturbed metric. In a perturbed metric, however, one could define the hypersurfaces of simultaneity in a different way, which would mean a different choice of gauge by which the perturbed and the non-perturbed spacetimes are related.  Nevertheless, considering it is desirable that fractal models be as close as possible to their Friedmann counterparts, and also considering mathematical simplicity, I shall take the specializations given by equations (\\ref{t30}) as the best modelling of a relativistic hierarchical (fractal) cosmology by Tolman's spacetime. Let us now write this model explicitly. Its metric is expressed as \\begin{equation} dS^2=dt^2-\\left( \\frac{{R'}^2}{\\cosh^2 r} \\right) dr^2 - R^2 (d \\theta^2+\\sin^2 \\theta d \\phi^2); \\ \\ \\ r \\ge 0, \\ R[r,t(r)] \\ge 0. \\label{t31} \\end{equation} The Einstein field equation for this metric may be written as an energy equation (Bondi 1947; paper I) \\begin{equation} \\frac{\\dot{R}^2}{2} - U(R) = E(r), \\label{t32} \\end{equation} where \\begin{equation} U(R) = \\frac{\\alpha r^p}{4R} \\label{t33} \\end{equation} is the effective potential energy, and \\begin{equation} E(r) = \\frac{1}{2} \\sinh^2 r \\label{t34} \\end{equation} is the total energy within $r$. The solution of equation (\\ref{t32}) is \\begin{equation} R = \\frac{ \\alpha r^p}{8 E(r)} \\left( \\cosh 2 \\Theta - 1 \\right), \\label{t35} \\end{equation} where $\\Theta$ is given by \\begin{equation} 4 \\left[ t(r) + \\beta_0 \\right] { \\left[ 2 E(r) \\right] }^{3/2} = \\alpha r^p \\left( \\sinh 2 \\Theta - 2 \\Theta \\right), \\label{t36} \\end{equation} $t(r)$ is the solution of the past radial null geodesic \\begin{equation} \\frac{d t}{d r} = - \\frac{R'}{ \\cosh r}, \\label{tt3366} \\end{equation} and the local density is expressed as \\begin{equation} \\rho = \\frac{4 p { \\left[ E(r) \\right] }^2}{ \\pi \\alpha r^{p+1} R' { \\left( \\cosh 2 \\Theta - 1 \\right) }^2 }. \\label{tt3377} \\end{equation}  The essentially new physical feature of the model above is its single difference from the open Friedmann one: the form of the function for the gravitational mass. That is given by $F = \\alpha r^p$, while in Friedmann this function must be $F = b_2 \\sinh^3 r$. Remembering that $\\alpha = 10^{-4} - 10^{-5}$ and $p=1 - 2.5$, the fractal metric (\\ref{t31}) appears to have a more rarefied dust  than its Friedmann equivalent. As fractal models are characterized by a power law nature of their average densities, with fractional exponents smaller than 3, it is hardly surprising that such models have a more rarefied distribution of mass.  \\subsection{Evolution of the Fractal Model and Comparison with the Spatially Homogeneous Case}  The constant $\\beta_0$ in equation (\\ref{t36}) gives the universal big bang time if we define ``now'' as $t=0$, and this fact allows us to study the evolution of the fractal model (\\ref{t31}) through the easily computable manner of simply varying $\\beta_0$. This investigation permits us to answer the question of whether or not the fractal features of the metric (\\ref{t31}) are present at different epochs. Bearing  this point in mind, I carried out simulations for different values of $\\beta_0$ in the interval $0 \\le r \\le 0.07$, but keeping $\\alpha = 10^{-4}$ and $p=1.4$ in all of them. I found out that the model under consideration remains fractal in integrations where $\\beta_0 =$ 1.5, 2, 2.5, 3.6, 4.5 and 6, with the resulting fractal dimensions being $D=$ 1.5, 1.5, 1.5, 1.4, 1.4 and 1.4, respectively. These results show that the metric (\\ref{t31}) effectively models a fractal distribution of dust at different epochs, with a remarkable constancy in $D$. All those integrations had $r = 0.07$, the final value of the integrating parameter, corresponding to $z \\cong 0.07$, although the luminosity distance varied  from $d_l \\cong 110$ Mpc when $\\beta_0 = 1.5$ to $d_l \\cong 450$ Mpc when $\\beta_0 = 6$. This variation may be physically explained as due to changes in the Hubble constant itself, whose values get bigger at earlier epochs. Finally, those results together with some simulations on different values of $p$ suggest it is possible to propose a simple, but very restricted, relationship between $p$ and $D$. For $0 \\le r \\le 0.07$, $1.5 \\le  \\beta_0 \\le 6$ and $ 1.4 \\le p \\le 2.5$ we can say approximately that $D = p \\pm 0.1$.  An interesting question about the model (\\ref{t31}) is to see how it would compare with the spatially homogeneous open Friedmann one. Looking at the results of the latter in figure 6 and the former in figure 20 we can qualitatively see that the absolute value of the difference in $\\log \\rho_v$ between the two models starts as zero, but increases rapidly. In the analytical study of the Einstein-de Sitter model presented in paper II it was shown that at the big bang singularity hypersurface the volume density $\\rho_v$ vanishes and the luminosity distance $d_l$ goes to infinity, and this effect is a consequence of the definition (\\ref{t10}) of the volume density: at the big bang the volume (\\ref{t9}) is infinity, but the total mass is finite. As already said at the end of \\S 4, we can therefore expect a similar asymptotic effect in both the open Friedmann case and the fractal model (\\ref{t31}), and this means that the difference in $\\rho_v$ between these two models should start decreasing after reaching a maximum in its increase. Actually, figure 6 already shows a sharp decrease of $\\rho_v$ in the open Friedmann case, but as the integrations of the model (\\ref{t31}) in figure 20 did not go far enough in the past, the results cannot show where this maximum might be. Nevertheless, based on this reasoning we can deduce that the observational relations of the fractal model (\\ref{t31}) appear to be asymptotically Friedmann when calculated along the past light cone.  \\subsection{Relations to Homothetic Self-Similarity}  Another topic which deserves some investigation is the relation, if any, between self-similarity due to fractals and self-similarity due to homothetic Killing vectors (Cahill \\& Taub 1971). There is some interest in this point because some attempts have been made to explain large-scale voids and clusters by self-similar perturbations of a Friedmann universe (Carr \\& Yahil 1990), and also because the first attempt to propose a workable relativistic hierarchical cosmology was done assuming a homothetic self-similar metric (Wesson 1978, 1979).  In general relativity a spacetime is called self-similar if all metric components can be put in a form in which they are functions of a single independent dimensionless variable which is a combination of the spacetime coordinates. Mathematically, this corresponds to the existence of homothetic Killing vectors, meaning that spherically symmetric similarity solutions are unchanged by coordinate transformations of the form $t \\rightarrow bt$, $r \\rightarrow br$, for any constant $b$ obeying the conformal transformation $g_{\\mu \\nu}(r,t) \\rightarrow \\frac{1}{b^2} g_{\\mu \\nu}(r,t)$ (Cahill \\& Taub 1971). Physically, spherically symmetric similarity spacetimes with $\\Lambda = 0$ contain no fundamental scales or dimensional constraints (Henriksen \\& Wesson 1978), and it would seem that these features make homothetic self-similarity a possible mathematical version of the scaling idea behind the empirical hierarchical clustering concept. However, homotheties, as defined in general relativitity, are basically geometrical features which will not necessarily translate themselves in observable quantities. In other words, the geometrical scaling features of the model do not necessarily mean that its observational relations are also scaling. For this reason it seems that homothetic self-similarity provides an unsatisfactory manner of modelling hierarchy.  It is beyond the aims of this work to make a general discussion about the possible relations between these two types of self-similarity. However it is of interest to investigate whether or not the fractal solutions presented in \\S 5 are homothetic. Recently Lemos \\& Lynden-Bell (1989) and Ponce de Leon (1991) studied homotheties in Tolman models and found that specific criteria are necessary to maintain the assumed similarity symmetry in the solutions. For models where $\\Lambda = 0$ (our case here), they showed that the first criterion is for the function $f(r)$ to be constant, that is, each comoving shell must have the same total energy. That immediately tells us that all fractal solutions of elliptic and hyperbolic type studied here, including the metric (\\ref{t31}), are not homothetic, leaving only the parabolic solutions still to investigate. The next criterion says that homothetic solutions with $f(r) = 1$ restrict the mass distribution to have the form $F(r) = \\mbox{constant} \\times r^{2 \\wp +3}$, where $\\wp$ is a constant, and this is the case in the solutions (\\ref{t24}) and (\\ref{t25}). The final requirement for homothetic symmetry to hold in our Tolman solutions demands that for $f=1$ and $\\wp \\not= 0$ the big bang hypersurface must be of the form $\\beta(r)= a + b r^{-\\wp}$, where $a$ and $b$ are constants. The solution (\\ref{t24}) satisfies these three requirements only if $p=3-2q$. The solution (\\ref{t25}) can have $\\beta(r)$ reduced to its form in equations (\\ref{t24}), but then the value $p=1$ must hold to satisfy the third requirement. Therefore, very restricted cases of the fractal solutions (\\ref{t24}) and (\\ref{t25}) have also homothetic self-similarity.  In conclusion, from what we have seen it does appear valid to say that fractal self-similarity is a much weaker requirement on the solutions than homotheties, and although we have reached this conclusion looking in more detail only at the Tolman spacetime, based on the self-similar requirements for both cases it seems reasonable to suppose that this conclusion may well be valid in general.  \\subsection{Fitting Condition}  In the previous section of this paper we dealt with the problem of finding a specific Tolman model which best represents the observed inhomogeneous distribution of galaxies, and in that respect it was concluded that the metric (\\ref{t31}) is the simplest one to achieve this aim. In other words, what was being sought was the optimal way of fitting the Tolman metric to the real lumpy large-scale structure of the universe. It was discussed in paper I that it may be desirable for us to have a Friedmann metric as background spacetime to the inhomogeneous Tolman region (or regions) used here to describe a fractal distribution of galaxies, and having found the specific forms for this inhomogeneous region an important question arises at once: what are the implications that the fractal metric (\\ref{t31}) and the hyperbolic solution (\\ref{t29}) bring to a possible Friedmann spacetime background? Answering this question is equivalent to finding a response to the ``fitting problem'' in cosmology, in the specific context of this work.  Ellis \\& Stoeger (1987) have outlined the fitting problem as being the search for an ideal Friedmann model which best fits another cosmological model that gives a realistic representation of the universe, including all inhomogeneities down to some specified length scale. In Ellis \\& Stoeger's words, ``the approach resembles that used in geodesy, where a perfect sphere is fitted to the pear-shaped earth; deviations of the real earth from the idealised model can then be measured and characterized''. Various ways in which this approach might be tried are discussed in detail by them, however here we shall restrict ourselves to the specific one outlined by Ellis \\& Jaklitsch (1989) where the matching between the Tolman and Friedmann spacetimes is interpreted as a fitting condition.  It was shown in paper I that the Darmois junction conditions between the two spacetimes under consideration require that $f=g'$ on the joining surface. Here $g= \\sin r$, $r$, $\\sinh r$ is the function which determines the curvature of the Friedmann model. As both the solution (\\ref{t29}) and the metric (\\ref{t31}) are of hyperbolic type with $f= \\cosh r$, there is no way of satisfying this condition when $r > 0$ unless we have $g = \\sinh r$. Therefore, the first response to the fitting problem in this context says that our Tolman fractal solutions imply an open Friedmann background model.  It was also shown in paper I that the matching between these two spacetimes severely restricts the gravitational mass inside the Tolman cavity. If $m(r) = F(r)/4$ is the gravitational mass of the Tolman region within a comoving radius $r$, and $\\overline{m}(x) = 4 \\pi \\mu a^3(t) g^3(x)/3$ is its Friedmann equivalent for a radius $x$ and dust density $\\mu$, the junction conditions demand \\begin{equation} m ( \\Sigma_0 ) = \\overline{m} ( \\Sigma_0 ), \\label{ultimaIII} \\end{equation} where $\\Sigma_0$ is the constant that defines the joining surface $r=x=\\Sigma_0$ between the two spacetimes. Therefore, as discussed by Ellis \\& Jaklitsch (1989), equation (\\ref{ultimaIII}) allows us to choose the Friedmann background model whose density is appropriate to our lumpy Tolman model. Let us see in more detail how this background spacetime can be specified.  In the open Friedmann model the local density is given by \\[ \\mu = \\frac{3 b_2}{16 \\pi a^3(t)}, \\] and as $F = \\alpha r^p$ in our fractal models, we can then write the equation (\\ref{ultimaIII}) as \\begin{equation} \\alpha {\\Sigma_0}^p = b_2 \\sinh ^3 \\Sigma_0. \\label{ultimaIV} \\end{equation} The value of the parameter $b_2$ is what we are seeking in order to determine precisely the Friedmann background and give a more accurate answer to the fitting problem in this context. Equation (\\ref{ultimaIV}) shows that $b_2$ is dependent on the other three parameters $\\alpha$, $p$, $\\Sigma_0$, and hence, there is a certain degree of flexibility in choosing the mass of the background spacetime. Thus, even when the interior region is determined by known values of $\\alpha$ and $p$, different values of $b_2$ are obtained according to exactly where the joining surface is located.  We can work out how the Friedmann background is in the case of the numerical integrations of the model (\\ref{t30}) shown in figure 20. We have $\\alpha = 10^{-4}$, $p=1.4$, and if we take $\\Sigma_0 = 0.07$ (the value where the numerical evaluation ends) we get $b_2 = 0.007$, $\\Omega_0 \\cong 0.002$ and $H_0 = 83$ km/s/Mpc. This low value for $\\Omega_0$ is in the lower limits of the interval where it has been reportedly measured. However, it is important to notice that in the approach used in this work no kind of dark matter was considered, but only the luminous matter associated with galaxies. Galactic luminous matter gives a value for $\\Omega_0$ of the same order of magnitude as the one found for the background model above (see White 1990, p. 38).  As a final remark, we have so far considered the interior Tolman region joining directly to the exterior Friedmann metric. That does not need to be always the case and we can envisage an interior region surrounded by one or more intermediary regions with higher or lower densities, in a scheme designed to model specific observational features. For example, a structure like the ``Great Wall'' (Geller \\& Huchra 1989; Ramella, Geller \\& Huchra 1992) could be modelled by an intermediary overdensity region before the background spacetime is reached, and with an underdensity interior fractal Tolman region (see Bonnor \\& Chamorro 1991 on how to join an underdensity Tolman region to an overdensity intermediary section). Such a modelling will obviously increase the value of $\\Omega_0$ for the fitted Friedmann background. This method, however, demands more detailed work in order to achieve a model where this sort of structure is precisely characterized. I shall not pursue further this study here.  \\subsection{Criticism of Criticisms of Fractal Cosmology}  As the final issue of this section, I shall discuss some of the objections raised by some authors against a fractal cosmology. The first type of criticism is contrary to a possible unlimited fractal pattern for the large-scale clustering of galaxies, and although some of the critical voices do accept fractals at small scales, their objections are usually based either on reasoning from the 2-point angular correlation function (Peebles 1989), or on supposed strong theoretical limitations of the standard Friedmannian cosmology. Therefore, in one way or another those authors see the strong need for a crossover to homogeneity on the fractal structure, at a scale yet to be agreed upon.  Criticisms based on the angular correlation function have been addressed by Coleman \\& Pietronero (1992) and I shall not discuss them here, although this point was briefly mentioned in paper II. The theoretical criticisms, on the other hand, must be addressed in this paper, and for this purpose I shall reproduce  here two quotations from Mart\\'{\\i}nez (1991) which well represent this point of view, although he is by no means the only one to raise such kind of objections. Mart\\'{\\i}nez  states that ``...~it should be noted that in the standard cosmology, the distribution of mass must tend to a non-zero finite density when averaged over large volumes'' \\footnote{ \\ The expression ``large volumes'' used in this quotation is imprecise, and can be interpreted as meaning either big local volumes or volumes which are big enough to be no longer considered as local. In a cosmological context the latter is more appropriate and from now on I shall assume the expression ``large volumes'' to mean non-local ones.}. From a relativistic point of view, the problem with this statement is its failure to specify where this average is supposed to be carried out. It is correct to say that if in a Friedmannian cosmology we make averages of density at spacelike hypersurfaces of constant $t$, those averages cannot be zero as this cosmology is spatially homogeneous. However, at large volumes such averages would be observationally irrelevant as astronomy in the electromagnetic spectrum is actually made along the backward null cone and {\\it not} at such spacelike hypersurfaces. Therefore, the statement above is only true in an unobservable situation, and considering that voids and clusters of galaxies were, and still are being, identified in the so-called redshift space, which lies on the past null cone, this is where the average of density must be carried out. Hence, the quotation above is inappropriate as an objection to a fractal structure for the distribution of galaxies. These two types of averages will only coincide locally, and it will depend on the model and the value assumed for the Hubble constant in order to establish what scales are local, although, in any case they will certainly differ at large volumes. In effect, in this imprecise formulation this statement may actually reinforce, or be taken by, a common misinterpretation of the standard model, which, due to inappropriate Newtonian analogies, confuses the model's geometry with its observable quantities .  Mart\\'{\\i}nez (1991) goes on and states that ``... a fractal universe without a crossover to homogeneity (...) implies a vanishing density for very large volumes and this idea cannot be accepted without creating important additional problems''. It was shown in paper II (and rediscussed in \\S 4) that if averages on density are made along the backward null cone, at the big bang singularity hypersurface the luminosity distance goes infinite and this average density is zero. That happens in the Einstein-de Sitter model, the most popular of the cosmological models, and no additional hypothesis or change in the metric was done to achieve this result. The point is, once more, where the average is made and which definition of density is adopted. Thus, again the statement is in fact an untenable objection since the standard cosmology does have a vanishing average density without any important additional problem. Having or not having zero average density in the model is just a question of interpretation.  It should be clearly understood that the above criticisms of Mart\\'{\\i}nez's (1991) statements are  made solely on the grounds of the standard Friedmannian cosmology, and there is no need whatsoever to mention any fractal hypothesis in order to show the impreciseness and inappropriateness of such statements. The important point being that even accepting the spatially homogeneous standard Friedmannian cosmology, this model tells us we would only be able to see, through our telescopes, its homogeneity locally.  Closely related to this point of local homogeneity is the issue, one could argue, of how we would understand in this context the reported uniform distribution of some deep samples like radio sources. In the first place it must be made very clear that the discussion made so far is aimed at showing that from a theoretical point of view there is no constraint to an unlimited fractal distribution, even from within a Friedmannian framework, but that does not imply the fractal system is indeed limitless, as an upper cutoff to homogeneity is not yet ruled out. \\footnote{ \\ The proposal for an upper cutoff to homogeneity in the fractal system appears to have been initially advanced by Ruffini, Song \\& Taraglio (1988), although Pietronero (1987) had already made a discussion about this issue.} Nonetheless, a simple calculation in the Einstein-de Sitter model, as presented in paper II, will show that a 30\\% decrease in $\\rho_v$ (from the value at present time $t=0$) occurs at $z~\\approx~0.1$ ($d_l~\\approx~500$ Mpc), and this means that even considering such high error in the determination of the volume density, this is, roughly speaking, the maximum approximate range where the homogeneity of this model could be observed. Beyond this range the homogeneity of the Einstein-de Sitter model would no longer be observed in the past light cone. Thus the first issue raised by this result is a problem of methodology: curvature effects occur in Friedmann models at much closer ranges than usually assumed, and this means that those surveys must consider in their data analysis expressions along the past light cone. Currently, calculations of observational relations where the backward null cone is taken into account is a very much neglected problem in cosmology. Secondly, if it is confirmed that in those deep surveys the distribution is really uniform, that would put the Friedmann model in even greater difficulties as it would appear to predict inhomogeneity in deeper ranges where this would not be observed. Thirdly, the sceptical viewpoint on this issue would be to argue that usually deeper observations are less precise than shallower ones, and previous claims of the so-called homogeneous ``fair-sample'' finally being observed did not stand once more refined and complete observations were made. Historically, the range at where the homogeneity is, or would be, finally reached has being pushed further and further away as more complete data become available and observational techniques improve, and so, the sceptic may say, we may not have necessarily seen the end of this story.  In addition to the points discussed above, a second kind of criticism to the fractal cosmology has been voiced by Peebles, Schramm, Turner \\& Kron (1991) in the following form. ``If the galaxy distribution had been observed to follow a pure scale-invariant fractal, (...) the closely thermal spectrum and isotropy of the cosmic background radiation in this highly inhomogeneous Universe would have been a deep puzzle''. First of all, it should be said that in an inhomogeneous model with a Friedmann background, the apparent discrepancy between the inhomogeneity of the model and the isotropy of the microwave background is not really an issue as the junction conditions already require that an overdensity must be compensated by an underdensity before the uniform region is reached, in order that the average densities will be the same (see paper~I, \\S 4). Nevertheless, the most important point is the result already obtained in paper II and extended in this paper for the other Friedmann models: the standard Friedmannian cosmology may be taken to be inhomogeneous depending on how we look at it. That means that the apparent contradiction between inhomogeneity and the isotropy of the microwave background may not be a contradiction at all. This is an essential point in order to put Peebles, Schramm, Turner \\& Kron's statement above into perspective, as we have already seen in this paper and in paper~II that at relatively modest luminosity distances and redshifts, even the standard spatially homogeneous Friedmannian cosmology becomes highly inhomogeneous on the past light cone because $\\rho_v$ departs considerably from its constant value in our constant time hypersurface the further into the past we look, and $\\rho$ is dependent on $r$. Considering that so far even the deepest all-sky redshift surveys have failed to reach the so-called ``fair sample'' where the homogeneity is supposed to be, waiting for us to discover it, perhaps it is about time to ask if the cosmic background radiation could be accommodated in a different cosmology, or wonder whether it is really a deep puzzle.  After all this discussion, we are left with an important point to consider. If the supposed theoretical need for a crossover to homogeneity in the fractal system is much weaker than previously thought, we have the question: is it really necessary? Since we have seen here and in paper II that even the Friedmann models do not seem to remain homogeneous along the past null geodesic (see figures 1, 2, 6 and 7), if the homogeneity of the standard model does not survive, where is the strong need for a crossover? In paper I it was assumed that the Tolman metric would eventually join a Friedmann background and, among other things, I argued the need for that was to make the model compatible with a different interpretation of the Copernican principle. However, in the light of inhomogeneity even in the Friedmann metric, it could be argued that from an observational point of view, that is, in calculations along our past null cone, if the relativistic fractal cosmology developed in this series of papers has or has not a Friedmann background might well be irrelevant."
    },
    "0807/0807.1956_arXiv.txt": {
        "abstract": "Recent observational and theoretical studies on the three-dimensional (3D) space motions of the Large and the Small Magellanic Clouds (LMC and SMC, respectively) have strongly suggested  that the latest proper motion measurements of the Magellanic Clouds (MCs) are consistent with their  orbital evolution models  in which the MCs have arrived in the Galaxy quite recently for the first time. The suggested orbital models appear to be seriously inconsistent with the tidal interaction models in which the Magellanic Stream (MS) can be formed as a result of the mutual tidal interaction between the MCs and the Galaxy for the last $\\sim 2$ Gyr. Based on orbital models of the MCs, we propose  that if the MCs have a common diffuse dark halo with the mass larger than $\\sim 2 \\times 10^{10} M_{\\odot}$, the MCs can not only have the present 3D velocities consistent with the latest proper motion measurements but also interact strongly  with each other and with the Galaxy for the last 2 Gyr. These results imply that if the observed proper motions of the MCs are true ones of the centers of mass for  the MCs, the common halo of the MCs would need to be considered in constructing  self-consistent MS formation models. We discuss whether the origin of the possible common halo can be closely associated  either with the past binary formation or with the MCs having been in a small group. ",
        "introduction": "Recent proper motion measurements of the MCs by the Advanced Camera for Surveys (ACS) on the  {\\it Hubble Space Telescope (HST)} have reported that the LMC and the SMC have significantly high Galactic tangential velocities ($367 \\pm 18$ km s$^{-1}$ and $301 \\pm 52$ km s$^{-1}$, respectively) and thus suggested that the MCs could be unbound from each other (Kallivayalil et al. 2006; K06). Piatek et al. (2008) have independently  analyzed  the same data sets as those used by K06 and confirmed the proper motions of the MCs derived from  K06. Besla et al. (2007, B07)  extensively investigated the long-term orbital evolution of the MCs by using the results of K06 and thereby suggested  that the MCs  have recently arrived in the Galaxy for the first time (``the first passage scenario'').   These observational and theoretical results on the 3D space motions of the MCs appears to be seriously inconsistent with the tidal interaction models in which the MS can be formed as a result of the strong tidal interaction between the MCs and the Galaxy (e.g., Gardiner \\& Noguchi 1996, GN): the MCs are required not only to keep their binary status but also to interact strongly with the Galaxy at least for the last $\\sim 2$ Gyr in the  models. Recently Bekki \\& Chiba (2008) have shown that the tidal interaction models consistent with the results by K06 can not reproduce well the observed  location of the MS on the sky, though they did not investigate all possible orbits  consistent with K06. Given that the tidal interaction models of the MS formation can explain not only the fundamental properties of the MS but also the presence of the leading arm features (GN), it is worth while to discuss whether the tidal interaction models consistent with K06 can be constructed by considering some new physical processes that have not been so far included in the previous models of the MC evolution.  The purpose of this Letter is to propose that if the MCs have a common diffuse dark halo, the halo can play an important  role in the long-term orbital evolution  of the MCs. Based on orbital models of the MCs with and without common halos, we investigate whether the MCs can keep their binary status at least for the last 2 Gyr for their  present  3D velocities consistent with K06. We demonstrate that some models with common halos can keep the binary status of the MCs and thus show strong tidal interaction between the LMC and the SMC and between MCs and the Galaxy for the last 2 Gyr. This work is inspired by recent numerical simulations  which have clearly shown the concurrent  accretion of multiple satellite systems onto galaxy-scale halos in a hierarchical galaxy formation scenario (e.g., Sales et al. 2007; Li \\& Helmi 2008; Ludlow et al. 2008). ",
        "conclusions": "The present study  has first shown that if the MCs have a common halo,  they can not only have the present 3D velocities consistent with K06 but also keep their binary status within  the last more than 2  Gyr in some models. It should be however stressed that the common diffuse halo needs to have (i) the present velocity ($|v|$) significantly smaller than that of the LMC  and (ii) the mass larger than $\\sim 2 \\times 10^{10}$ in order for the MCs to keep their binary status for the last $\\sim 2$ Gyr. These two  requirements would be  very hard to be confirmed directly by observations: differences in velocities between the center of the common halo and  those of the MCs and the mass of the halo can be inferred from numerical simulations on the binary formation of the MCs. Then, how could  they have formed  a common halo in the histories of the MCs  ?  We here suggest the following two scenarios for the common halo formation. The first is that the MCs might have dynamically coupled recently ($<4$ Gyr) to form a common halo: dynamical relaxation processes  of the two pre-existing halos of the MCs during binary galaxy formation can be  responsible for the common halo formation. The orbital evolution models including dynamical friction between the MCs by BC05 showed  that the MCs could  become dynamically coupled for the first time about $3-4$ Gyr ago. Recent cosmological N-body simulations of the pair galaxy formation based on a  $\\Lambda$ cold dark matter ($\\Lambda$CDM) cosmology have shown that the pair formation like the MCs can occur at $z<0.33$ corresponding to less than 3.7 Gyr ago for a canonical set of cosmological parameters (Ishiyama et al. 2008). These results imply that the common halo formation of the MCs might have happened recently ($<4$ Gyr ago). We suggest that the required larger  mass of the common halo (i.e., $M_{\\rm ch} \\ge 2\\times 10^{10} M_{\\odot}$) for the binary statue of the MCs in the last $\\sim 2$ Gyr can come from the stripped dark matter halos of the MCs at the epoch of their binary formation: their original masses are significantly larger than the present ones.  The second scenario  is that a small group of galaxies including the MCs in its central region fell onto the outer region of the Galaxy and then lost most of the halo and the group member galaxies via tidal stripping by the Galaxy: the MC system  with the common halo is the remnant of a destroyed group. Li \\& Helmi (2008) have investigated merging histories of subhalos in a Milky Way-like halo using high-resolution simulations based on  a $\\Lambda$CDM model and thereby demonstrated  that about one-third of the subhalos have been accreted in groups (see also Sales et al. 2007 and  Ludlow et al. 2008). The demonstrated higher incidence of the group infall appears to suggest that the MCs can  originate from a group thus that the second scenario is also viable. It is currently unclear which of the two scenarios are more consistent with other observations.    If the MCs really have a  common halo, then the common halo would have the following possible dynamical effects on the Galaxy and the LMC. Firstly, the MC system embedded in the common halo, which is more massive than the LMC,  can more strongly influence the outer part of the Galaxy than the LMC alone is demonstrated to be able to do (e.g., Tsuchiya 2002)  so that the observed HI warp of the Galaxy (e.g., Diplas \\& Savage 1991) can be more naturally explained in terms of the common halo scenario. Secondly, the common halo can weakly influence the disk of the LMC so that the combined tidal effect of the Galaxy, the SMC, and the common halo could  form a off-center bar that is more pronounced than the simulated one in the  last LMC-SMC interaction about 0.2 Gyr ago (Bekki \\& Chiba 2007). Thirdly, the common halo enables the LMC and the SMC to have their  stellar halos extended much beyond their optical radii.  The present study suggests  that it would be  difficult to observationally determine the 3D velocities of the MC system solely  from proper motion measurements of the MCs in the common halo scenario. Furthermore, it would be  even more difficult for theoretical and numerical works  to predict precisely the long-term orbital evolution of the MCs with a common halo owing to additional two (or more) parameters for physical properties of the  halo. If the observed proper motions of the MCs (K06) are really true ones of  {\\it the centers of mass for the LMC and  the SMC},  some new physical processes need to be incorporated into the tidal models of the MS formation for self-consistency. The hypothesized  common halo of the MCs in the present study  would be  just one of possible physical ingredients that need to be considered in constructing a more self-consistent model of the MS formation."
    },
    "0807/0807.0741_arXiv.txt": {
        "abstract": "{Shell galaxies are considered the debris of recent accretion/merging episodes. Their high frequency in low density environments suggests that such episodes could drive the secular evolution for at least some fraction of the early-type galaxy population.} {We present  XMM-Newton X-ray  observations of two shell galaxies, NGC~7070A and ESO~2400100, and far UV observations obtained with the Optical Monitor for these and for an additional shell galaxy, NGC 474, for which we also have near and far UV data from GALEX.  We aim at gaining insight on the overall evolution traced by their star formation history and by their hot gas content.} {The X-ray and the far UV data are used to derive  their X-ray spatial and spectral characteristics and their UV  luminosity profiles.  We use models developed ad hoc to investigate the age of the last episode of star formation from the (UV - optical) colors and line strength indices. } {The X-ray spatial and spectral analysis show significant differences in the two objects. A low luminosity nuclear source is the dominant component in NGC~7070A (log L$_X$=41.7 erg~s$^{-1}$ in the 2-10 keV band).  In ESO~2400100, the X-ray emission is due to a low temperature plasma with a contribution from the  collective emission of individual sources. In the Optical Monitor image ESO~2400100 shows a double nucleus, one bluer than the other. This probably results from  a very recent star formation event in the northern nuclear region. The extension of the UV emission is consistent with the optical extent for all galaxies, at different degrees of significance in different filters. The presence of the double nucleus, corroborated by  the (UV - optical) colors  and line strength indices analysis, suggests that ESO~2400100 is accreting a faint companion.  We explore the evolution of the X-ray luminosity during accretion processes with time. We discuss the link between the presence of gas and age, since gas is detected either before coalescence or several Gyr ($>3$) after. } {} ",
        "introduction": "In $\\Lambda$CDM cosmology, galaxies are assembled hierarchically over an extended period by mergers of smaller systems. In this framework, early-type galaxies showing {\\it fine structure} occupy a special position since they are a testimony of the effects of past merging/accretion events, and as such they fill the gap between on-going mergers and the relaxed elliptical galaxy population.  Among examples of {\\it fine structure}, shells are faint, sharp-edged stellar features \\citep{Malin83} that characterize a significant fraction ($\\approx$ 16.5\\%) of the field early--type galaxy population \\citep{Malin83, Schweizer92,Reduzzi96,Colbert01}.  Different scenarios for the origin of shells emerge from the rich harvest of simulations performed since their discovery in the early 80's, mostly involving galaxy-galaxy interactions.  These range  from merging/accretion  between galaxies of different morphological type or masses \\citep[mass ratios typically 1/10 - 1/100, ][]{Quinn, Dupraz86, Hernquist87a,Hernquist87b} to significantly weaker interaction events \\citep{Thomson90,Thomson91}.  A few models that invoke gas ejection due to the  central AGN or the power of supernovae \\citep{Fabian80, Williams85} do not associate the shell formation with environment.  Most models can reproduce qualitatively basic characteristics such as spatial distribution, frequency and shape of observed shell systems \\citep[see e.g.][and references therein]{Wilkinson00,Pierfederici04,Sikkema07}. In the accretion models, the more credited scenario for their formation,  shells are density waves formed  by  infall of stars from a companion. A major merger may also  produce shells \\citep{Barnes92,Hernquist92,Hernquist95}.  The fact that shells are frequently found in the field, rather then in the cluster environment, finds a direct explanation in the week interaction hypothesis, since the  galaxy-galaxy ``harassment'' within the cluster tends to destroy shells, while poorer environments are much better suited, because the group velocity dispersion is of the order of the internal velocity of the member galaxies \\citep[see e.g.][]{Aarseth, Barnes85, Merritt}. The class of the so-called ``internal models'' for shell formation  suggests that star formation within a giant shell is the result of shocked interstellar gas. In such a case shells are expected to be bluer than the parent galaxy, up to $\\approx$0.5 mag in U-B and B-V \\citep[see e.g.][]{Fabian80}. In a few cases this has been observed, \\citep[see e.g.][and reference therein]{Sikkema07} but attributed to multiple accretion events of galaxies of different intrinsic color.  In the framework of a hierarchical cosmology, all galaxies, including early-type galaxies, are expected to contain multi-epoch stellar populations.  Shell galaxies, among the early-type class of galaxies, are the ideal candidate to contain also a young stellar population since some simulations indicate a shell dynamical age of 0.5--2 Gyr \\citep{Hernquist87a,Nulsen89}. The star formation history of shell early-type galaxies has been analyzed by \\citet{Longhetti00} using line--strength indices. They show that shell-galaxies encompass the whole range of ages inferred from the H$\\beta$ vs. MgFe plane, indicating that among them recent and old interaction/acquisition events are equally probable. If shells are formed at the same time at which the ``rejuvenating'' event took place, shells ought to be long--lasting phenomena. Recently, \\citet{Rampazzo07} combining {\\it GALEX} far UV data and line--strength indices  show that the peculiar position of some shell galaxies in the (FUV-NUV) vs. H$\\beta$ plane could be explained in terms of a recent (1-2 Gyr old) rejuvenation episode.  A rejuvenation of stellar population requires the presence of fresh gas during the accretion event. Therefore, a study of the cold, warm and hot gas phases is important in order to consider many of the elements relevant  for the evaluation of shell galaxy evolution. \\citet{Rampazzo03} found that the warm ionized (H$\\alpha$) gas  and stars appear often decoupled suggesting an external acquisition of the gas, as predicted by merging models \\citep{Weil93}. At the same time, a set of observations showing a clear association between cold (\\hi/CO) gas and stars challenge present merging models which do not predict it unless cold gas behaves differently from the ionized gas \\citep{Schiminovich94,Schiminovich95,Charmandaris00,Balcells01}.  The bulk of the interstellar medium (ISM) in early-type galaxies emits in X-rays, and only comparatively small quantities are detected in the warm and cold phases of the ISM \\citep{Bregman92}. The hot ISM  is believed to build up primarily through stellar mass loss ($\\approx$ 1 M$_\\odot$/yr  in an old passively evolving elliptical galaxy with M$_B=-22$~mag) and Type Ia supernova ejecta  \\citep[$\\approx$0.03 M$_\\odot$/yr in a galaxy with M$_B=-22$~mag; see details in][]{Greggio05}. The two processes produce an almost identical mass in metals although the SNIa ejecta are dominated by Fe. \\citet{HB06}, using proper spectral fits to excellent Chandra data,  were able to better determine  the metal abundances of the X-ray emitting gas for a sample of early-type galaxies. By estimating  stellar abundances from  optical line strength indices, adopting simple stellar population models, they showed that the hot ISM and the stars have similar abundances.  The link between the stellar and ISM metallicities could be masked/contaminated by the accretion  of a gas-rich system. At the same time, the content of hot X-ray emitting gas could be  correlated with the ``age'' of the rejuvenation episode. Early-type galaxies with {\\it fine structure} (e.g. shells), which are considered {\\it bona fide} signatures of the ``dynamical youth'' of the galaxy, tend to be less X-ray luminous than more relaxed, ``mature'' ellipticals with little evidence of  {\\it fine structure}. These are also characterized by extended, and generally stronger, X-ray emission \\citep{Sansom00}. Simulations suggest that mergers \\citep[see e.g.][]{Cox06} or interactions \\citep[see e.g.][]{dercole00} among galaxies produce indeed a wide range of X-ray luminosities. One extreme example in this picture is  NGC 474, which shows  a well developed shell system and has an X-ray luminosity consistent with the low end of the expected emission from discrete sources \\citep{Rampazzo06}. The X-ray domain may further disclose an otherwise hidden AGN activity, that could also be a result of the merging episodes. Again an extreme example of nuclear activity in a merger remnant can be found in the newly discovered spectacular shell system in the elliptical host of the QSO MC2 1635+119 by \\citet{Canalizo07}.   So, whereas accretion/merging events are widely believed to be at the origin of shell galaxies, all the details such as: the age of the event and duration of the shell structure, the global secular evolution of the stellar and gas components of the host galaxy as well as the timing for triggering the AGN, its duty cycle and feedback, are far from being firmly established.  In light of the high fraction of shell galaxies in the field, interaction/accretion/merging events seem to have played a significant role in the evolution of the early-type class as a whole. At the same time, there is still  the general open  question of whether a link exists between shell galaxies and the early phases of merging processes (ULIRGs, AGN, E+A galaxies etc.) on one side and the general class of ``normal'' early-type  galaxies on the other.  In the above framework, we discuss the X-ray (XMM-{\\it Newton}) observations of two shell galaxies, NGC~7070A and ESO~2400100, taken from the \\citet{Malin83} compilation. We further present far UV XMM-{\\it Newton}  Optical Monitor (OM) observations of these objects, and we add {\\it GALEX} far UV observations of NGC~474, for which the results from XMM-Newton X-ray observations have been already presented in  \\citet{Rampazzo06}. Through far UV photometry we aim at inferring whether these galaxies have ongoing/recent star formation activity and its distribution across the galaxy. We finally aim to correlate the above information with those extracted from their hot gas content and properties in light of our current understanding of these components.  The plan of the paper is  the following. Section~2 describes the relevant properties of our sample gathered from the literature. Section~3  presents the X-ray and Far UV observations and data reduction. Results are presented in Section~4 and discussed in Section~5. H$_0$=75 km~s$^{-1}$ is used throughout the paper.   \\begin{table*} \\caption{Relevant photometric, structural and kinematic  properties} \\begin{tabular}{llllc} \\hline\\hline &  NGC~474     & NGC~7070A &ESO 2400100   & Ref. \\\\ \\hline Morphol. Type          &  (R')SA(s)0    &(R')SB(l)0/a & SAB0: pec  & [1] \\\\ Mean Hel. Sys. Vel. [km~s$^{-1}$] &2366$\\pm$16 &2391$\\pm$18 &3184$\\pm$14 & [2] \\\\ Adopted distance [Mpc]         &  32.5           & 31.9 & 42.4 & [3]\\\\ $\\rho_{(x,y,z)}$ [gal Mpc$^{-3}$] & 0.19            & & & [3] \\\\ Environment                             & Cetus-Aries  & & & [3] \\\\ cloud                                   &  52-12          & & & [3]\\\\ &                    & & &\\\\ {\\bf Apparent magnitude,}          &                    & & &\\\\ {\\bf colours, indices}:            &                    & & &\\\\ B$_T$                            & 12.36$\\pm$0.16& 13.35$\\pm$0.15&12.63$\\pm$0.24 & [2] \\\\ K$_T$ (2MASS)         & 8.555$\\pm$0.039 &9.130$\\pm$0.026 & 8.698$\\pm$0.028 & [1] \\\\ 6 cm [mJy]                             &            & 1.3   &         &   [5] \\\\ S$_{60\\mu m}$ (IRAS) [Jy] & & 0.26$\\pm$0.040 & & [1]\\\\ S$_{100\\mu m}$ (IRAS) [Jy] & &0.75$\\pm$0.077 & & [1]\\\\ Mg2                           &            &         & 0.28/0.21 & [6]\\\\ H$\\beta$                   &            &         & 1.54/2.79 & [6]\\\\ $\\Delta4000$                &            &         & 2.31/2.01 & [6]\\\\ H+K(CaII)                 &            &         & 1.18/1.24 & [6]\\\\ H$\\delta$/FeI              &            &         & 0.97/0.88 & [6]\\\\ &                  & & & \\\\ {\\bf Galaxy structure}:                  &                  & & & \\\\ Effective Surf. Bright. $\\mu_e$(B) &21.99$\\pm$0.31&22.48$\\pm$0.33 &21.84$\\pm$0.33 &[2]\\\\ Diam. Eff. Apert., A$_e$ [\\arcsec] &  64.6      & & &[2] \\\\ Average ellipticity      & 0.21    &0.28 & 0.46 & [2]\\\\ P.A. [deg]                  & 75   &6.5 & 132.5& [2]\\\\ Fine structure ($\\Sigma$) & 5.26     & & & [4]\\\\ &                           &  & &   \\\\ {\\bf Kinematic parameters}        &                        & & & \\\\ Vel.disp. $\\sigma_0$ stars [km~s$^{-1}$]  &163.9$\\pm$5.1& 101.0$\\pm$20.2 & 225/223&[1,6] \\\\ App.Max. rotation V$_{max}$ star [km~s$^{-1}$] & 30$\\pm$6 &0.0$\\pm$0.0 & & [2] \\\\ App.Max. rotation V$_{max}$ gas [km~s$^{-1}$] &158.4$\\pm$9.3 & & & [2] \\\\ &                           & & & \\\\ \\hline \\end{tabular} \\label{table1}  \\medskip References: [1] {\\tt NED http://nedwww.ipac.caltech.edu/}; [2] {\\tt HYPERLEDA http://leda.univ-lyon1.fr/}; [3] \\citet{Tully88} (H$_0$=75 km~s$^{-1}$~Mpc$^{-1}$); [4] \\citet{Sansom00}; [5] \\citet{Sadler89}; [6] \\citep{Longhetti99,Longhetti00} data refer respectively to the $a/b$ nuclei embedded in the envelope of ESO~2400100. \\end{table*} ",
        "conclusions": "We have studied the characteristics of  the X-ray emission of two shell systems, NGC~7070A and ESO~2400100 using XMM-Newton. We  also analyzed their far UV emission using the XMM-Newton Optical Monitor.  We include another shell galaxy NGC~474, for which we have already presented the X-ray characteristics in \\citet{Rampazzo06}, exploiting both original XMM-Newton OM images and {\\it GALEX} archival data.  The XMM-Newton  spatial and spectral analysis suggest that a nuclear source is the dominant component in NGC 7070A. In ESO~2400100, the emission is due to a low temperature plasma with a contribution from the collective emission of individual sources in the galaxy.  {\\it GALEX} data of NGC~474 show that the extension of NUV emission is comparable with that of the optical image,  while the FUV emission shows up only in the central regions of the galaxy. Also  in the UVW1 and UVM2 filters  NGC~474, ESO~2400100 and NGC~7070A have extensions similar to that of the optical image. XMM-OM UV images of ESO~2400100 show the presence of a double nucleus. The shape of the luminosity profiles suggest that a disk component is present, confirming the morphological classification provided by {\\tt NED} and the suggestion, from the kinematical study of \\citet{Sharples83}, that NGC 7070A is an S0 seen face-on.  This study further suggests that the prominent dust lane is not yet in equilibrium, indicating a recent accretion episode.  We model line--strength indices and the far UV - optical colors of the galaxies to infer the time elapsed from the last significant episode of star formation. From the comparison between  our galaxies and shell galaxies in the \\citet{Rampazzo07} sample we suggest that the (UV - optical) colors of NGC~474 and NGC~7070A are consistent with a recent burst of star formation. We argue that the double nucleus in ESO~2400100 is indicative of an ongoing accretion event. The combined UV and line-strength indices analysis suggests indeed a very recent star formation episode in the northern nucleus \\citep[see also simulations by ][]{Kojima97}.  Using the above estimates of the time elapsed from the last significant episode of star formation we investigate whether the evolutionary scheme for gas-rich systems that would lead to mature ellipticals proposed by \\citet{Brassington07} could be applied to the same early-type galaxies undergoing merging episodes.  The time range is very poorly sampled, with only 4 objects, but  we span almost the full range, since  \\eso\\ should be at $Time \\sim -200$ Myr, i.e.  before coalescence. We notice that shell galaxies are systematically underluminous relative to gas-rich systems at similar evolutionary stages, and that the $L_x/L_B$ or $L_x/L_K$ values vary only a factor of a few, compared to factors of 100. The only indication of a possible difference is in the gas content, which is present either before or several Gyr after coalescence. This is consistent with a picture in which several gigayears are required to refurbish a galaxy of a  hot gaseous halos after a merging responsible for its depletion \\citep{Sansom00, Osul01}.  To fully understand whether  shell galaxies are the precursors of  relaxed ellipticals, their more mature counterparts according to hierarchically evolutionary scenarios, we need  to better define the  dynamical and photometric time-scales of  the accretion/merging event and understand whether there is an evolution of their X-ray properties linked to their ``age\" based on a larger and better studied sample of objects. The study of the UV emission in early-type galaxies in connection with optical line-strength indices \\citep[see e.g.][]{Rampazzo07} could provide useful insight for timing their photometric evolution."
    },
    "0807/0807.0431_arXiv.txt": {
        "abstract": "A comprehensive statistical analysis of the broadband properties of EGRET blazars is presented. This analysis includes sources identified as blazars in the Sowards-Emmerd publications. Using this sample of 122 sources, we find that there is a relationship $L_\\gamma \\propto {L_r}^{0.77 \\pm 0.03} $ as well as a correlation between $\\alpha_{og}$ and $\\alpha_{ro}$, and a correlation between radio luminosity and $\\alpha_{og}$. Through the use of Monte Carlo simulations, we can replicate the observed luminosity relationship if a synchrotron self-Compton model is assumed. However, this relationship can not be replicated if an external Compton scattering model is assumed. These differences are primarily due to beaming effects.  In addition it has been determined that the intrinsic radio luminosity of the parent sample falls in the range  $10^{21} < L < 10^{30}\\, {\\rm Watts\\,Hz^{-1}}$ and that the bulk Lorentz factors of the source are in the range $ 1 < \\Gamma < 30 $, in a agreement with VLBI observations.  Finally, we discuss implications for GLAST, successfully launched in June 2008. ",
        "introduction": "During the lifetime of the Energetic Gamma-Ray Experiment Telescope (EGRET) instrument on board {\\it Compton Gamma-Ray Observatory} (CGRO) 271 sources were detected with 66 being confidently  identified as blazars in the Third EGRET Catalog \\citep{har99}. Many differing statistical analyses of these gamma ray detected blazars have been conducted, among them \\citet{fos98};\\citet{muc97}. Some of these analyses concentrated on the direct statistical relationship between luminosities in the gamma-ray band and radio bands \\citep{sal96,ste93,fan98}. Of these, most have used single dish radio data, but some have used VLBI fluxes \\citep{zho97,mat97}, all at various radio frequencies $>$ 1 GHz. In addition, \\citet{muc97} and \\citet{imp96} use Monte Carlo simulations to aid in interpreting these relationships. Though significant correlations are reported in all of these works, \\citet{muc97} show that in some cases these will result from a combination of variability and selection effects. To investigate the variability effects further, \\citet{zha01} have compared their statistical results using time averaged data for the entire sample to similar results for a restricted sample for which data were available during high and low states. They report a similar significant correlation in each case and show that using time averaged data tends to under estimate the underlying linear regression slope (as applied to the logarithmic data).  Various models have been invoked to explain the origin of the gamma-ray emission of blazars and specifically, the radio gamma-ray correlation. Among these are the synchrotron self Compton (SSC) model \\citep{blo96,ghi85} and various models in which the source of seed photons for scattering is from a source external to the jet (henceforth called ECS models for ``external Compton scattering''). \\citet{der93} use the accretion disk as the source of soft photons whereas others \\citep{ghi96} use broad-line region clouds as the source of soft photons.  Since the end of the EGRET mission there have also been several reanalyses of the significance of identifications of EGRET sources, particularly those of \\citet{mat97,mat01} and \\citet{sow03,sow04}.The \\citet{sow03,sow04} survey excludes sources with $|b|<10^{\\circ}$ (the Mattox papers exclude sources with $|b|<3^{\\circ}$), and thus may exclude additional sources that are thought to be blazars \\citep{sgu04}. Likewise, some work continues on identifying other individual sources, such as 3EG J0416+3650, possibly identified with 3C 111 \\citep{sgu06}. In the eight years since the end of the CGRO mission, there has been substantial modification to the identifications given in the last catalog published by the EGRET team \\citep{har99}, so we present a comprehensive statistical analysis including all potential identifications, using homogeneous criteria for inclusion in the sample ({\\S 2}). We have also included various statistical techniques of survival analysis for the inclusion of upper and lower limits in the data ({\\S 3}). We later discuss whether or not the results of the analyses are dependent on the precise source list. The approach used in understanding the physical implications of our results is to use a Monte Carlo technique \\citep{lis97} to generate multiple simulated samples under differing assumptions of overall theoretical or phenomenological model ({\\S 4,5}), distribution of physical parameters such as bulk Lorentz factor (e.g., Gaussian, power-law, etc.) and selection effects. Additional physical implications of observed correlations and implications for the Gamma Ray Large Area Space Telescope (GLAST), successfully launched in June 2008, are discussed in {\\S 6,7}. ",
        "conclusions": "After taking into account statistical tests and Monte Carlo analysis, we find the following: \\begin{enumerate} \\item For this sample of 122 gamma-ray blazars there is a strong correlation between radio and gamma ray luminosity which persists even after the effects of redshift and limits are taken into account. The correlation is of the form $L_{\\gamma} \\propto {L_r}^{0.77}$. This correlation remains with similar regression coefficients even when only the strongest 76 candidates are included in the sample. \\item There is a correlation between $\\alpha_{og}$ and $\\alpha_{ro}$ as well as a correlation between $L_r$ and $\\alpha_{og}$. Each correlation is consistent with the increasing dominance of the SED by gamma ray luminosity as the maximum relativistic electron energy decreases and the increasing importance of the ECS process at high radio luminosities and low maximum electron energy. \\item A detailed simulation of source statistics using Monte Carlo techniques shows that the relationship $L_{\\gamma} \\propto {L_r}^{0.77}$ can only be reproduced assuming the SSC model. Though the  assumed intrinsic physical model in the source frame is of the form $L_{\\gamma} \\propto L_r$, selection effects lead to a simulated sample with the relationship described above. This effect is much less significant when ECS is assumed as the underlying model.  However, upon also comparing the observed and simulated distributions of the luminosities, flux densities and redshifts, both the SSC and ECS models are only weakly consistent with the data, assuming linear dependence of intrinsic gamma ray luminosity on intrinsic radio luminosity. Taken together with the previous results, this would suggest that if either SSC or ECS is indeed responsible for the gamma ray emission, a more complex model is likely needed. In addition, effects of evolution and assumed cosmologies can be explored in more detail. \\end{enumerate} In the GLAST era, our findings can be clarified in the following ways. Acquiring more gamma-ray data, including, particularly, detections of new dim sources near the GLAST detection limit will clarify whether the correlations we find are caused by truncation effects due to the flux limits of our instruments, or whether this is due, at least in part, to physical causes, at least down to the new limit established by GLAST.  In order to predict what we might see with GLAST, we can extend our observed radio/gamma ray luminosity correlation down to lower luminosities than what are covered in Figure 1. A radio luminosity of $1.25 \\times 10^{24} {\\rm Watts\\,Hz^{-1}}$ would lead to a gamma ray luminosity of approximately $4.98 \\times 10^{13} {\\rm Watts \\, Hz^{-1}}$. We can then covert these luminosities to flux densities  for a range of potential redshifts. For z=0.1-1.1 the radio flux density would be in the range of 0.8-60 mJy and the gamma-ray flux density would be in the range, $3.3 \\times 10^{-14}$-$2.4 \\times 10^{-12}$  Jy. Greater values of redshift would lead to even lower flux density values. Most of the range of radio flux densities ($>$ few mJy) would be detectable by the Green Bank Telescope (GBT), with the lower limit detectable with the Very Large Array (VLA)\\citep{con08}. Assuming that the sensitivity of GLAST is approximately $2.8 \\times 10^{-13}$ Jy at 100 MeV, these sources would only be  be detected if  $z< 0.2$. A radio luminosity of $1.58 \\times 10^{23} {\\rm Watts \\, Hz^{-1}}$ corresponds to a gamma-ray luminosity of $1 x10^{13} {\\rm Watts \\, Hz^{-1}}$ . The radio flux density would fall in the range of 0.1-8 mJy for z=0.1-1.1. These sources could potentially be  radio detectable with  the VLA under optimal conditions. At the higher flux density limit it would also be possible to detect with GBT, especially at higher frequencies (i.e., 40-50 GHz)\\citep{min08}. However, if the regression slope were actually closer to 1, implying a d! irect proportionality between radio and gamma ray luminosities, then radio sources in the same luminosity, redshift and flux ranges discussed above would not be detected in gamma rays at all. In short, for the observed luminosity relationship for EGRET blazars to be appreciably extended down to lower luminosities with GLAST detections, then a large percentage of GLAST sources near the detection limit would have radio flux densities about 10 mJy or greater (and $z< 1$).  Several authors have stated that the predicted number of blazars to be detected by GLAST rougly matches with the sky density of flat spectrum radio sources down to 50 mJy at 5 GHz \\citep{pad07} and 65 mJy at 8.4 GHz \\citep{hea07}. However, our results show that a large number of GLAST sources would be detected at even lower radio flux densities if the luminosity relationship we observe for EGRET sources also holds up for GLAST sources. It is very likely that many new radio observations would be required in a! ny of the cases mentioned above, but especially for cases in w! hich the putative radio source may not even be in any previous catalog.  >From a theoretical perspective, confirmation of the previously determined correlation between gamma-ray and radio luminosities will lead to confirmation of SSC, though the precise agreement would have to be re-analyzed with these new data."
    },
    "0807/0807.2434_arXiv.txt": {
        "abstract": "We present limits on transit timing variations and secondary eclipse depth variations at 8 microns with the Spitzer Space Telescope IRAC camera. Due to the weak limb darkening in the infrared and uninterrupted observing, Spitzer provides the highest accuracy transit times for this bright system, in principle providing sensitivity to secondary planets of Mars mass in resonant orbits.  Finally, the transit data provides tighter constraints on the wavelength-dependent atmospheric absorption by the planet. ",
        "introduction": "The extremely precise 33 hour phase function measurement of HD~189733b \\cite[Knutson et al.\\ (2007)]{Knutson2007} observed with the 8 micron IRAC camera on the Spitzer Space Telescope yielded the most precise times of transit and secondary eclipse for any extrasolar planet, 6 and 24 seconds, respectively. This led us to propose to observe an additional six transits and eclipses of this system over time with the goal of measuring:  \\begin{itemize} \\item precise transit-timing (\\cite[Agol et al.\\ 2005]{Agol2005}, \\cite[Holman \\& Murray 2005]{Holman2005}) to search for the presence of resonant (or near-resonant) terrestrial-mass planets captured by migration (\\cite[e.g.\\ Mandell, Raymond \\& Sigurdsson 2007]{Mandell2007}) or on longer timescales precession of an eccentric orbit (\\cite[Miralda-Escud{\\'e} 2002]{MiraldaEscude2002}, \\cite[Heyl \\& Gladman 2007]{Heyl2007}; also Fabrycky and Wolfe, these proceedings); \\item variations in the depth of the secondary eclipses with time which might be caused by large-scale variable atmospheric features (\\cite[Rauscher et al. 2007]{Rauscher2007}; also Showmand and Dobbs-Dixon, these proceedings); \\item precise transit depth for comparison with atmospheric-absorption models to contrain the molecular composition (\\cite[e.g. Tinetti et al. 2007]{Tinetti2007}; also Tinetti et al., Fortney et al., and Hubeny et al., these proceedings); \\item improved system parameters for better characterization of the planet, host-star, and orbit properties (\\cite[Winn et al.\\ (2007)]{Winn2007}; also Winn, these proceedings). \\end{itemize}  For librating planets in a low-order mean motion resonance, the times of transit vary with an amplitude: \\begin{equation} \\delta t_{2:1} \\approx P_{trans} \\left({M_{pert} \\over M_{trans}}\\right) \\approx 3 {\\rm \\ min} \\left({P_{trans} \\over 3 {\\rm \\ day}}\\right) \\left({M_{pert} \\over M_{\\oplus}}\\right) \\left({M_{J} \\over M_{trans}}\\right), \\end{equation} and libration period of \\begin{equation} P_{2:1} \\approx P_{trans}\\left({M_* \\over M_{trans}}\\right)^{2/3} \\approx 150 {\\rm \\ days}\\left({P_{trans} \\over 3 {\\rm \\ day}}\\right) \\left({M_{*} \\over M_{\\odot}}\\right)^{2/3} \\left({M_{J} \\over M_{trans}}\\right)^{2/3}, \\end{equation} where $M_\\oplus, M_J, M_*$ are masses of the Earth, Jupiter, and host star, $P_{trans}$ is the period of the transiting planet, and the numbers have been estimated for the libration amplitude for planets starting on circular orbits with exact commensurability \\cite[(Agol et al.\\ 2005)]{Agol2005}; the actual value depends on the libration amplitude. This timescale requires observations separated by months with $\\approx$ seconds precision in timing, and in principle could be sensitive to sensitive to Mars mass planets. ",
        "conclusions": "The five transits we have observed are consistent with no transit timing variations greater than 5 seconds.  The five secondary eclipses are consistent with no eclipse depth variations at greater than the 10\\% level.  From just an analysis of the lightcurves we can place constraints on the eccentricity and longitude of periastron: $\\vert e \\cos{\\omega} \\vert = 0.0002 \\pm 0.0001$ and $\\vert e \\sin{\\omega}\\vert = 0.015 \\pm 0.012$; similar to the results of \\cite[Winn et al.\\ (2007)]{Winn2007}. We are collecting two more transits and two more secondary eclipses in Summer 2008; thus the results presented here are preliminary and require further more careful analysis."
    },
    "0807/0807.4481_arXiv.txt": {
        "abstract": " ",
        "introduction": "A major unsolved problem in theoretical physics is to reconcile the classical theory of general relativity with quantum mechanics. Of the numerous attempts, some have postulated new and so far unobserved ingredients, while others have proposed radically new principles governing physics at the as yet untested Planckian energy scale. Here we report on a much more mundane approach using only standard quantum field theory. In a sum-over-histories approach we will attempt to define a nonperturbative quantum field theory which has as its infrared limit ordinary classical general relativity and at the same time has a nontrivial ultra\\-vio\\-let limit. From this point of view it is in the spirit of the renormalization group approach, first advocated long ago by Weinberg \\cite{weinberg}, and more recently substantiated by several groups of researchers \\cite{reuteretc}. However, it has some advantages compared to the renormalization group approach in that it allows us to study (numerically) certain geometric observables which are difficult to handle analytically.  We define the path integral of quantum gravity nonperturbatively using the lattice approach known as {\\it causal dynamical triangulations} (CDT) as a regularization. In Sec.\\ \\ref{cdt} we give a short description of the formalism, providing the definitions which are needed later to describe the measurements. CDT establishes a nonperturbative way of performing the sum over four-geometries (for more extensive definitions, see \\cite{ajl4d,blp}). It sums over the class of piecewise linear four-geometries which can be assembled from four-dimensional simplicial building blocks of link length $a$, such that only {\\it causal} spacetime histories are included. The continuum limit of such a lattice theory should ideally be obtained as for QCD defined on an ordinary fixed lattice, where for an observable $\\cO(x_n)$, $x_n$ denoting a lattice point, one can measure the correlation length $\\xi(g_0)$ from \\beq\\label{1.1} -\\log (\\la  \\cO(x_n) \\cO(y_m) \\ra) \\sim |n-m|/ \\xi(g_0) + o(|n-m|). \\eeq A continuum limit of the lattice theory may then exist if it is possible to fine-tune the bare coupling constant $g_0$ of the theory to a critical value $g_0^c$ such that the correlation length goes to infinity, $\\xi(g_0) \\to \\infty$. Knowing how $\\xi(g_0)$ diverges for $g_0 \\to g_0^c$ determines how the lattice spacing $a$ should be taken to zero as a function of the coupling constants, namely \\beq\\label{1.2} \\xi(g_0) = \\frac{1}{|g_0-g_0^c|^\\n},~~~~~a(g_0) = |g_0-g_0^c|^\\n. \\eeq  The challenge when searching for a {\\it field theory} of quantum gravity is to find a theory which behaves in this way. The challenge is three-fold: (i) to find a suitable nonperturbative formulation of such a theory which satisfies a minimum of reasonable requirements, (ii) to find observables which can be used to test relations like \\rf{1.1}, and (iii) to show that one can adjust the coupling constants of the theory such that \\rf{1.2} is satisfied. Although we will focus on (i) in what follows, let us immediately mention that (ii) is notoriously difficult in a theory of quantum gravity, where one is faced with a number of questions originating in the dynamical nature of geometry. What is the meaning of distance when integrating over all geometries? How do we attach a meaning to local spacetime points like $x_n$ and $y_n$? How can we define at all local, diffeomorphism-invariant quantities in the continuum which can then be translated to the regularized (lattice) theory? -- What we want to point out here is that although (i)-(iii) are standard requirements when relating critical phenomena and (Euclidean) quantum field theory, gravity {\\it is} special and may require a reformulation of (part of) the standard scenario sketched above. We will return to this issue when we discuss our results in Sec.\\ \\ref{discussion}.  Our proposed nonperturbative formulation of four-dimensional quantum gravity has a number of nice features. Firstly, it sums over a class of piecewise linear geometries, which -- as usual -- are described without the use of coordinate systems. In this way we perform the sum over geometries directly, avoiding the cumbersome procedure of first introducing a coordinate system and then getting rid of the ensuing gauge redundancy, as one has to do in a continuum calculation. Our underlying assumptions are that 1) the class of piecewise linear geometries is in a suitable sense dense in the set of all geometries relevant for the path integral (probably a fairly mild assumption), and 2) that we are using a correct measure on the set of geometries. This is a more questionable assumption since we do not even know whether such a measure exists. Here one has to take a pragmatic attitude in order to make progress. We will simply examine the outcome of our construction and try to judge whether it is promising.  Secondly, our scheme is background-independent. No distinguished geometry, accompanied by quantum fluctuations, is put in by hand. If the CDT-regularized theory is to be taken seriously as a potential theory of quantum gravity, there has to be a region in the space spanned by the bare coupling constants where the geometry of spacetime bears some resemblance with the kind of universe we observe around us. That is, the theory should create dynamically an effective background geometry around which there are (small) quantum fluctuations. This is a very nontrivial property of the theory and one we are going to investigate in detail in the present piece of work. New computer simulations presented here confirm in a much more direct way the indirect evidence for such a scenario which we provided earlier in \\cite{emerge,semi}. They establish the de Sitter nature of the background spacetime, quantify the fluctuations around it, and set a physical scale for the universes we are dealing with. The main results of our investigation, without the numerical details, were announced in \\cite{agjl} (see also \\cite{desitter}).  The rest of the article is organized as follows. In Sec.\\ \\ref{cdt} we describe briefly the regularization method of quantum gravity named CDT and the set-up of the computer simulations. In Sec.\\ \\ref{S4} we present the evidence for an effective background geometry corresponding to the four-dimensional sphere $S^4$, i.e.\\ Euclidean de Sitter spacetime. Sec.\\ \\ref{effective} deals with the reconstruction of an effective action for the scale factor of the universe from the computer data, and in Sec.\\ \\ref{fluctuations} we analyze the quantum fluctuations around the ``classical'' $S^4$-solution. Sec.\\ \\ref{S3} contains an analysis of the geometry of the spatial slices of our computer-generated universe. In Sec.\\ \\ref{size} we determine the physical sizes of our universes expressed in Planck lengths and try to follow the flow of the gravitational coupling constant of the effective action under a change of the bare coupling constants of the bare, ``classical'' action used in the path integral. Finally we discuss the results, their interpretation and future perspectives of the CDT-quantum gravity theory in Sec.\\ \\ref{discussion}. ",
        "conclusions": "}  The CDT model of quantum gravity is extremely simple. It is the path integral over the class of causal geometries with a global time foliation. In order to perform the summation explicitly, we introduce a grid of piecewise linear geometries, much in the same way as when defining the path integral in quantum mechanics. Next, we rotate each of these geometries to Euclidean signature and use as bare action the Einstein-Hilbert action\\footnote{Of course, the full, effective action, including measure contributions, will contain all higher-derivative terms.} in Regge form. That is all.  The resulting superposition exhibits a nontrivial scaling behaviour as function of the four-volume, and we observe the appearance of a well-defined average geometry, that of de Sitter space, the maximally symmetric solution to the classical Einstein equations in the presence of a positive cosmological constant. We are definitely in a quantum regime, since the fluctuations of the three-volume around de Sitter space are sizeable, as can be seen in Fig.\\ \\ref{fig1}. Both the average geometry and the quantum fluctuations are well described in terms of the mini\\-superspace action \\rf{n5}. A key feature to appreciate is that, unlike in standard (quantum-)cosmological treatments, this description is the {\\it outcome} of a nonperturbative evaluation of the {\\it full} path integral, with everything but the scale factor (equivalently, $V_3(t)$) summed over. Measuring the correlations of the quantum fluctuations in the computer simulations for a particular choice of bare coupling constants enabled us to determine the continuum gravitational coupling constant $G$ as $G\\approx 0.42 a^2$, thereby introducing an absolute physical length scale into the dimensionless lattice setting. Within measuring accuracy, our de Sitter universes (with volumes lying in the range of 6.000-47.000 $\\ell_{Pl}^4$) are seen to behave perfectly semiclassically with regard to their large-scale properties.  We have also indicated how we may be able to penetrate into the sub-Planckian regime by suitably changing the bare coupling constants. By ``sub-Planckian regime\" we mean that the lattice spacing $a$ is (much) smaller than the Planck length. While we have not yet analyzed this region in detail, we expect to eventually observe a breakdown of the semiclassical approximation. This will hopefully allow us to make contact with attempts to use renormalization group techniques in the continuum and the concept of asymptotic safety to study scaling violations in quantum gravity \\cite{reuteretc}.  On the basis of the results presented here, two major issues suggest themselves for further research. First, we need to establish the relation of our effective gravitational coupling constant $G$ with a more conventional gravitational coupling constant, defined directly in terms of coupling matter to gravity. In the present work, we have defined $G$ as the coupling constant in front of the effective action, but it would be desirable to verify directly that a gravitational coupling defined via the coupling to matter agrees with our $G$. In principle it is easy to couple matter to our model, but it is less straightforward to define in a simple way a set-up for extracting the semiclassical effect of gravity on the matter sector. Attempts in this direction were already undertaken in the ``old'' Euclidean approach \\cite{js,newton}, and it is possible that similar ideas can be used in CDT quantum gravity. Work on this is in progress.  The second issue concerns the precise nature of the ``continuum limit''. Recall our discussion in the Introduction about this in a conventional lattice-theoretic setting. The continuum limit is usually linked to a divergent correlation length at a critical point. It is unclear whether such a scenario is realized in our case. In general, it is rather unclear how one could define at all the concept of a divergent length related to correlators in quantum gravity, since one is integrating over all geometries, and it is the geometries which dynamically give rise to the notion of ``length\".  This has been studied in detail in two-dimensional (Euclidean) quantum gravity coupled to matter with central charge $c \\le 1$ \\cite{corr2d}. It led to the conclusion that one could associate the critical behaviour of the matter fields (i.e.\\ approaching the critical point of the Ising model) with a divergent correlation length, although the matter correlators themselves had to be defined as non-local objects due to the requirement of diffeomorphism invariance. On the other hand, the two-dimensional studies do not give us a clue of how to treat the gravitational sector itself, since they do not possess gravitational field-theoretic degrees of freedom. What happens in the two-dimensional lattice models which can be solved analytically is that the only fine-tuning needed to approach the continuum limit is an additive renormalization of the cosmological constant (for fixed matter couplings). Thus, fixing the two-dimensional spacetime volume $N_2$ (the number of triangles), such that the cosmological constant plays no role, there are no further coupling constants to adjust and the continuum limit is automatically obtained by the assignment $V_2 = N_2 a^2$ and taking $N_2 \\to \\infty$. This situation can also occur in special circumstances in ordinary lattice field theory. A term like \\beq\\label{k5} \\sum_i c_1 (\\phi_{i+1}-\\phi_i)^2 + c_2(\\phi_{i+1}+\\phi_{i-1}-2\\phi_i)^2 \\eeq (or a higher-dimensional generalization) will also go to the continuum free field theory simply by increasing the lattice size and using the identification $V_d = L^d a^d$ ($L$ denoting the linear size of the lattice in lattice units), the higher-derivative term being sub-dominant in the limit. It is not obvious that in quantum gravity one can obtain a continuum quantum field theory without fine-tuning in a similar way, because the action in this case is multiplied by a dimensionful coupling constant. Nevertheless, it is certainly remarkable that the infrared limit of our effective action apparently reproduces -- within the cosmological setting -- the Einstein-Hilbert action, which is the unique diffeomorphism-invariant generalization of the ordinary kinetic term, containing at most second derivatives of the metric. A major question is whether and how far our theory can be pushed towards an ultraviolet limit. We have indicated how to obtain such a limit by varying the bare coupling constants of the theory, but the investigation of the limit $a \\to 0$ with fixed $G$ has only just begun.   \\subsection*"
    },
    "0807/0807.3742_arXiv.txt": {
        "abstract": "This paper presents a photometric and spectroscopic study of the bright blue eclipsing binary LMC-SC1-105, selected from the OGLE catalog as a candidate host of very massive stars ($\\geq 30\\,\\msun$). The system is found to be a double-lined spectroscopic binary, which indeed contains massive stars. The masses and radii of the components are $\\rm M_{1}=30.9\\pm1.0\\;\\msun$, $\\rm M_{2}=13.0\\pm0.7\\;\\msun$, and $\\rm R_{1}=15.1\\pm0.2\\;\\rsun$, $\\rm R_{2}=11.9\\pm0.2\\;\\rsun$, respectively. The less massive star is found to be filling its Roche lobe, indicating the system has undergone mass-transfer. The spectra of LMC-SC1-105 display the Struve-Sahade effect, with the \\ion{He}{1} lines of the secondary appearing stronger when it is receding and causing the spectral types to change with phase (O8$+$O8 to O7$+$O8.5). This effect could be related to the mass-transfer in this system. To date, accurate ($\\leq 10\\%$) fundamental parameters have only been measured for 15 stars with masses greater than 30~\\msun, with the reported measurements contributing valuable data on the fundamental parameters of very massive stars at low metallicity. The results of this work demonstrate that the strategy of targeting the brightest blue stars in eclipsing binaries is an effective way of studying very massive stars. ",
        "introduction": "\\label{section:intro} \\vspace{-0.4cm} The fundamental parameters of very massive stars ($\\geq 30\\,\\msun$) remain uncertain, despite the large impact massive stars have in astrophysics, both individually and collectively \\citep[see review by][]{Massey03}. The equations of stellar structure allow for stars with arbitrarily large masses, however the mechanisms to form massive stars \\citep[accretion and mergers; e.g.][]{Bally05} and the associated instabilities \\citep[see][and references therein]{Elmegreen00, Zinnecker07} are not well understood, hindering theoretical predictions on the existence of an upper limit on the stellar mass. The ``mass discrepancy'' problem, i.e. the disagreement between masses derived from parameters determined by fitting stellar atmosphere models to spectra and from evolutionary tracks \\citep[see e.g.][for a comparison]{Repolust04, Massey05}, still affects studies of single massive stars, even though significant progress has been made in both stellar atmosphere \\citep[see review by][]{Herrero07} and stellar evolution models \\citep[e.g.][]{Meynet03}. The parameters of single stars also suffer from suspected multiplicity, which in many cases cannot be determined.  Pismis 24-1 demonstrates this problem: its inferred evolutionary mass $>200\\,\\msun$ \\citep{Walborn02} contradicted the upper stellar mass limit of $\\sim 150\\,\\msun$ suggested by statistical arguments based on observations \\citep{Figer05,Oey05}. \\citet{Maiz-Apellaniz07} resolved it into a visual binary with the {\\it Hubble Space Telescope}, thereby removing the discrepancy. One of its components is also a spectroscopic binary, illustrating the systematic effects often accompanying ``single'' stars.  \\vspace{-0.2cm} The only model-independent way to obtain accurate fundamental parameters of distant massive stars and to resolve the ``mass discrepancy'' problem is to use eclipsing binaries \\citep[see review by][]{Andersen91}. In particular, double-lined spectroscopic binary systems exhibiting eclipses in their light curves are extremely powerful tools for measuring masses and radii of stars. Specifically, the light curve provides the orbital period, inclination, eccentricity, the fractional radii and flux ratio of the two stars. The radial velocity semi-amplitudes determine the mass ratio; the individual masses can be solved for by using Kepler's third law. Furthermore, by fitting synthetic spectra to the observed ones, one can infer the effective temperatures of the stars, solve for their luminosities and derive the distance \\citep[e.g.\\@][]{Bonanos06}. The most massive stars measured in eclipsing binaries are galactic Wolf-Rayet stars of WN6ha spectral type: NGC3603-A1 \\citep[M$_1=116 \\pm 31\\; \\msun$, M$_2=89 \\pm 16\\; \\msun$;][]{Schnurr08}, and WR~20a \\citep[M$_1=83.0 \\pm 5.0\\; \\msun$ and M$_2=82.0 \\pm 5.0\\; \\msun$][]{Rauw04, Bonanos04} in Westerlund~2, presenting a challenge for both stellar evolution and massive star formation models \\citep{Yungelson08, Zinnecker07} and raising the issue of the frequency and origin of ``binary twins'' \\citep{Pinsonneault06, Lucy06, Krumholz07}. Such systems are of particular interest, since massive binaries might be progenitors of gamma-ray bursts \\citep[e.g.\\@][]{Fryer07}, especially in the case of Population III, metal-free stars \\citep[see][]{Bromm06}.  Analogs of these heavyweight champions, if not more massive binaries, are bound to exist in the young massive clusters at the center of the Galaxy (Center, Arches, Quintuplet), in nearby super star clusters (e.g.\\@ Westerlund\\,1, R136), in Local Group galaxies (e.g.\\@ LMC, SMC, M31, M33) and beyond (e.g.\\@ M81, M83, NGC 2403). A systematic wide-ranging survey of these clusters and galaxies is currently underway. The goal is to provide data with which to test star formation theories, stellar atmosphere and stellar evolution models for both single and binary stars as a function of metallicity, and the theoretical predictions on the upper limit of the stellar mass. The adopted strategy involves two steps: a variability survey to discover eclipsing binaries in these massive clusters and nearby galaxies, which is followed by spectroscopy to derive parameters of the brightest -- thus most luminous and massive -- blue systems. However, characterizing massive stars requires the availability of 8-m class telescopes and high resolution near-infrared spectrographs (since massive stars in the Galaxy are extincted and extragalactic ones are faint) and has only become feasible in the past few years. \\citet{Bonanos07} demonstrated that this method efficiently finds massive candidates, by performing the first variability survey of the Westerlund~1 super star cluster and discovering 4 massive eclipsing binary systems.  Figure~\\ref{massradius} illustrates the extent of our knowledge of precise fundamental parameters of massive stars. It presents published mass-radius measurements from eclipsing binaries, accurate to better than $10\\%$ for the more massive component. The zero-age main sequences (ZAMS) at both Z=0.02 \\citep{Schaller92} and Z=0.008 \\citep{Schaerer93} are overplotted as a reference. The Galactic data are mainly taken from the compilations of \\citet{Andersen91} and \\citet{Gies03} with additions and updates from \\citet{Vitrichenko07}, but also \\citet{Gonzalez05} for the Large Magellanic Cloud (LMC), \\citet{Harries03} and \\citet{Hilditch05} for the SMC, and \\citet{Ribas05} for M31. A literature search was done to include all accurate measurements of stars in eclipsing binaries with masses $\\geq 30\\,\\msun$, which are presented in Table~\\ref{massradiusdata}. This Table is, to my knowledge, complete at present and consists of only 14 very massive stars with better than $10\\%$ mass-radius measurements, located in 3 galaxies. Of these, WR~20a and M33 X-7 \\citep{Orosz07} are the most massive and noteworthy. M33 X-7 contains a very massive $70.0 \\pm 6.9\\; \\msun$ O-type giant and a record-breaking $15.65\\;\\msun$ black hole, challenging current evolutionary models, which fail to explain such a large black hole mass. Without accurate measurements for a large sample of massive stars, theoretical models will remain unconstrained.  A survey to determine accurate parameters for several massive eclipsing binaries in the low metallicity (Z$=0.008$) LMC was undertaken, with the purpose of increasing the sample and improving our understanding of these rare systems. Several candidates were selected from the OGLE-II catalog of eclipsing binaries in the LMC \\citep{Wyrzykowski03} as the brightest systems with $B-V<0$. LMC-SC1-105, or OGLE053448.26-694236.4, is one of the brightest eclipsing binaries with $I_{max}=13.04$~mag, $V_{max}=12.97$~mag, $B_{max}=12.81$~mag and a preliminary semi-detached classification. This work presents the analysis of the follow-up observations obtained for LMC-SC1-105. The paper is organized as follows: \\S2 describes the spectroscopy and data reduction, \\S3 the spectral classification, \\S4 the radial velocity curve, \\S5 the light curve analysis, \\S6 the evolutionary status, and \\S7 the conclusion. ",
        "conclusions": "This paper presents accurate (to better than $5\\%$) fundamental parameters of LMC-SC1-105, one of the brightest blue eclipsing binary stars in the LMC found by the OGLE survey. The aim of this work is twofold: 1) to demonstrate that targeted surveys of the brightest blue eclipsing binaries in nearby galaxies do indeed select very massive stars and 2) to measure accurate parameters for one of these rare systems. The parameters of LMC-SC1-105 were determined from the light curves available from the OGLE and MACHO surveys and newly acquired high resolution spectroscopy that targeted quadrature phases, in part applying the strategy proposed by \\citet{Gonzalez05} to constrain the radial velocity curve with a small number of spectra. The system was found to contain a very massive main sequence primary ($30.9\\pm1.0\\,\\msun$) and a possibly evolved Roche lobe-filling secondary. The spectra display the Struve-Sahade effect, which is present in all the \\ion{He}{1} lines, causing the spectral classification to change with phase, and could be related to the mass transfer occurring in the system. LMC-SC1-105 could further be used as a distance indicator to the LMC. However, in addition to accurate radii, accurate flux (i.e. effective temperatures) and extinction estimates are necessary for accurate distances. Eclipsing binaries have been used to derive accurate and independent distances to the LMC \\citep[e.g.][]{Guinan98, Fitzpatrick03}, the Small Magellanic Cloud \\citep{Harries03,Hilditch05}, M31 \\citep{Ribas05} and most recently to M33 \\citep{Bonanos06}.  The accurate parameters determined herein for LMC-SC1-105 contribute valuable data on very massive stars, increasing the current sample of 14 very massive stars with accurate parameters to 15, which despite their importance remain poorly studied. Such data serve as an external check to resolve the ``mass discrepancy'' problem, as \\citet{Burkholder97} have shown, and to constrain stellar atmosphere, evolution and formation models. Further systematic studies of massive binaries in nearby galaxies are needed to extend the sample of 50 SMC eclipsing binaries \\citep{Harries03, Hilditch05} to higher masses and metallicities and populate the sparsely sampled parameter space (mass, metallicity, evolutionary state) with accurate measurements of their masses and radii. The method of targeting very massive stars in bright blue eclipsing binaries can therefore be employed towards this goal."
    },
    "0807/0807.3748_arXiv.txt": {
        "abstract": "Radiative cooling is central to a wide range of astrophysical problems. Despite its importance, cooling rates are generally computed using very restrictive assumptions, such as collisional ionization equilibrium and solar relative abundances. We simultaneously relax both assumptions and investigate the effects of photo-ionization of heavy elements by the meta-galactic UV/X-ray background and of variations in relative abundances on the cooling rates of optically thin gas in ionization equilibrium. We find that photo-ionization by the meta-galactic background radiation reduces the net cooling rates by up to an order of magnitude for gas densities and temperatures typical of the shock-heated intergalactic medium and proto-galaxies ($10^4\\,\\K \\la T \\la 10^6\\,\\K$, $\\rho / \\left<\\rho\\right> \\la 100$). In addition, photo-ionization changes the relative contributions of different elements to the cooling rates. We conclude that photo-ionization by the ionizing background and heavy elements both need to be taken into account in order for the cooling rates to be correct to order of magnitude. Moreover, if the rates need to be known to better than a factor of a few, then departures of the relative abundances from solar need to be taken into account. We propose a method to compute cooling rates on an element-by-element basis by interpolating pre-computed tables that take photo-ionization into account. We provide such tables for a popular model of the evolving UV/X-ray background radiation, computed using the photo-ionization package \\textsc{cloudy}. ",
        "introduction": "\\label{sec:intro} Dissipation of energy via radiative cooling plays a central role in many different astrophysical contexts. In general the cooling rate depends on a large number of parameters, such as the gas density, temperature, chemical composition, ionization balance, and the radiation field. In the absence of radiation, however, the equilibrium ionization balance depends only on the temperature. In that case the cooling rate in the low density regime, which is dominated by collisional processes, is simply proportional to the gas density squared, for a given composition. Thus, the cooling rates for a plasma in collisional ionization equilibrium (CIE) can be conveniently tabulated as a function of the temperature and composition (metallicity) of the gas \\cite[e.g.,][]{Cox1969,Raymond1976,Shull1982,Gaetz1983,Boehringer1989,Sutherland1993,Landi1999,Benjamin2001,Gnat2007,Smith2008}, and such tables are widely used for a large variety of problems.  Although it is convenient to ignore radiation when calculating cooling rates, radiation is generally important for the thermal and ionization state of astrophysical plasmas. For example, \\cite{Efstathiou1992} investigated the effect of the extragalactic UV background on the cooling rates for gas of primordial composition (in practice this means a pure H/He plasma) and found that including photo-ionization can suppress the cooling rates of gas in the temperature range $T\\sim 10^4 - 10^5\\,\\K$ by a large factor. Although the effects of radiation are often taken into account for gas of primordial composition, photo-ionization of heavy elements is usually ignored in the calculation of cooling rates (but see \\citealt{Leonard1998,Benson2002}).  In this paper we will investigate the dependence of cooling rates of gas enriched with metals on the presence of ionizing radiation, focusing on the temperature range $T\\sim 10^4 - 10^8\\,\\K$ and optically thin plasmas. We will show that, as is the case for gas of primordial composition \\cite[][]{Efstathiou1992}, photo-ionization can suppress the metallic cooling rates by a large factor. Moreover, the suppression of the cooling rate is significant up to much higher temperatures than for the primordial case.  We will also investigate the relative contributions of various elements to the cooling rates. If the relative abundances are similar to solar, then oxygen, neon, and iron dominate the cooling in the temperature range $T \\sim 10^4 - 10^7\\,\\K$. However, we will show that the relative contributions of different elements to the cooling rate are sensitive to the presence of ionizing radiation.  Although we will illustrate the results using densities and radiation fields that are relevant for studies of galaxy formation and the intergalactic medium (IGM), the conclusion that photo-ionization significantly reduces the cooling rates of enriched gas is valid for a large variety of astrophysical problems. For example, for $T \\sim 10^5\\,\\K$ and $T\\sim 10^6\\,\\K$ the reduction of metal-line cooling rates is significant as long as the dimensionless ionization parameter\\footnote{$U \\equiv \\Phi_{\\rm H}/(n_{\\rm H}c)$, where $\\Phi_{\\rm H}$ is the flux of hydrogen ionizing photons (i.e.,\\ photons per unit area and time), $n_{\\rm H}$ is the total hydrogen number density and $c$ is the speed of light.} $U \\ga 10^{-3}$ and $U \\ga 10^{-1}$, respectively. We will focus on the temperature range $10^4 - 10^8\\,\\K$ because gas in this temperature range is usually optically thin and because the effects of photo-ionization are generally unimportant at higher temperatures.  Tables containing cooling rates and several other useful quantities as a function of density, temperature, redshift, and composition, appropriate for gas exposed to the models for the evolving meta-galactic UV/X-ray background of \\cite{Haardt2001} are available on the following web site: \\texttt{http://www.strw.leidenuniv.nl/WSS08/}. The web site also contains a number of videos that illustrate the dependence of the cooling rates on various parameters.  This paper is organized as follows. In Section~\\ref{sec:method} we present our method for calculating element-by-element cooling rates including photo-ionization and we compare the limiting case of CIE to results taken from the literature. Section~\\ref{sec:photometals} shows how metals and ionizing radiation affect the cooling rates. Section~\\ref{sec:whim} demonstrates the importance for the low-redshift shock-heated IGM, which is thought to contain most of the baryons. In this section we also illustrate the effect of changing the intensity and spectral shape of the ionizing radiation. We investigate the effect of photo-ionization on the relative contributions of individual elements in Section~\\ref{sec:relabund} and we summarize and discuss our conclusions in Section~\\ref{sec:discussion}.  Throughout this paper we use the cosmological parameters from \\cite{Komatsu2008}: ($\\Omega_{\\rm m}, \\Omega_{\\Lambda}, \\Omega_{\\rm b}, h) = (0.279, 0.721, 0.0462, 0.701$) and a primordial helium mass fraction $X_{\\rm He} = 0.248$. Densities will be expressed both as proper hydrogen number densities $n_{\\rm H}$ and density contrasts $\\delta \\equiv \\rho_{\\rm b}/\\left <\\rho_{\\rm b}\\right >-1$, where $\\left < \\rho_{\\rm b}\\right >$ is the cosmic mean baryon density. The two are related by \\begin{equation} n_{\\rm H} \\approx 1.9 \\times 10^{-7}\\,\\cm^{-3} ~\\left (1+\\delta\\right )\\left (1+z\\right )^3 \\left ({X_{\\rm H} \\over 0.752}\\right ). \\end{equation} ",
        "conclusions": "\\label{sec:discussion} Radiative cooling is an essential ingredient of hydrodynamical models of a wide range of astrophysical objects, ranging from the IGM to (proto-)galaxies and molecular clouds.  While numerical simulations of objects with a primordial composition often compute non-equilibrium radiative cooling rates explicitly and sometimes even include the effect of ionizing background radiation, the treatment of cooling of chemically enriched material is typically much more approximate. For example, simulations of galaxy formation typically either ignore metal-line cooling altogether or include it assuming pure CIE. In addition, the abundances of all heavy elements are typically scaled by the same factor (the metallicity) (but see \\citealt{Martinez2008} and \\citealt{Maio2007} for recent exceptions). In this simplified treatment metal-line cooling depends only on temperature and metallicity, allowing straightforward interpolation from pre-computed two-dimensional tables.  We have used \\textsc{cloudy} to investigate the effects of heavy elements and ionizing radiation on the radiative cooling of gas with properties characteristic of (proto-)galaxies and the IGM, i.e., optically thin gas with densities $n_{\\rm H} \\la 1\\,\\cm^{-3}$ and temperatures $T\\ga 10^4\\,\\K$, assuming ionization equilibrium. We presented a method to incorporate radiative cooling on an element-by-element basis including photo-ionization by an evolving UV/X-ray background, using precomputed tables, which for heavy elements are functions of density, temperature, and redshift and for H\\&He (which must be considered together because they are important contributors to the free electron density) depend additionally on the He/H ratio. Using the 11 elements H, He, C, N, O, Ne, Mg, Si, S, Ca, and Fe, the redshift $z=0$ median absolute errors in the net cooling rate range from 0.33\\%, at $Z = 0.1 Z_{\\odot}$ to 6.1\\% for the extreme metallicity $Z= 10 Z_{\\odot}$, and the errors are smaller for higher redshifts.  The tables as well as some scripts that illustrate how to use them are available from the following web site: \\texttt{http://www.strw.leidenuniv.nl/WSS08/}. We also include tables for solar relative abundances which can be used if metallicity, but not the abundances of individual elements are known, as in equation (\\ref{eq:Zmethod}). This web site also contains a number of videos that may be helpful to gain intuition on the importance of various parameters on the cooling rates.  We confirmed that, assuming CIE, heavy elements greatly enhance the cooling rates for metallicities $Z \\ga 10^{-1}~Z_\\odot$ and temperatures $T \\la 10^7\\,\\K$. We demonstrated that this remains true in the presence of photo-ionization by the meta-galactic UV/X-ray background.  The background radiation removes electrons that would otherwise be collisionally excited, thus reducing the cooling rates. The effect is stronger for higher ionization parameters (i.e.,\\ higher radiation intensities or lower densities) and if the spectral shape of the radiation field is harder. Considering only the meta-galactic radiation field, which provides a lower limit to the intensity of the radiation to which optically thin gas may be exposed, the reduction of the metal-line cooling rates becomes important below $10^6\\,\\K$ for ionization parameters $U \\ga 10^{-1}$ and below $10^5\\,\\K$ for $U \\ga 10^{-3}$ (note that for the HM01 background $U=9\\times 10^{-7}/n_{\\rm H}$ and $2\\times 10^{-5}/n_{\\rm H}$ at $z=0$ and $z=3$, respectively).  As an example of the potential importance of including the effects of both photo-ionization and heavy elements, we considered the so-called warm-hot intergalactic medium (WHIM), which is thought to contain a large fraction of the baryons at redshifts $z < 1$. We demonstrated that the overdensities for which gas at typical WHIM metallicities ($Z\\sim 10^{-1}~Z_\\odot$) and temperatures ($T\\sim 10^5 - 10^7\\,\\K$) can cool within a Hubble time, can shift by an order of magnitude depending on whether photo-ionization and metal-line cooling are taken into account. Hence, photo-ionization of heavy elements may have important consequences for predictions of the amount of matter contained in this elusive gas phase.  Because chemical enrichment happens in a number of stages, involving a number of processes with different timescales, the relative abundances of the heavy elements varies with redshift and environment by factors of a few. Hence, computing cooling rates on an element-by-element basis rather than scaling all elements by the metallicity, will change the cooling rates by factors of a few. The difference is therefore typically somewhat smaller than the effect of neglecting metals or photo-ionization altogether, but still highly significant.  While it was known that different elements dominate the cooling for different temperatures in CIE, we showed that photo-ionization both shifts the peaks due to individual elements to smaller temperatures and reduces their amplitude. Note that since photo-ionization over-ionizes the gas, this effect is similar (but not equivalent to) that found in non-equilibrium calculations without ionizing radiation \\cite[e.g.,][]{Sutherland1993, Gnat2007}. Because the importance of photo-ionization depends on the ionization parameter, the relative contributions of individual elements exposed to a fixed ionizing radiation field depends also on the gas density.  Would dropping our assumption of ionization equilibrium have a large effect on the cooling rates?  Ionizing radiation results in a plasma that is overionized relative to its temperature. Its effect is therefore similar to that of non-equilibrium ionization following rapid cooling (i.e., if the cooling time is shorter than the recombination times of the ions dominating the cooling, see e.g.,\\ \\citealt{Kafatos1973,Shapiro1976}). We therefore anticipate that the effect of non-equilibrium ionization will be much smaller for our cooling rates than for those that assume CIE.  The assumptions that the gas is optically thin and exposed only to the meta-galactic background radiation are likely to be more important than the assumption of ionization equilibrium, particularly since non-equilibrium collisional cooling rates only differ from those assuming CIE by factors of a few or less \\cite[e.g.,][]{Schmutzeler1993,Sutherland1993,Gnat2007}. For column densities $N_{\\rm HI} > 10^{17}~\\cm^{-2}$ self-shielding becomes important and only part of the H-ionizing radiation will penetrate the gas cloud, which would particularly affect the cooling rates for $T\\la 10^5\\,\\K$. At higher temperatures line cooling is dominated by heavier ions, which can only be ionized by higher energy photons and which therefore remain optically thin up to much higher column densities. This is because the photo-ionization cross sections of H and He drop rapidly with increasing frequency for energies exceeding their ionization potentials. Moreover, for $T\\gg 10^4~\\K$ hydrogen is collisionally ionized to a high degree and consequently the optical depth for ionizing radiation will be significantly reduced.  It is, however, far from clear that high column densities would reduce the effect of radiation. For self-shielded clouds the cooling radiation may itself be trapped, providing a source of ionizing radiation even in the absence of an external one \\cite[e.g.,][]{Shapiro1976,Gnat2007}. Moreover, gas clouds with columns that exceed $10^{17}~\\cm^{-2}$ are on average expected to be sufficiently close to a galaxy that local sources of ionizing radiation dominate over the background \\citep{Schaye2006,Miralda2005}.  Ultimately, these issues can only be resolved if non-equilibrium cooling rates are computed including radiative transfer and if the locations of all relevant sources of ionizing radiation are known.  It will be some time before it is feasible to carry out such a calculation in, say, a cosmological hydrodynamical simulation. In the mean time, we believe that our element-by-element calculation of the equilibrium cooling rates for an optically thin gas exposed to the CMB and an evolving UV/X-ray background provides a marked improvement over earlier treatments. In future publications we will present cosmological, hydrodynamical simulations using these cooling rates."
    },
    "0807/0807.4442_arXiv.txt": {
        "abstract": "We use a combination of analytic tools and an extensive set of the largest and most accurate three-dimensional field theory numerical simulations to study the dynamics of domain wall networks with junctions. We build upon our previous work and consider a class of models which, in the limit of large number $N$ of coupled scalar fields, approaches the so-called `ideal' model (in terms of its potential to lead to network frustration). We consider values of $N$ between $N=2$ and $N=20$, and a range of cosmological epochs, and we also compare this class of models with other toy models used in the past. In all cases we find compelling evidence for a gradual approach to scaling, strongly supporting our no-frustration conjecture. We also discuss the various possible types of junctions (including cases where there is a hierarchy of them) and their roles in the dynamics of the network. Finally, we revise the Zel'dovich bound and provide an updated cosmological bound on the energy scale of this type of defect network: it must be lower than $10 \\, {\\rm keV}$. ",
        "introduction": "Introduction}  As the early universe expanded and cooled down, it is believed to have gone through a series of phase transitions, at which networks of topological defects must necessarily have formed \\cite{KIBBLE}. The type of defect that forms and its specific properties depend on the particular details of each symmetry breaking, and hence there is a wide range of possibilities, which will lead to correspondingly very different cosmological consequences \\cite{VSH}. Domain walls are known to be pathological except if they are very light \\cite{ZEL}, but it has been claimed \\cite{SOLID} that if a domain wall network is frozen in co-moving coordinates (or `frustrates', as is often colloquially put) then it can naturally explain the observational evidence that points to a recent acceleration of the universe.  It is clear that this scenario is subject to a number of observational constraints. For example the cosmic microwave background (CMB) data \\cite{WMAP5} severely constrains the characteristic scale of the network, $L$, which needs to be tiny in order not to give rise to exceedingly large CMB fluctuations. Also, recalling that the equation of state of a domain wall gas is given by \\begin{equation} w \\equiv \\frac{p}{\\rho} = -\\frac{2}{3} + v^2\\,, \\end{equation} with $v$ being the root-mean squared (RMS) velocity of the walls and that we require that \\begin{equation} w < -\\frac{1}{3}\\left(1+\\frac{\\Omega_m^0}{\\Omega_{DE}^0}\\right) \\lsim -\\frac{1}{2} \\end{equation} in order to accelerate the universe at the present time it is clear that only a non-relativistic value of the velocity (in fact, basically $v\\sim0$) would have any chance of working. The simplest domain wall models are known to reach a scaling regime (until they dominate the energy density of the universe), as first pointed out in \\cite{PRESS} and recently studied in detail in \\cite{AWALL,SIMS1,SIMS2}, and hence are obviously unable to satisfy these constraints. Nevertheless, it was thought that more complicated models, notably those having junctions, would eventually frustrate and therefore might conceivably be able to satisfy them.  In previous work \\cite{IDEAL1,IDEAL2} we have studied the dynamics of domain wall networks with junctions, and investigated in detail energy, geometrical and topological constraints on their properties. This led us to develop an ideal class of models which includes, in the large $N$ limit, what we called the \\textit{ideal model} (that is, the best possible candidate for frustration). A series of analytic and numerical arguments then led us to a \\textit{no frustration conjecture}: even though one can build (purely by hand, as was done in \\cite{BATTYE,CARTER,IDEAL2}) special lattices that would be locally stable against small perturbations, no such configurations are expected to ever emerge from any realistic cosmological phase transition. (A much simpler analysis of the evolution of non-interacting and entangled cosmic string networks using a velocity-dependent one-scale model had already led to a similar result, at least in four space-time dimensions \\cite{NONINT}.) Our high-resolution numerical simulations of various realizations of the ideal class and other models showed clear evidence of a gradual approach to scaling, which was subsequently confirmed in \\cite{BMFLAT} (though note that the latter only reports on numerical simulations in Minkowski space thus neglecting the role of the expansion).  Still, a fair criticism of our earlier work \\cite{IDEAL1,IDEAL2} would be that we only considered numerical simulations in two spatial dimensions, and the behavior could well be different in the three spatial dimensions of our universe. More recently \\cite{THIRD} we endeavored to eliminate this shortcoming and presented results of a series of massively parallel domain wall network numerical simulations of the `ideal' class of models in three spatial dimensions, which confirmed our earlier work, providing conclusive evidence for a gradual approach to scaling and hence strongly supporting our no frustration conjecture. The present article is the fourth in this series of papers, and its aims are to present the results of \\cite{THIRD} more thoroughly, to extend them in a number of different ways, and to discuss in some detail the corresponding cosmological implications.  In Sect. \\ref{swal} we introduce various approaches to the study of domain walls---the microscopic description, the analytic modeling and the building of phenomenological models---that we will be using in the rest of the paper. In particular we highlight some relevant scaling solutions for domain wall networks and describe the ideal class of models which will be our primary focus. In Sect. \\ref{sfield} we discuss several technical details of our massively parallel field theory numerical code (which uses an improved version of the algorithm of Press, Ryden and Spergel \\cite{PRESS}, henceforth referred to as the PRS algorithm). In particular we describe in some detail how we identify the domain walls and how we measure their velocities. Sect. \\ref{ssim} contains our main numerical results: we describe the outcome of our various sets of three-dimensional high-resolution simulations of the ideal class of models. We also discuss the relevance of the key findings and contrast the results with those of other models. This then leads us to discuss, in Sect. \\ref{smod}, the analytic modeling of our results, including the cosmologically important asymptotic limit of the ideal class of models. In Sect. \\ref{sjun} we present an analysis of the possible hierarchies of junctions in the ideal and other models, as well as a brief discussion of the role of the junctions on the dynamics of the networks. Finally Sect. \\ref{scon} has our conclusions, as well as some comments on the implications of our results for cosmological scenarios involving domain walls and a brief outline of future endeavors in this area. ",
        "conclusions": ""
    },
    "0807/0807.1098_arXiv.txt": {
        "abstract": "We present a study of the individual super star clusters (SSCs) in the low-metallicity galaxy SBS 0335-052 using new near-infrared and archival optical {\\it Hubble Space Telescope} observations.  The physical properties of the SSCs are derived from fitting model spectral energy distributions (SEDs) to the optical photometry, as well as from the H$\\alpha$ and Pa$\\alpha$ nebular emission.  Among the clusters, we find a significant age spread that is correlated with position in the galaxy, suggesting successive cluster formation occurred in SBS 0335-052 triggered by a large-scale disturbance traveling through the galaxy at a speed of $\\sim 35$ km s$^{-1}$.  The SSCs exhibit $I$-band ($\\sim 0.8~\\mu$m) and near-IR ($\\sim 1.6-2.1~\\mu$m) excesses with respect to model SEDs fit to the optical data.  We hypothesize that the $I$-band excess is dominated by a photoluminescent process known as Extended Red Emission; however, this mechanism cannot account for the excesses observed at longer near-IR wavelengths.  From the cluster SEDs and colors, we find that the primary origin of the near-IR excess observed in the youngest SSCs ($\\lesssim 3$ Myr) is hot dust emission, while evolved red supergiants dominate the near-IR light in the older ($\\gtrsim 7$ Myr) clusters.  We also find evidence for a porous and clumpy interstellar medium (ISM) surrounding the youngest, embedded SSCs: the ionized gas emission underpredicts the expected ionizing luminosities from the optical stellar continuum, suggesting ionizing photons are leaking out of the immediate vicinity of the clusters before ionizing hydrogen. The corrected, intrinsic ionizing luminosities of the two SSCs younger than $\\sim 3$~Myr are each $\\sim 5 \\times 10^{52}~{\\rm s}^{-1}$, which is equivalent to each cluster hosting $\\sim 5000$ O7.5 V stars.  The inferred masses of these SSCs are $\\sim 10^6 M_\\odot$. ",
        "introduction": "Ancient globular clusters are among the oldest objects known, almost as old as the universe itself with ages estimated in excess of 12~Gyr \\citep{Vandenberg96, Freeman02}.  As such, these clusters are valuable relics of the earlier universe, when violent and intense star formation was common.  Indeed, globular clusters are ubiquitous around massive galaxies today \\citep{Harris91, Brodie06}, and they must have been formed prodigiously in the primordial universe given that $\\gtrsim 90\\%$ may have been subsequently destroyed \\citep{Fall01,Whitmore07}.  However, the conditions required for the creation of globular clusters have puzzled astronomers for decades. For many years the prevailing belief was that globular clusters were simply formed by the gravitational collapse of density inhomogeneities in the early universe \\citep{Peebles68, Fall85}, and little was known about their early evolution. However, in the mid 1990's, observations using the {\\it Hubble Space Telescope} discovered extremely young, massive, and compact clusters in the local universe, the so-called ``super star clusters'' (SSCs) that have properties consistent with those expected for adolescent globular clusters \\citep[e.g.][]{Whitmore03}.  The discovery of SSCs precipitated a major shift in our understanding of the conditions required for globular cluster formation.  Although the evidence connecting young SSCs and ancient globular clusters is compelling, a critical issue remains: modern day SSCs and ancient globular clusters were formed in environments with very different metallicities.  We have not yet observationally constrained what effect low metal abundances had on the formation of massive star clusters. Much effort has gone into understanding the formation of the {\\it first} stars in the universe \\citep[e.g. those stars formed out of truly primordial material,][]{Bromm03,Santoro06,Tumlinson07}, and it is clear that metallicity has a critical role in this regime.  Even moderately low metal abundances may affect star formation in a number of ways, including dust formation, cooling and pressure, and the hardness of the resulting stellar spectra.  \\begin{figure*}[!t] \\begin{center} \\includegraphics[scale=0.61]{f1.jpg} \\caption{{\\it HST} images of the observations used in this work. \\label{allims}} \\end{center} \\end{figure*}  In an effort to investigate massive star cluster formation in an environment similar to that which might be found in primordial galaxies during the time ancient globular clusters were prolifically formed, we have targeted the galaxy SBS~0335-052 (E) for a detailed multi-wavelength study.  SBS~0335-052 is a remarkable blue compact dwarf galaxy at a distance of 54 Mpc (NED\\footnote{NASA/IPAC Extragalactic Database}) that is well-known for its extremely low oxygen abundance of $\\sim 1/40Z_\\odot$ \\citep{Izotov90,Melnick92,Izotov01}.  Unlike I~Zw~18 and SBS 0335-052W, which have slightly lower metallicities, SBS~0335-052 is undergoing a vigorous starburst with a star formation rate $\\gtrsim 1~M_\\odot$~yr$^{-1}$ and hosts extremely massive young clusters, making it ideal for this study \\citep{Thuan97, Thuan99, Hunt01, Dale01, Plante02, Hunt04, Houck04}.  Here we present new near-infrared and archival optical high-resolution observations of SBS 0335-052 with the goal of studying SSC formation and evolution at low metallicity.  Near-IR observations are particularly important for probing the dusty birth environments of embedded massive star clusters, and the wavelength range of $\\sim 1-2~\\mu$m is well-suited to assessing a cluster's evolutionary state because it samples the spectral energy distribution (SED) at the nexus of emission between dust and stellar light.  The {\\it HST} observations and photometry are presented in \\S\\ref{hstsec} and \\S\\ref{photsec}.  In \\S\\ref{propsec}, we discuss various properties of the SSCs. Specifically, we derive the physical properties of the SSCs from fitting model SEDs to the optical photometry, as well as from the H$\\alpha$, Pa$\\alpha$, and thermal radio nebular emission \\citep{Johnson08}.  In addition, we present evidence for a porous and clumpy ISM in the youngest embedded clusters in SBS 0335-052 which can account for the apparently discrepant extinction estimates found in the literature.  We also provide explanations for the observed {\\it I}-band ($\\sim 0.8~\\mu$m) and near-IR ($\\sim 1.6-2.1~\\mu$m) excesses.  Evidence for successive cluster formation is given in \\S\\ref{suc_sec}, and finally, our conclusions are summarized in \\S\\ref{conclusions}. ",
        "conclusions": "We have presented a multi-wavelength study of the super star clusters in SBS 0335-052, the lowest metallicity galaxy with a star formation rate $\\gtrsim 1~M_\\odot$~yr$^{-1}$ known.  We use new near-IR and archival optical {\\it HST} observations as well as free-free radio continuum measurements \\citep{Johnson08} of the galaxy to probe the stellar populations and the gas and dust components of the clusters.  The main results of our work are summarized below:  \\begin{enumerate}  \\item{There is a significant age spread between the SSCs in SBS 0335-052 that is correlated with position in the galaxy. The ages are in the range $\\lesssim 3$ Myr to $\\sim 15$ Myr with the youngest clusters located in the south and the oldest clusters in the north.  A large-scale disturbance, with a velocity of $\\sim 35$ km s$^{-1}$, appears to have triggered the successive cluster formation that we observe.}  \\item{We find evidence for a porous and clumpy ISM surrounding the youngest SSCs in SBS 0335-052 (1 and 2).  The measured ionizing luminosities from H$\\alpha$, Pa$\\alpha$, {\\it and} optically-thin, free-free radio emission \\citep{Johnson08} are lower than expected compared to the optical SEDs, suggesting a fraction of ionizing photons from the stellar continuum are escaping the strict confines of the clusters before ionizing hydrogen, that would have otherwise contributed to the measured ionized gas emission (i.e. H$\\alpha$, Pa$\\alpha$, and radio).  The $A_V$'s of SSCs 1 and 2 derived from SED fitting ($\\sim 0.5$ mag) also indicate the existence of low-extinction regions that provide relatively unobscured lines-of-sight into the embedded clusters.  The corrected, intrinsic ionizing luminosities of SSCs 1 and 2 are each $\\sim 5 \\times 10^{52}~{\\rm s}^{-1}$ (the equivalent of $\\sim 5000$ O7.5 V stars) and they each have total stellar masses of $\\sim 10^6 M_\\odot$.}  \\item{An $I$-band (F791W, $\\sim 0.8~\\mu$m) excess with respect to model SEDs is observed for all of the SSCs in SBS 0335-052.  The $I$-band magnitudes are $\\sim 0.5-0.7$ mag brighter than the model predictions.  A similar result was found in the young massive clusters in NGC 4449 \\citep{Reines08} and attributed to a photoluminescent process known as Extended Red Emission. We hypothesize that the same mechanism dominates the $I$-band excess in SBS 0335-052.}  \\item{All of the SSCs have red near-IR ($\\sim 1.6-2.1~\\mu$m) colors and large near-IR excesses with respect to model SEDs fit to the optical photometry.  The red near-IR colors, however, are quantitatively different (redder) in the youngest SSCs (1 and 2) compared to the older SSCs (3-6), clearly indicating a different origin for the red colors.  We have demonstrated that the near-IR light from the older clusters is dominated by evolved red supergiants, whereas the near-IR light from the youngest clusters is dominated by hot dust emission.  The young SSCs 1 and 2 are also the only radio-detected clusters in the galaxy \\citep{Johnson08}.}  \\item{We have proposed a scenario that can account for the apparently discrepant extinction estimates found in the literature for the starburst region encompassing the youngest SSCs (1 and 2) in SBS 0335-052.  Our picture is consistent with all of the emission (at all wavelengths) being associated with the optically visible SSCs 1 and 2, and is based on our calculated absorptions of the optical and IR recombination lines relative to the (thermal) free-free radio emission \\citep{Johnson08}. We find that an ISM containing dense dust clumps harbouring large grains ($\\sim 1-2~\\mu$m), with a dust clump covering factor of $\\sim 60\\%$, can account for the absorption of the optical/near-IR recombination lines relative to the radio, as well as the large extinctions derived from previous mid-IR observations.  Interclump regions containing diffuse dust can account for the low {\\it measured} extinctions derived from optical/near-IR observations.}  \\end{enumerate}  The main goal of this case study was to investigate the formation of massive star clusters in an extremely low metallicity environment that might be analogous to conditions in which the now ancient globular clusters were born.  In general, the formation of massive star clusters in SBS 0335-052 appears to be similar to that found in higher metallicity counterparts, e.g. Henize~2-10 and Haro~3 \\citep{Johnson03, Johnson04}.  These galaxies are all blue compact dwarfs (BCDs), and each of them currently hosts radio-detected massive star clusters with ages $\\lesssim 3$~Myr and masses of $\\sim 10^6 M_\\odot$, all of which are consistent with evolving into globular clusters in several billion years.  Perhaps even the low metal abundance in SBS 0335-052 is insufficiently extreme to reflect primordial star formation; in fact, models predict that the transition from metal-free massive Population III stars to ``normal'' (less massive) Population II occurs at $Z/Z_\\odot\\lesssim 10^{-5}$ \\citep{Mackey03,Bromm04}.  On the other hand, large metallicity fluctuations because of radiative feedback effects may enable ``hidden'' Population III star formation at slightly higher metal abundances \\citep{Tornatore07}.  Clearly more detailed studies need to be carried out on metal-poor star-forming galaxies in the Local Universe, and comparative samples need to be obtained in order to investigate interactions and trends among star cluster and dust formation, environment, and metal enrichment."
    },
    "0807/0807.1321_arXiv.txt": {
        "abstract": "The highly amplified magnetic fields suggested by observations of some supernova remnant (SNR) shells are most likely an intrinsic part of efficient particle acceleration by shocks. This strong turbulence, which may result from cosmic ray driven instabilities, both resonant and non-resonant, in the shock precursor, is certain to play a critical role in \\SC, \\NL\\ models of strong, \\CR\\ modified shocks. Here we present a Monte Carlo model of nonlinear diffusive shock acceleration (DSA) accounting for magnetic field amplification through resonant instabilities induced by accelerated particles, and including the effects of dissipation of turbulence upstream of a shock and the subsequent precursor plasma heating.  Feedback effects between the plasma heating due to turbulence dissipation and particle injection are strong, adding to the \\NL\\ nature of efficient DSA. Describing the turbulence damping in a parameterized way, we reach two important results: first, for conditions typical of supernova remnant shocks, even a small amount of dissipated turbulence energy ($\\sim 10\\%$) is sufficient to significantly heat the precursor plasma, and second, the heating upstream of the shock leads to an increase in the injection of thermal particles at the subshock by a factor of several. In our results, the response of the non-linear shock structure to the boost in particle injection prevented the efficiency of particle acceleration and magnetic field amplification from increasing. We argue, however, that more advanced (possibly, non-resonant) models of turbulence generation and dissipation may lead to a scenario in which particle injection boost due to turbulence dissipation results in more efficient acceleration and even stronger amplified magnetic fields than without the dissipation. ",
        "introduction": "Observations of young supernova remnants (SNRs) suggest that strong collisionless shocks can simultaneously place a large fraction of the shock ram kinetic energy into \\rel\\ protons \\citep[e.g.,][]{BE87,MD2001, WarrenEtal2005} and amplify the ambient turbulent magnetic field by large factors \\citep[e.g.,][]{Cowsik80, RE92, BambaEtal2003, BKV2003, VL2003, UA2008}. This coupling of diffusive shock acceleration (DSA) and magnetic field amplification (MFA) is critically important because the self-generated magnetic field largely determines the efficiency of DSA, the maximum particle energy a given shock can produce, and the \\syn\\ emission from radiating electrons.  The generation and dissipation of strong MHD turbulence in collisionless, multi-fluid plasmas is a complex process. Different \\NL\\ approaches to the modeling of the large scale structure of a shock undergoing  efficient cosmic ray acceleration \\citep[e.g.,][]{MV82,AB86,BL2001,AB2006,VEB2006,PLM2006,ZPV2008} have predicted the presence of strong MHD turbulence in the shock precursor. However, an exact modeling of the shock structure in a turbulent medium, including nonthermal particle injection and acceleration, requires a nonperturbative, self-consistent description of a multi-component and multi-scale system including the strong MHD-turbulence dynamics. While a number of analytic models describing resonant and non-resonant amplification and damping of magnetic fluctuations have been proposed, these generally rely on the quasi-linear approximation that the fluctuations are small compared to the background magnetic field, i.e., $\\Delta B \\ll B_0$ \\citep[e.g.,][]{Wentzel74,BT2005,Kulsrud2005}. No consistent analytic description of magnetic turbulence generation with $\\Delta B \\gtrsim B_0$ exists. For these reasons, numerical models with varying ranges of applicability have been proposed which offer a compromise between completeness and speed \\citep[e.g.,][]{Bell2004,AB2006,VEB2006,ZPV2008}.  In principle, the problem can be solved completely with few assumptions and approximations with particle-in-cell (PIC) simulations \\citep[e.g.,][]{Bell2004,Spitkovsky2008,NPS2008} or, in the assumption that electrons are not dynamically important, by hybrid models \\citep[e.g.,][]{WO1996, Giacalone2004}. However, modeling the \\NL\\ generation of \\rel\\ particles and strong magnetic turbulence in collisionless shocks is computationally challenging and PIC simulations will not be able to fully address this problem in \\nonrel\\ shocks for some years to come even though they can provide critical information on the plasma processes producing injection that can be obtained in no other way. In Appendix~\\ref{estimatesforpic} we outline the requirements that a PIC simulation must fulfill in order to tackle the problem of efficient DSA with non-linear MFA in SNR shocks.  In the \\mc\\ approximation we use here, the plasma interactions are parameterized allowing us to study coupled \\NL\\ effects between the extended shock precursor and the gas subshock. In particular, we investigate the \\NL\\ effects caused by upstream plasma heating due to magnetic field dissipation.  The importance of the dissipation of turbulence in the shock precursor can be illustrated by the following estimate. Suppose that in a shock wave of speed $u_0$, the turbulence is generated by the resonant cosmic ray (CR) streaming instability, so the energy density of the turbulence, $U_w$, evolves approximately as $u_0\\,dU_w(x)/dx=v_A\\,d\\Pcr(x)/dx$ \\citep[e.g.,][]{BL2001}, where $\\Pcr(x)$ is the CR pressure at position $x$ and $v_A$ is the \\Alf\\ speed. Ignoring all non-linear effects, the turbulence energy density at the shock positioned at $x=0$ is $U_w(0)=\\rho_0 u_0 \\vazero \\cdot \\Pcr(0)/(\\rho_0 u_0^2)$. The ratio $\\AccEff=\\Pcr(0)/(\\rho_0 u_0^2)$ characterizes the efficiency of acceleration and is typically assumed to be on the order of ten percent or more. In the above, $\\rho$ is the fluid density and the subscript ``0'' always indicates far upstream values. Suppose a fraction, $\\heatpar$, of this energy goes into heating of the thermal gas in the shock precursor so the energy density of the thermal plasma increases by $\\Delta U_{H}(0) = \\heatpar U_w(0)$ at $x=0$. Comparing $\\Delta U_{H}(0)$ with the internal energy density of the far upstream plasma, $\\epsilon_0$, we find \\begin{equation} \\label{upsilonequation} \\eta_H\\ = \\frac{\\Delta U_H(0)}{\\epsilon_0} \\approx \\heatpar \\AccEff \\frac{M_s^2}{M_A} \\ , \\end{equation} where $M_s$ is the sonic, and $M_A$ the \\Alf, Mach number. If $\\eta_H$ is large, the thermal gas in the shock precursor is strongly heated and this influences the subshock strength and the particle injection efficiency. In a non-linearly modified shock, a change in the injection efficiency causes the whole shock structure to change. For typical supernova remnant (SNR) parameters (e.g., $u_0 \\sim 3000$ \\kmps, $T_0 \\sim 10^4$ K, $n_0 \\sim 0.3$ \\pcc, and $B_0 \\sim 3$ \\muG), the ratio $M_s^2/M_A \\approx 250$, and values of $\\heatpar$ as low as a few to ten percent may be important.  Because existing analytical descriptions of MHD wave damping rely on the quasi-linear approximation $\\Delta B \\ll B_0$, which is inapplicable for strong turbulence, and because an exact description of this process in the framework of non-linear DSA is currently impossible (see Appendix~\\ref{estimatesforpic}), we propose a parameterization of the turbulence damping rate. In doing this, we are pursuing two goals. First, we make some predictions connecting cosmic ray spectra, turbulent magnetic fields and plasma temperatures, which, in principle, can be tested against high resolution X-ray observations in order to estimate the heating of the thermal gas by turbulence dissipation. And second, once heating is included in our simulation in a parameterized fashion, we will be ready to implement more realistic models of turbulence generation and dissipation as they are developed.  Our \\mc\\ simulation can be briefly summarized as follows \\citep[see][ and references therein for more complete details]{EJR90,JE91,VEB2006}. We describe particle transport in a plane shock by Bohm diffusion in a plasma flowing in the $x$-direction with speed $u(x)$. Particles move in small time steps as their local plasma frame momenta are `scattered' at each step in a random walk process on a sphere in momentum space. Some shock heated thermal particles are injected into the acceleration process when their history of random scatterings in the downstream region takes them back upstream.  These particles gain energy and some continue to be accelerated in the first-order Fermi mechanism. This form of injection is generally called `thermal leakage' and was first used in the context of DSA in \\citet{EJE1981} \\citep[see also][]{Ellison82}.  The magnetic field determining the random walk properties through the diffusion coefficient is the `seed' (interstellar) magnetic field after it has been amplified by a large factor by the CR streaming instability and compressed and advected downstream with the flow.  The streaming instability in the non-linear regime ($\\Delta B \\gg B_0$) is described by the traditional quasi-linear equations, where the instability driving term is the CR pressure gradient. These quasi-linear equations are extrapolated into the non-linear regime in a parameterized fashion for lack of a more complete analytic description. The magnetic turbulence generated by the instability is assumed to dissipate at a rate proportional to the turbulence generation rate, and the dissipated energy is pumped directly into the thermal particle pool.  An iterative scheme is employed to ensure the conservation of mass, momentum, and energy fluxes, thus producing a self-consistent solution of a steady-state, plane shock, with particle injection and acceleration coupled to the bulk plasma flow modification and to the magnetic field amplification and damping.   Our results show that even a small rate of turbulence dissipation can significantly increase the precursor temperature and that this, in turn, can increase the rate of injection of thermal particles. The \\NL\\ feedback of these changes on the shock structure, however, tend to cancel so that the spectrum of high energy particles is only modestly affected. ",
        "conclusions": "\\label{conclusions}  We have parameterized magnetic turbulence dissipation as a fraction of turbulence energy generation and included this effect in our \\MC\\ model of strongly \\NL\\ shocks undergoing efficient DSA. The energy removed from the turbulence goes directly into the thermal particle population in the shock precursor. The \\mc\\ simulation \\SCly\\ reacts to the changes in precursor heating and adjusts the injection of thermal particles into the DSA mechanism, as well as other \\NL\\ effects of DSA, accordingly.  Our two most important results are, first, that even a small rate ($\\sim 10$\\%) of turbulence dissipation can drastically increase the precursor temperature, and second, that the precursor heating boosts particle injection into DSA by a large factor. The increase in particle injection modifies the low-energy part of the particle spectrum but, due to \\NL\\ feedback effects, does not significantly change the overall efficiency or the high energy part of the spectrum. Both the precursor heating and modified spectral shape that occur with dissipation may have observable consequences.  The fact that the shock back-reaction to the increased injection prevents the acceleration efficiency from changing significantly is a clear consequence of the non-linear structure of the system. The boosted particle injection additionally smoothes the flow speed close to the subshock, which makes it harder for particles returning upstream to gain energy. As a result, the population of the high energy particles is not much changed (except for the decrease in the maximum particle momentum due to the reduction in the effective amplified magnetic field from dissipation) and, because those particles carry the bulk of the CR pressure $\\Pcr$ which drives the streaming instability, the amplification of the magnetic field is not strongly affected by the heating-boosted particle injection.  The parameterization we use here is a simple one and a more advanced description of the turbulence damping may change our results. In our model the energy drained from the magnetic turbulence, at all wavelengths, is directly `pumped' into the thermal particles. Superthermal particles only gain extra energy due to heating because the thermal particles were more likely to return upstream and get accelerated.  In a more advanced model of dissipation, where energy cascades from large-scale turbulence harmonics to the short-scale ones, the low energy CRs might gain energy directly from the dissipation. It is conceivable that cascading effects might increase the overall acceleration efficiency, the magnetic field amplification, and increase the maximum particle energy a shock can produce.  It is also possible that non-resonant turbulence instabilities play an important role in magnetic field amplification \\citep[e.g.,][]{PLM2006}.  This opens another possibility for the turbulence dissipation to produce an increase in the magnetic field amplification.  For instance, \\citet{BT2005} proposed a mechanism for generating long-wavelength perturbations of magnetic fields by low energy particles. If such a mechanism is responsible for generation of a significant fraction of the turbulence that confines the highest energy particles, then the increased particle injection due to the precursor heating may raise the maximum particle energy and, possibly, the value of the amplified magnetic field.  While our model is for the most part phenomenological as far as particle transport, injection and acceleration, magnetic field generation, and dissipation are concerned, it allows us to investigate the coupled nonlinear effects in a shock undergoing efficient DSA and MFA. The more efficient DSA is, the more basic considerations of momentum and energy conservation determine the shock structure and our model describes these effects fully."
    },
    "0807/0807.3597.txt": {
        "abstract": "With only six known examples, M-dwarf debris disks are rare, even though M dwarfs constitute the majority of stars in the Galaxy. After finding a new M dwarf debris disk in a shallow mid-infrared observation of NGC~2547, we present a considerably deeper \\textsl{Spitzer}-MIPS image of the region, with a maximum exposure time of 15~minutes per pixel. Among sources selected from a previously published membership list, we identify nine new M dwarfs with excess emission at 24~$\\mu$m tracing warm material close to the snow line of these stars, at orbital radii of less than 1~AU. We argue that these are likely debris disks, suggesting that planet formation is under way in these systems. Interestingly, the estimated excess fraction of M stars appears to be higher than that of G and K stars in our sample. ",
        "introduction": "Even though about 80~\\% of the stars in the Galaxy are M~dwarfs \\citep{lad06}, only very few of these are known to harbor debris disks. This is even more surprising since many M~dwarfs do show signs of protoplanetary disks in earlier stages of evolution with similar frequency as more massive stars (e.g., \\citealp{and05}, \\citealp{muz06}, \\citealp{lad06b}) and we now know seven M dwarfs that harbor extrasolar planets (most recently \\citealp{end08}, see references therein). While protoplanetary disks contain large amounts of primordial gas and dust out of which planets form, debris disks are older and entirely made of collisionally evolving planetesimals and dust particles. Debris disks, thus, are secondary products of the planet formation process and their existence in a given system implies previous formation of protoplanets. Compared to optically thick protoplanetary disks, debris disks are fainter and optically thin. For a recent observational and theoretical review of debris disks, see \\citet{mey07} and \\citet{mor07}.  Since the cold dust in debris disks is best studied at far-infrared (FIR) or submillimeter wavelengths, the \\textsl{Spitzer} Space Telescope, given its unprecedented FIR sensitivity, has revolutionized the research on the role of debris disks in star and planet formation. Most of the \\textsl{Spitzer} studies of debris disks have been carried out using the Multiband Imaging Photometer for Spitzer \\citep[MIPS,~][]{rie04}, operating at wavelengths of 24~$\\mu$m, 70~$\\mu$m, and 160~$\\mu$m.  Prior to our study, only six debris disks around M dwarfs were known. Early studies based on data obtained by IRAS, the Infrared Astronomical Satellite, reported far-infrared excess emission towards a small number of field M stars \\citep{tsi88,mul89,mat91}, but few of these results could be confirmed subsequently \\citep[e.g.,][]{son02,ria06}. The first firm detection of an excess was the disk of AU~Mic \\citep{kal04,liu04b}, a member of the $\\beta$~Pic moving group with an age of $\\approx$12~Myr. The discovery observation was made with a stellar coronograph detecting scattered light in near-infrared broad-band filters. AU~Mic was also found to have excess emission at 70~$\\mu$m, but not at 24~$\\mu$m \\citep{che05}. Subsequently, in submillimeter radio observations at 850~$\\mu$m, \\citet{liu04a} identified GJ~182, a member of the Local Association Group with an age of $\\approx$50~Myr, as harboring a debris disk. GJ~182 does not show a far-infrared excess in the \\textsl{Spitzer}-MIPS study of \\citet{che05}. In \\textsl{Spitzer}-MIPS data of NGC~2547, a cluster at an age of 30--40~Myr, \\citet{you04} detected the third M~star debris disk (source 23 in our Table~\\ref{tabsigexcl}, see also Teixeira et al., \\textsl{in prep.}). TWA~7 was the fourth object reported; it is a member of the TW Hya association with an age of $<6$~Myr and excess emission reported at 24~$\\mu$m, 70~$\\mu$m, and submillimeter radio wavelengths \\citep{low05,mat07}. One field M dwarf (GJ~842.2) was found to have a debris disk \\citep[in submillimeter observations,][]{les06}. Finally, \\citet{gor07} identified one additional M~dwarf debris disk in NGC~2547 at 24~$\\mu$m, listed as source 22 in our Table~\\ref{tabsigexcl}.  Only three sources (TWA~7, source 22, and source 23) have been identified as far-infrared excess sources at 24~$\\mu$m, while the other sources have been found at either 70~$\\mu$m  (AU Mic) or submillimeter radio wavelengths. Since material of quite different distances from the central star is probed with these two techniques, debris disks identified in the mid-infrared and in submillimeter radio are not necessarily directly comparable to one another.  Studying a sample of 62 field M dwarfs estimated to be older than 1~Gyr, \\citet{gau07} did not find any indication of excess emission at 24~$\\mu$m or 70~$\\mu$m indicative of debris disks (see also \\citealp{pla05,ria06}). In contrast to these \\textsl{Spitzer} results, \\citet{les06} discuss two sets of (sub-)millimeter data, including previously published observations, and find three detections of excess emission in a sample of 23 M~dwarfs with ages ranging from 20--200~Myr, corresponding to an excess fraction of $13^{+6}_{-8}$~\\%, which may not be significantly different from the disk fractions of earlier spectral types.  After identifying an M~dwarf debris disk in earlier \\textsl{Spitzer} observations of NGC~2547 (source 23, see above), we performed ten times longer \\textsl{Spitzer}-MIPS 24~$\\mu$m observations of the same region in order to find out whether this was a singular object or part of a larger population. In the remainder of the introduction, we summarize the current knowledge about NGC~2547 before describing the observations in Section~\\ref{sec_obs}. We present the results in Section~\\ref{sec_results}, including remarks on the 70~$\\mu$m and 160~$\\mu$m data as well as variability at 24~$\\mu$m, discuss these results in Section~\\ref{sec_disc} and close with a summary in Section~\\ref{sec_summ}. ",
        "conclusions": "\\label{sec_disc}  There are three different possibilities for the interpretation of the observed excess emission \\citep[e.g.,][]{stru06}. It could be due to remnant primordial material, due to a debris disk with continuously replenished dust, or it could be debris produced by a recent catastrophic collision of planetesimals, i.e., debris not in (quasi-) equilibrium. The dust in debris disks is produced and replenished by collisions, and then removed by processes such as the Poynting-Robertson (PR) drag, radiation pressure and stellar winds. Direct radiation pressure does not play a role for the low-luminosity stars that we consider \\citep{pla05}. The balance of dust production and these removal processes determines the amount of observable dust.  %===> Which radius? In order to discuss the physical processes, we first need to determine which region around a central star is probed by the 24~$\\mu$m (excess) emission. As an estimate, we consider for a given stellar luminosity the orbital radius of a blackbody which is at a temperature that puts the peak of its SED at 24~$\\mu$m. For a dwarf star with a luminosity of $L$=0.01~$L_\\odot$, the corresponding radius is only 0.5~AU.  %===> PR drag and timescale -> debris disk With the orbital radius of the dust particles constrained, we can now determine the lifetime of dust particles under the influence of the PR drag, which should yield an upper limit of the dust removal timescale \\citep[e.g.,][]{dom03}. The PR lifetime of a dust particle in orbit around a central star scales with particle size and density as well as the square of its orbital radius and is inversely proportional to the luminosity of the central star \\citep{bur79}. For micron-sized particles with a density of $\\rho=2.5$~g\\,cm$^{-3}$ \\citep{div93,che01} located at a distance of 1~AU from a central star with a luminosity of $L=0.01\\,L_\\odot$, the lifetime is still on the order of $1.7\\times10^5$ years. This value is ten times longer for 10~$\\mu$m-sized particles.  Since the M~dwarfs discussed here have an age of 30--40~Myr, the particles would need to be replenished and thus cannot be primordial in order to still be observable at that age, indicating the debris disk nature of these systems and ruling out the first possibility mentioned in the beginning of this section. The lifetimes do not approach the cluster age until an orbital radius of 13~AU. Since other dust removal processes like stellar winds potentially operate at the same time, the derived timescale is in fact an upper limit.  This result thus suggests that at an age of 30--40~Myr, M~dwarfs can have warm inner debris disks. Given that we know already seven M dwarfs that harbor extrasolar planets orbiting close to their host stars, this finding reinforces the idea that planets form around M stars.  With M~star debris disks more frequent than previously thought, chance occurrences like catastrophic collisions of big planetesimals may not have to be invoked for their explanation, ruling out the remaining third possibility.  %===> Remarks on radius It is interesting to note that a blackbody with a peak of its SED at 24~$\\mu$m has a temperature of 120~K, not far from the temperature of water sublimation, 153~K, defining the ``snow line'' \\citep{hay81}, separating the inner region of rocky planet formation from the outer region of icy planet formation and playing a role in the definition of ``habitable zones''. \\footnote{We can estimate a permissable range of dust temperatures for the dust we detect at 24\\,$\\mu$m. Comparing the ratio of an assumed 3$\\sigma$ excess at 8\\,$\\mu$m and the observed 24\\,$\\mu$m excess with that expected for blackbody emission suggests an approximate upper limit of $\\sim260$~K for a typical source without an 8\\,$\\mu$m excess. As a strict lower limit, we note that the M-dwarf debris disk around AU~Mic which does not exhibit 24\\,$\\mu$m excess emission has a dust temperature of $\\sim40$~K \\citep{ria06}.} For M stars, the habitable zone lies at orbital radius of about 0.1~AU \\citep{sca07}. Indeed, seven M dwarfs are now known to harbor one or several extrasolar planets at orbital radii of 0.02 to 2.3~AU (most recently \\citealp{end08}, see references therein), with masses spanning 0.016--2.1~$M_J$. The five planets with the smallest semi-major axes ($<0.1$~AU) have masses of $<0.07$~$M_J$. During the early evolution of an M dwarf, the snow line moves inwards due to the decreasing luminosity of the central star, enabling the formation of icy protoplanets within the first 1~Myr at orbital radii of 1--4~AU, not taking into account subsequent migration \\citep{ken07}. The snow lines of M~dwarfs in NGC~2547 have already moved further in when compared to their initial location, so that rocky protoplanets may have previously formed at orbital radii larger than the current snow line.  %===> Comp to submm Given that some of the few previously known M dwarf debris disks have been found at submillimeter wavelengths, it is important to keep in mind that material traced in the far infrared is quite different from material traced in the submillimeter radio range. The latter corresponds to material at temperatures of $<20$~K at orbital radii of several hundred AU \\citep[e.g., ][]{les06}. Thus, different regions are probed by far-infrared and submillimeter observations and they cannot be directly compared.  %===> SEDs We show the photometric spectral energy distributions of the five M~dwarfs with the strongest excess emission at 24~$\\mu$m in Fig.~\\ref{sedplots}. Their 24~$\\mu$m emission is about an order of magnitude stronger than the expected photospheric flux. To obtain approximate physical properties from the crudely constrained photometric SEDs, we performed fits using disk models by \\citet{rob07}. For one of our best-constrained disks, source 23, the fits suggest that slightly more than 1\\% of the total luminosity is due to the disk, even though the central object is a star with a luminosity of only $L=0.028$~$L_\\odot$. To within an order of magnitude, the resulting dust disk masses are $\\sim10^{-9}$~$M_\\odot$, corresponding to only a few percent of one lunar mass, but in some cases are barely constrained. As a major source of uncertainty, these mass estimates scale with the assumed dust properties which would be much better constrained with (sub-) millimeter data. For all but the densest regions of the disk, the model assumes small dust particles slightly larger than those in the diffuse ISM, and therefore, the mass estimates are a lower limit. Reviewing the submillimeter detections of M~dwarf debris disks, \\citet{les06} quote masses of one to 13 lunar masses for three debris disks in the M0 range. In marked contrast, \\citet{che05} use \\textsl{Spitzer}-MIPS data to derive a minimum dust mass for the debris disk of AU~Mic of only 10$^{-4}$ lunar masses, assuming a particle size of 0.2~$\\mu$m. This mass difference may indicate that submillimeter and far-infrared--detected debris disks are indeed quite different. However, a comparison of both methods towards the same objects would be needed to confirm this conjecture.  %===> age effect? All of the M~dwarf members of NGC~2547 can be assumed to be of the cluster age, 30--40~Myr, possibly indicating that at this age, the evolutionary stage of M~dwarfs makes it easier to detect such debris disks. For debris disks around stars of earlier spectral types, there have already been several studies concerning age trends. In a study of main-sequence A stars older than 20~Myr, \\citet{rie05} find that the magnitude of excess at 24~$\\mu$m declines with a $t_0/t$ dependence, with $t_0\\approx 150$~Myr. Studying the double cluster h and $\\chi$ Persei in context with this result and other observations, \\citet{cur08a} conclude that the debris disk excess emission of A~stars peaks at about the age of this cluster, 13~Myr. On longer timescales, the evolution is less clear. While \\citet{tri08} did not find a clear age trend studying FIR excess emission of solar-type field stars with ages ranging from 200~Myr to 10~Gyr, \\citet{mey08}, also studying solar-type stars, found an excess rate decreasing from 18\\% at ages 3--30~Myr to 2\\% at ages 0.3--3~Gyr. \\citet{hil08} find that the 70~$\\mu$m excess emission of 328 FGK stars, tracing cooler dust compared to 24~$\\mu$m, peaks at ages of 30--100~Myr.  For a solar-mass star, the debris for rocky planet formation peaks after a few million years for material at 1~AU \\citep{ken04,ken05}. To scale this result for an M dwarf, we can use a simple relation for the disk mass, $M_{\\rm disk} \\propto M_{\\rm star}$, and the collision timescale, $t_{\\rm coll} \\propto M_{\\rm star}^{-1/2}$ \\citep{ken08}. So, if the M dwarf is three times less massive than the Sun, we would expect planet formation to take about five times longer around an M dwarf, i.e., 10--20 Myr, not taking into account minor complications due to the moving snow line. This indicates that the M dwarfs in NGC~2547 are beyond but still close to the maximum predicted debris production. Especially with regard to stars of earlier spectral types, it is important to keep in mind that this argument only allows to compare the evolutionary timescale of debris at an orbital radius of 1~AU. For A stars, the radius probed by 24~$\\mu$m observations is beyond 20~AU, and the timescales at these radii may be very different.  %===> known constraints on evolution Currently, the early debris-disk evolution of M dwarfs is poorly constrained from an observational point of view. \\citet{wei04} did not find significant mid-infrared excess (12 and 18~$\\mu$m) towards 16 members of the young, nearby TW Hydrae association and conclude that apparently any planet formation in the terrestrial planet region was rapidly completed. In a study targeting stars at age similar to the age of NGC~2547, \\citet{mam04} observed members of the 30-Myr-old Tucana-Horologium Association, including two M dwarfs, at a wavelength of $\\sim$10~$\\mu$m (\\textsl{N} band). They do not find excess emission towards the two M dwarfs, and their quoted uncertainties would allow excess emission at only $\\sim$10\\% of the expected photospheric flux level. Looking at our SEDs in Fig.~\\ref{sedplots}, if the excess emission is modeled correctly, it would be far more difficult to detect at 10$\\mu$m. The M dwarfs studied by \\citet{pla05}, and not found to exhibit excess emission at 11.7~$\\mu$m, include one young source and nine more objects with ages of $>$600~Myr. The single young star is GJ~3305, a member of the $\\beta$~Pic moving group at an age of $\\sim$12~Myr. \\citet{gau07} found that field M~dwarfs older than 1~Gyr do not show far-infrared excess emission. Due to the nearby location of these sources ($d<20$~pc), the absolute sensitivity of these observations is much higher than that of our dataset. This upper limit for the evolutionary timescale of M~dwarf debris disks indicates that planet formation around M dwarfs has ceased by this time.  %===> earlier spectral types We finally consider the excess fractions of stars with earlier spectral types. While \\citet{gor07}, restricting the discussion to their completeness limit, found 30\\ -- 45\\% of the B--F members to show excess emission at 24~$\\mu$m, we interestingly find virtually no G~star excess emission at this wavelength. As pointed out above, while we find excess emission towards K stars, we may not be complete throughout their \\textsl{R-K} range, and many of the ``highly probable'' K members may be background giants. Even so, the 24~$\\mu$m excess rate for both G and K stars, using our conservative criterion for excess sources, is at most a few percent. In the most comprehensive sample of field stars to date, \\citet{tri08} find statistically indistinguishable excess rates for A, F, G, and K stars, with  an average age of the solar-type sample of 5~Gyr. The fact that we do not find excess emission towards G~stars can possibly be understood using an argument similar to the one used in our discussion of M dwarf debris disks: given the cluster age, the solar-type stars are far beyond their peak in dust production. In marked contrast, the excess rate for M dwarfs is surprisingly high, even though it is a lower limit for two different reasons: first, our 24~$\\mu$m detections of M dwarfs are not complete and second, we use a conservative criterion to define excess sources. It remains unclear, however, to what degree the excess rates of G, K, and M stars are skewed due to contamination by non-members."
    },
    "0807/0807.2578_arXiv.txt": {
        "abstract": "{The rate at which galaxies grow via successive mergers is a key element to understand the main phases of galaxy evolution.} {We measure the evolution of the fraction of galaxies in pairs and the merging rate since redshift $z \\sim 1$ assuming a ($H_0=70 km s^{-1} Mpc^{-1}, \\Omega_{M}=0.3$ and $\\Omega_{\\Lambda}=0.7$) cosmology.} {From the VIMOS VLT Deep Survey we use a sample of 6464 galaxies with $I_{AB} \\leq 24$ to identify 314 pairs of galaxies, each member with a secure spectroscopic redshift, which are close in both projected separation and in velocity.} {We estimate that at $z\\sim 0.9$, $10.9 \\pm 3.2\\%$ of galaxies with $M_B(z) \\leq -18-Qz$ ($Q=1.11$) are in pairs with separations $\\Delta r_p \\leq 20h^{-1}\\ kpc$, $\\Delta v \\leq 500$ km/s, and with $\\Delta M_B \\leq 1.5$, significantly larger than $3.8 \\pm 1.7 \\%$ at $z \\sim 0.5$; thus, the pair fraction evolves as $(1+z)^m$ with $m = 4.73 \\pm 2.01$. For bright galaxies with $M_B(z=0) \\leq -18.77$, the pair fraction is higher and its evolution with redshift is flatter with $m=1.50 \\pm 0.76$, a property also observed for galaxies with increasing stellar masses. Early-type pairs (dry mergers) increase their relative fraction from $3\\%$ at $z \\sim 0.9$ to $12\\%$ at $z \\sim 0.5$. The star formation rate traced by the rest-frame [OII] $EW$ increases by $26 \\pm 4\\%$ for pairs with the smallest separation $r_p \\leq 20h^{-1}kpc$. Following the prescription to derive merger timescales of Kitzbichler \\& White (2008) we find that the merger rate of $M_B(z) \\leq -18-Qz$ galaxies evolves as $N_{mg}=(4.96 \\pm 2.07)\\times 10^{-4}) \\times (1+z)^{2.20 \\pm 0.77}  mergers\\ Mpc^{-3} Gyr^{-1}$.} {The merger rate of galaxies with $M_B(z) \\leq -18-Qz$ has significantly evolved since $z \\sim 1$ and is strongly dependent on the luminosity or stellar mass of galaxies. The major merger rate increases more rapidly with redshift for galaxies with fainter luminosities or stellar mass, while the evolution of the merger rate for bright or massive galaxies is slower, indicating that the slow evolution reported for the brightest galaxies is not universal. The merger rate is also strongly dependent on the spectral type of galaxies involved. Late-type mergers were more frequent in the past, while early-type mergers are more frequent today, contributing to the rise in the local density of early-type galaxies. About $20\\%$ of the stellar mass in present day galaxies with $log(M/M_{{\\odot}}) \\geq 9.5$ has been accreted through major merging events since $z = 1$. This indicates that major mergers have contributed significantly to the growth in stellar mass density of bright galaxies over the last half of the life of the Universe.} ",
        "introduction": "In the current hierarchical structure formation paradigm, the mass assembly in galaxies proceeds via a process of coalescence between increasingly more massive dark matter halos. This halo merging tree history can be quantified by a  halo merger rate, measuring the growth of mass per average mass in a representative volume of the Universe. However, these models do not directly predict a growth of galaxy mass via mergers (\\cite{moore}), and the actual contribution of mergers to the evolution of galaxies remains poorly predicted.  Merging two galaxies is potentially a very powerful process. It is possible that during major merger events, i.e. mergers where the two components have more or less the same mass, disks could be transformed into spheroidals, as predicted using detailed simulations (\\cite{combes}; \\cite{mihos}; \\cite{conselice06}). It is also expected that major merger events profoundly modify the spectrophotometric properties of the galaxies involved, for instance triggering a burst of star formation (e.g. \\cite{P05}). Galaxies in the process of merging are observed, however the contribution of this process to the evolution of the global galaxy population is not yet precisely constrained. Indirect evidence for merging is also inferred from other galaxy properties like the luminosity or mass function. The luminosity of the red bulge dominated population of galaxies is measured to increase since $z\\sim1$, part of which could be produced by mergers (\\cite{ilbert06}). it seems possible that the increase in the density of intermediate mass early-type galaxies since $z\\sim1$ may be happening at the expense of late-type galaxies involved in merging events (\\cite{tresse}). Merging is therefore potentially a very important physical phenomenon which could drive the evolution of galaxies along cosmic time. The average numbers of merger events needed to build a typical $M_*$ galaxy, the contribution of mergers to the mass growth of galaxies, or the identification of a prefered time in the life of the Universe when mergers were more frequent, are all important elements to help towards our understanding of galaxy evolution. It is then crucial to quantify the contribution of merging to the evolution process and its impact on  important quantities like the cosmic star formation rate (e.g. \\cite{bouwens}; \\cite{tresse}; \\cite{woods}) or the global stellar mass density (e.g. \\cite{Arnouts}; \\cite{bundy}; \\cite{pozzetti}).  To estimate the contribution of mergers to the formation and evolution of galaxies is not a trivial task. In the nearby Universe merger events can be identified aposteriori from perturbed morphologies, wisps, tails, and other peculiar signatures seen at low surface brightness. Only recently volume complete measurements of the merger rate in the nearby Universe are becoming available. In the Millennium catalogue, especially tailored to a volume complete identification of merging events, de Propris et al. (2007) use the relative velocity measured from spectroscopic redshifts to confirm true galaxy pairs in the process of merging. They find that the merger fraction is 2\\% at a mean redshift of 0.06, refining earlier estimates based on pair fraction (\\cite{P00}; \\cite{P02}). At higher redshifts, searching for evidence for past mergers becomes increasingly difficult, because the residual signatures of mergers often have a too low surface brightness. At redshifts $z\\geq0.3$, it is therefore easier to search for 'apriori mergers', encounters that are likely to lead to a merger event, rather than to look for 'aposteriori' signs of past mergers. When two galaxies are close together in space, and depending on their relative velocities, gravity is acting to bring them closer for a bound system that will merge. A measure of the merging frequency is then to count galaxy pairs with a separation and velocity difference such as they are likely to be gravitationally bound and destined to merge. By selecting pairs of galaxies with similar magnitudes and hence approximately with similar masses, one can focus on major merger events. They are able to significantly contribute to the mass assembly, to modify morphologies, as well as to significantly alter the star and gas content of the incoming galaxies. Assuming that a dynamically bound system of two galaxies will most likely evolve into one more massive galaxy, one can then derive the merger rate from the pair count. A major uncertainty of this estimator is the timescale upon which a merger will be completed. N-body simulations are then used to provide reasonable estimates of the merger timescales (\\cite{conselice06}; Kitzbichler \\& White, 2008).  Cold dark matter simulations show that the evolution of the dark matter halos merger rate follows a power law  $N_{mg} = N_{mg,0}(1+z)^m$ where $N_{mg,0} = N_{mg}(z=0)$ is the local value, and $m$ parameterizes the evolution. While some simulations predict that $m$ should have $2.5 \\leq m \\leq 3.5$ (Gottl\\\"\\o ber, 2001), measuring $m$ directly from galaxy samples is an important step to understand the evolution of galaxies. Many observational attempts have been carried out to track the evolution of the merger rate as a function of redshift (e.g. \\cite{burkey}; \\cite{carlberg94}; \\cite{yee}; \\cite{P97}; \\cite{lefevre00}; \\cite{P00}; \\cite{P02}; \\cite{conselice03}; \\cite{lin04}; \\cite{kartaltepe07a}; \\cite{lin08}; \\cite{lotz08}; \\cite{kampczyk}). Even though, $m$ remains poorly constrained with $0 \\leq m \\leq 6$, meaning either no evolution of the merger rate with cosmic time, or a strong evolution. Part of this large range of values can be understood as coming from the different criteria used to identify merger candidates, or the photometric band used to identify pairs (\\cite{bundy}). Furthermore, comparing measurements at low and high redshifts from different surveys is complicated due to the different selection functions used. At redshifts $z>0.3$, most pair counts so far have been performed from a measurement of the number of pairs observed on deep images, with either a photometric redshift of the galaxies (e.g. \\cite{conselice03}; \\cite{lotz08}), or a spectroscopic redshift of one of the galaxies in the pair (e.g. \\cite{P97}; \\cite{lefevre00}). The effect of contamination by galaxies projected along the line of sight producing false pairs is then estimated from galaxy counts, and the observed pair fraction is corrected to get an estimate of the true pair fraction. As redshift increases, projection effects become increasingly important making it difficult to estimate the true pair fraction, creating a fondamental uncertainty in the measurement of $m$. At $z\\sim1$ a galaxy with a luminosity $L_*$ has a 40\\% probability to have a galaxy with a similar magnitude but at a different redshift projected within an apparent radius of $20 h^{-1} kpc $ (\\cite{lefevre00}).  To overcome these limitations, the most secure method to identify a physical pair of galaxies is to obtain a velocity measurement of each galaxy in the pair, enabling to identify pairs of galaxies which are most likely to be gravitationnaly bound. Only recently samples with spectroscopic redshifts for both galaxies in a pair are becoming available (\\cite{lin07}, \\cite{lin08}). In this paper we use for the first time a complete redshift survey to $z \\sim 1$ and as faint as $I_{AB}=24$ to securely identify pairs with {\\it both} galaxies having a spectroscopic redshift. We use the VIMOS VLT Deep Survey (VVDS) (\\cite{lefevre05a}), to search for galaxy pairs and to derive the pair fraction and the merger rate evolution.  We present the galaxy sample and the methodology to build a pair sample in Section 2, we derive the pair fraction evolution in Section 3, and we examine the spectrophotometric properties of galaxies in pairs in Section 4. We compute the merger rate in Section 5. We evaluate the fraction of the stellar mass involved in mergers since $z\\sim1$ in Section 6, and conclude in Section 7. We adopt a $H_0=70 km s^{-1} Mpc^{-1}$, $\\Omega_{\\lambda}=0.7$ and $\\Omega_{m}=0.3$ cosmology throughout this work and magnitudes are given in the AB system. ",
        "conclusions": "\\label{discuss}   Our results can be summarized as follows:  (i) We find that $3.8 \\pm 1.7$, $9.4 \\pm 2.8$, and $10.9 \\pm 3.2\\ \\%$ of galaxies with $M_B(z) < -18 - Q(z)$ at $z \\sim 0.5, 0.7$ and $0.9$ respectively, are in pairs of galaxies with luminosities $\\Delta M_B \\leq 1.5$ and separations less than $20h^{-1}\\ kpc$.  (ii) The evolution of the pair fraction with redshift is strongly dependent on the absolute luminosity or stellar mass of the brighter galaxy in the pair: it evolves more slowly for brighter or more massive galaxies than for faint galaxies. Using the VVDS alone, the pair fraction of galaxies with  $M_B(z) < -18 -Q(z)$ is found to strongly evolves with redshift as $\\propto (1+z)^m$ with $m=4.46 \\pm 0.81$ for separations of $(100h^{-1}\\ kpc,500 km/s)$, while for brighter galaxies with $M_B(z) < -18.77 -Q(z)$, we find a slower evolution with $m=3.18 \\pm 1.34$. Combining VVDS data with low redshift measurements from de Propris et al. (2007), Patton et al. (2000) and Patton et al. (2002), and taking $r_p^{max}=\\ 20h^{-1}\\ kpc$, we similarly find $m=1.50 \\pm 0.76$ for bright galaxies with $M_B(z=0) \\leq -18 + 5log(h) \\sim -18.77$ and $m=4.73 \\pm 2.01$ for the fainter $M_B(z=0) \\leq -18$ sample. In addition, the evolution of the pair fraction is found to be stronger with $m = 3.13 \\pm 1.54$ for less massive galaxies with $log(M/M_{\\odot}) \\geq 9.5$, than for more massive galaxies with $log(M/M_{\\odot}) \\geq 10$ for which we find $m = 2.04 \\pm 1.65$. Low mass pairs are therefore contributing more to the evolution of the pair fraction than high mass pairs.  (iii) The star formation rate of close pairs is enhanced at separations $r_p \\leq 150h^{-1}\\ kpc$. We find that the mean $EW(OII)$ in close pairs are larger by $26 \\pm 4\\%$ than the one derived for galaxies with larger separations.  (iv) The evolution of the pair fraction is stronger for late-type pairs with $m_{late} = 4.74 \\pm 0.81$, than for early-type pairs with $m_{early} = 1.44 \\pm 0.93$. Late-type pairs are therefore contributing significantly more to the observed evolution of the pair fraction than early-type pairs in our $I_{AB} \\leq 24$ sample.  (v) Using the merging timescale from Kitzbichler \\& White (2008), we find that the merger rate increases from $\\sim 12.3 \\times 10^{-4}$ to $\\sim 19.4\\times 10^{-4}\\ $mergers$\\ h^{3}\\ Mpc^{-3}\\ Gyr^{-1}$ from $z=0.5$ to $z=0.9$. The merger rate of galaxies with $M_B(z) < -18 -Q(z)$ evolves as $N_{mg}=(4.96 \\pm 2.07) \\times 10^{-4}) \\times (1+z)^{2.20 \\pm 0.77}$. Similarly to the pair fraction, we find that the merger rate evolves faster for fainter or less massive galaxies, with $m_{mg}=2.20 \\pm 0.77$ and $2.38 \\pm 1.57$ respectively, than for brighter or more massive galaxies with $m_{mg}=1.57 \\pm 0.44$ and $1.27 \\pm 1.67$ respectively. The merger rate is evolving more strongly for late-type mergers than for early-type mergers.    We conclude that the observed evolution of the pair fraction and merger rate in our $I_{AB} \\leq 24$ sample is mostly driven by low mass late-type galaxies, while the pair fraction and merger rate of high mass early-type galaxies remains roughly constant since $z \\sim 1$. Therefore, the pair fraction or the merger rate are not universal numbers but rather are dependent on the luminosity or stellar mass, and on the spectral type of galaxies involved. Our finding that bright or massive galaxies experience a lower merger rate and a lower evolution of the merger rate extends to higher redshifts the results found in the local Universe by Patton \\& Atfield (2008). Taking into account this pair fraction and merger rate dependancy on galaxy luminosity and spectral type offers a first step to reconcile apparently inconsistent observations.Lotz et al. (2008) find a slow or no evolution of the merger rate and claim that they disagree with previous studies. When taking into account that their result is derived from bright $M_B \\leq -19.94 - 1.3 \\times z$ galaxies, their result is consistent with other studies like Conselice et al. (2003) or Le F\\`evre et al. (2000) which have been analysing fainter samples.  The dependency of the merger rate and its evolution on luminosity or stellar mass is indeed a prediction from the latest simulations using advanced semi-analytic models as described in Kitzbichler and White (2008). At the limiting magnitudes or stellar masses of our sample, Kitzbichler and White (2008) predict that the merger rate decreases and evolves more slowly for galaxy samples with increasing luminosity or stellar mass, similar to the trend observed in our sample.  The star formation rate is significantly enhanced in merging pairs with a net star formation increase of $\\sim 25\\%$ for these galaxies. Nevertheless, it accounts for only $12\\%$ to $3\\%$ of the global galaxy population from redshift $z \\sim 1$ to $z \\sim 0$ which is not sufficient to counteract the strong fading of the global star formation rate observed since $z\\sim1$. This may indicate that the gas reservoir of massive and intermediate mass galaxies has already been depleted at redshifts $z \\sim 1$, in agreement with their observed peak in star formation at $z\\sim 3.5$ (e.g. Tresse et al., 2007). It is then apparent that the decreasing SFR since $z\\sim 1$ is regulated by other physical processes like gas availability in the intergalactic medium, or feedback.  Major merging events are largely dominated by pairs of late or mixed type galaxies, but while early-type mergers represent about $15\\%$ of the merging events of bright galaxies at $z \\sim 1$, they become approximately $25\\%$ of all mergers at $z \\sim 0.5$, which is in good agreement with previous results on dry mergers (e.g. \\cite{lin08}). This indicates that major mergers are efficient in lowering the number density of intermediate mass late-type galaxies to build up more early-type galaxies. We confirm that merging is one of the important physical processes driving galaxy evolution, with the observed galaxy merger rate undoubtedly closely linked to the hierarchical build up of dark matter galaxy halos, with a rapid mass accretion phase of massive halos since $z\\sim1$ (Abbas et al., 2008). Our finding that $\\sim20\\%$ of the stellar mass in present day massive galaxies has experienced a major merger since $z \\sim 1$ is an indication that major mergers are significantly contributing to the observed evolution of the stellar mass density since $z\\sim1$ (Bundy et al., 2005; Arnouts et al., 2007)."
    },
    "0807/0807.0057_arXiv.txt": {
        "abstract": "We present three-dimensional magnetohydrodynamic (MHD) simulations of superbubbles, to study the importance of MHD effects in the interpretation of images from recent surveys of the Galactic plane. These simulations focus mainly on atmospheres defined by an exponential density distribution and the \\citet{1990ARA&A..28..215D} density distribution. In each case, the magnetic field is parallel to the Galactic plane and we investigate cases with either infinite scale height (constant magnetic field) or a constant ratio of gas pressure to magnetic pressure.  The three-dimensional structure of superbubbles in these simulations is discussed with emphasis on the axial ratio of the cavity as a function of magnetic field strength and the age of the bubble.  We investigate systematic errors in the age of the bubble and scale height of the surrounding medium that may be introduced by modeling the data with purely hydrodynamic models. Age estimates derived with symmetric hydrodynamic models fitted to an asymmetric magnetized superbubble can differ by up to a factor of four, depending on the direction of the line of sight. The scale height of the surrounding medium based on the Kompaneets model may be up to 50\\% lower than the actual scale height.  We also present the first ever predictions of Faraday rotation by a magnetized superbubble based on three-dimensional MHD simulations.  We emphasize the importance of MHD effects in the interpretation of observations of superbubbles. ",
        "introduction": "\\label{intro}  The combined stellar wind and supernova ejecta of groups of O and B stars blow large bubbles in the interstellar medium. The largest of these bubbles, with size scales of 100 pc to 1 kpc are commonly referred to as superbubbles. The basic structure of a superbubble consists of a hot low-density interior, the cavity, surrounded by a cool shell of swept-up interstellar medium. The continuous formation and dissipation of superbubbles is an important factor in the energy balance of the interstellar medium, and determines the locations of different phases of the interstellar medium on large scales \\citep{mckee1977}. Compression of the interstellar medium in the shell may increase cooling and trigger the formation of a new generation of stars.  Also, the ability of large superbubbles to break out of the disk of a galaxy and initiate an outflow of chemically enriched plasma and ionizing radiation from the disk into the halo has a profound influence on the evolution of galaxies.  Several examples of well-defined superbubbles have been identified in the Galaxy (e.g. Heiles 1984; Maciejewski et~al. 1996; Normandeau et al. 1996; Heiles 1998; Ehlerov\\'a \\& Palous 1999; Callaway et~al. 2000; Reynolds et~al. 2001; McClure-Griffiths et al. 2002, Pidopryhora et~al. 2007). These have been studied in detail, thanks to their relative proximity. Observations with parsec-scale resolution of neutral and ionized gas reveal important details about the interaction between the hot ejecta and the interstellar medium.  New high-resolution surveys of Galactic atomic hydrogen (\\HI) emission \\citep{taylor2003,mcclure2005,stil2006} have provided unprecedented images with morphological and kinematic information of Galactic superbubbles \\citep{1996Natur.380..687N,mcclure2003}.  Physically interesting parameters are usually derived from the observations by means of analytic models that assume spherical symmetry \\citep{1975ApJ...200L.107C,1977ApJ...218..377W} or axial symmetry \\citep{Kompaneets1960,1999ApJ...516..843B}.  In this paper we investigate the importance of MHD effects on physical quantities derived from observed superbubbles. The effect of the Galactic magnetic field is difficult to model because it introduces anisotropy in the medium, which requires three-dimensional numerical simulations. The first three-dimensional magnetohydrodynamic simulations of superbubbles were presented by \\citet{1998MNRAS.298..797T}, who discussed the importance of the magnetic field in the break-out of superbubbles from the Galactic disk. \\citet{1998MNRAS.298..797T} also described significant departures from spherical and axial symmetry in the shape of a magnetized bubble resulting from the interaction of the expanding superbubble with the Galactic magnetic field. The shape and the size of a superbubble depend on the strength and the geometry of the Galactic magnetic field as much as they depend on the density distribution of the ambient interstellar medium. The expanding superbubble in turn redefines the geometry of the interstellar medium and the Galactic magnetic field in a volume several hundred parsecs across.  \\citet{korpi1999} and \\citet{deavillez2005} performed thee-dimensional MHD simulations that include the evolution of a superbubble in a supernova-driven turbulent multi-phase interstellar medium. The super bubbles in these simulations show significant departures form symmetry because of inhomogeneities in the medium in which the super bubble expands. \\citet{korpi1999} and \\citet{deavillez2005} found that a super bubble can break out into the halo in such a medium. Although these simulations provide valuable insight in the dynamics of superbubbles in a multi-phase interstellar medium, it is in general difficult to relate these simulations to specific observed super bubbles (see however Fuchs et~al. 2006 for a simulation of the Local Bubble).  Astrophysical parameters derived from observations, such as the age and the energy of a superbubble, or the density distribution of the ambient medium have relied on symmetric analytic models that do not include a magnetic field.  The departure from axial symmetry imposed by the magnetic field introduces systematic errors that have not been considered before. In this paper we present three-dimensional MHD simulations of superbubbles, and we explore the errors introduced by commonly used methods to determine basic parameters from observations. In particular, we study the axial ratio of the wind-blown cavity for different times and magnetic field configurations.  The smaller simulation volume and time span used in our simulations compared to \\citet{1998MNRAS.298..797T} are more tailored to the latitude coverage of the Galactic plane surveys. We also calculate the first images of Faraday rotation by a magnetized superbubble derived from our simulations, and emphasize the importance of such simulations to make meaningful predictions in this area.  We describe frequently used analytic models \\S 2 and detail our numerical setup and methods in \\S 3.  Numerical results and the analysis of magnetic effects on derived parameters are presented \\S 4. The specific case of the W4 supperbubble is discussed in \\S 5. Faraday rotation by magnetized superbubbles is discussed in \\S 6, and conclusions are presented in \\S 7. ",
        "conclusions": "\\label{conc}   We present three-dimensional MHD simulations of a superbubble evolving in a magnetized medium. In these simulations, we assume two different atmospheric models, exponential and Dickey \\& Lockman (1990), and two different magnetic field configurations, constant magnetic field and constant $\\beta$, for varying values of the magnetic field strength. With these simulations we aim to study the importance of MHD effects on the interpretation of observed superbubbles.  As noted before by \\citet{1998MNRAS.298..797T}, a superbubble in a magnetized medium becomes significantly elongated along the magnetic field. We find that the axial ratio of the bubble in the Galactic plane is $\\sim$ 0.8 after 5 Myr and $\\sim 0.6$ after 12.5 Myr, depending on the age of the bubble and the strength of the magnetic field but not on the vertical structure of either the magnetic field or the gas. The elongated shape of the bubble in the Galactic plane may lead to significant errors in the age determination of magnetized superbubbles from observations using symmetric hydrodynamic models. The derived age depends on the location of the observer. The ratio of the smallest age to the largest age derived by observers assuming axial or spherical symmetry is found to be proportional to the square of the axial ratio of the bubble in the Galactic plane, which creates a discrepancy in the age of up to a factor $\\sim$4.  We have analyzed systematic errors in the age of the bubble and the scale height of the ambient medium introduced by fitting the Kompaneets model to a magnetized superbubble looking along the magnetic field lines. The scale height may be underestimated by 30\\% to 50\\% and the age by 50\\%.  In particular, we investigated the curiously small scale height of the interstellar medium near the W4 superbubble found by \\citet{1999ApJ...516..843B}. We re-analyzed HI data from the Canadian Galactic Plane Survey and found that the density of the ambient medium $n_{H}\\approx 2 \\rm{cm}^{-3}$ which is smaller than previously thought. This lower density helps to diminish the apparent age discrepancy between the W4 bubble and the star cluster OCL 352. However, our analysis of the systematic errors introduced by fitting the Kompaneets model shows that they are not big enough to explain the scale height of 25 pc found by \\citet{1999ApJ...516..843B} for the W4 region.  We use the MHD simulations to predict the rotation measure distribution of superbubbles based on three-dimensional MHD simulations, and emphasize the importance of such simulations to make these predictions. As expected, the appearance of a magnetized superbubble depends on the perspective of the observer. If the observer looks along the magnetic field lines, the largest rotation measures are seen at the intersection of the shell with the Galactic plane. The rotation measure is increased at larger distances from the Galactic plane. If an observer in the Galactic plane looks perpendicular to the magnetic field lines, the rotation measures are much smaller, but most importantly most of the structure in rotation measure appears in projection on the low-density cavity, and not on the shell surrounding it.  The simulations and analysis presented in this paper highlight the importance of three-dimensional MHD simulations of superbubbles evolving in the Galactic magnetic field to the interpretation of new high-resolution images of the Galactic plane at radio wavelengths from the International Galactic Plane Survey."
    },
    "0807/0807.2863_arXiv.txt": {
        "abstract": "We perform a set of high-resolution, fully self-consistent dissipationless $N$-body simulations to investigate the influence of cold dark matter (CDM) substructure on the dynamical evolution of thin galactic disks. Our method combines cosmological simulations of galaxy-sized CDM halos to derive the properties of substructure populations and controlled numerical experiments of consecutive subhalo impacts onto initially-thin, fully-formed disk galaxies. We demonstrate that close encounters between massive subhalos and galactic disks since $z \\sim 1$ should be common occurrences in {\\LCDM} models. In contrast, extremely few satellites in present-day CDM halos are likely to have a significant impact on the disk structure. One typical host halo merger history is used to seed controlled $N$-body experiments of subhalo-disk encounters. As a result of these accretion events, the disk thickens considerably at all radii with the disk scale height increasing in excess of a factor of $2$ in the solar neighborhood. We show that interactions with the subhalo population produce a wealth of distinctive morphological signatures in the disk stars, many of which resemble those being discovered in the Milky Way (MW), M31, and in other disk galaxies, including: conspicuous flares; bars; low-lived, ring-like features in the outskirts; and low-density, filamentary structures above the disk plane. We compare a resulting dynamically-cold, ring-like feature in our simulations to the Monoceros ring stellar structure in the MW. The comparison shows quantitative agreement in both spatial distribution and kinematics, suggesting that such observed complex stellar components may arise naturally as disk stars are excited by encounters with subhalos.  These findings highlight the significant role of CDM substructure in setting the structure of disk galaxies and driving galaxy evolution. ",
        "introduction": "The currently favored cold dark matter (CDM) paradigm of hierarchical structure formation (e.g., \\cite[Blumenthal et al. 1984]{Blumenthal_etal84}), predicts significant dark matter halo substructure in the form of small, dense, self-bound {\\it subhalos} orbiting within the virialized regions of larger host halos (e.g., \\cite[Klypin et al. 1999]{Klypin_etal99}). Observational probes of substructure abundance thus constitute fundamental tests of the CDM model. Due to the fact that most subhalos associated with galaxy-sized host halos lack of a significant luminous component, a constraint on the amount of substructure in these systems may be obtained via their gravitational influence on galactic disks. If there is a considerable subhalo population, it may produce strong tidal effects and induce distinctive gravitational signatures which might be imprinted on the structure and kinematics of the host galactic disk. Thus, establishing the role of substructure in shaping the fine structure of galactic disks may prove fundamental in informing our ideas about global properties of galaxy formation and evolution.  Significant theoretical effort has been devoted to quantifying the resilience of galactic disks to infalling satellites (e.g., \\cite[Quinn \\& Goodman 1986; Velazquez \\& White 1999; Font et al. 2001; Gauthier et al. 2006; Read et al. 2008; Villalobos \\& Helmi 2008]{Quinn_Goodman86,Walker_etal96,Velazquez_White99,Font_etal01, Gauthier_etal06,Read_etal08,Villalobos_Helmi08}). Despite their usefulness, most earlier investigations suffered basic shortcomings that limited their applicability. For example, some considered encounters of single satellites with galactic disks, a set-up which is at odds with CDM predictions of multiple, nearly contemporaneous accretion events. Other studies made ad hoc assumptions about the orbital parameters and internal structures of the infalling systems. Consequently, it remains uncertain whether these earlier investigations faithfully captured the responses of galactic disks to halo substructure in a cosmological context.  Here we address this issue using a hybrid approach that combines cosmological simulations to derive the merger histories of galaxy-sized CDM halos with controlled numerical experiments of consecutive subhalo impacts onto $N$-body realizations of fully-formed disk galaxies. We demonstrate that accretion histories of the kind expected in {\\LCDM} models are capable of severely perturbing galactic disks and generating a wealth of distinctive signatures in the structural and kinematic properties of disk stars. The resulting morphological features are similar to those being discovered in the Milky Way (MW), M31, and in other disk galaxies. We also show that satellite-disk interactions produce dynamically-cold, ring-like features around galactic disks that are quantitatively similar to the Monoceros ring in the MW. This suggests that such observed stellar structures may arise naturally as a result of subhalo-disk encounters, which can excite perturbations in {\\it disk stars}. These findings imply that detailed observations of galactic structure may be able to distinguish between competing cosmological models by determining whether the detailed structure of galactic disks is as excited as predicted by the CDM paradigm. ",
        "conclusions": ""
    },
    "0807/0807.3674_arXiv.txt": {
        "abstract": "High-redshift, dust-obscured galaxies -- selected to be luminous in the radio but relatively faint at 850\\,$\\mu$m -- appear to represent a different population from the ultra-luminous submillimeter- (submm-) bright population. They may be star-forming galaxies with hotter dust temperatures or they may have lower far-infrared luminosities and larger contributions from obscured active galactic nuclei (AGN).  Here we present observations of three $z\\sim2$ examples of this population, which we term {\\it submm-faint radio galaxies -- SFRGs} in CO(3--2) using the IRAM Plateau de Bure Interferometer to study their gas and dynamical properties. We estimate the molecular gas mass in each of the three SFRGs ($8.3\\times10^{9}$M$_\\odot$, $<5.6\\times10^{9}$M$_\\odot$ and $15.4\\times10^{9}$M$_\\odot$, respectively) and, in the case of \\rga, a dynamical mass by measurement of the width of the CO(3--2) line ($8\\times10^{10} \\csc^2i$\\,M$_\\odot$).  While these gas masses are substantial, on average they are 4$\\times$ lower than submm-selected galaxies (SMGs). Radio-inferred star formation rates ($<{\\rm SFR_{\\rm radio}}>=970$\\,M$_{\\odot}\\,$yr$^{-1}$) suggest much higher star-formation efficiencies than are found for SMGs, and shorter gas depletion time scales ($\\sim$11\\,Myr), much shorter than the time required to form their current stellar masses ($\\sim$160\\,Myr; $\\sim$10$^{11}$\\,M$_\\odot$). By contrast, SFRs may be overestimated by factors of a few, bringing the efficiencies in line with those typically measured for other ultraluminous star-forming galaxies and suggesting SFRGs are more like ultraviolet- (UV-)selected star-forming galaxies with enhanced radio emission.  A tentative detection of \\rga\\ at 350\\,$\\mu$m suggests hotter dust temperatures -- and thus similar gas-to-dust mass fractions -- as the SMGs. We conclude that SFRGs' radio luminosities are larger than would naturally scale from local ULIRGs given their gas masses or gas fractions. ",
        "introduction": "Submm surveys have provided an efficient probe of star-formation activity in ultraluminous infrared (IR) galaxies (ULIRGs, $>$10$^{12}$\\,L$_\\odot$) in the distant Universe (e.g., Smail, Ivison \\& Blain 1997; Hughes et al.\\ 1998; Barger et al.\\ 1998), with bright submm emission providing unambiguous evidence of massive quantities of dust, heated predominantly by young stars rather than AGN (e.g.,  Chapman et al.\\ 2003a; Alexander et al.\\ 2005; Menendez-Delmestre et al.\\ 2007; Valiante et al.\\ 2007; Pope et al.\\ 2008).  Before the availability of the Atacama Large Millimeter Array (ALMA), confusion will continue to limit the sensitivity of current submm surveys. As a result, many ULIRGs fall below the detection limits due to variations in their spectral energy distributions (SEDs) -- usually parameterized in terms of dust temperature ($T_{\\rm d}$) -- meaning that entire populations of star-forming galaxies may have been missed by submm surveys.  For a fixed far-infrared luminosity (FIR), a galaxy with a higher $T_{\\rm d}$ will be weaker in the submm at 850$\\mu$m than a galaxy with a lower $T_{\\rm d}$. Specifically, raising $T_{\\rm d}$ from the canonical $\\sim$35~K for SMGs to 45~K will result in a factor $\\sim$10$\\times$ drop in 850$\\mu$m flux density (Blain 1999; Chapman et al.\\ 2004). These galaxies should, though, be accessible in the radio waveband, regardless of their specific SEDs, since the radio correlates with the integrated FIR emission (Helou et al.\\ 1985) with a small $\\sim0.2$dex dispersion and no observable dependence on SED type. However,  there is potential for large AGN contaminations in the radio, as has often been the case with mid-IR selection of $z>1$ ULIRGs (e.g., Houck et al.\\ 2005; Yan et al.\\ 2005, 2007; Sajina et al.\\ 2007; Weedman et al.\\ 2006a, 2006b; Desai et al.\\ 2006) and the facilities required to provide the high-resolution, multi-frequency radio data needed to decontaminate the samples (e.g., Ivison et al.\\ 2007a) are not yet available.  Substantial populations of apparently star-forming galaxies at $z\\sim2$ have been uncovered through deep 1.4-GHz radio continuum observations, many of which are not detected at submm wavelengths with the current generation of instruments (Barger et al.\\ 2000; Chapman et al.\\ 2001, 2003b, 2004a -- hereafter C04).  These galaxies are luminous in the radio and spectroscopy suggests that star formation is powering their bolometric output (there is little or no sign of high-ionization emission lines, characteristic of AGN in their UV/optical spectra).  These galaxies could, in principle, span a range in properties from deeply obscured AGN to far-IR-luminous starbursts. In the latter case, one would expect a different SED from a typical SMG -- a higher $T_{\\rm d}$ for instance.  These {\\it submm-faint radio sources} have a large volume density at $z\\sim2$, even larger than the SMGs (Haarsma et al.\\ 1998; Richards et al.\\ 1999; Chapman et al.\\ 2003a, C04; Barger et al.\\ 2007). There are $\\rho = 2 \\times 10^{-5}$\\,Mpc$^{-3}$ radio sources with L$_{\\rm 1.4 GHz}>10^{31}$ ergs\\,s$^{-1}$\\,Hz$^{-1}$ at $z\\sim2$ compared with $\\rho = (6.2\\pm2.3) \\times 10^{-6}$~Mpc$^{-3}$ for SMGs brighter than 5\\,mJy at 850\\,$\\mu$m at the same epoch (Chapman et al.\\ 2003b, 2005). As essentially all of these SMGs form a subset of these radio sources (Chapman et al.\\ 2005; Pope et al.\\ 2006; cf.\\ Ivison et al.\\ 2002) this implies $\\sim14\\times10^{-6}$~Mpc$^{-3}$ luminous radio sources remain undetected at submm wavelengths. Understanding the exact properties of these galaxies is therefore of great importance. If they are all forming stars at the rates implied by their radio luminosities, they would triple the observed SFR density (SFRD) at $z\\sim2$. By contrast, if their radio luminosity comes from a mix of star formation and AGN, they have less impact on the global SFRD but they increase the highly obscured AGN fraction at these epochs (e.g., Daddi et al.\\ 2007a; Casey et al.\\ 2008) and contribute substantially to black hole growth.  Together with other observations, the redshifted cooling emission lines of CO allow us to assess and compare the energy source of SFRGs with that of SMGs and other distant star-forming galaxies via measurements of their gas and dynamical masses. In this paper, we present the results of a pilot study with the \\pdbilong\\ to detect molecular gas in SFRGs through the rotational CO(3--2) line emission. In \\S~2 we describe the sample properties and observations both with \\pdbishort\\ and other facilities. Section \\S~3 presents the CO(3--2) detections and limits obtained from the \\pdbishort\\ observations, \\S~4.1 estimates gas properties, star-formation rates and efficiencies, and \\S~4.2 compares the SFRGs to other galaxy populations. Finally \\S~5 discusses the results and places them in a broader galaxy evolution context. Throughout we assume a cosmology with $h=0.7, \\Omega_\\Lambda = 0.72, \\Omega_M = 0.28$ (e.g., Hinshaw et al.\\ 2008).   \\begin{figure} \\centering \\includegraphics[width=5.1cm,angle=0]{f1a.ps} % \\includegraphics[width=4.7cm,angle=-90]{f1c.ps} % \\includegraphics[width=5.1cm,angle=0]{f1b.ps} % \\includegraphics[width=4.7cm,angle=-90]{f1d.ps} % \\vskip0.4cm \\caption{{\\bf top panels:} CO(3--2) spectra for the two candidate detections. The spectra are shown smoothed with a 50~km/s boxcar filter, and with respect to the zero velocity offsets defined from the H$\\alpha$ emission line redshift (red dashed line). The best-fit Gaussian profile is shown for the emission line in \\rga, along with the Uv-Inferred redshift from inter-stellar absorption lines (blue dashed line). {\\bf bottom panels:} velocity-averaged spatial maps of CO emission, from $-$1500 to $-$800\\,km\\,s$^{-1}$ (\\rga) where contours are from $-$1 to 5$\\sigma$ in steps of $\\sigma$ (0.05\\,mJy\\,beam$^{-1}$). \\rgc\\ does not represent a formal CO detection. We measure a significance of 3.2$\\sigma$ integrating over the full band). The field of view, on a side, is 35\\arcsec, with the size of the beam shown to the lower left.  Both CO emitters lie exactly at the radio source position to within 1\\arcsec.} \\label{Pic} \\end{figure} ",
        "conclusions": "$\\bullet$ We conclude that the radio luminosities of these SFRGs are higher for their overall mass (gas plus stellar) than for the SMGs (given that if the radio-SFRs are over-estimated for one class, they could well be for both).  $\\bullet$ We note that SMGs in general, and also these SFRGs, are outliers of the stellar mass-SFR correlation (Daddi et al.\\ 2007a), probably due to the higher efficiency in forming stars for a similar stellar mass and CO luminosity. Lower gas masses in the SFRGs would imply even higher SFEs than the SMGs.  By contrast, if the SFRs are significantly over-estimated by the radio, or the CO-to-H$_2$ conversion were significantly different from SMGs, then SFRGs could have similar efficiencies to typical ULIRGs. Together with the apparent low-efficiency star-forming (U)LIRGs from Daddi et al.\\ (2008), the SMGs and SFRGs with SFRs several to 10$\\times$ larger than the Daddi et al.\\ galaxies suggest roughly equal numbers of stars being formed in both high- and low-efficiency modes at $z\\sim2$. Massive galaxies are being built in an impressive variety of modes in the $z=1-3$ peak star-formation period.  $\\bullet$ If the radio-inferred SFRs are correct, then these SFRGs are more efficient star formers than SMGs, and cannot obviously be interpreted as {\\em scaled up} versions of local ULIRGs as Tacconi et al.\\ (2006, 2008) have argued is the case for SMGs.  The SFRGs' radio luminosities are larger than would naturally scale from local ULIRGs given the gas masses or gas fractions. These observed gas masses and star-formation properties may be typical of the SFRG population and further work is justified to explore this population with improved statistics.  $\\bullet$ Our results underscore the fact that ultraluminous galaxies in the high-redshift Universe have been discovered with a wide range of star-forming efficiencies, the SFRGs apparently being one extreme. Massive galaxies are likely being built in a variety of modes in the $z=1-3$ peak star-formation period."
    },
    "0807/0807.1447_arXiv.txt": {
        "abstract": "{} {To investigate the chemical relations between complex organics based on their spatial distributions and excitation conditions in the low-mass young stellar objects IRAS~16293-2422 ``A'' and ``B''.} {Interferometric observations with the Submillimeter Array have been performed at 5$\\arcsec\\times3\\arcsec$\\ (800$\\times$500~AU) resolution revealing emission lines of HNCO, CH$_3$CN, CH$_2$CO, CH$_3$CHO and C$_2$H$_5$OH. Rotational temperatures are determined from rotational diagrams when a sufficient number of lines are detected.} {Compact emission is detected for all species studied here. For HNCO and CH$_3$CN it mostly arises from source ``A'', CH$_2$CO and C$_2$H$_5$OH have comparable strength for both sources and CH$_3$CHO arises exclusively from source ``B''. HNCO, CH$_3$CN and CH$_3$CHO have rotational temperatures $>$200~K implying that they arise from hot gas. The $(u,v)$-visibility data reveal that HNCO also has extended cold emission, which could not be previously determined through single dish data. } {The relative abundances of the molecules studied here are very similar within factors of a few to those found in high-mass YSOs. This illustrates that the chemistry between high- and low-mass objects appears to be relatively similar and thus independent of luminosity and cloud mass. In contrast, bigger abundance differences are seen between the ``A'' and ``B'' source. For instance, the HNCO abundance relative to CH$_3$OH is $\\sim$4 times higher toward ``A'', which may be due to a higher initial OCN$^-$ ice abundances in source ``A'' compared to ``B''. Furthermore, not all oxygen-bearing species are co-existent, with CH$_3$CHO/CH$_3$OH an order of magnitude higher toward ``B'' than ``A''. The different spatial behavior of CH$_2$CO and C$_2$H$_5$OH compared with CH$_3$CHO suggests that successive hydrogenation reactions on grain-surfaces are not sufficient to explain the observed gas phase abundance of the latter. Selective destruction of CH$_3$CHO may result in the anti-coincidence of these species in source ``A''. These results illustrate the power of interferometric compared with single dish data in terms of testing chemical models.} ",
        "introduction": "The envelopes of some low-mass protostars contain many complex organic molecules\\footnote{In this paper molecules are considered complex if they contain more than four atoms.} \\citep{blake1994,dishoeck1995,cazaux2003,bottinelli2004b,bottinelli2004,bottinelli2007,jorgensen2005b,jorgensen2005a,sakai2006,sakai2007}. This raises the question whether these are low-mass versions of ``hot cores'', chemically very rich environments in high-mass star forming regions that are thought to have their origin in grain-mantle evaporation and the subsequent rapid gas phase reactions. The low-mass counterpart is sometimes called a ``hot corino''. The presence of warm material has long been suggested from modeling of the SEDs of these sources \\citep[e.g.,][]{adams1987,jorgensen2002,shirley2002} and has been firmly established by molecular excitation studies \\citep{blake1994,dishoeck1995,ceccarelli2000} and by resolved interferometric imaging \\citep[see e.g.,][]{chandler2005,bottinelli2004,jorgensen2005}. However, in some sources the emission peaks offset from the continuum source \\citep[e.g.,][]{chandler2005}. This offset is in disagreement with the ``hot corino'' hypothesis for low-mass stars in which complex molecules evaporate through passive heating and is more in favor of other explanations such as the presence of disks or outflows that create shocks in the envelope. Currently, there is an ongoing debate on whether the emission of complex organics comes from ``hot corinos'' or from other types of regions. Also, the extent to which the observed organics are first generation molecules created in the ice or second generation produced in the gas is still an open question. The aim of this work is to map the emission of complex organics to address the latter question, namely to determine their most likely formation mechanism.  One method to study chemical links between species is to study molecular abundances through single-dish surveys in a large number of sources \\citep{vdtak2000b,vdtak2003,ikeda2001,bisschop2007a}. An alternative method is to look for spatial correlations by interferometric observations of a single source through which it is possible to distinguish compact and extended emission as well as the exact location of the compact emission. The species that are studied here are the nitrogen-bearing species, HNCO and CH$_3$CN and the oxygen-bearing species CH$_2$CO, CH$_3$CHO, and C$_2$H$_5$OH that are commonly found toward or surrounding hot cores. The molecules CH$_2$CO, CH$_3$CHO and C$_2$H$_5$OH are proposed to be linked by successive hydrogenation on the surfaces of grains \\citep{tielens1997}.  The source studied in this paper, IRAS~16293-2422, is a very well studied and chemically rich low-mass YSO. It is a binary with its main components, named ``A'' and ``B'', separated by 5\\arcsec \\citep[corresponding to 800~AU at 160~pc][]{mundy1992}. From single dish observations and modeling, \\citet{dishoeck1995}, \\citet{ceccarelli1999,ceccarelli2000} and \\citet{schoier2002} concluded that emission from several organic molecules arises from a compact region. The data for some species such as H$_2$CO and CH$_3$CN are significantly better fitted if a ``jump'' in the abundance at 80--90~K due to grain-mantle evaporation is assumed in a spherical circum-binary envelope \\citep{ceccarelli2000,schoier2002}. Through interferometric observations it is possible to distinguish between the emission from the cold extended envelope and more compact emission as well as the peak location: is the emission coming from source ``A'' or ``B''?  Previous observations have shown that some complex species are much more abundant toward one source than the other, such as CH$_3$OCHO which is more prominent toward the ``A'' source \\citep[][]{bottinelli2004,huang2005,kuan2004,remijan2006}.  The nitrogen-bearing species HNCO and CH$_3$CN have previously been detected toward IRAS~16293-2422 through single dish observations \\citep{dishoeck1995,cazaux2003}. They often have high rotational temperatures in star forming regions implying that they are present in hot gas \\citep{olmi1993,zinchenko2000,cazaux2003,bisschop2007a}. This is also the case for the oxygen-bearing molecule C$_2$H$_5$OH \\citep{ikeda2001,bisschop2007a}. In contrast, CH$_2$CO and CH$_3$CHO are often detected in high-mass sources with low rotational temperatures \\citep{ikeda2002,bisschop2007a}. For CH$_3$CHO higher rotational temperatures are also found, but these can be due to the {\\it b}-type transitions that are radiatively pumped and do not represent the actual kinetic temperature of the gas \\citep{nummelin2000,turner1991}. Recent laboratory experiments by \\citet{bisschop2007c} have shown that it is possible to explain the gas phase abundances of C$_2$H$_5$OH if a solid state formation route through CH$_3$CHO hydrogenation is assumed. The observed absence of CH$_3$CHO in hot gas combined with its detection in cold ices \\citep{keane2001,gibb2004}, however, implies that it must be destroyed at higher ice temperatures, before evaporation commences, or directly after evaporation in the gas phase. A comparison of the excitation properties as well as interferometric observations in which the spatial distribution of the species can be determined are good tools to further elucidate the chemical relations between these species.  This paper is structured as follows: \\S~\\ref{obs} presents the Submillimeter Array (SMA) observing strategy as well as maps and spectra; \\S~\\ref{an} explains the analysis of the interferometric observations in the $(u,v)$-plane and the fitting of rotational diagrams; \\S~\\ref{results} presents the results of the rotational diagram and flux analysis; \\S~\\ref{disc} discusses the relative abundances of the different species with respect to each other as well as the implications for the chemistry; \\S~\\ref{sum} summarizes the main conclusions. ",
        "conclusions": "\\label{sum}  We have performed an interferometric study of the complex organic species HNCO, CH$_3$CN, CH$_2$CO, CH$_3$CHO and C$_2$H$_5$OH toward the low-mass protostar IRAS~16293-2422 with the SMA. Previously published data from \\citet{kuan2004} are used to determine abundances relative to CH$_3$OH. The main conclusions of this work are:  \\begin{itemize} \\item The emission from both HNCO and CH$_3$CN is compact and is seen toward both sources in the binary, with only 10--20\\% arising from source ``B''. Additionally, the lowest excitation line of HNCO shows extended emission suggestive of its presence in a cold extended envelope. The relatively higher abundances with respect to CH$_3$OH in source ``A'' may originate from higher initial abundances of OCN$^-$ in the ice. \\item For CH$_2$CO and C$_2$H$_5$OH only one line is detected due to compact emission. For CH$_2$CO these lines are detected with similar strength toward both sources and C$_2$H$_5$OH is clearly detected in source ``B'', but due to line-blending it is difficult to determine the flux for source ``A''. Compact hot emission for CH$_3$CHO is detected only toward source ``B''. Comparison with previous single dish observations by \\citet{cazaux2003} suggests that a cold extended component is present as well. If CH$_2$CO, CH$_3$CHO and C$_2$H$_5$OH are related through successive hydrogenation on the surfaces of grains, the same spatial behavior is expected for all three species. Since this is not observed, it suggests that hydrogenation reactions on grain surfaces alone cannot account for the observed gas phase abundance ratios. The difference between the two IRAS~16292-2422 sources can be explained if CH$_3$CHO would be selectively destroyed in source ``A'' right before or after grain-mantle evaporation. \\end{itemize} The discussion in this paper demonstrates the strength of clearly resolved interferometric observations for studies of the chemistry in star forming regions. First, molecules that have both extended and compact emission can easily be identified based on $(u,v)$-visibility analysis. In single dish observations these components will be averaged together giving rotational temperatures in between that of the hot and cold component. Additionally, compact emission arising from different components in the same single-dish beam can be separated, which leads in some cases to abundance ratios that can vary over an order of magnitude for different sources as exemplified in Table~\\ref{comp} for HNCO and CH$_3$CHO. For HNCO the presence of a cold extended component could not be derived from the single dish observations, whereas for CH$_3$CHO it is the hot component that is not detected. Further interferometric studies are needed to elucidate the chemical relations between complex organics."
    },
    "0807/0807.3442_arXiv.txt": {
        "abstract": "We estimate the local number density of sources of ultra-high-energy cosmic rays (UHECRs) based on the statistical features of their arrival direction distribution. We calculate the arrival distributions of protons above $10^{19}$ eV taking into account their propagation process in the Galactic magnetic field and a structured intergalactic magnetic field, and statistically compare those with the observational result of the Pierre Auger Observatory. The anisotropy in the arrival distribution at the highest energies enables us to estimate the number density of UHECR sources as $\\sim 10^{-4}~{\\rm Mpc}^{-3}$ assuming the persistent activity of UHECR sources. We compare the estimated number density of UHECR sources with the number densities of known astrophysical objects. This estimated number density is consistent with the number density of Fanaroff-Reily I galaxies. We also discuss the reproducability of the observed {\\it isotropy} in the arrival distribution above $10^{19}$ eV. We find that the estimated source model cannot reproduce the observed isotropy. However, the observed isotropy can be reproduced with the number density of $10^{-2}$-$10^{-3}~{\\rm Mpc}^{-3}$. This fact indicates the existence of UHECR sources with a maximum acceleration energy of $\\sim 10^{19}$ eV whose number density is an order of magnitude more than that injecting the highest energy cosmic rays. ",
        "introduction": "\\label{introduction}  The origin of ultra-high-energy cosmic rays (UHECRs) has been highly unknown in spite of prolonged effort to construct larger UHECR observatories and to detect more events \\cite{nagano00}. In cosmic ray spectrum, a sharply spectral steepening at around $10^{20}$eV has been predicted theoretically by interactions with photopion production with cosmic microwave background (CMB) photons, known as Greisen-Zatsepin-Kuz'min (GZK) steepening \\cite{greisen66,zatsepin66}. The feature of this steepening observed by the High Resolution Fly's Eye (HiRes) \\cite{abbasi04a,abbasi08} and the Pierre Auger Observatory (PAO) \\cite{yamamoto07,abraham08} implies that astrophysical sources are much more dominant than top-down sources (for a review, see Ref. \\cite{bhattacharjee00}) at the highest energies, though physical reasons for the extension of the energy spectrum beyond the GZK energies reported by the Akeno Giant Air Shower Array (AGASA) have been not understood yet \\cite{takeda98,takeda03}. Several objects to accelerate particles up to $10^{20}$ eV have been suggested, but there has been little observational evidence which is a UHECR source (see Ref. \\cite{bhattacharjee00,torres04} and reference therein).   The arrival distribution of UHECRs has information on the distribution of their sources, also including that on intergalactic and the Galactic magnetic field (IGMF and GMF). The AGASA and PAO reported statistically significant anisotropy at small angular scale in the arrival distribution \\cite{takeda99,pao07,pao08}. The small-scale anisotropy implies point-like sources, and enables us to constrain the number of nearby UHECR sources. The source number density, $n_s$, is one of important parameters to investigate the nature of UHECR sources. Comparing the estimated number density with the number densities of known astrophysical objects, we can constrain the object-type of UHECR sources. Several authors have estimated as $n_s \\sim 10^{-5}$-$10^{-6}~{\\rm Mpc}^{-3}$ using the published AGASA data above $4 \\times 10^{19}$ eV assuming the same injection rate over all sources \\cite{yoshiguchi03,blasi04,kachelriess05,takami06,takami07}, and we also constrained the source number density as $10^{-4}~{\\rm Mpc}^{-3}$, assuming a model that the injection rate is proportional to the luminosity of galaxies \\cite{takami07}.   The PAO reported the correlation between the arrival directions of UHECR above $5.7 \\times 10^{19}$ eV and the positions of nearby active galactic nuclei (AGNs) listed in the 12th Veron-Cetty \\& Veron catalog \\cite{veron06} within the angular scale of $3.1^{\\circ}$ \\cite{pao07,pao08}. However, most of the PAO-correlated AGNs are Seyfert galaxies and LINERs, which have much weaker activity than radio-loud AGNs like Fanaroff-Reily II (FR II) galaxies \\cite{moskalenko08}. To understand whether these galaxies with weak activity are really UHECR sources or not, another information on UHECR sources is required.   The PAO estimated a lower limit of the UHECR source number as 61 based on anisotropy in the arrival distribution of their detected events, simply assuming Poisson statistics and not taking into account UHECR propagation \\cite{pao08}. This is certainly a lower limit of the number of the sources, but is not quantity which can be compared with astronomical observables because we do not know what radius they are included in. Their number density is an observable. Taking into account UHECR propagation, we can also estimate plausible value of it, not {\\it limit}.   In this study, we simulate the arrival distribution of UHECRs above $10^{19}$ eV, taking into account their propagation process in intergalactic and the Galactic space. We extract an anisotropy signal from the simulated arrival distribution and then compare this signal with that observed by the PAO above $5.7 \\times 10^{19}$ eV to estimate the source number density. The reports for the correlation by the PAO with nearby large-scale structure \\cite{pao07,pao08,kashti08,takami08c} show that it plays a crucial role on the arrival distribution of UHECRs. Thus, we adopt the models of the IGMF and source distribution which reproduce the local structure actually observed around the Milky Way, developed in our previous work \\cite{takami06}. We also discuss the isotropy in the arrival distribution at $\\sim 10^{19}$ eV and estimate the source number density for the lower energies. The composition is assumed to be pure protons.   This paper is structured as follows. In section \\ref{methods}, our calculation method for calculating UHECR arrival distribution and a statistical method to estimate the source number density are explained. In section \\ref{results}, the calculation results are shown and we discuss the results and conclude in section \\ref{discussion}. ",
        "conclusions": "\\label{discussion}  In this study, we estimated the number density of UHECR sources based on the statistical features of the arrival distribution of UHECRs observed by the PAO. We simulated the arrival distributions of protons above $10^{19}$ eV, taking into consideration their propagation process in the Galactic and intergalactic space. The IGMF model adopted was associated with the observed large-scale structure. Comparing the simulated arrival distributions with the data observed by the PAO statistically, we estimated the number density which can best reproduce the observational result. The anisotropy signal above $5.7 \\times 10^{19}$ eV led to $n_s \\sim 10^{-4}~{\\rm Mpc}^{-3}$ which is consistent with that of FR I galaxies, which are a candidate to accelerate protons up to $10^{20}$ eV. We also focused on isotropy in the arrival distribution at around $10^{19}$ eV, and then found $n_s \\sim 10^{-2}$-$10^{-3}~{\\rm Mpc}^{-3}$, which is one or two order of magnitude larger than that estimated from the anisotropy.   In this calculation, cosmological evolution of UHECR sources is not considered. Since protons with $\\sim 10^{19}$ eV can reach the earth from sources at the distance of 1Gpc ($z \\sim 0.25$), which is comparable with the energy-loss length of such protons by Bethe-Heitler pair creation, it might be possible that the cosmological evolution could reproduce the isotropy even if the local number density of UHECR sources is comparable with $\\sim 10^{-4}~{\\rm Mpc}^{-3}$. However, a typical evolution factor is $(1+z)^3$, and therefore the number density at $z \\sim 0.25$ is only twice more than the local one. Such a small factor could not change the estimated number density by two order of magnitude less. The difference between the two estimated number densities is significant.   The difference can be interpreted as the evidence of the existence of UHECR sources which can accelerate protons up to $\\sim 10^{19}$ eV assuming the persistent activity of UHECR sources. $n_s \\sim 10^{-2}$-$10^{-3}~{\\rm Mpc}^{-3}$ is comparable with the number densities of bright galaxies or Seyfert galaxies. The proton acceleration up to $10^{19}$ eV might be common in the universe. We could also interpret this difference as transient activity to emit UHECRs. If bursting sources, like GRBs, are assumed, apparent number density, which corresponds to $n_s$ estimated in this study, depends on the threshold of UHECR energies because the dispersion of the time-delay is larger at lower energies \\cite{miralda96}. However, quantitative discussion on this possibility exceeds the scope of this study. It is next study of ours.   In order to reproduce isotropy in the arrival distribution of UHECRs above $10^{19}$ eV, sources within 5 Mpc were artificially neglected. In these sources, Cen A, the nearest radio-loud AGN, is involved. The isotropy at around $10^{19}$ eV implies that it is not UHECR sources in our persistent source model. However, radio-loud AGNs have been plausible site for particle acceleration up to the highest energy \\cite{torres04}, and the two of the highest energy events of the PAO are correlated with the position of Cen A within $\\sim 3^{\\circ}$. Whether Cen A is really a UHECR source or not is a key to understand the mechanism to generate the highest energy cosmic rays.   If Cen A is not a UHECR source, we can interpret that the source of cosmic rays arriving in the direction of Centaurus is behind Cen A and more distant, so that the strong anisotropy is not generated. The 2 events towards Cen A are also positionally correlated with NGC 5090 with $\\sim 3^{\\circ}$, a radio galaxy with the distance of $\\sim 40$ Mpc. This might be a real source though the activity is weak like the other PAO-correlated AGNs. These events are also towards Centaurus cluster whose distance is about 40 Mpc. An idea to generate UHE particles is cluster accretion shock \\cite{kang96,inoue05,inoue07}. However, this scenario can accelerate protons up to $\\sim 10^{19}$ eV.   If Cen A is a UHECR source, we propose several possibilities not to generate the anisotropy. One is UHECR composition. In the case of heavy dominated composition at energies above $10^{19}$ eV, the trajectories of UHECRs are deflected $Z$ times more than those of protons, and the discussion on the isotropy in this paper is not applied. Several composition measurements imply the existence of some fraction of heavy elements for all large uncertainty on hadronic interaction models in extensive air shower \\cite{unger07,glushkov07}. The other is that the UHECR production is transient. For a bursting source, the energy of cosmic rays observed at present is in narrow energy range since the time-delay of cosmic rays depends on their energies \\cite{miralda96}. In this brief picture, the 2 events are arrived from Cen A, and the lower energies will come in the future. The detailed discussion on the two possibilities is our near future plan.   The estimated number densities have some uncertainty because of the small number of detected events. The uncertainty to determine the number density at the highest energies can be reduced by increasing observed events \\cite{takami07} and significant anisotropy would be observed in the arrival distribution at $\\sim 10^{19}$ eV in the near future \\cite{waxman97,kashti08}. The dramatically increase of detected events in the near future will provide us more useful information on UHECR sources.   \\subsubsection*{Acknowledgements:} The work of H.T. is supported by Grants-in-Aid from JSPS Fellows. The work of K.S. is supported by Grants-in-Aid for Scientific Research provided by the Ministry of Education, Science and Culture of Japan through Research Grants S19104006."
    },
    "0807/0807.1196.txt": {
        "abstract": "%context heading (optional) {The Cepheid period-luminosity (PL) relation is unquestionably one of the most powerful tools at our disposal for determining the extragalactic distance scale. While significant progress has been made in the past few years towards its understanding and characterization both on the observational and theoretical sides, the debate on the influence that chemical composition may have on the PL relation is still unsettled.} %aims heading (mandatory) {With the aim to assess the influence of the stellar iron content on the PL relation in the $V$ and $K$ bands, we have related the $V$-band and the $K$-band residuals from the standard PL relations of Freedman et al. (2001) and Persson et al. (2004), respectively, to [Fe/H].} %method {We used direct measurements of the iron abundances of 68 Galactic and Magellanic Cepheids from FEROS and UVES high-resolution and high signal-to-noise spectra.} %results heading (mandatory) {We find a mean iron abundance ([Fe/H]) about solar ($\\sigma$ = 0.10) for our Galactic sample (32 stars), $\\sim$ -0.33 dex ($\\sigma$ = 0.13) for the Large Magellanic Cloud (LMC) sample (22 stars) and $\\sim$ -0.75 dex ($\\sigma$ = 0.08) for the Small Magellanic Cloud (SMC) sample (14 stars). Our abundance measurements of the Magellanic Cepheids double the number of stars studied up to now at high resolution. The metallicity affects the $V$-band Cepheid PL relation and metal-rich Cepheids appear to be systematically fainter than metal-poor ones. These findings depend neither on the adopted distance scale for Galactic Cepheids nor on the adopted LMC distance modulus. Current data do not allow us to reach a firm conclusion concerning the metallicity dependence of the $K$-band PL relation. The new Galactic distances indicate a small effect, whereas the old ones support a marginal effect. } %Conclusions {Recent robust estimates of the LMC distance and current results indicate that the Cepheid PL relation is not Universal. } { ",
        "introduction": "Since the dawn of modern astronomy the Cepheid Period-Luminosity (PL) relation is a key tool in determining Galactic and extragalactic distances. In spite of its fundamental importance, the debate on the role played by the chemical composition on the pulsation properties of Cepheids is still open, with different theoretical models and observational results leading to markedly different conclusions.  From the theoretical point of view pulsation models by different groups lead to substantially different results. Linear models (e.g. Chiosi et al. 1992; Sandage et al. 1999; Baraffe \\& Alibert 2001), based on nonadiabatic radiative models, suggest a moderate dependence of the PL relation on the metallicity. The predicted change at log(P) = 1 is less than 0.1 mag for metal abundances ranging from the SMC (Z=0.004) to the Galaxy (Z =0.02), independent of wavelength. Nonlinear convective models (e.g. Bono et al. 1999; Caputo et al. 2000; Caputo 2008) instead predict a larger dependence on the same interval of metallicity: the change is 0.4 mag in $V$, 0.3 mag in $I$ and 0.2 mag in $K$, again at log(P) = 1. Moreover, the predicted change in these latter models is such that metal-rich Cepheids are fainter than metal-poor ones, at variance with the results of the linear models. Fiorentino et al. (2002) and, more recently, Marconi, Musella \\& Fiorentino (2005) investigations, also based on nonlinear models, suggest that there may be also a dependence on the helium abundance.  On the observational side, the majority of the constraints comes from indirect measurements of the metallicity, mostly in external galaxies, such as oxygen nebular abundances derived from spectra of H II regions at the same Galactocentric distance as the Cepheid fields (e.g. Sasselov et al. 1997; Kennicutt et al. 1998; Sakai et al. 2004). These analyses indicate that metal-rich Cepheids are brighter than metal-poor ones (hence at variance with the predictions of nonlinear convective models), but it is important to note that the results span a disappointingly large range of values (see Table~1 and Fig.~1).  More recently Macri et al. (2006) found, by adopting a large sample of Cepheids in two different fields of NGC~4258 and the [O/H] gradient based on H II regions provided by Zaritsky et al. (1994), a metallicity effect of $\\gamma=-0.29\\pm0.09$ mag/dex. This galaxy has been adopted as a benchmark for estimating the metallicity effect, since an accurate geometrical distance based on water maser emission is available (Herrnstein et al. 2005). However, in a thorough investigation Tammann et al. (2007) suggested that the flat slope of the Period-Color relation of the Cepheids located in the inner metal-rich field could be due to a second parameter, likely helium, other than the metal abundance. Furthermore, Bono et al. (2008) found, using the new and more accurate nebular oxygen abundances for a good sample of H II region in NGC~4258 provided by D\\`iaz et al. (2000), a shallower abundance gradient. In particular, the new estimates seem to suggest that both the inner and the outer field might have a mean oxygen abundance very similar to LMC. They also found a very good agreement between predicted and observed Period-Wesenheit ($V,I$) relation. Nonlinear convective models predict for this relation a metallicity effect of $\\gamma=+0.05\\pm0.03$ mag/dex.  On the basis of independent distances for 18 galaxies based on Cepheid and on the Tip of the Red Giant Branch, Tammann et al. (2007) found a small metallicity effect ($\\gamma=-0.017\\pm0.113$ mag/dex). On the other hand, Fouqu\\'e et al. (2007) using a sample of 59 Galactic Cepheids whose distances were estimated using different methods -- HST trigonometric parallaxes (Benedict et al. 2007), revised Hipparcos parallaxes (van Leeuwen et al. 2007), infrared surface brightness method (Fouqu\\'e \\& Gieren 1997), and interferometric Baade-Wesselink method (Kervella et al. 2004), zero-age-main-sequence fitting of open clusters (Turner \\& Burke 2002) -- found no significant difference between optical and Near-Infrared (NIR) slopes of Galactic and LMC Cepheids (Udalski et al. 1999; Persson et al. 2004).   \\begin{figure}[h] \\centering \\includegraphics[height=0.35\\textheight, width=0.40\\textwidth]{5661f1.ps} \\vspace*{0.5truecm} \\caption{ Comparison of recent results for the metallicity sensitivity of Cepheid distances. FM1990: Freedman \\& Madore (1990); G1994: Gould (1994); Ko1997: Kochanek (1997); S1997: Sasselov et al (1997); Ke1998: Kennicutt et al (1998); F2001: Freedman et al (2001); U2001: Udalski et al (2001); C2002: Ciardullo et al (2002);  Sa2004: Sakai et al (2004); Sto2004: Storm et al (2004); Gro2004: Groenewegen et al (2004); M2006: Macri et al. (2006); Sah2006: Saha et al. (2006); Be2007: Benedict et al. (2007); Fo2007: Fouqu\\'e et al. (2007); Tam2007: Tammann et al. (2007); B2008: Bono et al. (2008). See Table~1. } \\label{Fig1} \\end{figure}  An alternative approach is to measure directly the metal content of Cepheid stars, which, so far, has been attempted only by few studies, primarily focused on stars of our own Galaxy (Luck \\& Lambert, 1992; Fry \\& Carney 1997; Luck et al. 1998; Andrievsky et al. 2002a,b,c; Luck et al. 2003; Andrievsky et al. 2004). Fry \\& Carney (1997, hereafter FC97), for instance, have derived iron and $\\alpha$-element abundances for 23 Galactic Cepheids from high resolution and high signal-to-noise spectra. They found a spread in [Fe/H] of about 0.4 dex, which they claim is real. Using approximately half of their sample, the stars belonging to clusters or associations, they have made a preliminary evaluation of metallicity effects on the zero point of the PL relation, finding that metal-rich Cepheids are brighter than metal-poor ones. Thus, finding a result similar to the studies based on indirect measurements of the metallicity.  The impressive observational effort carried out by Andrievsky and collaborators (Andrievsky et al. 2002a, 2002b, 2002c; Luck et al. 2003; Andrievsky et al. 2004; Kovtyukhet al. 2005b) has, instead, taken advantage of high resolution spectra of 130 Galactic Cepheids (collected with different instruments at different telescopes) in order to determine their chemical composition and study the Galactic abundance gradient. The sample covers a range of Galactocentric distances from 4 to 14 kpc. The emerging picture can be best described by a relatively steep gradient (about -0.14 dex kpc$^{-1}$) for Galactocentric distances less than 7 kpc, followed by a much shallower slope ($\\approx$ -0.03 dex kpc$^{-1}$) between 7 and 10 kpc, a discontinuity at approximately 10 kpc and a nearly constant metallicity of about -0.2 dex towards larger Galactocentric distances, out to about 14 kpc. In relation to our work, it is important to note that Andrievsky and collaborators did not investigate the effects of the chemical composition on the Cepheid PL relation: on the contrary, they used the PL relation to determine the distances of their stars.  Outside our Galaxy, Luck \\& Lambert (1992, hereafter LL92) have studied 10 Cepheids in the Magellanic Clouds (MCs). Five are in the Large MC (LMC) and five in the Small MC (SMC). For the former sample, they found a mean [Fe/H] of -0.36 dex with a dispersion of 0.3 dex, while for the latter one the mean [Fe/H] is -0.60 dex with a rather small dispersion of less than 0.15 dex. A more recent study by Luck et al. (1998, hereafter L98) on 10 LMC Cepheids and 6 SMC Cepheids, 4 of which in common with LL92, confirmed the mean [Fe/H] value in the LMC (-0.30 dex), found very little evidence of a significant metallicity dispersion in the LMC (contrary to LL92, but similarly to the SMC), and slightly revised downwards the mean [Fe/H] of the SMC (-0.74 vs -0.60 found by LL92).  \\begin{table*}[!ht] %\\begin{sidewaystable*} %\\begin{minipage}[t][180mm]{\\textwidth} \\label{table:1} \\caption{Overview of recent results for the metallicity sensitivity of Cepheid distances. In the first column is listed the variation of the distance modulus $\\mu$ per dex of metallicity, the negative sign indicates that the true distance is longer than the one obtained neglecting the effect of the metallicity. In the second column is listed the elemental abundance used as reference for the metallicity. The third and fourth columns give the method and the reference of the different studies. See also Fig.~1.} %\\centering \\begin{tabular}{ l c p{9cm} l} \\hline $\\delta \\mu / \\delta [M/H]$ & ~ & Method & Reference \\\\ (mag/dex) & ~ & ~ \\\\ \\hline\\hline -0.32 $\\pm$ 0.21 & [Fe/H] & Analysis of Cepheids in 3 fields of M31 ($BVRI$ bands) & Freedman \\& Madore (1990) \\\\ -0.88 $\\pm$ 0.16 & [Fe/H] & Comparison of Cepheids from 3 fields of M31 and LMC ($BVRI$ bands) & Gould (1994) \\\\ -0.40 $\\pm$ 0.20 & [O/H] & Simultaneous solution for distances to 17 galaxies ($UBVRIJHK$ bands) & Kochanek (1997) \\\\ -0.44$_{-0.20}^{+0.10}$ & [O/H] & Comparison of EROS observations of SMC and LMC Cepheids ($VR$ bands) & Sasselov et al. (1997) \\\\ -0.24 $\\pm$ 0.16 & [O/H] & Comparison of HST observations of inner and outer fields of M101 & Kennicutt et al. (1998) \\\\ -0.12 $\\pm$ 0.08 & [O/H] & Comparison of 10  Cepheid galaxies with Tip of the Red Giant Branch distances & Kennicutt et al. (1998) \\\\ -0.20 $\\pm$ 0.20 & [O/H] & Value adopted for the HST Key Project final result & Freedman et al. (2001) \\\\ 0 & [Fe/H] & OGLE result comparing Cepheids in IC1613 and MC ($VI$ bands) & Udalski et al. (2001) \\\\ 0 & [O/H] & Comparison of Planetary Nebula luminosity function distance scale and Surface Brightness fluctuation distance scale & Ciardullo et al. (2002) \\\\ -0.24 $\\pm$ 0.05 & [O/H] &  Comparison of 17 Cepheid galaxies with Tip of the Red Giant Branch distances & Sakai et al. (2004) \\\\ -0.21 $\\pm$ 0.19 & [Fe/H] &  Baade-Wesselink analysis of Galactic and SMC Cepheids ($VK$ bands) & Storm et al. (2004) \\\\ -0.23 $\\pm$ 0.19 & [Fe/H] &  Baade-Wesselink analysis of Galactic and SMC Cepheids ($I$ band) &  Storm et al. (2004) \\\\ -0.29 $\\pm$ 0.19 & [Fe/H] &  Baade-Wesselink analysis of Galactic and SMC Cepheids ($W$ index) &  Storm et al. (2004) \\\\ -0.27 $\\pm$ 0.08 & [Fe/H] &  Compilation from the literature of distances and metallicities of 53 Galactic and MC Cepheids ($VIWK$ bands) & Groenewegen et al. (2004) \\\\ -0.39 $\\pm$ 0.03 & [Fe/H] &  Cepheid distances to SNe Ia host galaxies & Saha et al. (2006)\\\\ -0.29 $\\pm$ 0.09 & [O/H]  & Cepheids in NGC~4258 and [O/H] gradient from Zaritsky et al. (1994)  & Macri et al. (2006)\\\\ -0.10 $\\pm$ 0.03 & [Fe/H] & Weighted mean of Kennicutt, Macri and Groenewegen estimates & Benedict et al. (2007)\\\\ -0.017 $\\pm$ 0.113 & [O/H] & Comparison between Cepheid and TRGB distances for 18 galaxies & Tammann et al. (2007)\\\\ 0  & [Fe/H] & Comparison between the slopes of Galactic and LMC Cepheids  & Fouqu\\'e et al. (2007)\\\\ +0.05 $\\pm$ 0.03  & [Fe/H] & Predicted Period-Wesenheit ($V,I$) relation & Bono et al. (2008)\\\\ \\hline \\hline \\end{tabular} %\\vfill %\\end{minipage} %\\end{sidewaystable*} \\end{table*}  Finally, we mention two studies that followed slightly different approaches. Groenewegen et al. (2004) have selected from the literature a sample of 37 Galactic, 10 LMC and 6 SMC Cepheids for which individual metallicity estimates and $BVIK$ photometry were known. Their work aimed at investigating the metallicity dependence of the PL relation using individual metallicity determinations as well as good individual distance estimates for Galactic Cepheids. They inferred a metallicity effect of about -0.27 $\\pm$ 0.08 mag/dex in the zero point in $VIWK$, in the sense that metal-rich Cepheids are brighter than the metal-poor ones (see Table~1 and Fig.~1, for a comparison with other studies). Also Storm et al. (2004) discussed the effect of the metallicity on the PL relation using 34 Galactic and 5 SMC Cepheids, for which they determined accurate individual distances with the Baade-Wesselink method. Assuming an average abundance for the SMC Cepheids of [Fe/H]= -0.7 and solar metallicity for the Galactic ones, they determined, in a purely differential way, the following corrections: -0.21 $\\pm$ 0.19 for the $V$ and $K$ bands, -0.23 $\\pm$ 0.19 for the $I$-band and -0.29 $\\pm$ 0.19 for the Wesenheit index W. These agree well with Groenewegen et al. (2004).  Despite these ongoing observational efforts, it is important to underline that none of the observational studies undertaken so far has directly determined elemental abundances of a large sample of Cepheids in order to explicitly infer the metallicity effect on the PL relation, taking advantage of a sample that has been homogeneously analysed.  The novelty of our approach consists exactly in this, i.e. in the homogeneous analysis of a large sample of stars (68) in three galaxies (the Milky Way, the LMC and the SMC) spanning a factor of ten in metallicity and for which distances and $BVJK$ photometry are available. Preliminary results based on a sub-sample of the data discussed here were presented by Romaniello et al. (2005). Here we present the results about the iron content for the complete sample. In a forthcoming paper we will discuss the $\\alpha$-elements abundances.  The paper is organised as follows. The data sample is presented in Section 2. In Section 3 we thoroughly describe the data analysis and how we determine the metallicity of our stars. We compare our iron abundances with previous results in Section 4. The dependence of the PL relation on [Fe/H] is discussed in Section 5. Finally, Section 6 summarizes our concluding remarks.  \\begin{table*} \\label{table:2} \\caption{Pulsation phases ($\\phi$) and intrinsic parameters of the Galactic Cepheids. Both AP Pup and AX Vel were not included in the analysis of the metallicity effect because accurate distance estimates are not available in literature. In particular, for AP Pup we only listed the apparent mean magnitudes. In the last column is listed the duplicity status according to Szabados (2003): B - spectroscopic binary, Bc - spectroscopic binary that needs confirmation, O - spectroscopic binary with known orbit, V - visual binary} \\centering \\renewcommand{\\footnoterule}{} \\begin{tabular}{l l c c c c c c c c c} \\hline\\hline ID & $\\log P$ & $\\phi$ & $\\mu_{Old}$  & $E(B-V)_{Old}$ & $\\mu_{New}$  &  $E(B-V)_{New}$ & $M_{B}$ & $M_{V}$ & $M_{K}$ &  Duplicity \\\\ \\hline  l  Car  &     1.5509 & 0.580  &    8.99$^{d}$ &   0.170$^{d}$   &  8.56$^{b}$ & 0.147$^{a}$ &   -4.17\t&    -5.28 \t &  -7.53    &  \\ldots \\\\ U  Car  &     1.5891 & 0.490  &   10.97$^{d}$ &   0.283$^{d}$   & 10.87$^{a}$ & 0.265$^{a}$ &   -4.50\t&    -5.41 \t &  -7.44    &  B \\\\ V  Car  &     0.8259 & 0.375  &    9.84$^{e}$ &   0.174$^{h}$   & 10.09$^{a}$ & 0.169$^{a}$ &   -2.54\t&    -3.24 \t &  -4.86    &  B \\\\ WZ  Car  &     1.3620 & 0.745  &   12.92$^{d}$ &   0.384$^{d}$   & 12.69$^{a}$ & 0.370$^{a}$ &   -3.80\t&    -4.58 \t &  -6.52    &  \\ldots \\\\ V  Cen  &     0.7399 & 0.155  &    9.18$^{d}$ &   0.289$^{c}$   &  8.91$^{a}$ & 0.292$^{a}$ &   -2.41\t&    -2.99 \t &  -4.49    &  \\ldots \\\\ KN  Cen  &     1.5319 & 0.867  &   13.12$^{d}$ &   0.926$^{d}$   & 12.84$^{a}$ & 0.797$^{a}$ &   -4.63\t&    -5.46 \t &  -7.59    &  B \\\\ VW  Cen  &     1.1771 & 0.967  &   12.80$^{d}$ &   0.448$^{d}$   & 12.76$^{a}$ & 0.428$^{a}$ &   -2.93\t&    -3.85 \t &  -6.08    &  B \\\\ XX  Cen  &     1.0395 & 0.338  &   11.11$^{d}$ &   0.260$^{d}$   & 10.90$^{a}$ & 0.266$^{a}$ &   -3.19\t&    -3.91 \t &  -5.58    &  B \\\\ $\\beta$ Dor  &  0.9931 & 0.529  &    7.52$^{c}$ &   0.040$^{c}$   &  7.50$^{b}$ & 0.052$^{a}$ &   -3.16\t&    -3.91 \t &  -5.57    &  \\ldots \\\\ $\\zeta$ Gem  &  1.0065 & 0.460  &    7.78$^{c}$ &   0.010$^{c}$   &  7.81$^{b}$ & 0.014$^{a}$ &   -3.16\t&    -3.94 \t &  -5.72    &  V \\\\ GH  Lup  &     0.9675 & 0.031  &   10.05$^{e}$ &   0.364$^{h}$   & 10.25$^{a}$ & 0.335$^{a}$ &   -2.77\t&    -3.66 \t &  -5.54    &  B \\\\ T  Mon  &     1.4319 & 0.574  &   10.82$^{d}$ &   0.209$^{c}$   & 10.71$^{a}$ & 0.181$^{a}$ &   -4.16\t&    -5.15 \t &  -7.25    &  O \\\\ S  Mus  &     0.9850 & 0.266  &    9.81$^{e}$ &   0.147$^{h}$   &  9.57$^{a}$ & 0.212$^{a}$ &   -3.48\t&    -4.10 \t &  -5.62    &  O \\\\ UU  Mus  &     1.0658 & 0.865  &   12.59$^{d}$ &   0.413$^{d}$   & 12.41$^{a}$ & 0.399$^{a}$ &   -3.12\t&    -3.86 \t &  -5.70    &  \\ldots \\\\ S  Nor  &     0.9892 & 0.343  &    9.91$^{d}$ &   0.189$^{c}$   &  9.87$^{a}$ & 0.179$^{a}$ &   -3.23\t&    -4.00 \t &  -5.77    &  B \\\\ U  Nor  &     1.1019 & 0.422  &   10.72$^{d}$ &   0.892$^{d}$   & 10.46$^{a}$ & 0.862$^{a}$ &   -3.14\t&    -3.90 \t &  -5.72    &  \\ldots \\\\ X  Pup  &     1.4143 & 0.232  &   12.36$^{e}$ &   0.443$^{h}$   & 11.64$^{a}$ & 0.402$^{a}$ &   -3.57\t&    -4.38 \t &  -6.34    &  \\ldots \\\\ AP  Pup  &     0.7062 & 0.109  &  \\ldots$^{f}$ &   \\ldots$^{f}$  &\\ldots$^{f}$ &\\ldots$^{f}$ &    7.37\t&     6.78\t &   5.26    &  B \\\\ AQ  Pup  &     1.4786 & 0.436  &   12.52$^{d}$ &   0.512$^{c}$   & 12.41$^{a}$ & 0.518$^{a}$ &   -4.53\t&    -5.35 \t &  -7.27    &  B \\\\ BN  Pup  &     1.1359 & 0.397  &   12.95$^{d}$ &   0.438$^{c}$   & 12.93$^{a}$ & 0.416$^{a}$ &   -3.55\t&    -4.33 \t &  -6.14    &  \\ldots \\\\ LS  Pup  &     1.1506 & 0.012  &   13.55$^{d}$ &   0.478$^{d}$   & 13.39$^{a}$ & 0.461$^{a}$ &   -3.60\t&    -4.37 \t &  -6.18    &  B \\\\ RS  Pup  &     1.6174 & 0.944  &   11.56$^{d}$ &   0.446$^{c}$   & 11.30$^{a}$ & 0.457$^{a}$ &   -4.71\t&    -5.69 \t &  -7.81    &  \\ldots \\\\ VZ  Pup  &     1.3649 & 0.816  &   13.08$^{d}$ &   0.471$^{c}$   & 12.84$^{a}$ & 0.459$^{a}$ &   -3.93\t&    -4.63 \t &  -6.31    &  \\ldots \\\\ KQ  Sco  &     1.4577 & 0.446  &   12.36$^{c}$ &   0.839$^{c}$   & 12.23$^{g}$ & 0.869$^{a}$ &   -4.05\t&    -5.11 \t &  -7.55    &  \\ldots \\\\ EU  Tau  &     0.3227 & 0.414  &   10.27$^{c}$ &   0.170$^{c}$   & 10.27$^{c}$ & 0.170$^{d}$ &   -2.26\t&    -2.74 \t &  -4.05    &  Bc\\\\ SZ  Tau  &     0.4981 & 0.744  &    8.73$^{c}$ &   0.290$^{c}$   &  8.55$^{a}$ & 0.295$^{a}$ &   -2.38\t&    -2.93 \t &  -4.33    &  B \\\\ T  Vel  &     0.6665 & 0.233  &    9.80$^{d}$ &   0.281$^{c}$   & 10.02$^{a}$ & 0.289$^{a}$ &   -2.24\t&    -2.88 \t &  -4.47    &  B \\\\ AX  Vel  &     0.5650 & 0.872  &   10.76$^{f}$ &   0.224$^{h}$   &\\ldots$^{f}$ & 0.224$^{a}$ &   \\ldots&    \\ldots \t &  \\ldots   &  \\ldots \\\\ RY  Vel  &     1.4496 & 0.704  &   12.02$^{d}$ &   0.562$^{c}$   & 11.73$^{a}$ & 0.547$^{a}$ &   -4.23\t&    -5.05 \t &  -6.96    &  \\ldots \\\\ RZ  Vel  &     1.3096 & 0.793  &   11.02$^{d}$ &   0.335$^{c}$   & 10.77$^{a}$ & 0.299$^{a}$ &   -3.78\t&    -4.61 \t &  -6.56    &  \\ldots \\\\ SW  Vel  &     1.3700 & 0.792  &   11.00$^{d}$ &   0.349$^{c}$   & 11.88$^{a}$ & 0.344$^{a}$ &   -4.02\t&    -4.83 \t &  -6.75    &  \\ldots \\\\ SX  Vel  &     0.9800 & 0.497  &   11.44$^{e}$ &   0.250$^{h}$   & 11.41$^{g}$ & 0.263$^{a}$ &   -3.33\t&    -3.95 \t &  -5.49    &  \\ldots \\\\ \\hline \\hline \\multicolumn{8}{l}{$^a$ Fouqu\\'e et al. (2007).}\\\\ \\multicolumn{8}{l}{$^b$ Benedict et al. (2007).}\\\\ \\multicolumn{8}{l}{$^c$ Groenewegen et al. (2004).}\\\\ \\multicolumn{8}{l}{$^d$ Storm et al. (2004).}\\\\ \\multicolumn{8}{l}{$^e$ Laney \\& Stobie (1995) and Groenewegen (2004).}\\\\ \\multicolumn{8}{l}{$^f$ Not included in the analysis of the metallicity effect.}\\\\ \\multicolumn{8}{l}{$^g$ Groenewegen (2008).}\\\\ \\multicolumn{8}{l}{$^h$ Fernie et al. (1995).}\\\\ \\end{tabular} \\end{table*}  \\begin{table*} \\label{table:3} \\caption{Pulsation phases ($\\phi$) and intrinsic parameters of the Magellanic Cepheids. Periods ($\\log P$), apparent mean magnitudes and reddenings come from Laney \\& Stobie (1994). The mean $K$-band magnitudes were transformed into the 2MASS photometric system using the transformation provided by Koen et al. (2007).} \\centering \\begin{tabular}{l r r r r r c} \\hline\\hline ID & $\\log P$ & $\\phi$ & $B_0$ & $V_0$ & $K_0$ & $E(B-V)$\\\\ \\hline \\multicolumn{7}{c}{LMC} \\\\ \\hline HV 877   &  1.654 & 0.682 & 14.06 & 12.98 &  10.77  &  0.12 \\\\ HV 879   &  1.566 & 0.256 & 14.12 & 13.15 &  11.03  &  0.06 \\\\ HV 971   &  0.968 & 0.237 & 14.86 & 14.24 &  12.68  &  0.06 \\\\ HV 997   &  1.119 & 0.130 & 14.94 & 14.19 &  12.37  &  0.10 \\\\ HV 1013  &  1.382 & 0.710 & 14.39 & 13.46 &  11.41  &  0.11 \\\\ HV 1023  &  1.425 & 0.144 & 14.48 & 13.51 &  11.45  &  0.07 \\\\ HV 2260  &  1.112 & 0.144 & 15.19 & 14.43 &  12.67  &  0.13 \\\\ HV 2294  &  1.563 & 0.605 & 13.19 & 12.45 &  10.74  &  0.07 \\\\ HV 2337  &  0.837 & 0.861 &\\ldots & \\ldots&  13.27  &  0.07 \\\\ HV 2352  &  1.134 & 0.201 & 14.49 & 13.84 &  12.25  &  0.10 \\\\ HV 2369  &  1.684 & 0.136 & 13.15 & 12.29 &  10.38  &  0.10 \\\\ HV 2405  &  0.840 & 0.037 & \\ldots& \\ldots&  13.43  &  0.07 \\\\ HV 2580  &  1.228 & 0.119 & 14.33 & 13.67 &  11.92  &  0.09 \\\\ HV 2733  &  0.941 & 0.411 & 14.85 & 14.34 &  13.00  &  0.11 \\\\ HV 2793  &  1.283 & 0.917 & 14.49 & 13.58 &  11.75  &  0.10 \\\\ HV 2827  &  1.897 & 0.880 & 13.19 & 12.03 &   9.80  &  0.08 \\\\ HV 2836  &  1.244 & 0.059 & 14.85 & 14.02 &  12.04  &  0.18 \\\\ HV 2864  &  1.041 & 0.055 & 15.16 & 14.42 &  12.77  &  0.07 \\\\ HV 5497  &  1.997 & 0.321 & 12.73 & 11.63 &   9.43  &  0.10 \\\\ HV 6093  &  0.680 & 0.024 & 15.74 & 15.16 &  13.71  &  0.06 \\\\ HV 12452 &  0.941 & 0.860 & 15.25 & 14.60 &  12.83  &  0.06 \\\\ HV 12700 &  0.911 & 0.342 & 15.62 & 14.87 &  13.12  & -0.01 \\\\ \\hline \\multicolumn{7}{c}{SMC} \\\\ \\hline HV 817   &  1.277 & 0.298 &   14.13  & 13.59  &  12.12  &  0.08 \\\\ HV 823   &  1.504 & 0.873 &   14.46  & 13.60  &  11.58  &  0.05 \\\\ HV 824   &  1.818 & 0.315 &   13.06  & 12.27  &  10.33  &  0.03 \\\\ HV 829   &  1.931 & 0.348 &   12.61  & 11.81  &   9.92  &  0.03 \\\\ HV 834   &  1.866 & 0.557 &   12.95  & 12.14  &  10.20  &  0.02 \\\\ HV 837   &  1.631 & 0.822 &   13.95  & 13.10  &  11.11  &  0.04 \\\\ HV 847   &  1.433 & 0.500 &   14.40  & 13.66  &  11.83  &  0.08 \\\\ HV 865   &  1.523 & 0.108 &   13.55  & 12.93  &  11.21  &  0.06 \\\\ HV 1365  &  1.094 & 0.184 &   15.39  & 14.79  &  13.20  &  0.07 \\\\ HV 1954  &  1.223 & 0.847 &   14.13  & 13.62  &  12.12  &  0.07 \\\\ HV 2064  &  1.527 & 0.279 &   14.28  & 13.50  &  11.61  &  0.07 \\\\ HV 2195  &  1.621 & 0.135 &   13.85  & 13.07  &  11.09  & -0.02 \\\\ HV 2209  &  1.355 & 0.822 &   13.99  & 13.42  &  11.84  &  0.04 \\\\ HV 11211 &  1.330 & 0.516 &   14.36  & 13.64  &  11.83  &  0.06 \\\\ \\hline \\hline \\end{tabular} \\end{table*} ",
        "conclusions": "We have directly measured the iron abundances for 68 Galactic and Magellanic Cepheids from FEROS and UVES high resolution and high signal-to-noise spectra. We have used these measurements to assess the influence of the stellar iron content on the Cepheid PL relation in the $V$ and in the $K$ band. In order to do this we have related the $V$-band and the $K$-band residuals from the standard PL relations of Freedman et al. (2001) and Persson et al. (2004), respectively, to [Fe/H]. Differently from previous studies, we can constrain the PL relation using Cepheids with known distance moduli and chemical abundances, homogeneously measured, that cover almost a factor of ten in metallicity.  For our Galactic sample, we find that the mean value of the iron content is solar ($\\sigma$ = 0.10, see Fig.~6), with a range of values between -0.18 dex and +0.25 dex. For the LMC sample, we find that the mean value is about $\\sim$ -0.33 dex ($\\sigma$ =0.13, see Fig.~6), with a range of values between -0.62 dex and -0.10 dex. For the SMC sample, we find that the mean value is about $\\sim$ -0.75 dex ($\\sigma$ = 0.08, see Fig.~6), with a range of values between -0.87 and -0.63.  We have compared our results with the analyses of FC97, Andrievsky et al. (2002a, 2002b, 2002c) and Luck et al. (2003) for the Galactic Cepheids and LL92 and L98 for the Magellanic Clouds. Regarding the Galactic sample, our results are marginally more in agreement with Andrievsky's values than with FC97 and the differences, on average, appear rather small. Considering the Magellanic Cepheids, we have a poor agreement with LL92, which could be in part accounted for by different analytical tools and data quality. Our data are in better agreement with L98 results and the spread of our iron abundances in the LMC and SMC is similar to the one they reported. We note that the mean metallicity that they found with their complete sample (-0.30 dex and -0.74 dex) is in very good agreement with our results.  Our main results concerning the effect of the iron abundance on the PL relation are summarized in Fig.~8 and Fig.~9 (bottom panels) and they hold for a LMC distance modulus of 18.50. In Fig.~10 is also showed the comparison in the $V$-band with the empirical results of Kennicutt et al. (1998) in two Cepheid fields of M101 (open circles and solid line). The main findings can be summarized as follows:  \\begin{itemize} \\item The $V$-band PL relation does depend on the metal abundance. This finding is marginally affected by the adopted distance scale for the Galactic Cepheids and by the LMC distance.  \\item Current data do not allow us to reach a firm conclusion concerning the dependence of the $K$-band PL relation on the metal content. The use of the most recent distances for Galactic Cepheids (Benedict et al. 2007; Fouqu\\`e et al. 2007; van Leeuwen et al. 2007) indicates a mild metallicity effect. On the other hand, the use of the old distances (Storm et al. 2004) suggest a vanishing effect.  \\item Residuals based on a canonical LMC distance ($\\mu_{LMC}=18.5$) and on the most recent distances for Galactic Cepheids present a well defined effect in the $V$-band. The metal-poor and the metal-rich bin are $\\approx 2 \\sigma$ and $\\approx 9 \\sigma$ from the null hypothesis. Moreover, the two metallicity bins differ at the $3\\sigma$ level.  \\item By assuming a \"short\" LMC distance ($\\mu_{LMC}=18.3$) the residuals present a strong metallicity dependency in the zero-point of the $V$-band PL relation. By assuming a \"long\" LMC distance ($\\mu_{LMC}=18.7$) we found a strong metallicity effect when moving from metal-poor to metal-rich Cepheids. This indicates a significant change in the slope and probably in the zero-point. The findings based on the \"short\" and on the \"long\" LMC distance are at odds with current empirical and theoretical evidence, suggesting a smaller metallicity effect.  \\item Metal-rich Cepheids in the $V$-band are systematically fainter than metal-poor ones. This evidence is strongly supported by the canonical, the \"short\" and the \"long\" LMC distance.  \\end{itemize}  The above results together with recent robust LMC distance estimates indicate that the behaviours based on the canonical distance appear to be the most reliable ones.  In order to constrain on a more quantitative basis the metallicity dependence of both zero-point and slope of the optical PL relations is required a larger number of Cepheids covering a broader range in metal abundances. Moreover, for each metallicity bin Cepheids covering a broad period range are required to reduce the error on the residuals and to constrain on a quantitative basis the fine structure of the PL relation in optical and NIR photometric bands."
    },
    "0807/0807.0731_arXiv.txt": {
        "abstract": "In this paper we present preliminary results from cosmological simulations of modified gravity in the dark matter sector. Our results show improvements over standard cold dark matter cosmology. The abundance of low-mass haloes in the modified gravity model fit observations better than the conventional theory, while the differences of the modified density fluctuation power spectrum differs from the standard, $\\Lambda$CDM power spectrum are small enough to make these two models observationally indistinguishable at large scales. ",
        "introduction": "Over the past 30 years cosmology became a successful and mature empirical science. The growing accuracy of astronomical observations gave the possibility to test theoretical models. Past decade has established the \"standard cosmological model\" dubbed also as the $\\Lambda$CDM cosmology, or the \"concordance model\". In this picture the evolution of the expanding universe, dominated by non-relativistic (cold) dark mater particles is governed by Einstein's gravity with a positive cosmological constant. Although the standard model is very successful in explaining many observations it still has some difficulties. The physical nature of the dark matter (DM) sector is still unknown. Moreover, the mass distribution in the real universe (in particular, the emptiness of voids) differs significantly from the predictions of the $\\Lambda$CDM model. Modifications of the standard gravity theory have been proposed to deal with this impasse. In this paper we show preliminary results of our work based on one of such modifications. ",
        "conclusions": "Our results show that the modification of standard gravity we consider can close the gap between some properties of standard $\\Lambda$CDM cosmology simulations and today's high resolution astronomical observations. Deviations from the standard power spectrum are negligible, and they agree with large scale observations. On the other hand, our simulations of scalar interaction show desired clustering properties on small scales. While the court on the physical nature of dark matter particles is still out it is scientifically justified to probe various theoretical possibilities. In this paper we have shown that the long-range scalar DM interaction has the potential to improve the standard cosmological model."
    },
    "0807/0807.0441_arXiv.txt": {
        "abstract": "We calculate the rate of in-fall of stellar matter on an accretion disk during the collapse of a rapidly rotating massive star, and estimate the luminosity of the relativistic jet that results from accretion on to the central black hole.  We find that the jet luminosity remains high for about $10^2$ s, at a level comparable to the typical luminosity observed in gamma-ray bursts (GRBs).  The luminosity then decreases rapidly with time for about $\\sim10^3$ s, roughly as $\\sim t^{-3}$; the duration depends on the size and rotation speed of the stellar core.  The rapid decrease of the jet power explains the steeply declining x-ray flux observed at the end of most long duration GRBs.  Observations with the Swift satellite show that, following the steep decline, many GRBs exhibit a plateau in the x-ray lightcurve (XLC) that lasts for about 10$^4$ s.  We suggest that this puzzling feature is due to continued accretion in the central engine.  A plateau in the jet luminosity can arise when the viscosity parameter $\\alpha$ is small, $\\sim10^{-2}$ or less. A plateau is also produced by continued fall-back of matter -- either from an extended stellar envelope or from material that failed to escape with the supernova ejecta. In a few GRBs, the XLC is observed to drop suddenly at the end of the plateau phase, while in others the XLC declines more slowly as $\\sim t^{-1}-t^{-2}$.  These features arise naturally in the accretion model depending on the radius and mean specific angular momentum of the stellar envelope.  The total energy in the disk-wind accompanying accretion is found to be about 10$^{52}$ erg.  This is comparable to the energy observed in supernovae associated with GRBs, suggesting that the wind might be the primary agent responsible for the explosion.  The accretion model thus provides a coherent explanation for the diverse and puzzling features observed in the early XLC of GRBs.  It might be possible to use this model to invert gamma-ray and x-ray observations of GRBs and thereby infer basic properties of the core and envelope of the GRB progenitor star. ",
        "introduction": "Observations of gamma-ray bursts (GRBs) carried out by the NASA Swift satellite in the last two years have shown that the $\\gamma$-ray prompt emission turns off abruptly after about a minute.  The abrupt shutoff is evidenced by the rapidly declining x-ray flux (t$^{-3}$ or faster) from about 80s to 300s, which joins smoothly the prompt GRB lightcurve when extrapolated back in time and spectral band (Tagliaferri et al. 2005; Nousek et al. 2006; O'Brien et al. 2006). The abrupt decline suggests that the activity at the center of the explosion declines very rapidly with time after about a minute of more or less steady activity.  At the same time, Swift observations have provided overwhelming evidence that the GRB central engine continues operating for hours and perhaps even days. There are two independent indications for this phenomenon.  First, we often see a phase in the early x-ray lightcurve of long-GRBs, from $\\sim 10^3$s to 10$^4$s, during which the flux declines slowly with time.  We will refer to this as a ``plateau'' in the lightcurve.  There are many proposals to explain the plateau. The models are, however, highly constrained by a lack of correlation between x-ray and optical features for many GRBs.  Therefore, a successful proposal must invoke continuing activity of the central engine to produce the x-ray plateau (e.g., Zhang 2006; Panaitescu 2007, and references therein), whereas the simultaneous optical emission may be produced in the afterglow.  Second, in roughly a third of the observed GRBs, the x-ray flux is seen to increase suddenly and then to drop precipitously in what are referred to as ``flares.''  In some cases the flux during these flares increases by a factor of $\\sim 10^2$ on a time scale $\\delta t\\ll t$ (Burrows et al. 2005; Chincarini et al. 2007), which cannot be explained in terms of a density inhomogeneity in the external medium (Nakar \\& Granot, 2006).  Variable activity in the central engine is a more natural explanation.  A sudden drop in the flux --- e.g., in the case of GRB 070110 the x-ray lightcurve was nearly flat for 20 ks and then fell off as $t^{-8}$ (Troja et al. 2007) --- is also not possible to understand other than as the result of highly variable central engine activity.  The above conflicting requirements, viz., (i) a sudden drop in activity at the end of the main GRB ($t\\ \\lta 10^2$ s), (ii) continued steady activity during an extended plateau ($t\\sim10^4$ s), and (iii) occasional dramatic flares, are challenging for models of the central engine.  In this paper we consider the currently most popular model of GRBs, which postulates ultra-rapid accretion of gas on to a newly-formed black hole or neutron star (Narayan, Paczy\\'nski \\& Piran 1992; Popham, Woosley \\& Fryer 1999; Narayan, Piran \\& Kumar 2001). The accretion disk may be the result of (i) gas fall-back after a hypernova, as in the collapsar model (Woosley 1993; Paczy\\'nski 1998; MacFadyen \\& Woosley 1999), which is considered relevant for long duration GRBs, or (ii) during the merger of a double neutron star or neutron star-black hole binary (Eichler et al. 1989; Narayan et al. 1992), which is currently considered the most likely explanation for short duration GRBs. We attempt to reconcile the accretion model for collapsars with Swift observations of long duration GRBs.  In \\S~2, we make use of a reasonably realistic model of the progenitor of a collapsar and estimate the rate at which gas is added to the accretion torus at the center.  The calculation is based on a crude free-fall model for the collapsing star and the results are to be taken in the spirit of an order of magnitude estimate. Then, in \\S~3, we use this model to study the variation of accretion power with time.  We show that the model naturally reproduces both the sudden shutoff of the prompt GRB emission and the extended plateau in the lightcurve.  We speculate on possible scenarios for producing flares and suggest an explanation for the hypernova explosion associated with some GRBs.  We conclude in \\S~4 with a discussion and summary. ",
        "conclusions": "This work was motivated by two basic questions posed by the extensive and excellent Swift observations of GRBs: (1) Why does the $\\gamma$-ray/x-ray flux undergo a sharp decline (flux decreasing as $t^{-3}$ or faster) about one minute after the start of the burst, even though the central engine itself is apparently active for hours? (2) After the sharp decline, how does the power from the engine remain nearly constant for a period of $\\sim10^4$ s to produce the long plateau that is observed in the x-ray lightcurves of many GRBs?  We have attempted to answer these questions in the context of an accretion model of GRBs in which the central engine is powered by accretion on to a BH, and the GRB luminosity is proportional to the power in the resulting relativistic jets.  We employ the pre-collapse stellar model of Woosley \\& Heger (2006), with the density structure and rotation profile shown in Fig. 2.  We investigate the post-collapse accretion activity of this model from $\\sim$1 s to $\\sim10^5$ s and compare the model predictions with GRB observations. We compute the mass in-fall rate during the collapse of the massive star, and the viscous accretion rate on to the central BH using ideas from Narayan et al. (2001) and Kohri et al. (2005).  We then estimate the relativistic jet luminosity from the mass accretion rate and the BH spin, using a prescription proposed by McKinney (2005).  We find that the jet has a luminosity of about $10^{51} ~{\\rm erg\\,s^{-1}}$, lasting for about 10-20s (Fig. 4), which is of order the power and duration of a typical long GRB.  The luminosity arises from the fall-back and accretion of the outer half of the core of the progenitor star; the material here has sufficient angular momentum to go into orbit, whereas the material from the inner half of the core collapses directly to form the BH.  When the outer layers of the stellar core -- where the density falls off rapidly with $r$ -- are accreted via the ADAF process, the jet luminosity drops rapidly. This provides a straightforward and natural explanation for the steep decline of the early x-ray lightcurve observed by Swift.  Fluence of gamma-ray bursts can be used to constrain the rotation rate in the core of their progenitor stars to within a factor of 10 or better; too large or too small $\\Omega$ in the core results in a small jet luminosity (\\S3.4), and causes the temporal decay of the lightcurve at the end of the prompt emission, for a few hundred seconds, to be either too shallow or too steep (fig. 5).  In our model, the plateau in the x-ray lightcurve requires continued accretion to power the jet.  One possibility is that the viscous time scale in the disk is so long that it takes a time $\\sim10^4$ s for the material in the disk to accrete.  A viscosity parameter $\\alpha\\ \\lta\\ 10^{-2}$ is required.  In this scenario, although mass fall-back ceases in a few hundred seconds, accretion continues on for a few hours (Fig. 6).  While the model succeeds in producing an extended plateau, it causes the fall-off of the early GRB lightcurve to be less steep -- more like $t^{-2}$ than $t^{-3}$ -- (Fig. 6) in conflict with observations.  The decline at the end of the plateau is also fairly shallow, inconsistent with observations of some GRBs.  A second and, in our opinion, more likely scenario is that there is continued mass fall-back for the entire duration of the plateau (Fig. 7).  In this case, the accretion time scale itself is short (i.e., $\\alpha$ is large), so the lightcurve reflects the mass fall-back rate.  One possibility is that the fall-back is due to material that fails to be ejected by the supernova explosion.  This idea has two problems.  First, it is hard to see why there should be an extended period of a fairly constant fall-back rate as required by the observations, whereas we expect the fall-back rate to vary much more steeply as $t^{-5/3}$ (Chevalier 1989).  Second, it is hard to see how we can have a sudden cutoff in the fall-back at the end of the plateau as seen in several GRBs.  Another possibility is that the progenitor star has a core-envelope structure, with the core producing the early GRB and the envelope producing the plateau.  The core must have a radius of $\\sim {\\rm few}\\times10^{10}$ cm to explain the duration of the GRB and the envelope must have a radius of $\\sim {\\rm few}\\times10^{11}$ cm to explain the duration of the plateau.  There should be a large density (or $j$) contrast between the core and the envelope, in order to explain the sharp cutoff of the prompt emission, and the density profile in the envelope must be fairly shallow, $\\rho\\sim r^{-2}$, in order to obtain a shallow plateau.  Depending on the rotation profile of the star, various kinds of lightcurves --- including ones in which we have a very rapid cutoff of the x-ray luminosity at the end of the plateau --- are possible.  This scenario is therefore capable of explaining almost all observed cases.  GRBs that do not have a plateau in their lightcurve are also easily explained; these presumably had progenitors with only a core and no envelope (like the model shown in Fig. 2).  A nice feature of this model is that we could, in principle, use GRB observations to deduce the density and rotation structure of the progenitor star.  On the other hand, it is not clear that evolved massive stars do have the kind of core-envelope structure we need to explain a typical GRB x-ray plateau.  (We are not aware of any pre-supernova models in the literature with the required properties.)  An important implication of the accretion model of GRB central engines is that the accretion flows are advection-dominated and thus have strong outflows/winds. We have estimated the total energy in the wind for the pre-collapse stellar model of Fig. 2 and find it to be about 2x10$^{52}$ erg.  This is sufficient to explode the star, and might explain the observed energetics of supernovae associated with GRBs.  A generic prediction of the late fall-back model for the x-ray plateau is that brighter GRBs should have a weaker (lower luminosity) and shorter duration plateau.  The reason is that a stronger GRB, with its stronger jet and wind, is likely to eject more of the stellar envelope during the main burst. Indeed, recent simulations of relativistic jet-induced supernovae support this prediction, with more luminous explosions expelling more of the stellar envelope and leaving less material available for accretion (eg. Tominaga 2007).  In order to test this prediction, we have looked at a sample of bursts with known redshift and isotropic equivalent luminosities (Butler \\& Kocevski 2007).  From this set of bursts, we have selected two subsets, those with distinct x-ray plateaus (GRBs 070110, 060614, 050315, 060607A, 060729, 070810A) and those clearly lacking x-ray plateaus (GRBs 071020, 070318, 061007, 050922C, 050826, 070411). Consistent with the prediction of the fall-back model, we find that the average peak isotropic equivalent luminosity (per frequency interval) of the subset of bursts with distinct x-ray plateaus is $\\sim$ 4 times lower than that of the subset of bursts lacking an x-ray plateau.  Moreover, using the same sample of bursts, we find that the average peak isotropic equivalent luminosity (per frequency interval) of bursts with an x-ray plateau is lower than that of bursts without an x-ray plateau at the $>$ 10\\% level of significance, assuming normal distributions for the luminosities of each of these populations of bursts.  A related prediction is that the plateau should be absent, or at least weak, in those cases where we see a bright supernova event associated with a GRB.  The idea is that a bright supernova implies powerful ejection, and there should be less material available for accretion.  We have only one well observed case of a GRB-supernova association in the Swift sample of bursts (GRB 030329), and the x-ray lightcurve did not have a plateau.  While this observation is consistent with the late fall-back model, we note that it is just a single object and therefore the result is not very significant.  It is interesting to note that x-ray plateaus are not seen for short duration GRBs (see Nakar, 2007, for an excellent review).  This is consistent with our model.  If short GRBs are the result of the merger of double neutron star binaries (the currently popular model), then there is no material in an extended envelope in the progenitor to produce late fall-back.  The model described in this paper is obviously incomplete.  We have not provided any quantitative explanation for the x-ray flares seen during the plateau phase (and even later) in many GRBs.  The only qualitative idea we have offered is that the flares reflect an instability in the accretion disk.  We have also simplified the model considerably by postulating a direct proportionality between the jet power at the point where it is launched from the BH and the observed luminosity.  Several factors could seriously modify this relation. First, the tunneling of the jet through the stellar material may be inefficient, and so the power that escapes from the surface of the star may be a small (and variable) fraction of the jet power at the base. Second, the efficiency with which the escaping jet power is converted to radiation (the physics of which is poorly understood) may be variable.  Finally, the beaming of the jet may be different for the prompt GRB and the x-ray plateau, and may even vary during the plateau.  We have ignored these complications in the interests of simplicity.   \\noindent{\\bf Acknowledgment:} We thank Stan Woosley for providing his stellar model and for very useful comments on the draft. We are grateful to Craig Wheeler for a number of excellent suggestions that improved the paper. This work is supported in part by a NSF grant (AST-0406878), and NASA Swift-GI-program."
    },
    "0807/0807.5052_arXiv.txt": {
        "abstract": "{ {\\em Context:} There is increasing need for good algorithms for modeling the aggregation and fragmentation of solid particles (dust grains, dust aggregates, boulders) in various astrophysical settings, including protoplanetary disks, planetary- and sub-stellar atmospheres and dense molecular cloud cores. Here we describe a new algorithm that combines advantages of various standard methods into one. \\\\ {\\em Aims:} The aim is to develop a method that 1) can solve for aggregation and fragmentation, 2) can easily include the effect and evolution of grain properties such as compactness, composition, etc., and 3) can be built as a coagulation/fragmentation module into a hydrodynamics simulation where it 3a) allows for non-`thermalized' non-local motions of particles (e.g.~movement of particles in turbulent flows with stopping time larger than eddy turn-over time) and 3b) focuses computational effort there where most of the mass is. \\\\ {\\em Methods:} We develop a Monte-Carlo method in which we follow the ``life'' of a limited number of representative particles. Each of these particles is associated with a certain fraction of the total dust mass and thereby represents a large number of true particles which all are assumed to have the same properties as their representative particle. Under the assumption that the total number of true particles vastly exceeds the number of representative particles, the chance of a representative particle colliding with another representative particle is negligibly small, and we therefore ignore this possibility. This now makes it possible to employ a statistical approach to the evolution of the representative particles, which is the core of our Monte Carlo method. \\\\ {\\em Results:} The method reproduces the known analytic solutions of simplified coagulation kernels, and compares well to numerical results for Brownian motion using other methods. For reasonably well-behaved kernels it produces good results even for moderate number of swarms.} ",
        "introduction": "Dust particle aggregation is a very common process in various astrophysical settings. In protoplanetary disks the aggregation of dust particles forms the very initial step of planet formation (see e.g.\\ Dominik et al.~\\citeyear{dominikblum:2007}). It also modifies the optical properties of the disk, and it has influence on the chemistry and free electron abundance in a disk (Sano et al.~\\citeyear{sanomirama:2000}; Semenov et al.~\\citeyear{semenovwiebe:2004}; Ilgner \\& Nelson \\citeyear{ilgnernelson:2006a}). The appearance and evolution of a protoplanetary disk is therefore critically affected by the dust aggregation process. In sub-stellar and planetary atmospheres the aggregation of dust particles and the coagulation of fluid droplets can affect the structure of cloud layers. It can therefore strongly affect the spectrum of these objects and influence the local conditions within these atmospheres. The process of aggregation/coagulation and the reverse process of fragmentation or cratering are therefore important processes to understand, but at the same time they are extremely complex.  Traditional methods solve the Smoluchowski equation for the particle mass distribution function $f(m)$, where $f(m)$ is defined such that $f(m)dm$ denotes the number of particles per cubic centimeter with masses in the interval $[m,m+dm]$. This kind of method has been used in many papers on dust coagulation before (e.g.\\ Nakagawa et al.~\\citeyear{nakanakahayashi:1981}; Weidenschilling \\citeyear{weidenschilling:1984,weid1997}; Schmitt et al.~\\citeyear{schmitthenningmucha:1997}; Suttner \\& Yorke \\citeyear{suttneryorke:1999}; Tanaka et al.~\\citeyear{tanakahimemoida:2005}; Dullemond \\& Dominik \\citeyear{duldom:2005}; Nomura \\& Nakagawa \\citeyear{nomuranaka:2006}). Methods of this kind are efficient, but have many known problems. First of all a coarse sampling of the particle mass leads to systematic errors such as the acceleration of growth (Ohtsuki et al.~\\citeyear{ohtsuki:1990}). High resolution is therefore required, which may make certain problems computationally expensive. Moreover, if one wishes to include additional properties of a particle, such as porosity, charge, composition etc, then each of these properties adds another dimension to the problem. If each of these dimensions is sampled properly, this can quickly make the problem prohibitively computationally expensive. Finally, the traditional methods are less well suited for modeling stochastic behavior of particles unless this stochastic behavior can be treated in an averaged way. For instance, in protoplanetary disks if the stopping time of a particle is roughly equal to the turbulent eddy turn-over time, then the velocity of a particle with respect to the gas is stochastic: at the same location there can exist two particles with identical properties but which happen to have different velocities because they entered the eddy from different directions (see e.g.\\ the simulations by Johansen et al.~\\citeyear{johansen:2006}).  To circumvent problems of this kind Ormel et al.~(\\citeyear{ormelmonte:2007}) have presented a Monte Carlo approach to coagulation. In this approach the particles are treated as computational particles in a volume which is representative of a much larger volume. The simulation follows the life of $N$ particles as they collide and stick or fragment. The collision rates among these particles are computed, and by use of random numbers it is then determined which particle collides with which.  The outcome of the collision is then determined depending on the properties of the two colliding particles and their relative impact velocity. This method, under ideal conditions, provides the true simulation of the process, except that random numbers are used in combination with collision rates to determine the next collision event. This method has many advantages over the tradiational methods. It is nearly trivial to add any number of particle properties to each particle. There is less worry of systematic errors because it is so close to a true simulation of the system, and it is easy to implement.  A disadvantage is that upon coagulation the number of computational particles goes down as the particles coagulate. Ormel et al.~solve this problem by enlarging the volume of the simulation and hence add new particles, but this means that the method is not very well suited for modeling coagulation within a spatially resolved setting such as a hydrodynamic simulation or a model of a protoplanetary disk.  It is the purpose of the present paper to present an alternative Monte Carlo method which can quite naturally deal with extremely large numbers of particles, which keeps the number of computational particles constant throughout the simulation and which can be used in spatially resolved models. ",
        "conclusions": "We have shown that our representative particle method for aggregation of particles in astrophysical settings works well for standard kernels. It has the usual advantages of Monte Carlo methods that one can add particle properties easily and without loss of computational speed. Moreover, it naturally conserves the number of computational elements, so there is no need to ``add'' or ``remove'' particles. Each representative particle represents a fixed portion of the total mass of solids.  Our method may have various possible interesting extensions and applications. Here we speculate on a few of these. For instance, the fact that each representative particle corresponds to a fixed amount of solid mass makes the method ideal for implementation into spatially resolved models such as hydrodynamic simulations of planetary atmospheres or protoplanetary disks. We can then follow the exact motion of each representative particle through the possibly turbulent environment, and thereby automatically treat the stochastic nature and deviation from a Boltzmann distribution of the motion of particles with stopping times of the same order as the turbulent eddy turn-over time. It is necessesary, however, to assure that a sufficiently large number of representative particles is present in each grid cell of the hydrodynamic simulation. For large scale hydrodynamic simulations this may lead to a very large computational demand for the coagulation computation, as well as for tracking the exact motion of these particles. If strong clumping of the particles happens, however, much of the ``action'' anyway happens in these ``clumps'', and it may then not be too critical that other grid cells are not sufficiently populated by representative particles. This, however, is something that has to be experimented.  Our representative particle method can in principle also be used to model the sublimation and condensation of dust grains. If a particle sublimates then the representative particle becomes simply an atom or molecule of the vapor of this process. It will then follow the gas motion until the temperature becomes low enough that it can condense again. Other representative particles which are still in the solid phase may represent physical particles that can act as a condensation nucleus. Finally, in our method the properties of the particle can not only change due to collisions, but we can easily implement other environmental factors in the alteration of particle properties.  There are two main drawbacks of the method. First, it only works for large particle numbers, i.e.\\ it cannot treat problems in which individual particles start dominating their immediate environment. Ormel's method and its expected extension do not have this problem. Secondly, the method cannot be accelerated using implicit integration, while Brauer's method can.  All in all we believe that this method may have interesting applications in the field of dust aggregation and droplet coagulation in protoplanetary disks and planetary atmospheres.   \\begin{acknowledgement} We wish to thank Frithjof Brauer, Anders Johansen, Patrick Glaschke, Thomas Henning, J\\\"urgen Blum, Carsten G\\\"uttler, Carsten Dominik and Dominik Paszun for useful comments. We also thank the anonymous referee for very useful comments that helped to improve the paper and for pointing us to a problem in our presentation of the solution to the particle number conservation paradox. \\end{acknowledgement}"
    },
    "0807/0807.2499_arXiv.txt": {
        "abstract": "{\\mbox{BD$-6\\degr$1178} identified with the infrared source IRAS\\,05238$-$0626 is shown for the first time to be a spectroscopic binary (SB2) by analyzing the high-resolution spectra taken with the NES echelle spectrograph of the 6-m telescope. The components of the binary have close spectral types and luminosity classes: F5\\,IV--III and F3\\,V. The heliocentric radial velocities are measured for both components at four observing moments in 2004--2005. Both stars have close rotation velocities, which are equal to 24 and 19\\,km/s. We do not confirm the classification of BD$-6\\degr$1178 as a supergiant in the transition stage of becoming a planetary nebula. BD$-6\\degr$1178 probably is a young pre-MS stars. It is possibly a member of the 1c subgroup of the Ori\\,OB1 association.}  \\titlerunning{\\it Spectroscopic binary BD${ -6\\degr}$1178} \\authorrunning{\\it Klochkova \\& Chentsov} ",
        "introduction": "In this paper we continue to publish the results of our spectroscopy of stars with IR excesses (see [\\cite{Klochkova1995, KSPV, IRAS23304, AFGL2688, IRAS20000}] and references therein for the main results). BD$-6\\degr$1178 is an optical component of the infrared source IRAS\\,05238$-$0626 (with galactic coordinates l=208.9$\\degr$, b=$-21.8\\degr$). This object is considered to be a candidate to protoplanetary nebula (PPN) according to the observed excess of radiation in the 12--60\\,$\\mu$m wavelength region and its position on the IR colour--colour diagram [\\cite{Garcia-1990, Reddy, Fujii}]. Recall that according to modern concepts (see, e.g., [\\cite{Block}]), objects observed at the short-lived evolutionary stage of a young planetary (protoplanetary) nebula are intermediate-mass stars evolving away from the asymptotic giant branch (AGB) toward the stage of a planetary nebula. The initial main-sequence (MS) masses of these stars lie in the 3--8\\,${\\mathcal M}_{\\odot}$ mass interval. During the AGB stage these stars have lost much mass in the form of a powerful stellar wind, and as a result, at the PPN stage the stars have the form of degenerate carbon-oxygen nuclei with typical masses of about 0.6\\,${\\mathcal M}_{\\odot}$ surrounded by expanding gas-and-dust shell. The astronomers are interested in studying PPNs, first because they allow one to study stellar-wind driven mass loss and second, because they offer a unique opportunity of observing the result of the stellar nucleosynthesis, mixing, and dredge-up of products of nuclear reactions that occurred during the preceding evolution of the star.  About a dozen objects overabundant in heavy metals synthesized via neutronization of iron nuclei under the conditions of low neutron density (\\mbox{$s$-process}) have been found among the PPN-candidate studied. An analysis of the properties of PPNs showed that the expected overabundances of $s$-process elements are observed only the atmospheres of \\mbox{C-rich} stars whose IR spectra contain an emission at 21\\,$\\mu$m [\\cite{Klochkova1995, IRAS20000, Klochkova1997, Winckel}]. However, the overwhelming majority of PPNs exhibit neither carbon (\\mbox{O-rich} stars) nor heavy-element overabundance (see, e.g., [\\cite{Klochkova1995, QYSge, IRAS19475}]). The correlation found between the excess of heavy elements in the star's atmosphere and the peculiarity of the IR spectrum of the envelope of the star remains unexplained and hence a further increase of the sample of PPN objects studied is needed.   Currently, we know little about BD$-6\\degr$1178. Its sky coordinates for the epoch of 2000 are: $\\alpha$=05$^{\\rm h} 26^{\\rm m} 19.8^{\\rm s}$, $\\delta$=$-6\\degr 23^{\\rm '} 57^{\\rm ''}$. The $V$- and $B$-band apparent magnitudes are equal to $V$\\,=\\,$10.52^m$ and \\mbox{$B$\\,=\\,$10.96^m$ [\\cite{Fujii}]}. Some evidence for the photometric variability of the star was found: according to the NSVS catalog [\\cite{Wozniak}], the mean magnitude of the star in a close-to-$R$-filter passband varies in the interval 10.78--$10.87^m$ and its standard deviation is about $0.01^m$. Modeling of the spectral energy distribution based on the multicolor photometry in the visual and near IR yields effective-temperature values T$_{eff}$ ranging from 8000\\,K [\\cite{Garcia-1990}] to 7400\\,K [\\cite{Fujii}], which corresponds to late A --- early F-subclasses.  As for spectroscopic observations, only low-resolution ($\\approx$5\\,\\AA/pixel) spectra have been published so far for BD$-6\\degr$1178. These spectra yielded the following estimates for the spectral type: F2\\,II [\\cite{Reddy}], F4 [\\cite{Suarez}], and F5 [\\cite{Torres}]. In view of the aforesaid, it becomes evident from these results that further detailed study of the optical spectrum of the star is needed. In this paper we report the results of our numerous high-resolution spectroscopic observations of BD$-6\\degr$1178 made with the 6-m telescope of the Special Astrophysical Observatory of the Russian Academy of Sciences (SAO RAS). The aim of this study is to perform two-dimensional spectral classification, search for spectroscopic variability, analyze the velocity field in the star's atmosphere and envelope, and refine its evolutionary status. In Section\\,\\ref{observ} the methods of observation and reduction are described; in Section\\,\\ref{results} we present and analyze the observational data obtained, and in Section\\,\\ref{conclus} we briefly sums up the main results. ",
        "conclusions": "The results of our quantitative spectral classification of BD$-6\\degr$1178 led us to conclude that it is a double-line spectroscopic binary. Both components are F-type stars: F5\\,IV--III\\,+\\,F3V. We measured the heliocentric velocities of both components of the binary at four time moments. We found no grounds to classify BD$-6\\degr$1178 as a post-AGB star. The coordinates of BD$-6\\degr$1178 and its heliocentric distance 450\\,pc allow it to be suspected of a membership in the Ori\\,OB1 association. Thus BD$-6\\degr$1178 may be a young pre-MS object of the Galactic disk."
    },
    "0807/0807.2450_arXiv.txt": {
        "abstract": "We present Keck adaptive optics imaging of the L4+L4 binary \\hdbin\\ along with archival \\HST\\ and Gemini-North observations, which together span $\\approx$70\\% of the binary's orbital period.  From the relative orbit, we determine a total dynamical mass of 0.109$\\pm$0.002~\\Msun\\ (114$\\pm$2~\\Mjup).  The flux ratio of \\hdbin\\ is near unity, so both components are unambiguously substellar for any plausible mass ratio.  An independent constraint on the age of the system is available from the primary \\hdprim\\ (G2V, [M/H]~=~0.0).  The ensemble of available indicators suggests an age comparable to the Hyades, with the most precise age being \\hdage\\ based on gyrochronology.  Therefore, \\hdbin\\ is now a unique benchmark among field L and T~dwarfs, with a well-determined mass, luminosity, and age.  We find that substellar theoretical models disagree with our observations.  (1) Both components of \\hdbin\\ appear to be overluminous by a factor of $\\approx$2--3$\\times$ compared to evolutionary models.  The age of the system would have to be notably younger than the gyro age to ameliorate the luminosity disagreement.  (2) Effective temperatures derived from evolutionary models for HD~130948B and C are inconsistent with temperatures determined from spectral synthesis for objects of similar spectral type.  Overall, regardless of the adopted age, evolutionary and atmospheric models give inconsistent results, which indicates systematic errors in at least one class of models, possibly both. The masses of \\hdbin\\ happen to be very near the theoretical mass limit for lithium burning, and thus measuring the differential lithium depletion between B and C will provide a uniquely discriminating test of theoretical models.  The potential underestimate of luminosities by evolutionary models would have wide-ranging implications; therefore, a more refined age estimate for \\hdprim\\ is critically needed. ",
        "introduction": "More than a decade after their discovery, brown dwarfs continue to offer key insights into the astrophysics governing of some of the lowest temperature products of star formation.  Brown dwarfs in the field are particularly useful as probes of very cold atmospheres.  For instance, the atmospheres of extrasolar planets are very difficult to study directly due to their intrinsic faintness and proximity to very bright stars.  However, brown dwarfs are typically found in relative isolation, and their atmospheres are subject to the same processes (e.g., dust formation and sedimentation, and non-equilibrium molecular chemistry) that are at work in their much less massive plantary counterparts.  Despite the broad relevance of brown dwarfs, their fundamental properties remain poorly constrained by observations.  In particular, very few direct mass measurements are available for brown dwarfs.  To date, a total of six objects have been identified as unambiguously substellar \\citep[$M<0.072$~\\Msun\\ at solar metallicity;][]{2000ARA&A..38..337C} via dynamical mass measurements with precisions ranging from 6--9\\%: both components of the T5.0+T5.5 binary 2MASS~J15344984$-$2952274AB \\citep{2008arXiv0807.0238L}, both components of the young M6.5+M6.5 eclipsing binary 2MASS~J05352184$-$0546085 in the Orion Nebula \\citep{2006Natur.440..311S}, and two tertiary components of hierarchical triples in which the primaries are M stars, GJ~802B \\citep{gl802b-ireland} and Gl~569Bb \\citep{2001ApJ...560..390L, 2004astro.ph..7334O, 2006ApJ...644.1183S}.  Direct mass measurements of brown dwarfs are critical for empirically constraining substellar evolutionary models.  Since brown dwarfs have no sustainable source of internal energy, they follow a mass--luminosity--age relation, rather than the simpler mass--luminosity relation for main-sequence stars. Thus, mass measurements alone cannot fully constrain theoretical models, although mass and luminosity measurements of brown dwarfs in coeval binary systems can offer stringent tests of theoretical models \\citep[e.g.,][]{2008arXiv0807.0238L}.  To fully constrain evolutionary models, systems with independent measurements of the mass, age, and \\Lbol\\ (or one of the much less observationally accessible quanities \\Teff\\ or $R$) are required.  Such systems are quite rare, but they represent the gold standard among ``benchmark'' brown dwarfs.  \\citet{2002ApJ...567L.133P} discovered the L~dwarf binary \\hdbin\\ in a hierarchical triple configuration with the young solar analog \\hdprim\\ (G2V) using the curvature adaptive optics (AO) system Hokupa`a on the Gemini North Telescope on 2001 February 24 UT.  The L~dwarfs are separated from each other by $\\lesssim$0$\\farcs$13, and they lie 2$\\farcs$6 from the primary G star.  \\hdbin\\ has been the target of AO-fed slit spectroscopy with NIRSPEC on the Keck II Telescope (1.15--1.35~\\micron) and IRCS on the Subaru Telescope (1.5--1.8~\\micron, 1.95--2.4~\\micron).  \\citet{2002ApJ...567L..59G} used the latter spectra to determine the spectral types of the B and C components, both L4$\\pm$1, via spectral template matching.  These are consistent with the less precise NIRSPEC $J$-band spectral types of dL2$\\pm$2, which are on the \\citet{1999AJ....118.2466M} system, found by \\citet{2002ApJ...567L.133P} for both HD~130948B and C.  We present here a dynamical mass measurement for \\hdbin\\ based on Keck natural guide star adaptive optics (NGS AO) imaging of \\hdbin, as well as an analysis of {\\sl Hubble Space Telescope} (\\HST) and Gemini archival images.  In addition to an independent age estimate (\\hdage, see \\S \\ref{sec:age}), the primary star provides a wealth of information about the system.  \\citet{2005ApJS..159..141V} measured a solar metallicity for \\hdprim\\ ([M/H]~=~0.00, [Fe/H]~=~0.05), which is important since metallicity can play a significant role in shaping the spectra of brown dwarfs \\citep[e.g.,][]{2005astro.ph..9066B,2006liu-hd3651b}.  Most importantly, the distance to \\hdprim\\ has been measured very precisely by \\Hipparcos, with a revised parallax of 55.01$\\pm$0.24~mas \\citep{2007hnrr.book.....V}, corresponding to a distance of $d$~=~18.18$\\pm$0.08~pc.  Thus, the distance is measured to an exquisite precision of 0.44\\%, which is invaluable since the error in the dynamical mass scales as 3$\\times$ the distance error (i.e., the 0.44\\% error in distance translates into a 1.3\\% error in mass).   HD~130948BC can thus serve as both an ``age benchmark'' and ``mass benchmark'' system in studying brown dwarfs.  In the literature, the term benchmark is often applied to any readily observable unique or extreme objects, but here we specifically use the term to refer to systems for which fundamental properties may be directly determined. \\citet{2006MNRAS.368.1281P} highlighted the value of systems where the age and composition of substellar objects can be independently determined, e.g., from a stellar or white dwarf companion, and \\citet{2008arXiv0807.0238L} described an equivalent use of systems with dynamical mass measurements.  Essentially, since brown dwarfs follow a mass--luminosity--age relation, the measurement of either mass or age in addition to the measured luminosity allows any other quantity to be fully specified using evolutionary models.  This approach can be extremely useful, for example, by offering precise determinations of \\Teff\\ and \\logg, which can then be compared directly to atmospheric models.  Of course, the measurement of mass, age, {\\em and} luminosity offers a direct test of evolutionary models, which is possible with \\hdbin. ",
        "conclusions": "We have determined the orbit of the young L4+L4 binary \\hdbin\\ using relative astrometry of the system spanning 7 years of its 10-year orbital period.  The astrometric measurements and their uncertainties were extensively tested through Monte Carlo simulations.  The fitted orbital parameters and revised \\Hipparcos\\ parallax give a total dynamical mass of 0.109$\\pm$0.002~\\Msun.  The precision in mass is 2\\%, with nearly equal contributions to the uncertainty from the 1.7\\% error in the best-fit orbit and the 1.3\\% error in mass from the \\Hipparcos\\ parallax error.  For any plausible mass ratio, both components of \\hdbin\\ are unambiguously substellar.  \\hdbin\\ has the most precise mass determination for a brown dwarf binary to date.  The primary star \\hdprim\\ offers an independent constraint on the age of the system from various indicators: rotation, chromospheric activity, isochrone fitting, X-ray emission, and lithium depletion. The ensemble of all available age indicators is consistent with an age for \\hdprim\\ similar the Hyades (625~Myr).  For example, its rotation period is inconsistent with ages much younger than the Hyades and its chromospheric activity is inconsistent with ages much older than the Hyades.  Our preferred age estimate is \\hdage, derived from the gyrochronology formalism of \\citet{2007ApJ...669.1167B} and \\citet{mam08-ages}.  With a measured mass, luminosity, and age, \\hdbin\\ provides the first direct test of the luminosity evolution predicted by theoretical models for substellar field dwarfs.  Both the Tucson models \\citep{1997ApJ...491..856B} and Lyon models \\citep[DUSTY;][]{2000ApJ...542..464C} underpredict the luminosities of HD~130948B and C given their masses and age.  The discrepancy is quite large, about a factor of 2 for the Lyon models and a factor of 3 for the Tucson models.  In order to explain this discrepancy entirely, model radii would have to be underpredicted by 30--40\\%.  The age of \\hdprim\\ would need to be $\\approx$0.4~Gyr younger than we have estimated in order to resolve this discrepancy.  This is inconsistent with the preferred gyro age but can be accommodated by other age indicators; a more refined age estimate for \\hdprim\\ is critically needed.  Since the mass of \\hdbin\\ is more precisely determined than its age, we have used the mass with the individual bolometric luminosities to infer all other properties (age, \\Teff, etc.) from evolutionary models.  We use a Monte Carlo approach to compute model-inferred quantities, and we are careful to account for covariance between the observational errors, the most notable of which is the correlation of the luminosities of the two components through their measured flux ratio.  Because we use mass and \\Lbol\\ to derive model-inferred properties, any potential systematic errors in luminosity evolution will be reflected in the model-inferred quantities.  For example, the very precise model-inferred ages for \\hdbin\\ (0.41$^{+0.04}_{-0.03}$~Gyr from Tucson models; 0.45$^{+0.05}_{-0.04}$~Gyr from Lyon models) are self-consistent, but they are inconsistent with the independent age estimate for \\hdprim\\ (\\hdage).  Lacking measured radii for \\hdbin, we have used evolutionary models to derive effective temperatures.  Given the mass and luminosity of each component, evolutionary models predict effective temperatures of $\\approx$1900--2000~K.  Alternatively, given the mass of each component and age of the primary star, evolutionary models predict effective temperatures of $\\approx$1600--1700~K.  (The disagreement between these two temperature ranges is just a reflection of the systematic errors in luminosity evolution.)  Spectral synthesis using atmospheric models gives temperatures of 1700--1800~K for objects of similar spectral type to \\hdbin\\ \\citep[L3--L5;][]{2008ApJ...678.1372C}.  Using evolutionary models and the measured luminosity ratio gives $\\Delta$\\Teff~=~90~K.  Resolved spectroscopy of HD~130948B and C has previously shown that they have indistinguishable spectral types, so this rather large temperature difference may indicate that spectral type does not hold a one-to-one correspondence with \\Teff\\ mid-L~dwarfs, even for two coeval objects. Better spectral types for the two components of \\hdbin\\ are needed to address this apparent discrepancy.  Comparing the different effective temperature determinations for \\hdbin\\ on the H-R diagram shows that the evolutionary models, atmospheric models, and observational data cannot be simultaneously brought into consistency with each other, regardless of the adopted age of the system.  Thus, systematic errors in some combination of the atmospheric and/or evolutionary models are needed to explain the observed discrepancy.  The best current age estimate indicates that both evolutionary and atmospheric models harbor systematic errors. Further evaluation of the disagreement between models and the data requires a refined age estimate for \\hdprim.  Resolved multi-band spectroscopy of \\hdbin\\ is also needed to reduce the uncertainties in the atmospheric model effective temperatures by direct spectral synthesis fitting.  We also find large discrepancies when comparing the observed near-infrared colors of \\hdbin\\ to the Lyon models.  This suggests that using color-magnitude diagrams to infer the properties of field L~dwarfs from evolutionary models will lead to large errors in the resulting quantities (e.g., mass and/or age).  For example, if we inferred the ages and masses of the components of \\hdbin\\ from the model $J-K$ or $H-K$ color-magnitude diagrams, we would derive masses that are $\\approx$20--30\\% smaller than observed and ages $\\approx$2$\\times$ younger than the age of the primary star (\\hdage).  One novel aspect of using \\hdbin\\ to constrain theoretical models is the application of the ``binary lithium test'', originally proposed by \\citet{2005astro.ph..8082L}.  This is made possible by the fortuitous circumstance that the components of \\hdbin\\ are very near the mass limit for lithium burning.  As a consequence, the Lyon and Tucson evolutionary models, which are almost indistinguishable in their predictions of substellar bulk properties, give very different predictions for the amount of primordial lithium remaining in the B and C components.  Thus, resolved optical spectroscopy to detect the lithium doublet at 6708~\\AA\\ would provide a very discriminating test of the evolutionary models.  Such a constraint is significant in that it directly tests the properties of fully convective substellar interiors (e.g., the core temperature) and/or the lithium reaction rates.  \\hdbin\\ is the only system currently known for which such an empirical calibration of lithium burning is possible.  Substellar theoretical models are in sore need of empirical validation as they have been employed for more than a decade to interpret observations of field dwarfs.  Given the independent constraints on the age and composition provided by a stellar companion, dynamical mass measurements for triple systems like HD~130948ABC provide the most challenging tests of substellar theoretical models.  However, substellar companions to stars are quite rare \\citep[$\\approx$1$\\pm$1\\%, e.g.,][]{2001AJ....121.2189O, 2004AJ....127.2871M, 2005AJ....130.1845L, 2007ApJS..173..143B, 2007ApJ...670.1367L}, and even more rare are substellar binary companions that yield dynamical mass measurements in a reasonable time frame.  When the stellar companion is a bright star like \\hdprim, a wealth of additional information is available, the most important of which is a very precise \\Hipparcos\\ distance measurement since this is the limiting factor in the precision of the dynamical mass.  Stars bright enough to enable seismological measurements can yield the most stringent (10--20\\%) age determinations possible \\citep[e.g.,][]{2005A&A...434.1085C,2008ApJ...673.1093B}.  Thus, \\hdbin\\ represents a rare class of benchmark systems for which the most precise mass and age determinations are possible.  Our observations of \\hdbin\\ indicate that substellar models currently harbor significant systematic errors.  The potential underestimation of \\Lbol\\ by evolutionary models has far-reaching implications.  For example, such models have been used to determine the low-mass end of the intial mass function and to predict the radii of extrasolar planets.  Obtaining measurements for more systems like \\hdbin\\ over a broad range of mass, luminosity, and age will be critical in understanding and resolving the discrepancies that have been revealed between observations and theoretical models."
    },
    "0807/0807.0039_arXiv.txt": {
        "abstract": "We discuss the cosmological degeneracy between the Hubble parameter $H(z)$, the age of the universe and cosmological parameters describing simple variations from the minimal $\\Lambda$CDM model. We show that independent determinations of the Hubble parameter $H(z)$ such as those recently obtained from ages of passively evolving galaxies% , combined with Cosmic Microwave Background data (WMAP 5-years), provides stringent constraints on possible deviations from the $\\Lambda$CDM model. In particular we find that this data combination constrains at the 68\\% (95\\%) \\textit{c.\\,l.}\\,\\,the following parameters: sum of the neutrino masses $\\sum m_\\nu < 0.5\\,(1.0)$ eV, number of relativistic neutrino species $N_{\\rm rel} = 4.1^{+0.4}_{-0.9}\\,(^{+1.1}_{-1.5})$, dark energy equation of state parameter $w = -0.95 \\pm 0.17\\,(\\pm\\,0.32)$, and curvature $\\Omega_k = 0.002 \\pm 0.006\\,(\\pm\\,0.014)$\\,, in excellent agreement with dataset combinations involving Cosmic Microwave Background, Supernovae and Baryon Acoustic Oscillations. This offers a valuable consistency check for systematic errors. ",
        "introduction": "The recent measurements of cosmic microwave background (CMB) anisotropies and polarization \\cite{spergelwmap2,dunkleywmap08}, alone or in combination with other cosmological data sets,  have provided  confirmation of the standard cosmological model and an accurate determination of some of its key parameters.  In particular, the new determination of the age of the Universe $13.68 \\pm 0.13$ Gyrs improves by an order of magnitude previous determinations from, e.g., cosmochronology  of long-lived radioactive nuclei \\cite{thorium} and population synthesis of the oldest stellar populations \\cite{Jimenez96,dunlop96,chaboyerkrauss03} and by a factor of 2 previous determinations from CMB data.  With cosmological parameters so tightly constrained within the framework of the standard flat-$\\Lambda$CDM model, it is important however to constrain possible deviations from the standard cosmological model. Beyond the primordial parameters describing the shape of the primordial power spectrum and late-time parameters such as the optical depth to the last scattering surface, CMB observations so far constrain directly parameters such as \\cite{kosowsky} the angular size distance to last scattering combined with the sound horizon at decoupling, the baryon-to-photon ratio and the redshift of matter radiation equality. This implies that, for models beyond the standard flat-$\\Lambda$CDM,  CMB data alone still show large degeneracies among ``derived\"\\footnote{Note that the name \"derived parameters\" has been sometimes used in the literature with a slightly different emphasis, denoting parameters such as the bias $b$, $\\sigma_8$, etc..} cosmological parameters such as  the matter density parameter $\\Omega_m$, the curvature $\\Omega_k$,  the dark energy equation of state paramenter $w$,  the  effective number of relativistic neutrino species $N_{\\rm eff}$,  the sum of neutrino masses $\\sum m_{\\nu}$ and  the Hubble parameter  $H_0$. For example  (see e.g., \\cite{Eisensteinwhite04, debernardisetal08, de Bernardis:2007bu}),  departures from the standard model described by a deviation from  $3$ neutrino species, can arise from  the decay of dark matter particles \\cite{bonometto98, lopez98,  hannestad98, kaplinghat01}, quintessence \\cite{bean}, exotic models \\cite{unparticles}, and additional hypothetical relativistic particles. This affects the matter-radiation equality yielding, even for a flat, cosmological constant-dominated  model,  a degeneracy between $N_{\\rm eff}$, $H_0$ and  $\\Omega_m$. A departure from dark energy  being described by a cosmological constant (i.e. a component with equation of state $w\\neq-1$), yields a  different   angular size distance to last scattering, and thus degeneracy between $w$, $H_0$ and $\\Omega_m$ even for a flat universe. Finally, relaxing the flatness assumption yields the so-called ``geometric degeneracy\" (between age or $H_0$ and $\\Omega_m$ and  $\\Omega_{\\Lambda}$).  In order to go beyond the concordance $\\Lambda$CDM model parameters determination, one needs extra data sets that probe different physics and are affected by different systematics. In this work we concentrate on the measurements of $H(z)$ using passively evolving red-envelope galaxies and how using them helps to constrain cosmological parameters dropping the assumption of the concordance model.  In particular we show that the recent determinations of $H_0$ from the HST key project \\cite{hstkey} and $H(z)$ provided by \\cite{svj} (SVJ; based on \\cite{data} and references therein) can provide, when combined with CMB and other cosmological data, new and tighter constraints on deviations from the standard $\\Lambda$CDM model, as first shown in \\cite{de Bernardis:2007bu, data}. This approach  of combining different data-sets to constrain parameters that are otherwise poorly constrained,  is called ``concordance approach\".  While it is a very powerful approach, the same ``concordance\" approach is used to test data sets for systematic errors.   It is therefore important to consider enough data sets to  have an  over-constrained  problem and as diverse data sets as possible, relying on different physics and affected by different systematics. Only in this case, if all data sets agree, one can be confident that the systematic errors are safely below the statistical errors and that the cosmological constraints are robust.  After obtaining constraints on deviations from the simple $\\Lambda$CDM model obtained with WMAP 5-years data and $H(z)$ measurements, we compare them with those obtained from the combination of WMAP 5-years with Baryon Acoustic Oscillations and Supernovae data. We find good agreement between the two approaches. We conclude that any possible systematic effect in the non-CMB data sets is below the statistical errors, and that there is no evidence for a deviation from the % flat-$\\Lambda$CDM model, thus offering  support to the standard cosmological model. ",
        "conclusions": "Measurements of $H(z)$ constrain the age of the Universe at different redshifts and thus break the CMB-only degeneracy between the age of the universe today ($t_0$) and the parameters describing deviations from the $\\Lambda$CDM model. As shown in table~\\ref{tab:ages}, the age of the Universe as constrained by WMAP5-only data, is very sensitive to the presence of some of these parameters, especially to the possibility of having a background of $N_{\\rm eff}$ relativistic particles (with $N_{\\rm eff}$ not fixed to $3.04$) or allowing a non-zero curvature (see first row, 3rd and 5th columns of table~\\ref{tab:ages}). This is because many models which deviates from the standard $\\Lambda$CDM but are consistent with CMB data, are not a good fit to the $H(z)$ data. Some illustrative examples are shown in Figure 1. The combination WMAP5+H significatively reduces the degeneracy between $t_0$ and some of the ``extra\" parameters, thus improving the constraints on the age of the Universe in those models, for example, by almost a factor of 3 for the case with $N_{\\rm eff}$ not fixed to 3.04, and almost a factor of 5 for the case of non-zero curvature (see same columns but second row, in table~\\ref{tab:ages}).  \\begin{table}[htb] \\caption{\\label{tab:ages} Determination of the age of the Universe ($68.3 \\%$ c.l.) in several different cosmological models for WMAP 5-years data alone and WMAP5+H data. } \\begin{center} \\begin{tabular}{|c||c|c|c|c|c|} \\hline AGE (Gyr)    & $\\Lambda$CDM & $\\Lambda$CDM+$N_{\\rm eff}$ & $\\Lambda$CDM +$\\sum m_{\\nu} $  & $\\Lambda$CDM+$\\Omega_k$ & wCDM \\\\ \\hline & & & & &\\\\ WMAP 5-years  & $13.69 \\pm 0.13$ & $12.08 \\pm 1.29$ & $14.06\\pm 0.27$  &$16.32 \\pm 1.76$   & $13.74 \\pm 0.34$    \\\\ & & & & &\\\\ WMAP5+H  & $13.65^{+0.14}_{-0.10}$  &  $12.87^{+0.61}_{-0.31}$  &  $13.81^{+0.24}_{-0.14}$  &  $13.61^{+0.29}_{-0.44}$ &  $13.67^{+0.24}_{-0.08}$  \\\\ & & & & &\\\\ \\hline \\end{tabular} \\end{center} \\end{table}  In Figure \\ref{lcdm_mnu} we explore the resulting constraints on the neutrino properties. In all cases the dotted line is the WMAP 5-years only result, the dashed line is WMAP5+HST and the solid line is WMAP5+H. We find that the combination WMAP5+H constrains the sum of neutrino masses to be $\\sum m_{\\nu} < 0.48$ eV and $< 0.93$ eV at the 68\\% and 95\\% confidence levels, respectively, thus improving the WMAP-only constraints by 50\\%. The constraints on the effective number of neutrino species is $N_{\\rm rel} = 4.1^{+0.4}_{-0.9}\\,(^{+1.1}_{-1.5})$ at the 68\\% (95\\%) confidence level and $N_{\\rm rel} > 2.2 $ at better than the 99\\% confidence level. For comparison, note that Ref.~\\cite{de Bernardis:2007bu} obtained $N_{\\rm rel} = 4.0 \\pm 1.2$ at the 68\\% confidence level and $N_{\\rm rel} > 1.8$ at better than 99\\% confidence level, from WMAP 3-years data combined with the same $H(z)$ measurements we use in this paper. Therefore, the improvement we obtain in the constraints of $N_{\\rm rel}$, simply reflects the improvement of the CMB alone constraints in the WMAP 5-years data (compare e.g. the age constraints in table 1 of \\cite{de Bernardis:2007bu} and our table 1). Note that our results are compatible (within the 1-2$\\sigma$ regions) with the constraints on $N_{\\rm eff}$ obtained~\\cite{SimhaSteigman} from BBN alone or combined with other data sets, as well as have statistical errors of the same order (see table 1 from ~\\cite{SimhaSteigman}).   \\begin{figure}% \\hspace{-0.55cm}\\includegraphics[width=8cm, angle=0]{lcdm+curv_WMAP5.eps} \\hspace{-0.1cm}\\includegraphics[width=8cm, angle=0]{cdm+wlam_WMAP5.eps} \\caption{\\textit{Left}: Constraints on the curvature of a $\\Lambda$CDM model from WMAP 5-years alone (dotted line), WMAP5+HST (dashed line) and WMAP5+H (solid line). With the $H(z)$ measurements, the curvature is constrained to 0.002 $\\pm$ 0.006 ($\\pm$ 0.014) at the $68\\%$ ($95\\%$) confidence level. The WMAP 5-years only line shows the well known geometric degeneracy. \\textit{Right}: Constraints on the  dark energy equation of state parameter from WMAP 5-years alone (dotted line), WMAP5+HST (dashed line) and WMAP5+H (solid line). With the $H(z)$ measurements we obtain $w = -0.95 \\pm 0.17\\,(\\pm\\,0.32)$ at the 68\\% (95\\%) confidence level.} \\label{fig:olcdm} \\end{figure}    The left panel of Figure \\ref{fig:olcdm} shows the constraints on the geometry of the Universe. The WMAP5+H combination yields  $\\Omega_k=0.002 \\pm 0.006\\,(\\pm\\,0.014)$ at the 68\\% (95\\%) confidence levels, thus breaking the geometric degeneracy.  In  the right panel of Figure \\ref{fig:olcdm}, we report the constraints on the dark energy equation of state parameter (asssumed constant). The WMAP5+H combination yields $w = -0.95 \\pm 0.17\\,(\\pm\\,0.32)$ at the  68\\% (95\\%) confidence level, which improves the WMAP 5-years only constraints by a factor $\\sim 70\\%$. While the WMAP 5-years constraint has a hard prior on the Hubble constant $H_0<100$ km/s/Mpc which imposes a lower limit on $w$, the WMAP5+HST and WMAP5+H combinations are insensitive to this prior.   In table \\ref{tab:compare} we compare the WMAP5+H constraints  on deviations from the $\\Lambda$CDM model, with those obtained by the combination WMAP 5-years with Baryon Acoustic Oscillation data (BAO) \\cite{percivaletal07} and with Supernovae as obtained by \\cite{komatsuwmap08}.    \\begin{table}[htb] \\caption{\\label{tab:compare} Cosmological constraints  at $68$\\% ($95 \\%$) c.l.\\,\\,on the extra parameters characterizing deviations of the standard $\\Lambda$CDM model% , comparing their values as extracted from WMAP 5-years only, WMAP5+BAO+SN and WMAP5+H.} \\begin{center} \\begin{tabular}{|c||c|c|c|c|c|c|} \\hline Parameter & WMAP 5-years only &  WMAP5+BAO+SN    & WMAP5+H\\\\ \\hline & & &\\\\ $N_{\\rm eff}$ & $> 2.3\\,\\, (95 \\%)$ & $4.4^{+1.5}_{-1.5}\\,^{(*)}$ & $4.10^{+0.37}_{-0.94} (^{+1.12}_{-1.50})$\\\\ & & &\\\\ $\\sum m_\\nu$ & $< 1.3$ eV (95 \\%) & $< 0.61$ eV (95 \\%) & $< 0.93$ eV (95 \\%) \\\\ & & &\\\\ $w$ & $^{> -2.37}_{< -0.68}\\,\\, (95 \\%)$ & $-0.972_{-0.060}^{+0.061} (_{-0.138}^{+0.112})$ & $-0.945_{-0.155}^{+0.194} (_{-0.350}^{+0.311})$\\\\ & & &\\\\ $\\Omega_k$ & $^{< +0.017}_{> -0.063}\\,\\, (95 \\%) $ & $-0.0052^{+0.0064}_{-0.0064} (^{+0.0137}_{-0.0123})$ & $0.002^{+0.0059}_{-0.0059} (^{+0.012}_{-0.018})$\\\\ & & & \\\\ \\hline \\end{tabular} \\end{center} $^{(*)}$with HST prior \\end{table}  \\begin{figure}% \\hspace{-0.0cm}\\vspace{0.2cm}\\includegraphics[width=7.5cm, height=6.2cm, angle=0]{contour_w_omega_k.eps} \\hspace{0.2cm}\\includegraphics[width=7.5cm, height=5.8cm, angle=0]{wmapfig.eps} \\caption{\\textit{Left}: Constraints in the $\\Omega_k$-$w$ plane from WMAP 5-years alone (purple), WMAP5+H (blue 68\\% and 95\\% c.l.). For comparison we show (right) the wmap team's WMAP 5-years only constraints (black), WMAP5+BAO (yellow) and WMAP5+Supernovae (red) (see \\cite{komatsuwmap08}). The differences in the WMAP 5-years only constraints are due to different choice of priors. Most notably, different boundaries on the $H_0$ prior are used: $0.4<h<1$ (left) vs $h<1$ (right), and on $w$ (on the left panel there is an additional  prior $w<-0.3$).} \\label{fig:owcdm} \\end{figure}  Finally, we  consider a model which deviates from the standard $\\Lambda$CDM by two parameters: curvature is allowed to vary and the equation of state of dark energy (allowed to cluster), is assumed to be constant but not fixed to $w=-1$. When running the WMAP 5-years only Markov chain we use different priors from those used in \\cite{dunkleywmap08}, the most important differences being on $h$ and $w$: we use $0.4<h<100$ and $-2.5<w<-0.3$, and a flat prior on the angular size distance to the last scattering surface rather than a flat prior on $\\Omega_{\\Lambda}$. In Figure \\ref{fig:owcdm} we show  how the addition of $H(z)$ data helps break the  degeneracies. For comparison, on the right hand panel of Figure \\ref{fig:owcdm} we show the constraints obtained by  \\cite{komatsuwmap08} from the combination WMAP5+BAO and WMAP5+SN. As already noted by e.g. \\cite{weinberg70, polarski,ClarksonCortesBassett07}, measurements of $H(z)$ are crucial to break degeneracies between the curvature and dark energy properties.  We conclude that the addition of Hubble parameter determinations at different redshifts, break the CMB-only degeneracies arising in models that allow deviations from the simple flat $\\Lambda$CDM model. We find constraints on the number of effective neutrino species, the sum of neutrino masses, the curvature of the universe and the equation of state parameter for dark energy. These constraints are comparable to those obtained from the combination of WMAP 5-years with Supernovae and Baryon Acoustic Oscillations \\cite{komatsuwmap08}. This ``concordance\" approach shows that systematic errors in non-CMB data sets are smaller than the statistical errors and offers further support to the simple flat-$\\Lambda$CDM model. Finally note that future BAO surveys will have the potential to constrain H(z) with \\% accuracy in several z bins, i.e. they will constrain not only dark energy models but other deviations from the standard $\\Lambda$CDM model, as explored in this paper."
    },
    "0807/0807.1515_arXiv.txt": {
        "abstract": "We have computed theoretical models of circumstellar disks for the classical Be stars $\\kappa$ Dra, $\\beta$ Psc, and $\\upsilon$ Cyg. Models were constructed using a non-LTE radiative transfer code developed by \\citet{sig07} which incorporates a number of improvements over previous treatments of the disk thermal structure, including a realistic chemical composition.  Our models are constrained by direct comparison with long baseline optical interferometric observations of the H$\\alpha$ emitting regions and by contemporaneous H$\\alpha$ line profiles. Detailed comparisons of our predictions with H$\\alpha$ interferometry and spectroscopy place very tight constraints on the density distributions for these circumstellar disks. ",
        "introduction": "Classical Be (or B-emission) stars are rapidly rotating, hot stars with optical spectra that show hydrogen emission lines, and, frequently, emission lines from singly ionized metals. It has been recognized since the days of \\citet{str31} that recombination in a flattened disk of circumstellar gas can reproduce the basic features of the spectroscopically observed H$\\alpha$ line profiles of Be stars. However, it is now generally accepted that this model is too simplistic. The detailed mechanism(s) that creates and maintains Be star disks remains unclear \\citep{por03, owo03}, although rapid rotation of the central B star certainly plays a role \\citep{tow04}.  In addition to this lack of a successful dynamical model, the evolutionary status of Be stars and related stars (such as B[e] stars) is not well understood, and possible evolutionary connections between these groups are ambiguous. These problems are compounded by the additional uncertainty in their rotation rates \\citep{cra05, tow04, owo03} and lack of understanding of how rotation and evolution interplay. The suggestion by \\citet{tow04} that the rotation rates of Be stars may be systematically underestimated, combined with new interferometric data, has rekindled interest in finding suitable dynamical models for these stars.  Until recently, the bulk of the observational evidence for the presence of disks surrounding Be stars has been spectroscopic and polarimetric in nature. The spectroscopic evidence includes the profiles of the emission lines observed in the optical and IR spectra, as well as continuum excesses at infrared wavelengths.  The net linear polarization observed in Be stars arises via electron scattering in the non-spherical distribution of gas, as first suggested by \\citet{coy69}.  However, interferometric observations that spatially resolve the disks in several wavelength regimes are becoming available, \\citep[see for example][]{che05, mca05, tyc06}.  Interferometric observations, in concert with other observables, allow the disk physical properties to be determined with greater accuracy and this may contribute to the development of more successful dynamical models.  \\citet{tyc06} collected long-baseline, optical, interferometric observations of the classical Be stars $\\gamma\\;$Cas and $\\phi\\;$Per at the Navy Prototype Optical Interferometer (NPOI) and interpreted these observations using simple models of the intensity distribution on the sky, such as uniform disks, rings, or disks with a Gaussian intensity distribution.  In this paper, we extend the work of \\citet{tyc06} by computing detailed theoretical models for the intensity distribution on the sky produced by the central B star and disk of a Be system. We then compare these models to observations collected at the NPOI.  The theoretical models are computed with the {\\sc bedisk} code developed by \\citet{sig07}. Assuming a density distribution of the disk gas, the code can compute the temperatures in the disk given the energy input from the central stars' photoionizing radiation field, by enforcing radiative equilibrium. As the disk density model contains several adjustable parameters, values for these parameters can then be extracted from the match to observations.  We present new observations and detailed models for the Be stars $\\kappa\\;$Dra, $\\beta\\;$Psc, and $\\upsilon\\;$Cyg. We also use near-contemporaneous observations of the H$\\alpha$ spectral lines for these stars as additional constraints on the density model. As we shall demonstrate, the additional spectroscopic constraint of detailed line profiles is often an important ingredient in selecting among models consistent with the interferometric observations. The overall goal of this study is to place tight constraints on physical conditions in the circumstellar regions of these Be stars. ",
        "conclusions": "In this work, disk density models were obtained for the classical Be stars $\\kappa\\;$Dra, $\\beta\\;$Psc, and $\\upsilon\\;$Cyg by matching the observed interferometric H$\\alpha$ visibilities with Fourier transforms of synthetic images produced by theoretical models. It was demonstrated that the additional constraint obtained by the (co-temporal) observed H$\\alpha$ line profile was critical to selecting a density model. This was particularly evident in the case of $\\kappa$ Dra. We have demonstrated that in the case of these 3 stars, only by additional observational constraints can the degeneracy in the interferometric best-fitting models be removed.  The best fit models to both the interferometric observations and the observed H$\\alpha$ line profiles for the three stars are summarized in Table~\\ref{summary}. The range of $n$ is consistent with previous determinations based on other diagnostics (such as IR excess see -- \\citet{wat86}).  We note that the range of ``base\" densities, $\\rho_o$, varies by over two orders of magnitude between the three stars.  \\begin{table} \\begin{center} \\caption{Summary of Best-Fit Models} \\label{summary} \\begin{tabular}{lcr} \\tableline\\tableline Star&$n$& $\\rho_o$\\\\ &&g/cm$^3$\\\\ \\tableline $\\kappa\\;$Dra&4.2& $1.5 \\times 10^{-10}$\\\\ $\\beta\\;$Psc&2.1&$3.0 \\times 10^{-12}$\\\\ $\\upsilon\\;$Cyg&4.0&$8.0 \\times 10^{-10}$\\\\ \\tableline \\end{tabular} \\end{center} \\end{table}"
    },
    "0807/0807.1665_arXiv.txt": {
        "abstract": "We present the first results of a study where we determine the metallicity distribution function in the Galactic disks as a function of height above the Galactic plane. Observations in the Str\\\"{o}mgren photometric system enables us to identify the dwarf stars and derive metallicities for them. The resulting metallicity distribution functions at 0.5 and 2.0 kpc above the Galactic plane are significantly broader and more metal-rich than is anticipated from standard models such as the Besan\\c{c}on model.  Our results can be explained by invoking a smaller scale height and larger local normalisation for the thick disk than is commonly used in the models. These results are compatible with recent determinations of the thick disk scale height based e.g. on SDSS data.  The age of the stellar populations as a function of height above the Galactic plane is also investigated by studying the turn-off colour and metallicity.  We tentatively find that at 2.0 kpc above the Galactic plane there exist an intermediate age population. ",
        "introduction": "The Milky Way galaxy hosts two stellar disks.   The properties of the disks appear to differ in that the thick disk is older and more metal-poor than the thin disk. Much of our knowledge of these properties for the disks is acquired through the study of local stellar samples (e.g. \\cite[Bensby et al. 2005]{Bensby05}, \\cite[Fuhrmann 2008]{fuhrmann08}, \\cite[Reddy et al. 2006]{reddy06}).   The stellar sample studied here lies along the line-of-sight towards the Draco dwarf spheroidal galaxy, located at $l=86.4$ and $b=34.7$, Fig. \\ref{figure1}. Along this line-of-sight the Milky Way disk stars are approximately at the same galactocentric distance as the sun. This means that we are fairly insensitive to any radial gradients that might exist in the disk, and can assume that any gradients found are basically due to changes in the properties of the stellar populations in the $z$-direction.  This study uses the Str\\\"{o}mgren photometric system which is made up of four intermediate band filters ($uvby$) that are strategically located so that information can be obtained about the amount of line blanketing and the size of  the Balmer jump.  \\cite[Str\\\"{o}mgren (1963)]{Stromgren63} devised two indices, $m_1$ and $c_1$, that measures these effects. While the $m_1$ index is metallicity sensitive, the $c_1$ index is sensitive to surface gravity and we use it to separate distant giant stars, belonging to the dwarf spheroidal galaxy, from Galactic dwarf stars and sub-giant stars, belonging to the thin disk, thick disk, and halo, respectively.  \\begin{figure}[t] \\begin{center} \\scalebox{0.42}{ % \\includegraphics{Plots/dlos11.eps} \\vspace*{-1.0 cm} } \\caption{Schematic figure showing the observed line-of-sight (LOS).} \\label{figure1} \\end{center} \\end{figure} ",
        "conclusions": "Using Str\\\"{o}mgren photometry to derive the stellar parameters of dwarf stars in the Milky Way along the line-of-sight towards the Draco dwarf spheroidal galaxy we have studied the properties of the stellar populations in the Milky Way disks and find that  \\begin{itemize} \\item The metallicity distribution functions at 0.5 kpc and 2.0 kpc above the Galactic plane are broader than predicted by the Besan\\c{c}on model and the observed metallicity gradient as a function of height above the Galactic plane is not as steep as predicted. \\item Preliminary investigations, using a simple toy model, indicates that our observations favour models with smaller thick disk scale heights and higher local normalisation such as found by e.g. \\cite[Chen et al. (2001)]{Chen2001}. \\item By studying the turn-off colour and metallicity as a function of height above the Galactic plane, we find an indication of an intermediate age population located 2 kpc above the Galactic plane. \\end{itemize}  Further modelling of the data is being undertaken, but we believe that these preliminary results will hold."
    },
    "0807/0807.2793_arXiv.txt": {
        "abstract": "{Analytic methods to investigate periodic orbits in galactic potentials.} {To evaluate the quality of the approximation of periodic orbits in the logarithmic potential constructed using perturbation theory based on Hamiltonian normal forms.} {The solutions of the equations of motion corresponding to periodic orbits are obtained as series expansions computed by inverting the normalizing canonical transformation. To improve the convergence of the series a resummation based on a continued fraction may be performed. This method is analogous to that looking for approximate rational solutions (Prendergast method).} {It is shown that with a normal form truncated at the lowest order incorporating the relevant resonance it is possible to construct quite accurate solutions both for normal modes and periodic orbits in general position.} {} ",
        "introduction": "In his book on Dynamical Astronomy, Contopoulos (\\cite{c1}) encourages to investigate higher-order versions of the Prendergast (\\cite{prende}) method to solve non-linear differential equations. The original method was applied by Contopoulos \\& Seimenis (\\cite{contos}, hereafter CS90) to periodic orbits in the logarithmic potential and consists in approximating the exact solution with rational trigonometric functions. Even though the trigonometric series used in the rational approximation are truncated {\\bf at} the first non-trivial order, in CS90 is shown that the quality of the fit to the exact result is quite good over a wide range of energy and ellipticity. On this basis it is natural to presume that higher-order truncations would improve the quality of the prediction.  However, even the simplest version of the Prendergast (\\cite{prende}) method has two problematic aspects: 1) the choice of the dominant harmonic in the trigonometric series has to be made on the basis of some knowledge about the orbit type under study; 2) the determination of the coefficients in the series, which depend on the parameters of the system and on initial conditions, originates from a non-linear algebraic system the solution of which must in general be performed numerically. This second aspect spoils the approach of much of its simplicity, even more if we aim at higher order truncations and consider the growth of the number of unknown coefficients.  In this paper we would like to explore the link between the Prendergast-Contopoulos approach and the approximation of orbital solutions found with a resonant normal form. The motivation for this study stems from the idea of rooting a simplified version of the rational solution method into the frame of a modified normalization algorithm in order to devise a completely analytical approach. In fact it has recently proposed (Pucacco et al. \\cite{PBB}) to exploit a resummation technique based on continued fractions to speed up the convergence of series obtained in the framework of normal form perturbation theory. This technique is able to extend the quality of predictions concerning the instability of normal modes and consequent bifurcations of families of boxlets (Belmonte et al. \\cite{BBP2}).  In analogy with  CS90 we apply this approach to investigate periodic orbits in the logarithmic potential (Binney \\& Tremaine \\cite{BT}). We find analytical solutions of the equations of motion for the normal modes and the main low-order boxlets (`loops' and `bananas'). By inverting the normalizing transformation of coordinates, these solutions are either in the form of standard truncated power series or in a rational form constructed by a continued fraction truncated at the same order of the series. Knowing the `normal form' approximating the system under study, the procedure of creating those solutions is straightforward and does not require any further approximation or numerics.  We show that the analytic rational solutions obtained in this way offer a degree of reliability comparable, where data are available, to those of the semi-analytic treatment based on the Prendergast-Contopoulos approach. Both loops and bananas are quite well reconstructed in shape and dimension. We extend the analysis in CS90 also to check the energy conservation along the boxlets: it turns out that, whether for normal modes energy is conserved within a few percent, for loops and bananas, at this level of approximation, it is not easy to go below 10\\%.  The plan of the paper is as follows: in Section 2 we briefly recall the method to construct normal forms for the logarithmic potential, relegating to the Appendix the explicit expressions of the 1:1 and 1:2 Hamiltonian and generating function. In Section 3 we analyze the approximation of the major-axis orbit and in Sections 4 and 5 we do the same respectively for the loop and banana families. In Section 6 we sketch the conclusions. ",
        "conclusions": "the two methods are almost equivalent for what concerns the precision of the analytic prediction when performed to the same order (Contopoulos \\& Seimenis \\cite{contos}). However, the normal form perturbation expansions, even if computationally heavy, are completely algorithmic and analytic at every stage, whereas the explicit evaluation of the coefficients in the rational expansions require the numerical solution of non-linear systems. We have shown that the analytic rational solutions obtained in this way offer a degree of reliability such that both loops and bananas are quite well reconstructed in shape and dimension. We have extended the analysis in Contopoulos \\& Seimenis (\\cite{contos}) also to check the energy conservation along the boxlets: it turns out that, whereas for normal modes energy is conserved within a few percent, for loops and bananas, at this level of approximation, it is not easy to go below 10\\%.  On the theoretical side, the general usefulness of rational solutions can be explained in the light of the better convergence performance of truncated continued fractions. These come into play as a resummation method of the series expansions produced in the usual way in the normalization approach. The generality of this setting allows us to envisage analogous results in the case of higher-order resonances and the corresponding higher commensurable boxlets.  In addition to the formal and algorithmic improvements, we remark on the relevance of this work also in relation with specific problems of galactic dynamics. The study of orbits in non-axisymmetric potentials is usually performed numerically; however, an exhaustive study with conventional integration methods is costly and difficult to interpret (Touma \\& Tremaine \\cite{TT}). The availability of simple and accurate analytical recipes can be quite useful in several contexts in which periodic orbits and boxlets play an important role: we mention the study of the parameter space of non-axisymmetric discs (Zhao, Carollo \\& de Zeeuw \\cite{ZZ}, Zhao \\cite{Zhao}) and that of the orbit structure around massive black holes in galactic nuclei (Sridhar \\& Touma \\cite{ST}). Even more promising seems to be the possibility of getting accurate solutions for periodic orbits in the triaxial case with and without rotation, for which the analysis is still at the level of the first-order averaging method applied to the 1:1:1 resonance by de Zeeuw (\\cite{dz})."
    },
    "0807/0807.0934_arXiv.txt": {
        "abstract": "The acceleration of interstellar pick-up ions as well as solar wind species has been observed at a multitude of interplanetary (IP) shocks by different spacecraft. The efficiency of injection of the pick-up ion component differs from that of the solar wind, and is strongly enhanced at highly oblique and quasi-perpendicular shock events. This paper expands upon previous work modeling the phase space distributions of accelerated ions associated with the shock event encountered on day 292 of 1991 by the Ulysses mission at 4.5 AU.  As in the prior work, a kinetic Monte Carlo simulation is employed here to model the diffusive acceleration process. This exposition presents recent developments pertaining to the incorporation into the simulation of the diffusive characteristics incurred by field line wandering (FLW), according to the work of Giacalone and Jokipii. Resulting ion distributions and upstream diffusion scales are presented and compared with Ulysses data.  For a pure field-line wandering construct, it is determined that the upstream spatial ramp scales are too short to accommodate the HI-SCALE flux increases for 200 keV protons, and that the distribution function for \\teq{H^+}  somewhat underpopulates the combined SWICS/HI-SCALE spectra at the shock.  This contrasts our earlier theory/data comparison where it was demonstrated that diffusive transport in highly turbulent fields according to kinetic theory can successfully account for both the proton distribution data near the shock, and the observation of energetic protons upstream of this interplanetary shock, using a single turbulence parameter. The principal conclusion here is that, in a FLW scenario, the transport of ions across the mean magnetic field is slightly less efficient than is required to effectively trap energetic ions within a few Larmor radii of the shock layer, and thereby sustain acceleration at levels that match the observed distributions. This highlights the contrast between ion transport in highly turbulent shock environs and remote, less-disturbed interplanetary regions. ",
        "introduction": "\\label{sec:Introduction} There is bountiful evidence for efficient particle acceleration at collisionless shocks in the heliosphere, and interplanetary (IP) shocks provide interesting and useful test cases for shock acceleration theories.  For observations of energetic ion populations at IP shocks in the pre-Ulysses era, see, for example, \\cite{SvA74,DPK81,Tan88}.  A leading hypothesis is that in these systems, both solar wind and pick-up ions are energized to above 200 keV by the mechanism of diffusive shock acceleration. This process forms the focus of this paper, which employs the kinematic Monte Carlo technique of Ellison and Jones (e.g., \\cite{EJE81,JE91,EBJ95}) to model ion diffusion and convection in the turbulent shock environs, and describes the acceleration of particles that are injected directly from the thermal and pick-up ion populations. Upstream plasma quantities are input from observational data, and downstream quantities are determined using the full MHD Rankine-Hugoniot relations. The Monte Carlo method was employed by Ellison et al. \\cite{EMP90} to perform the first successful theory/data comparison for the quasi-parallel portion of the Earth's bow shock, using proton, \\teq{He^{++}} and other ion distributions from the AMPTE experiment.  A similar theory/data comparison was enunciated for IP shocks in Baring et al. \\cite{BOEF97}, where impressive agreement was found between the simulation predictions and spectral data obtained by the Solar Wind Ion Composition Spectrometer (SWICS) aboard Ulysses, in the case of two highly-oblique shocks observed early in the Ulysses mission.  Such agreement was possible only with the assumption of diffusion fairly near the Bohm limit.  Summerlin \\& Baring \\cite{BS05,SB06} extended this program to explore the role of pick-up ions in the acceleration process at an IP shock detected by Ulysses' SWICS and HI-SCALE instruments at around 4.5 AU, as reported in \\cite{Gloeck94}.  Phase space distributions from the simulations, in the case of large angle scattering, were compared with the data, yielding acceptable fits for the proton populations and upstream flux increases, using standard prescriptions for the injected pick-up ion distribution. This paper further extends our investigative program by exploring whether or not diffusion due to field line wandering can model Ulysses' observations for this shock equally well. ",
        "conclusions": "\\label{sec:conclusion}  This paper has modeled measurements of the phase space distributions for protons observed by the Ulysses instruments SWICS and HI-SCALE in the Day 292, 1991 shock using a Monte Carlo simulation of diffusive shock acceleration. Two diffusion models were employed, large angle scattering appropriate for highly turbulent regions, and one approximating the characteristics of field line wandering as determined by Giacalone and Jokipii \\cite{GJ99}. In both, the injection of pick-up protons dominates that of solar wind protons in this highly oblique shock. It was found that while there was consistency between the LAS model and the data for energetic protons above speeds around 600 km/sec, outlined at length in \\cite{SB06}, the FLW scenario explored here has some difficulty in generating superthermal \\teq{H^+} in the power-law domain with the requisite efficiency.  This is attributed to the comparative lack of cross-field diffusion (\\teq{\\kappa_\\perp /\\kappa_\\parallel \\lesssim 10^{-3}}) in the FLW model, which is implemented for variances \\teq{\\sigma^2\\lesssim 0.1 B^2} commensurate with field data for the quiet, interplanetary solar wind.  Because a decline of \\teq{\\lambda/r_g} with ion momentum is implicit in the FLW model, upstream ramp diffusion scales are accordingly very short at 200 keV energies, much shorter than the observed pre-shock scale in HI-SCALE data, unless large values of \\teq{\\eta_0 \\equiv\\lambda_0/r_{g0}\\gtrsim 130} for suprathermal ions are adopted.  Summerlin \\& Baring \\cite{SB06} observed that the LAS model could consistently explain this upstream ramp.  Investigation of the turbulence parameter space in the FLW implementation is ongoing, to assess whether it absolutely requires higher \\teq{\\kappa_\\perp /\\kappa_\\parallel} in order to successfully model the data.  However, the fact that the diffusive LAS model works so well in coupling the spectral and spatial properties suggests that diffusion in highly turbulent fields is an integral aspect of the acceleration process at this shock."
    },
    "0807/0807.0105_arXiv.txt": {
        "abstract": "We present the results of a thorough study of the nucleus of the Sgr dwarf spheroidal galaxy (Sgr dSph) and of the bright globular cluster M54 (NGC~6715) that resides within the same nucleus (Sgr,N).  We have obtained accurate radial velocities and metallicity estimates for 1152 candidate Red Giant Branch stars of Sgr and M54 lying within $\\sim 9\\arcmin$ from the center of the galaxy, from Keck/DEIMOS and VLT/FLAMES spectra of the infrared Calcium II triplet. Using both velocity and metallicity information we selected two samples of 425 and 321 very-likely members of M54 and of Sgr,N, respectively. The two considered systems display significantly different velocity dispersion profiles: M54 has a steeply decreasing profile from $r=0\\arcmin$, where $\\sigma\\simeq 14.2$ km/s, to $r\\simeq 3\\mcnd5$ where it reaches $\\sigma\\simeq 5.3$ km/s, then it appears to rise again to $\\sigma\\simeq 10$ km/s at $r\\sim 7\\arcmin$. In contrast Sgr,N has a uniformly flat profile at $\\sigma\\simeq 9.6$ km/s over the whole $0\\arcmin \\le r\\le 9\\arcmin$ range. Using data from the literature we show that the velocity dispersion of Sgr remains constant at least out to $r\\sim 100\\arcmin$ and there is no sign of the transition between the outer flat-luminosity-profile core and the inner nucleus in the velocity profile. These results - together with a re-analysis of the Surface Brightness profile of Sgr,N and a suite of dedicated N-body simulations - provide very strong support for the hypothesis that the nucleus of Sgr formed independently of M54, which probably plunged to its present position, coincident with Sgr,N, because of significant decay of the original orbit due to dynamical friction. ",
        "introduction": "The occurrence of stellar nuclei at the photometric center of several dwarf elliptical galaxies has been the subject of statistical investigations since the pioneering study by \\citet{bts85} and \\citet{bts87}, and nucleated dwarf ellipticals (dE,N) have become a generally recognized and well studied class of galaxies in their own right \\cite[see][and references therein]{GW,ferbin}. In the last few years the results from systematic studies performed with the instrumentation on board of the Hubble Space Telescope (HST), has revolutionized the field, providing strong evidence supporting the possibility of an intimate connection between the process of nucleation and the process of galaxy formation as a whole. In particular:  \\begin{itemize}  \\item It was generally accepted that stellar nuclei occurred in a significant but minor fraction of dwarf elliptical galaxies \\cite[$\\sim$ 25\\%][]{bts85}; on the other hand, the high resolution HST analysis of a large and well selected sample of Virgo dEs by \\citet[][hereafter C06]{cote6} has revealed that the fraction of dE,Ns ($f_n$) can be as high as 66\\% $\\le f_n\\le$ 87\\%. Strong support to a significant upward revision of $f_n$ comes also from the recent ground-based study by \\citet{grant}. Hence, nucleation appears as the natural status of dEs instead of an exceptional occurrence (see C06 for details and discussion).  \\item An HST study of a large sample of Sa-Sd galaxies by \\citet{bok} found that a similarly high fraction of stellar nuclei is found also among these late type spirals, $f_n\\simeq$ 77\\%. In general, the most recent studies agree in finding a fraction of nucleation larger than 50\\% in any kind of galaxy \\cite[see, for example][and references therein]{carollo,balc,fercot,bok07}, with the only exception of those brighter than $M_B\\simeq -20.3$ mag \\cite[][hereafter FC06]{fercot}. Hence, nucleation seems to occur very frequently in both dwarf and giant galaxies \\citep[see][and references therein, for a deeper discussion of the dwarf/giant dichotomy]{GG}.   \\item Stellar galactic nuclei are found to obey the same scaling relation that links Supermassive Black Holes (SBH) with their host galaxies \\cite[][C06, FC06]{wh}. It has been suggested that ``{\\em ...SBHs and stellar nuclei are nothing but complementary incarnation of Central Massive Objects - they likely share a common formation mechanism and follow a similar evolutionary path...}'' (FC06). In this view, the accumulation of a compact overdensity of baryons at the very center would be a common process during galaxy formation, while its further evolution into a SBH or a stellar nucleus would depend on various ''local'' circumstances.  \\item The similarity between stellar nuclei, the brightest globular clusters and the recently discovered Ultra Compact Dwarf (UCD) galaxies \\citep{ucd} clearly suggests possible links among these classes of stellar systems \\citep{freeman,lotz1,lotz4,bmr94,hase,lucky,bok07}. In particular it has been suggested either that nuclei can form by merging of pre-existing globular clusters \\cite[][and references therein]{lotz1}, or that some objects currently classified as globular clusters may be in fact galactic nuclei, the only relics of the tidal shredding of their host galaxies that once orbited the Milky Way and/or M31 \\cite[see][and references therein]{macsyd,lucky,BS}, or even that most/all globulars were nuclei of Galactic building blocks \\cite[i.e. dwarf galaxies, see][BS, and references therein]{freeman,frebland, bok07}.  \\end{itemize}  Within the scenario briefly outlined above, the crucial relevance of any independent observational fact that may provide some insight on the process of nucleation is clearly manifest. The mechanisms for the formation of stellar nuclei that have been suggested until now belong to two broad classes: (a) the orbital decay and merging at the center of the parent galaxies of pre-existing stellar systems (star clusters), and (b) the accumulation of gas (of various origins) in the very center of a galaxy and its subsequent transformation into stars \\cite[see][for a detailed description of various proposed flavors of these classes of models]{grant}.  At present, all the observational constraints we have in hand come from the study of the integrated properties of samples of distant galaxies/nuclei. The nearest galaxies that are known to host a stellar nucleus are M32, M33 and NGC205, in the Local Group. Even in the most favorable case of NGC205, the structure of the nucleus can be studied in some detail but it can be resolved into individual stars only partially, even with HST imaging \\citep{n205}. In this context the case of the Sagittarius dwarf spheroidal galaxy \\cite[Sgr dSph,][]{iba94} constitutes an absolute {\\em unicum} \\cite[][hereafter M05a]{lnuc}.  The Sgr dSph is a dwarf satellite ($L\\sim 2-5 \\times 10^7 L_{\\sun}$) of the Milky Way \\citep{iba97} whose main body is at $D_{\\sun}=26.3\\pm 1.8$ kpc from us and $R_{GC}\\simeq 18.7$ kpc from the center of the Galaxy \\citep{ltip}. It is devoid of gas \\citep{burton} and its stellar content is dominated by an intermediate-age \\cite[$\\sim 6-8$ Gyr old, see][and references therein]{mic_con} metal-rich \\cite[$\\langle Z\\rangle\\sim$ \\slantfrac{2}{5} $Z_{\\sun}$,][and references therein]{luves} population, while $\\simeq$ 10\\% of old and metal poor stars are also present \\citep{lbhb}. The Sgr dSph is currently being destroyed by the Galactic tidal field \\citep{iba97}. The tidally stripped stars form a huge and coherent tidal tail system that can be observed over the whole celestial sphere \\cite[Sgr Stream, see][and references therein]{ibata01,maj}. There are four globular clusters within the main body of the galaxy \\citep{iba94,DCA,paul} and probably several others that have been stripped from the main body and are now associated with the Sgr Stream \\citep{mic3,mic3b,carra}.  Particularly noticeable, however, is the very bright \\cite[$M_V=-10.0$,][]{harris} metal-poor ($[Fe/H]=-1.55$) \\citep{brown} globular cluster M54 (NGC6715), located exactly at the center of the light distribution of the Sgr galaxy. For this reason, immediately after the discovery of the galaxy it was suggested that M54 could be the actual nucleus of Sgr dSph \\citep{bm95,sl95}. Some doubt on this hypothesis was advanced by \\citet{DCA} based on the circumstance that the integrated color of M54 is much bluer than that of its host galaxy, while the opposite was generally believed to occur in dE,Ns; however \\citet{sl97} noticed that there were exceptions to this rule. In fact, the latest studies suggest that the opposite may be true \\cite[][see also C06]{lotz4}. Since the mean metallicity of M54 and of the population dominating the Sgr galaxy differ by one full dex, their Red Giant Branch (RGB) stars (and in some cases also their Horizontal Branch stars, HB) can be easily discriminated from each other in a Color Magnitude Diagram (CMD); at the same magnitude the RGB of Sgr is much redder than that of M54. Selecting genuine Sgr stars in this way \\citet{ls00} and \\citet{maj} independently found an overdensity of Sgr stars that appear concentric with M54 and have a similar spatial scale. Using the same technique and a very large optical photometry database M05a demonstrated that Sgr has actually a nucleus of metal rich stars (in the following we will refer to this structure as the {\\em nucleus} of the Sagittarius galaxy, {\\em Sgr,N} for brevity, if not otherwise stated). This stellar structure has the same center as M54, to within the uncertainties, but it displays a different surface brightness profile, suggesting a different origin from the cluster. The strong incompleteness in the innermost $\\sim 10\\arcsec$ prevented M05a from obtaining accurate estimates of the structural parameters of the Sgr nucleus: they derived $-10.0 \\le M_V\\le -7.6$, $\\mu_V(0) \\la 18.5$ mag/arcsec$^2$, and, with a tentative fit of a \\citet{king} model $r_c\\la 0.21 \\arcmin$ and $r_t\\sim 17\\arcmin$ (where C=log($r_t/r_c$)$\\sim$ 1.9; we will refine these parameters in Sect.~2, below). They also found that the observed properties of the nucleus of the Sgr galaxy were fully compatible with those of known nuclei of dwarf ellipticals. Moreover, the recent study by \\citet{siegel} has shown that while the stellar budget of Sgr,N is dominated by the same intermediate-age population found in the surrounding galaxy, there is clear evidence of more recent and repeated episodes of star formation \\cite[not observed in extra-nuclear fields,][]{mic_con}, pointing to a history of subsequent phases of re-accumulation of gas followed by star formation bursts, very similar to what is observed by \\citet{rossa} in a large sample of distant galactic nuclei.  In any case, the key result of M05a is that Sgr is a nucleated galaxy independently of the presence of M54; even if one were able to remove the cluster from the galaxy by magic, a nucleus made of typical Sgr stars would still be there. To explain the strict spatial coincidence of M54 and the nucleus, M05a proposed two possibilities: (a) the nucleus formed {\\em in situ} from processed Sgr gas that has fallen to the bottom of the galactic potential well, and M54 was (independently) driven into the same place by dynamical friction; (b) M54 was born in the present place (or was driven there as in case a) and formed a co-spatial overdensity of Sgr stars by capturing them (or the gas from which they formed) from the surroundings. The key difference between the two hypothesis can be summarized by the questions {\\em ``did M54 provide the mass seed around which the metal-rich nucleus of Sgr was assembled? is M54 the main contributor of the mass budget in the innermost $\\sim 100$ pc of the Sgr galaxy?''}, that enclose the basic point of hypothesis (b).  Hypothesis (a) was found to be fully compatible with all the available information, while it was not possible to consider case (b) in more depth. In a recent paper, \\citet{mibau} studied in detail the process of the capture of ``field'' stars from an intervening star cluster and they concluded that, even in the most favourable cases, the fraction of stars captured in a Hubble time is negligible, less that $10^{-4}$ of the cluster mass. The range of conditions studied by \\citet{mibau} includes the ``M54 within Sgr'' case, hence their conclusions are fully applicable here. Nevertheless, the results by \\citet{mibau} do not exclude the possibility that M54 operates as a collector of processed gas at the center of Sgr.  However it seems quite reasonable to expect that if case (b) is true, the stars of the Sgr,N structure should share the same kinematics as M54 stars, or, at least, their kinematics should be compatible with a mass-follows-light model, or, finally, their kinematics should be different from that of extra-nuclear Sgr stars, as they would be bound to M54 and would orbit within its potential \\citep{gil}. Note that M54 and Sgr,N are enclosed within the innermost $10\\arcmin$ of a galaxy whose core radius and limiting radius are as large as $r_c=224\\arcmin\\pm 10\\arcmin$ and $r_t=1801\\arcmin\\pm 112\\arcmin$, respectively \\citep{maj}. Hence, even if the main body of Sgr is clearly undergoing tidal disruption the kinematics of the nuclear region should be unaffected and is consequently expected to provide a trustworthy insight on the inner mass distribution \\citep{munoz}.  To follow-up these ideas and to obtain a deeper insight into this unique case of a galactic nucleus that can be studied {\\em in vivo} on a star by star basis, we performed an extensive spectroscopic survey, principally aimed at the study of the kinematics of Sgr,N and M54. Here we present the main results of this survey. The plan of the paper is as follows. As a preliminary step we re-analyse the Surface Brightness profile of Sgr,N with new data in Sect.~2. In Sect.~3 we describe the observation and data reduction of our spectroscopic survey, we test the reliability and the accuracy of our measures, the observed metallicity distribution is briefly discussed and, finally, the selection of the samples is described in detail. Sect.~4 is devoted to the kinematical analysis of the M54 and Sgr,N samples and in Sect.~5 we briefly present the results of a suite of N-body simulations aimed at the study of the orbital decay of M54 within the Sgr galaxy. Finally, in Sect.~6 we summarize and discuss our results. \\vskip0.9truecm ",
        "conclusions": "We present the results of a large radial velocity survey of stars in the nucleus of the Sgr galaxy and in the globular cluster M54, which lies at the center of the nucleus itself. From high signal to noise Keck-DEIMOS and VLT-FLAMES spectra, we obtained accurate radial velocities ($\\epsilon_{V_r} \\simeq 2$ km/s) and   metallicities from the Ca triplet for 1152 candidate RGB/RC stars of M54/Sgr,N. Selecting by the position in the CMD, by radial velocity and by metallicity, we obtained two reliable and clean samples of 321  Sgr,N members and 425  M54 stars and we used these samples, further cleaned from $3\\sigma$ velocity outliers, to study the kinematics of the two systems.  In support of this analysis we used archival HST/ACS images to refine the estimates by M05a of the structural parameters of Sgr,N.  We also performed a suite of N-body simulations of the orbital decay (driven by dynamical friction) of a massive cluster into a host dwarf spheroidal galaxy, mimicking the M54 - Sgr system. We focused our attention on the possibility to realize a phase-space match between the infalling cluster and the density cusp of the host NFW halo as close as that observed with M54 and Sgr,N.  The main results of our analyses can be summarized as follows:  \\begin{enumerate}  \\item We have obtained new, much more reliable estimates of the total luminosity and of the typical size of Sgr,N, confirming that it is significantly less luminous than M54 and that it has a different luminosity distribution with respect to the cluster.  \\item The systemic radial velocities of M54 and Sgr,N are identical within $\\sim \\pm 1$ km/s. Coupling this result with those by M05a, it appears that the two systems coincide in phase-space to within the observational errors, at least for what concerns position in space and radial velocity.  \\item The velocity dispersion profile of M54 is in good agreement with the theoretical profile of the King model that best fits its surface brightness, at least over a range of 30 core radii ($r\\le 3.5\\arcmin$), but possibly also beyond this radius (see Sect.~4.2). In particular, the velocity dispersion drops from $\\sigma=14.2$ km/s at $r=0$, to $\\sigma\\simeq 5.3$ km/s at $r\\sim 3.5\\arcmin$. The observed velocity profile strongly suggests that M54 behaves as an ordinary, self-bound globular cluster, at least in its innermost and densest region. In our view, this indicates that even if at the present epoch the motion of its stars is driven by the potential produced by the overall mass distribution within Sgr,N (see below), its velocity profile keeps memory of its original nature, that is an ordinary massive globular cluster orbiting within the Sgr galaxy. A turn-over of the profile to $\\sigma=8.8$ km/s around the tidal radius of the cluster may suggest that some degree of tidal disruption is occurring in the outskirts of the cluster; however the alternative hypothesis that the sample is contaminated by metal-poor Sgr stars at large radius is also viable and seems more likely, at the present stage (Sect.~4.2).  \\item The velocity dispersion profile of Sgr,N is consistent with being flat at $\\sigma\\simeq 10 $ km/s over the whole considered radial range ($r\\le 9\\arcmin$). The profile is hardly compatible with any realistic King model roughly reproducing the observed surface brightness profile of Sgr,N. More generally, {\\em the fact that the velocity profile is flat in a radial range in which the surface density declines by a factor of $>1000$ is strongly uncompatible with any mass-follows-light model} \\citep{gil}. Conversely, a realistic NFW model reasonably reproduces the observed velocity profile and, in principle, does not conflict with the presence of an overdensity of baryons at its center, like Sgr,N. However, it has to be recalled that the NFW model does not make any definite prediction on the SB profile of the embedded stellar nucleus and the compatibility of the observed SB profile of Sgr,N with the NFW model that that fits the velocity profile is not established.  \\item M54 and Sgr,N have definitely different kinematical properties. In particular, the velocity dispersion profiles are very different: in the radial range $1.5\\arcmin < r < 6.5\\arcmin$ the statistical significance of the difference in the velocity distribution is very large.  \\item We have provided observational evidence that the velocity dispersion profile of Sgr remains flat from $r\\simeq 1\\arcmin$ to $r\\simeq 100\\arcmin$ and that there is no apparent transition in the velocity profile corresponding to the onset of the stellar nucleus, at $r\\la 10\\arcmin$. This fact as well as those listed above strongly suggest that Sgr,N and M54 had independent origins, as we would expect that, if the cluster provided the mass seed to collect the Sgr gas that later formed the metal-rich nucleus, Sgr,N stars would have shown a declining velocity dispersion profile, compatible with a mass-follows-light distribution.  \\item Our N-body simulations that follow the orbits of massive clusters (1-2 $\\times 10^6 M_{\\sun}$, representing M54) within different NFW halos (representing the Sgr galaxy) show that for a large range of initial distances and relative velocities, the orbit of the cluster decays completely by dynamical friction within 3 Gyr, at most. Moreover, at the end of the simulations, the cluster is perfectly concentric with the cusp of the host halo (within the resolution of the simulation) and the difference in average velocity is always less than $\\sim 2$ km/s. Hence, the observed phase-space coincidence between M54 and Sgr,N can be naturally explained by the ``dynamical friction hypothesis'' (M05a).  \\item According to FC06 the mass of Central Massive Objects, independently if they are Black Holes or stellar nuclei, is  $\\simeq \\frac{3}{1000}$ of the mass of the host galaxy ($\\frac{M_{CMO}}{M_{gal}}=0.003$). Using our estimates for  the total mass of Sgr,N and M54 we obtain the following estimates for the total mass of the Sgr galaxy, depending on whether we assume is the M54, Sgr,N or the sum of the two  (as we would do if we observed the system at the distance of Virgo cluster) is the Central Massive Object: $M_{Sgr}=1.0\\times10^9~M_{\\sun}$, $M_{Sgr}=2.1\\times10^9 ~M_{\\sun}$, and $M_{Sgr}=3.0\\times10^9~M_{\\sun}$, respectively. These values are in very good agreement with  the independent estimates by \\citet{maj} that ranges from $M_{Sgr}=5.8\\times10^8 ~M_{\\sun}$ to $M_{Sgr}=6.9\\times10^9~M_{\\sun}$. Hence both Sgr,N and M54 singularly or taken together have a mass compatible with being the CMO of the Sgr galaxy.  \\item Both Sgr,N and M54, as well as the combination of the two, when placed in the $M_V$ vs. log$r_h$ diagram lie in a region that is populated by globular clusters {\\em and} galactic stellar nuclei. They are also compatible with the Color-Magnitude relation of nuclei shown by FC06 and with the $M_V$ vs. $\\sigma$ relation satisfied by globular clusters, nuclei and UCDs \\cite[see][]{geha,evsti}, as shown in Fig.~\\ref{mvsigma} Hence the structure and dynamics of Sgr,N, M54 and their combinations are fully compatible with other galactic nuclei.  \\item A detailed study of the mass profile of Sgr,N and M54 using the Schwarzschild method \\cite[see][and references therein]{rix,ven} is currently ongoing and the results will be presented elsewhere (Ibata et al., in preparation).  \\end{enumerate}  These findings lend very strong support to the scenario proposed by M05a to explain the M54/Sgr,N system: the nucleus of the galaxy formed {\\em in situ}, at the bottom of the potential well of the Sagittarius galaxy; the {\\em globular cluster} M54 was independently driven to the same site by dynamical friction. As a complement of the above conclusions, it must be recalled that the stellar population that dominates Sgr,N formed several Gyrs later than M54 \\citep[see][and references therein]{paul,sl97,ls00,mic_con,siegel}. In the present context this does not appear particularly relevant, but it should be kept in mind that M54 can have reached the very center of Sgr {\\em after} Sgr,N formed its stars \\cite[within the last $\\sim 5-9$ Gyr][]{mic_con}, or even {\\em before} this, depending on the birthplace of the cluster within the Sgr galaxy. In any case it is very likely that both the cluster and the processed Sgr gas were {\\em independently} driven - by different mechanisms - to the bottom of the potential well of the Sgr galaxy, i.e. to the center of its Dark Matter halo. A conclusive proof that M54 was driven to its present position by dynamical friction could be provided by the successful detection of genuine extra-tidal stars at large distances from the cluster center as envisaged in Sect.~4.3, while, as said their non-detection will not disprove the above scenario.  As a possible alternative to this view (or as an extreme version of it) it can be conceived that M54 formed at the bottom of the overall potential well of the Sgr galaxy since the beginning. In this framework M54 and Sgr,N are just the results of two subsequent episodes of star formation both occurring at the very center of Sgr, the second from  enriched gas that was infalling on very different orbits with respect to the gas that formed M54, thus resulting in different final stellar kinematics. We regard this possibility as much less likely with respect to the ``dynamical-friction'' scenario depicted above. First, if M54 is considered as part of the (metal-poor) field population of Sgr, its presence would be at odd with the global metallicity gradient observed by \\citet{alard}, \\citet{sdgs2}, \\citet{sl97} and others in Sgr, as well in all dwarf spheroidal galaxy studied to date \\cite[more metal rich and younger populations are preferentially found at the center of dSph's][]{harbeck}. While dE nuclei have been indicated as a possible exceptions to this trend, the ingestion of large metal poor globular clusters in pre-existing metal rich nuclei seems one possible natural way to reconcile the generally observed metallicity gradients and the presence of nuclei that are bluer than their parent galaxy \\citep{lotz4}. Second, as both M54 and Sgr,N would have {\\em formed} from gas falling into the same potential well, the reason for the different kinematics remains to be explaned, while it is a natural outcome if M54 formed elsewhere as a classical globular cluster.  \\begin{figure} \\plotone{mvsigma.ps} \\caption{Sgr,N and M54 (filled triangles) in the $M_V$ vs. $\\sigma$ plane \\cite[see][]{evsti}. Both systems lie on the locus common to globular clusters, dE nuclei, Dwarf-Globular Transition Objects \\cite[DGTO][]{hase}, and, possibly, UCDs. The same is true for their combination (Sgr,N+M54, open triangle), i.e. for a system having luminosity equal to the sum of the luminosities of M54 and Sgr,N and having $\\sigma$ as estimated from an integrated spectrum obtained with a $1\\arcsec$ slit from the distance of the Virgo cluster of galaxies ($\\sigma\\simeq 12.4$ km/s). Data for Galactic globulars are from \\citet[][$M_V$]{macsyd} and \\citet[][$\\sigma$]{tad}; for M31 globulars we took integrated magnitudes from the Revised Bologna Catalog \\citep{rbc} and $\\sigma$ from \\citet{djor}; data for dE and dE,N are from \\citet{geha,geha2}; UCDs data are from \\citet{ucd} and \\citet{evsti}; data for DGTO are from \\citet{hase}. Some remarkable bright clusters have been labeled, for reference \\cite[see][]{lucky}. \\label{mvsigma}} \\end{figure}     \\subsection{The process of galaxy nucleation}  If we take for demonstrated the above concluding remarks, we can ask what we have learned about the process of galaxy nucleation from the case studied here. Concerning the two main mechanisms that have been put forward in the literature, i.e. (a) formation of the nucleus by infall of globular cluster(s) to the center of the galaxy, or (b) {\\em in situ} formation by accumulation of gas at the center of the potential well and its subsequent conversion into a stellar overdensity \\citep[see Sect. 1. and][]{grant,cote6}, the main conclusion that can be drawn from the case of Sgr is that both channels are viable and actually both have been at work ``simultaneously'' in Sgr.  The present analysis has shown that a stellar nucleus made of the typical material that dominates the baryonic mass budget of Sgr is present in this galaxy, independent of M54, as it display the same flat velocity dispersion profile as the whole core of Sgr, much different from that of the cluster.  In a likely scenario, the enriched gas from previous generations of Sgr stars accumulated at the bottom of the overall potential well of the galaxy, until star formation transformed it into a stellar nucleus whose surface brightness is $\\ga 100$ times larger than in the surrounding Sgr core, a substructure within a larger galaxy. On the other hand, M54 is a (relatively) ordinary massive, old and metal poor globular cluster. Independent of its birthplace within the early Sgr galaxy, its mass and the density of the surrounding medium of the host galaxy drove it to the bottom of the Sgr potential well, by dynamical friction. During its trip to the densest central region of Sgr, the dense cluster managed to survive the tidal force of the host galaxy, hence it reached the present position  as a (partially?) self-bound stellar system.  While the mass budget at the center of Sgr is probably dominated by the underlying DM halo, M54 dominates the overall light distribution : an observer taking photometry and/or spectra of the unresolved nucleus of Sgr from a distant galaxy (say, a galaxy in the Virgo cluster) would find that the object looks like a bright and blue globular cluster; the integrated velocity dispersion would not show any peculiar feature revealing the composite nature of the observed nucleus (see Fig.~\\ref{mvsigma}, above). Probably, it would be impossible to disentangle the two systems from the integrated light\\footnote{It would be interesting to check if there is some spectral feature that may reveal the composite nature of the ``system'', in the present case. We plan to do this in the future by combining properly scaled synthetic spectra representing the light output of M54 and Sgr,N.}.   Finally, the Sgr case seems to support the observed ubiquity of nuclei (see Sect.~1. and C06). The Sgr galaxy was able to form a sizable nucleus ``twice'' and with two different formation channels: if either of the two channels had not been viable for some reason, the galaxy would have ended up with a nucleus in any case. The fact that Sgr is the only case of galaxy with a clear nucleus among those classified as dwarf spheroidals in the Local Group may suggest that the progenitor of Sgr was in fact a brighter dE or disc galaxy that has been transformed into a dSph by the interaction with the Milky Way \\cite[][M05b, and references therein]{maj,mayer}   \\subsection{Suggestions for further investigations}  The results presented in this paper suggest several interesting lines of research that we did not follow up for practical reasons. However we feel that it is worth briefly mentioning some of them here, as a possible starting point for future studies.  \\begin{itemize}  \\item The results presented by \\citet{read}, \\citet{goerdt} and \\citet{sanchez}, suggest that the complete decay of a massive cluster to the very center of the host galaxy is much easier and faster if a central density cusp is present. It is even possible that a central cusp is actually {\\em required} to bring a cluster down to the very bottom of the overall potential well. If this is the case, the position of M54 within Sgr,N would support the existence of a central cusp in actual DM halos, a point that has been questioned by several authors \\cite[see][and references therein]{sanchez}. The ``complete infall'' of a massive cluster may need a NFW cusp and, simultaneously, it may transform the cusp into a core by transferring orbital energy and momentum to the surrounding ``medium'', thus possibly providing a self-regulating mechanism that simultaneously prevent the further decay of other clusters \\citep[that, in general cases, would be quite difficult, given the expected decay times, see][and references therein]{herna,oh,read,sanchez,goerdt}. In this context, it may be worth to recall also the work by \\citet{stri}, whose results militates against the presence of a large-size core in the Fornax dSph, and by \\citet{boy} on the resilience of cuspy halos in major mergers. These ideas seems worthy of detailed theoretical follow up.    \\item It has been suggested several times that bright and metal poor globular clusters may be of cosmological origin \\citep[see][and references therein]{BS,sills}. If this is the case they may be very intimately linked to the earliest phases of formation of galaxies and there may have been plenty of opportunities for most of them to become the nuclei of some dwarf galaxy. The tidal stress that they suffered during their infall to the center of their host galaxies may be at the origin of the larger half-light sizes that are observed in those of them that have been suggested as possible nuclear remnants of ancient dwarf galaxies \\citep{frebland,lucky,macsyd}. This kind of scenario may be explored with dedicated N-body simulations, possibly including gas and stars. It would be interesting also to consider in detail the results presented here in relation with the scenario for the origin of globular clusters recently discussed by \\citet{bok07}.  \\item All the analysis presented in this paper have been performed within the standard Newtonian gravitation theory and dynamics. It may be worth to consider the observational scenario emerged from this study also in the framework of Modified Newtonian Dynamics paradigms \\citep[MOND, see][and references therein]{mil,sand}, even if it may not necessarily be an ideal case. The transition between ordinary Newtonian regime and MOND regime occurs around $r\\simeq 4.5\\arcmin - 6.5\\arcmin$, depending on the actual stellar mass of M54+Sgr,N, that is in a range covered by our data.   \\item There is little doubt that the final fate of the Sgr galaxy will be its complete tidal disruption. Once the large scale stellar body of Sgr will be completely dispersed into the Galactic halo, the final remnant of this (once) relatively large galaxy would be a faint nucleus embedding a bright globular cluster. An observer lacking any knowledge of the origin of this object would conclude that it is a very bright and peculiar globular cluster, dominated by a metal poor population ($[Fe/H]\\sim -1.5$, possibly with some spread) but also including a small fraction ($\\sim$ 10\\%) of metal rich stars (with average $[Fe/H]\\sim -0.4$). Also the abundance patterns would appear different: for instance, the metal poor (M54) stars would appear as moderately $\\alpha$-enhanced \\citep{brown}, while metal rich stars (Sgr,N) would have solar or sub-solar $[\\alpha/Fe]$ ratios \\citep{luves}. The radial velocities would reveal that the metal rich and metal poor stars have {\\em the same systemic velocity}, but {\\em different velocity dispersion profiles} and {slightly different density profiles}, possibly with some rotation in the metal poor component. However the dominance of metal-poor (M54) stars would be so high that there would be no hint of the presence of a Dark Matter component. The half-light radius of the system would appear slightly larger with respect to ordinary globulars \\citep{macsyd,lucky}. Most of these likely characteristics of the future remnant of Sgr seem to have a counterpart in the widely studied and mysterious stellar system $\\omega$ Centauri \\citep[see][and references therein]{norris,omega,ele1,ele2,ele3,ven}, that has been proposed as the possible remnant of a nucleated dwarf elliptical since \\citet{freeman}. While there are also noticeable differences between the two cases, the analogy seems very intriguing and potentially powerful\\footnote{It is interesting to recall that $\\omega$ Cen shares with Sgr some remarkable chemical peculiarities. In particular, stars of comparable metallicity in the two systems are strongly enhanced in s-process elements and the ratio of heavy s-process to light s-process elements [hs/ls] is very similar. Furthermore, $\\omega$ Cen and Sgr are, up to now, the only stellar systems known to have deficient [Cu/Fe] ratios with respect to the trend in the Galactic disk and halo \\citep{mcw}}. It is possible that at least some of the observational features of $\\omega$ Centauri that appear so difficult to explain may find their natural place within a scenario like the one described above.   \\end{itemize}"
    },
    "0807/0807.4510_arXiv.txt": {
        "abstract": "The energy region spanning from $\\sim 10^{17}$ to $\\lesssim 10^{19}$ eV is critical for understanding both, the Galactic and the extragalactic cosmic ray fluxes. This is the region where the propagation regime of nuclei inside the Galactic magnetic environment changes from diffusive to ballistic, as well as the region where, very likely, the most powerful Galactic accelerators reach their maximum output energies.  In this work, a diffusion Galactic model is used to analyze the end of the Galactic cosmic ray spectrum and its mixing with the extragalactic cosmic ray flux. In particular, we study the conditions that must be met, from the spectral and composition points of view, by the Galactic and the extragalactic fluxes in order to reproduce simultaneously the total spectrum and elongation rate measured over the transition region by HiRes and Auger. Our analysis favors a mixed extragalactic spectrum in combination with a Galactic spectrum enhanced by additional high energy components, i.e., extending beyond the maximum energies expected from regular supernova remnants. The two additional components have mixed composition, with the lowest energy one heavier than the highest energy one. The potential impact on the astrophysical analysis of the assumed hadronic interaction model is also assessed in detail. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0807/0807.1932_arXiv.txt": {
        "abstract": "A growing number of observations indicate that magnetic fields are present among a small fraction of massive O- and B-type stars, yet the origin of these fields remains unclear.  Here we present the results of a VLT/FORS1 spectropolarimetric survey of 15 B-type members of the open cluster NGC 3766.  We have detected two magnetic B stars in the cluster, including one with a large field of nearly 2 kG, and we find marginal detections of two additional stars.  There is no correlation between the observed longitudinal field strengths and the projected rotational velocity, suggesting that a dynamo origin for the fields is unlikely.  We also use the Oblique Dipole Rotator model to simulate populations of magnetic stars with uniform or slightly varying magnetic flux on the ZAMS.  None of the models successfully reproduces our observed range in $B_\\ell$ and the expected number of field detections, and we rule out a purely fossil origin for the observed fields. ",
        "introduction": "Magnetic fields have long been inferred in solar and late-type stars and are produced by a dynamo mechanism in the outer convection zones.  Studies of hot star structure and evolution have long neglected magnetic fields since the radiative envelopes of hot stars are not expected to produce them.  Yet, an increasing number of observations have revealed that magnetic fields are present in $\\sim 5$\\% of hot stars \\citep{petit2008}, even though the origin of such fields is uncertain \\citep{neiner2007}.  Various models have proposed that they may be generated in the convective cores, although no known mechanism can transport them to the surface in a timescale consistent with observations \\citep{charbonneau2001, macgregor2003}.  Other models have proposed that the radiative envelopes somehow maintain a dynamo process near the stellar surface (e.g.\\ \\citealt{spruit2002}; \\citealt{macdonald2004}).  Finally, magnetic fields in hot stars may be a relic from stellar formation, hence the present day field is a ``fossil'' of the primordial field.  A growing number of studies have begun to investigate magnetic fields among massive stars in open clusters.  Main-sequence (MS) and pre-main sequence (PMS) magnetic OB stars have recently been detected in the Orion Nebula Cluster \\citep{petit2008}, NGC 6611, and NGC 2244 \\citep{alecian2008} with a variety of ages, rotational periods, and chemical abundances.  Currently, the most widescale investigation of magnetic, early-type cluster members considers 258 A- and B-type stars in more than 40 open clusters \\citep{bagnulo2006, landstreet2008}.  These studies offer strong potential to infer the origin of the magnetic fields in massive stars.  If the magnetic fields are relics of the molecular cloud from which the stars formed, then a population of coeval OB stars should offer widespread evidence of the initial field.  A past or present dynamo likely depends on other physical conditions within the stars, especially rotation, and would likely be less widespread among cluster members.  Observations currently suggest a fossil origin for the magnetic fields in massive stars.  \\citet{landstreet2008} find that the strongest fields are associated with stars near the zero-age main-sequence (ZAMS) and that the field strength decays over time. For all but the strongest fields, the decay is consistent with magnetic flux conservation as the stars increase in radius during the MS lifetime.  \\citet{landstreet2008} also find that some of the strongest fields are found in the most slowly rotating stars, and they find no correlation between rotation and magnetic field strength as expected for a magnetic dynamo in hot stars.  However, there is no evidence that widespread magnetic braking contributes to the loss of angular momentum in B-type stars.  The rotational velocity of B stars has been observed to decline with age \\citep{huang2006}, but no more than expected for their increase in radius during the MS lifetime and angular momentum losses due to stellar winds.  The open cluster NGC 3766 is between 14.5 and 25 Myr in age (WEBDA; \\citealt{moitinho1997}; \\citealt{tadross2001}), with a reddening $E(B-V) = 0.22 \\pm 0.03$ and at a distance of 1.9 to 2.3 kpc \\citep{mcswain2008}.  In our spectroscopic study of this cluster \\citep{mcswain2008}, we measured projected rotational velocities, $V \\sin i$, for 37 members.  We also measured the effective temperatures, $T_{\\rm eff}$, and polar surface gravities, $\\log g_{\\rm polar}$, for 42 cluster members.  More recently, we have analyzed spectra of additional cluster members using the same technique, and these will be presented in a forthcoming work (McSwain et al., in prep).  In these works, we have identified several B-type members with unusually high or low helium abundances that are excellent candidates for the presence of magnetic fields.  Since NGC 3766 offers a number of likely magnetic field candidates as well as a chance to study the rotational and evolutionary dependence of hot star magnetic fields, it is an excellent target to search for magnetic B stars.  For these reasons, we have conducted a spectropolarimetric survey of 15 B-type stars in NGC 3766.  In this work, we present the definite detection of two magnetic stars in the cluster, including one strong field of almost 2 kG.  We also identify two stars with marginal magnetic field detections that are worth further investigation.  The observations and magnetic field measurements are discussed in Section 2.  In Section 3, we discuss the possibility that the stellar magnetic field strength is correlated with rotation and/or MS evolution.  We also present a method to test whether the observed magnetic field distribution is consistent with a fossil field origin.  Our conclusions are summarized in Section 4. ",
        "conclusions": "We have detected 2 definite and 2 possible magnetic B-type stars in our survey of 15 members of the open cluster NGC 3766.  Both of our definite $B_\\ell$ detections are found in stars with abnormal He abundances, and we find a marginal $B_\\ell$ signature in two other cluster members that should be confirmed by higher resolution spectropolarimetric observations.  Our observations provide only a single snapshot of $B_\\ell$ among the 15 stars in our survey.  Further time-resolved observations of these B-type stars are necessary to determine the intensity and topology of the fields and investigate the spatial distribution of orientation angles throughout the region.  We do not observe any trends with the detection of $B_\\ell$ with $V \\sin i$ or $\\log g_{\\rm polar}$, suggesting neither rotation nor evolution along the MS are major factors influencing the presence of magnetic fields.  If the magnetic fields are produced by a magnetic dynamo, a correlation between $B_\\ell$ and $V \\sin i$ would be expected.  To examine the possibility that the detected fields have a fossil origin, we used the Oblique Dipole Rotator model to simulate populations having a uniform or slightly varying magnetic flux on the ZAMS.  Assuming conservation of the field strength with MS evolution, we compared our 4 definite and marginal detections to the modeled fossil field populations.  None of our simple fossil models successfully reproduces both our observed range in $B_\\ell$ and the expected number of field detections.  Therefore we tentatively rule out a purely fossil origin for the magnetic stars in NGC 3766.  Our results are also consistent with the bimodal distribution of magnetic field strengths found by \\citet{auriere2007}, who proposed a critical magnetic threshold for the stability of large scale magnetic fields among early type stars.  This or some other, unknown mechanism may contribute to the magnetic fields detected in NGC 3766."
    },
    "0807/0807.1103_arXiv.txt": {
        "abstract": "AM-1, at $\\sim$120 kpc, and Pal~14, at $\\sim$70 kpc, are two of the most distant Galactic globular clusters known. We present {\\it Hubble Space Telescope} WFPC2 photometry of AM-1 and Pal~14 that reveals unprecedented depth and detail in the color-magnitude diagrams of these two clusters. Absolute and relative age measurements confirm that both are younger than the inner halo globular cluster M~3 by 1.5--2 Gyr assuming all three clusters have similar compositions. Thus AM-1 and Pal~14 join Pal~3, Pal~4, and Eridanus (studied by Stetson et al.) as distant Galactic globular clusters with red horizontal branches and young ages relative to the inner halo. Within the context of the entire body of research on the ages of second parameter globular clusters, the observed correlation between age and horizontal branch morphology suggests that age is the best candidate for the second parameter. However, this conclusion is tempered by the lack of precise chemical abundance determinations for a significant fraction of second parameter globular clusters. ",
        "introduction": "The second parameter phenomenon, a long-standing problem in Galactic astronomy, recognizes that the horizontal branch (HB) morphology of Galactic globular clusters (GCs) cannot be explained by the first parameter, metallicity, alone: at least one more factor, a `second parameter,' must be involved.  In a landmark study, \\citet{sz78} demonstrated that GCs inside the solar circle (which they deemed the inner halo) exhibit a tight correlation between [Fe/H] and HB morphology while those that lie outside the solar circle (the outer halo) show no such well-defined trend.  This in turn led \\citet{sz78} to theorize that the formation of the Galaxy was a prolonged, complex process.  While the second parameter phenomenon has informed our understanding of galaxy formation, the identity of the second parameter remains a topic of some debate. A number of candidates have been suggested including age, helium abundance, CNO abundances, cluster central density, and stellar rotation (see \\citet{fp96} for a review). It is clear that its influence grows with distance from the Galactic center \\citep{sz78,zi80,ldz94}. To gain a better understanding of the second parameter it is logical to carefully examine the most distant GCs.  Only six known GCs inhabit the Galaxy at distances greater than 50 kpc \\citep{ha96} \\footnote{References to the Harris catalog actually refer to the February 2003 revision that can be found at http://www.physics.mcmaster.ca/$\\sim$harris/mwgc.dat.} and five of the six are prime examples of second parameter GCs, i.e. metal poor GCs with red HBs.  All six of the GCs with $\\rgc >$ 50 kpc: Pal~14, NGC~2419, Eridanus, Pal~3, Pal~4, and AM-1 (listed in order of increasing $\\rgc$) have been observed by the Hubble Space Telescope (HST) using the Wide Field Planetary Camera 2 (WFPC2).  \\citet{ha97} presented a HST/WFPC2 color-magnitude diagram (CMD) of NGC~2419---the only one with $\\rgc >$ 50 kpc and a blue HB---and concluded that it is roughly coeval with the old, metal poor cluster M~92. In contrast, \\citet[hereafter S99]{st99} presented HST/WFPC2 CMDs of Pal~3, Pal~4, and Eridanus and concluded that these clusters are younger than relatively nearby clusters of similar metallicities (M~3 and M~5) by $\\sim$1.5--2 Gyr assuming the inner and outer halo GCs have similar abundances.  The two remaining GCs, AM-1 and Pal~14, are the subject of this paper. AM-1 was classified as a GC by \\citet{am79} though it had been observed previously and named ESO 201-10 by \\citet{ho75}. \\citet{am79} estimated the distance to AM-1 at 300 kpc and, although it has since been revised downward to $\\sim$120 kpc \\citep{aa84,or85,ma89,hi06}, it is still the most distant Galactic GC known.  Age estimates for AM-1 have proved uncertain to date because none of the published photometry has reached the main sequence. Pal~14, at $\\rgc \\sim$70 kpc, was discovered by van den Bergh in 1958 and later classified as a GC by \\citet{ar60}; it is also known as AvdB. \\citet{sa97} published the deepest CMD of Pal~14 to date, reaching more than a magnitude fainter than the main sequence turn off (MSTO), and concluded that Pal~14 is 3--4 Gyr younger than inner halo GCs with similar metallicities.  We describe the observations and data reduction in $\\S$2 and present the WFPC2 CMDs of AM-1 and Pal~14 in section $\\S$3.  Available information regarding the reddening and metallicity of both clusters is discussed in $\\S$4 followed by relative and absolute age analyses in $\\S$5.  The results are discussed and put into context in $\\S$6 and the paper concludes with a summary in $\\S$7. ",
        "conclusions": "Color magnitude diagrams of AM-1 and Pal~14 from HST/WFPC2 F555W--F814W photometry were presented. The CMDs reveal unprecedented depth, reaching $\\sim$4 mag below the MSTO in both clusters. AM-1 and Pal~14 have very similar CMD morphologies.  AM-1 shows a more densely populated RGB and HB, albeit with very few bright red giants, and an impressive display of blue stragglers.  Pal~14 has a sparsely populated RGB and HB which is likely due to its large size (compared to the area covered by WFPC2) and low density.  Three different age measurement techniques reveal that AM-1 and Pal~14 are $\\sim$1.5--2 Gyr younger than M~3 assuming all three GCs have similar compositions. Isochrone fitting suggests that if AM-1 and Pal~14 have similar abundances, then Pal~14 is about 0.5 Gyr younger than AM-1.  S99 found that Pal~3, Pal~4 and Eridanus are also $\\sim$1.5--2 Gyr younger than inner halo GCs of comparable metal abundance.  Thus all five of the GCs with $\\rgc >$ 50 kpc and red HBs are 1.5--2 Gyr younger than GCs with similar metallicities in the inner halo (to the extent that assumptions regarding similar abundances are valid).  The variation of the second parameter phenomenon with $\\rgc$ has been known for some time, but the identity of the second parameter itself remains a topic of active debate in the literature, even some 40 years after it was discovered.  When combined with previously published results on the ages of clusters that exhibit the second parameter effect, the present study provides further evidence that age is a strong candidate for the second parameter. However, there are a significant number of globular clusters important to identifying the second parameter that have not been accurately age-dated.  Of equal importance is the precise determination of abundances in these clusters.  Improvements to both the total metal abundance and the abundance ratios in these clusters will lead to more appropriate comparisons with well-studied, nearby GCs and theoretical models and, ultimately, to a better understanding of the second parameter phenomenon and the formation history of the Galactic halo."
    },
    "0807/0807.1335_arXiv.txt": {
        "abstract": "{We show that expanding or contracting Kasner universes are unstable due to the amplification of gravitational waves (GW).  As an application of this general relativity effect, we consider a pre-inflationary anisotropic geometry characterized by a Kasner-like  expansion, which is  driven dynamically towards inflation by a scalar field. We investigate  the evolution of linear metric fluctuations around this background, and calculate  the amplification of the long-wavelength GW of a certain polarization during the anisotropic expansion (this effect is absent for another GW polarization, and for scalar fluctuations). These GW are superimposed to the usual tensor modes of quantum origin from inflation, and are potentially observable if the total number of inflationary e-folds exceeds the minimum required to homogenize the observable universe only by a small margin. Their contribution  to the temperature anisotropy  angular power spectrum decreases with the multipole $\\ell$ as $\\ell^{-p}$,  where  $p$ depends on the slope of the  initial GW power-spectrum. Constraints on the long-wavelength GW can be translated into limits on  the total duration of inflation and the initial GW amplitude. The instability of  classical GW (and zero-vacuum fluctuations of gravitons) during Kasner-like expansion (or contraction) may have other interesting applications. In particular, if GW become non-linear, they can significantly  alter the geometry before the onset of inflation.} ",
        "introduction": "The inflationary stage of the very early universe  explains the dynamical origin of  the observed isotropic and homogeneous FRW geometry. The patch of the FRW geometry covers the cosmological horizon and beyond if inflation lasted \\begin{equation}\\label{efold} N=62-\\ln \\left(\\frac{10^{16}GeV}{V^{1/4}} \\right)+\\Delta \\ , \\end{equation} e-folds  or longer. Here $V$ is the potential energy of the inflation, and  $\\Delta$ is a correction from the (p)reheating stage after inflation, which is not essential for our discussion. Chaotic inflationary models, associated with a large energy ($\\sim$ GUT scale)  of $V^{1/4} \\sim 10^{16}$GeV, predict a very large number of inflationary e-folds, $N \\gg 62$. Long-lasting inflation erases all classical anisotropies and inhomogeneities of the pre-inflationary stage. However, scalar and tensor vacuum fluctuations during inflation lead to almost scale free post-inflationary scalar and tensor metric inhomogeneities  around our   smooth   observable FRW patch.  In particular, the amplitude of the gravitational waves generated from the vacuum fluctuations during inflation is proportional to $V^{1/2}$,  $h_k \\simeq\\frac{H}{2\\pi M_p} \\sim  \\frac{V^{1/2}}{M_p^2}$ (where $M_p$ is the reduced Planck mass).  There are significant efforts to measure the $B$-mode of $\\Delta T/T$ polarizations, since this  will provide a direct probe  of the scale of inflation. The current $95 \\%$ C.L. limits on $r$ (ratio of the tensor to scalar amplitudes of cosmological fluctuations) $r \\lta 0.43$ (WMAP-only) and $ r \\lta 0.2$ (WMAP plus acoustic baryon oscillation, plus supernovae) \\cite{WMAP} shall be improved to $r \\lta 0.1$ by the Planck mission \\cite{planck}, to $r \\lta 0.05$ by the ${\\rm C}_\\ell$over \\cite{clover}, EBEX \\cite{ebex}, and Spider \\cite{spider} experiments (see \\cite{epic} for the study of a mission that can improve over these limits). While these limits imply a detection in the case of high energy inflation, a number of  other inflationary models, including many of the string theory constructions have lower energy, and therefore lead to GW of much smaller amplitude, which are virtually unobservable through  $B$ mode polarization \\footnote{Future gravitational waves astronomy may  allow to probe $r$ up to the level $10^{-6}$ with BBO \\cite{BBO} or ultimate DECIGO \\cite{decigo} direct detection experiments.}.  In anticipation of the null signal observation of the primordial GW from inflation, it is worth  thinking about other implementations of this result for the theory of inflation, besides putting limits on the energy scale $V^{1/4}$. There are models of inflation (including many string theory inflationary models) where the total number of e-folds,  $N$,  does not exceed the minimum   (\\ref{efold}) by a large number. If the extra number of e-folds $\\Delta N$ beyond   (\\ref{efold})  is relatively small then pre-inflationary inhomogeneities of the geometry are not erased completely, and their residuals can be  subject to observational constraints. In the context of this idea, in this paper we suggest an additional  mechanism to have observable gravitational waves associated with inflation. These gravitational waves are very different from the GW generated from the vacuum fluctuations during inflation. Firstly, they are the residual tensor inhomogeneities from the pre-inflationary stage. Secondly, they can be of a classical, rather than quantum, origin. Thirdly, while their initial amplitude  and spectrum are given by the initial conditions, they are significantly affected by the number of ``extra'' e-folds $\\Delta N$. Therefore, observational limits on gravity waves result in constraints on a combination of $\\Delta N$ and of the initial amplitude.  The choice of the initial geometry of the universe before inflation is wide open. In principle, one may assume an arbitrary geometry with significant tensor inhomogeneities  component, and much smaller scalar inhomogeneities. This choice is, however, very artificial. A much more comfortable choice of the pre-inflationary stage will be a generic anisotropic Kasner-like geometry with small inhomogeneities around it. The origin of the anisotropic universe with the scalar field can be treated with  quantum cosmology, or can be embedded in the modern context of the tunneling in the string theory landscape. In fact, a Kasner-like (Bianchi I) space was a rather typical choice in previous papers on pre-inflationary geometry, see e.g. \\cite{pre}. Most of the works on an anisotropic pre-inflationary stage aimed to investigate how the initial anisotropy is diluted by the dynamics of the scalar field towards inflation \\cite{starob83}.  The formalism of linear fluctuations about an anisotropic geometry driven by a scalar field toward inflation was constructed only recently \\cite{Gumrukcuoglu:2006xj,Pereira:2007yy,Gumrukcuoglu:2007bx,Pitrou:2008gk}. Besides the technical aspects of calculations of cosmological  fluctuations, there is a substantial conceptual  difference between computations in the standard inflationary setting and in the anisotropic case. For an isotropic space undergoing inflationary expansion, all the modes have an oscillator-like time-dependence at sufficiently early times, when their frequency coincides with their momentum. One can therefore  use   quantum initial conditions  for these modes. This is no longer the case for an expansion starting from an initial Kasner singularity. In this case, a range of modes, which can potentially be observed today (if $\\Delta N$ is not too large), are not oscillating initially and therefore cannot be quantized on the initial time hyper-surface; as a consequence, there is an issue in providing the initial conditions for such modes. For this reason  we will adopt another perspective, namely, we will consider generic small classical inhomogeneities around the homogeneous background, as an approximation to the more  generic anisotropic and inhomogeneous cosmological solution.    Equipped with this philosophy, we consider an anisotropic expanding universe filled up by the  scalar field with a potential $V(\\phi)$ which is typical for the string theory inflation. We add generic linear metric fluctuations about this geometry. The evolution of these fluctuations is by itself an interesting academic subject. However, it acquires a special significance in the context of the GW signals from inflation, because of a new effect that we report here of amplification of long-wavelength GW modes during the Kasner expansion. This growth terminates when a mode enters the ``average'' Hubble radius (the average of that for all the three spatial directions), or, for larger wavelength modes, when the background geometry changes from anisotropic Kasner  to isotropic inflationary expansion. We perform explicit computations in the case of an isotropy of two spatial directions. In this case the computation becomes much more transparent and explicitly $k$ dependent. Fluctuations for arbitrary $a ,\\, b ,\\, c$ were considered in the formalism of \\cite{Pereira:2007yy,Pitrou:2008gk}, where the $k$ dependence is not explicit. We verified that our  results agree with \\cite{Pereira:2007yy,Pitrou:2008gk} in the axisymmetric $b=c$ limit. We find that only one  of the two GW polarizations undergoes significant  amplification. Therefore, even if we assume for simplicity equi-partition of the amplitudes of   the three inhomogeneous physical modes of the system (the scalar and the two GW polarization) at the initial time, the final spectra that will be frozen at large scales in the inflationary regime will be very different from each other, in strong contrast to what is obtained in the standard inflationary computations.  This result can have different consequences, that we explore in the present work. Suppose that the growing GW mode is still linear (but significantly exceeds other modes) when the space becomes isotropic. Then, we can have significant yet linear classical GW fluctuations at the beginning of inflation, say of amplitude $\\lta {\\cal O}(1)$. If the modes which correspond to the largest scales that we can presently observe left the horizon after the first $\\Delta N$ e-folds of inflation, their amplitude decreased by the factor $e^{-\\Delta N}$ in this period. If $\\Delta N$ is relatively small, say $\\sim {\\cal O}(10)$ the freeze out amplitude of these GW modes would be $\\sim {\\cal O}(10^{-6})$. The angular spectrum of these GW will rapidly decrease  as  the multipole number $\\ell$ grows, since smaller angular scales are affected by modes which spend more time inside the horizon during the inflationary stage.  Suppose instead that the growing GW mode becomes non-linear before the onset of inflation. In this case the background geometry departs from the original onset.  Besides the phenomenological signatures, it is interesting to study the origin of the amplification of the GW mode. It turns out that the effect of GW amplification is related to the anisotropic Kasner stage of expansion. Therefore we will separately study  GW in the pure expanding Kasner cosmology. For completeness, we also include the study of GW in a contracting Kasner universe, which is especially interesting due to the universality of anisotropic Kasner approach to singularity.  The plan of the paper is the following. In Section~\\ref{sec:background} we discuss the evolution of the anisotropic universe driven by the scalar field towards inflation. In Section~\\ref{sec-linear} we briefly review the formalism of the linear fluctuations in the case of a scalar field in an anisotropic geometry, paying particular attention to the GW modes. In Section~\\ref{sec:decoupled} we compute the amplification of one of the two GW modes that takes place at large scales in the anisotropic era. In Section~\\ref{sec:pair} we discuss instead the evolution of the other two physical modes of the system. In Section~\\ref{sec:kasner} we study the evolution of the perturbations in a pure Kasner expanding or contracting Universe. In Section~\\ref{sec:obs} we return to the cosmological set-up, and we compute the contribution of the GW polarization amplified during the anisotropic stage to the CMB temperature anisotropies. In particular, by requiring that the power in the quadrupole does not exceed the observed one, we set some limits on the initial amplitude of the perturbations  vs. the duration of the inflationary stage. In Section~\\ref{sec:sum} we summarize the results and list some open questions following from the present study, which we plan to address in a future work. ",
        "conclusions": "\\label{sec:sum}    We found a new effect of instability of the gravitational waves in an expanding and contracting Kasner geometries. We demonstrated the effect for a particular choice of the Kasner exponents $(-1/3, 2/3, 2/3)$, but we expect it is rather generic. This particular choice allows to simplify the description of GW in such a way that their wave equations is manifestly depending on the momenta. For the contracting Kasner geometry,  we found that   our unstable  GW mode is identical to the unstable large-scale inhomogeneous mode first identified in  \\cite{LK63}-\\cite{BKL82} (for arbitrary Kasner exponents ($p_1,p_2, p_3$)). Backreaction of this  unstable mode is conjectured to alter the Kasner exponents, differently in different spatial patches, depending on the spatial profile of the initial GW.  The Kasner geometry is a rather universal asymptotic solution for many interesting situations: it describes approach towards  black holes or cosmological singularities (including higher dimensional and supersymmetric cases \\cite{Damour:2000hv}), and it describes generic anisotropic expansion from singularity prior to inflation.   All of these situations are inevitably accompanied by quantum fluctuations of the gravitons), and, possibly, also by  classical gravitational waves. There is a long list of questions arising in connection with a new effect of the gravity waves instability in anisotropic geometry. We have to investigate how the instability growth depends on the Kasner exponents $p_1,p_2, p_3$. It will be interesting to understand if the instability of the classical GW  also results in the instability of the graviton zero vacuum fluctuations. It will also be worth to understand the impact of the GW instability on the structure of the singularity inside the black hole and cosmological singularity. It is also interesting to investigate the backreaction of the GW instability on the contracting and expanding anisotropic geometries depending on the initial GW profile.  In this paper we considered the effect of GW instability in the context of the anisotropic pre-inflationary stage. The transition from Kasner expansion to inflation terminates the effect of GW instability but leaves  classical GW signal as the initial conditions for inflation. If inflation does not last  very long, then the residual GW can contribute to the CMB temperature anisotropies. Since GW polarizations and the scalar mode of cosmological fluctuations behave differently during anisotropic pre-inflation, we can consider only the leading  contribution, namely, $H_\\times$ mode of the GW polarization. In this paper we calculated its contribution to the total $\\Delta T/T$ anisotropy angular power spectrum. The angular power spectrum of the signal decreases with $\\ell$ in a power-law manner and depends on the initial spectrum of the classical GW fluctuations,  $C_\\ell \\sim 1/ \\ell^p$, where the exponent $p$ depends on the  power-spectrum of the  classical GW at the initial hypersurface. It is interesting to note that this result is qualitatively similar to the result of \\cite{linde08}, where the impact of GW from the our-universe-bubble nucleation was considered.  The  signal from $H_\\times$ is rather anisotropic, and there is an interesting question about the anisotropy of its multipole structure $\\langle a_{\\ell m} \\, a_{\\ell' m'}^* \\rangle$, which is intriguing in connection with an apparent alignment of the low multipoles of $\\Delta T/T$.  While such an analysis is beyond the aims of the present work, we have estimated the initial conditions (initial amplitude of the GW, versus the duration of inflation) which can lead to potentially observable effects. We leave a more extended analysis to future investigation.        \\vspace{1cm}  {\\bf \\large Acknowledgements}  \\bigskip  \\noindent We thank J.R.~Bond, C.R.~Contaldi, T.~Damour,  I.~Khalatnikov,  A.~Linde, C.~Pitrou, M.~Sasaki, A.~Starobinsky,  J.P.~Uzan and J.~Weinwright for useful discussions and correspondence. The work of A.E.G. and M.P. was partially supported by  the DOE grant DE-FG02-94ER-40823. LK was supported by NSERC and CIFAR.    \\appendix"
    },
    "0807/0807.4456_arXiv.txt": {
        "abstract": "{}% {We identify and characterize low-mass stars in the $\\sim$\\,3\\,Myr old \\tr region by means of a deep \\chandra X-ray observation, and study their optical and near-IR properties. We compare X-ray activity of \\tr stars with known characteristics of Orion and Cygnus\\,OB2 stars.}{We analyzed a 88.4 ksec \\chandra\\,ACIS-I observation pointed at the center of \\tr. Because of diffuse X-ray emission, source detection was performed using the PWDetect code for two different energy ranges: 0.5-8.0\\,keV and 0.9-8.0\\,keV. Results were merged into a single final list. We positionally correlate X-ray sources with optical and 2MASS catalogues. Source events were extracted with the IDL-based routine ACIS-Extract. X-ray variability was characterized using the Kolmogorov-Smirnov test and spectra were fitted by using XSPEC. X-ray spectra of early-type, massive stars were analyzed individually.}{Our list of X-ray sources consists of 1035 entries, 660 of which have near-IR counterparts and are probably associated with \\tr members. From near-IR color-color and color-magnitudes diagrams we compute individual masses of stars and their \\Av values. The cluster median extinction is \\Av=3.6 mag, while OB-type stars appear less absorbed, having \\Av=2.0 mag. About 15\\% of the near-IR counterparts show disk-induced excesses. X-ray variability is found in 77 sources, and typical X-ray spectral parameters are N$_{\\rm H}\\sim$5.37$\\times$10$^{21}$ cm$^{-2}$ and kT$\\sim$1.95 keV, with 1$\\sigma$ dispersions of 0.45 dex and 0.8 keV, respectively. OB stars appear, softer with a median kT$\\sim$0.65 keV. The median X-ray luminosity is 6.3$\\times$10$^{30}$ \\ergs\\,, while variable sources show a larger median \\Lx value of 13$\\times$10$^{30}$ \\ergs. OB-stars have an even higher median \\Lx of 80$\\times$10$^{30}$ \\ergs, about 10 times that of the low-mass stars.}{The \\tr region has a very rich population of low-mass X-ray emitting stars. An important fraction of its circumstellar disks survive the intense radiation field of its massive stars. Stars with masses 1.5-2.5 M$_\\odot$ display X-ray activity similar to that of stars in Cyg\\,OB2 but much less intense than observed for Orion Nebula Cluster members.} ",
        "introduction": "\\label{intro}  The Carina nebula region (NGC\\,3372) is one of the most massive star formation regions of the Galaxy. It is associated with a giant H{\\sc ii} region spanning about 4 deg$^2$ of the sky and being bisected by a prominent V-shaped dark gas and dusty lane. This prominent young structure is not as compact as some of the other young galactic clusters, but seemingly related to a spiral feature. In this direction, we are looking almost tangentially to the now recognized Carina-Sagittarius spiral arm, at the edge of a giant molecular cloud extending over about 130 pc and with a content in excess of 5$\\times$10$^5$ solar masses \\citep{1988ApJ...331..181G}. The concentration of massive stars (i.e. M$\\geq$20 M$_\\odot$) interacts with the parent giant molecular cloud of the region, leading to triggered star formation events on intermediate to lower masses \\cite[e.g.][]{2004MNRAS.351.1457S}.\\\\ This region harbors several open clusters and/or star concentrations (Trumpler 14, 15 and 16; Collinder 228 and 232; Bochum 10 and 11) containing more than 60 known O-type stars \\citep{1995RMxAC...2...57F}. Large cavities within the giant molecular cloud are supposed to be carved out by the Tr\\,14 and 16 open clusters, which contain most of massive stars of the region. In particular Tr\\,16 includes three rare main-sequence O3 stars, the Wolf-Rayet (WR) star HD93162 and the famous luminous blue variable (LBV) \\etacar. There is a historical controversy about the distance and age of Tr\\,14 and Tr\\,16 \\citep{1995RMxAC...2...51W}. For instance, from extensive spectroscopy and photometry \\cite{1993AJ....105..980M} find 3.2 kpc for both clusters. However, photometric studies are strongly affected by differential extinction in the region and peculiar reddening, and so the derived distance  are different. An example is the \\cite{2004A&A...418..525C} work, who for different R=\\Av/E(B-V) values (3.48 and 4.16 for Tr\\,16 and Tr\\,14 regions), compute distances of 4.0 kpc and 2.5 kpc, respectively. A more reliable distance (2250$\\pm$180 pc), was derived from proper motion and Doppler velocities of the expanding \\etacar Homunculus using HST-STIS\\footnote{Data from Hubble Space Telescope (HST) with the Space Telescope Imaging Spectrograph (STIS)} observations \\citep{1997ARA&A..35....1D}. Recent work \\citep{2003MNRAS.339...44T} derives a common distance DM=12.14 (2.7 kpc) and an age between $\\sim$1\\,Myr and 3\\,Myr, for Tr14 and Tr16, respectively. For this study, we adopt for \\tr a distance of 2250 pc and an age of 3\\,Myr. This young age agrees with the \\cite{2000ApJ...532L.145S} results, who report the existence of several embedded IR sources where star formation might be active. Also, \\cite{2001ApJ...549..578D} confirm clear evidence of pre-main sequence (PMS) stars in the region, while \\cite{2001MNRAS.327...46B} have identified two compact H{\\sc ii} regions possibly linked to very young O-type stars. Finally, \\cite{2004MNRAS.355.1237H} report a compact cluster of infrared PMS-stars in Tr\\,16.  Of the existing methods to identify young stellar populations, the use of X-ray emission is perhaps the least biased \\citep{2002ApJ...574..258F}. While in main-sequence (MS) stars, from late A to M dwarfs, X-rays are believed to originate from the hot coronal gas that is heated by stellar dynamo magnetic fields \\citep{1987ApJ...315..687M}, for late type Pre-MS stars (T-Tauri stars (TTSs)) X-ray emission is attributed to solar-like coronal activity but elevated by a factor of 10$^3$-10$^4$ \\citep{1999ARA&A..37..363F}. Several authors suggested the possibility of detecting early pre-main sequence (PMS) objects through their hard X-ray emission escaping the highly obscured regions (see \\cite{1992AJ....104..758W, 1997PASJ...49..461K, 1997ApJ...486L..39H, 2002ApJ...579L..95H}). Recently, X--ray surveys have been successful in identifying the young and pre-MS population in star-forming regions, including: $i-$ deeply embedded Class I young stellar objects (YSOs), $ii-$ low-mass T-Tauri PMS stars, $iii-$ intermediate-mass Herbig Ae/Be PMS stars, $iv-$ zero-age MS stars. Moreover, X-ray emission from low-mass pre-MS stars usually exhibits a strong variability that helps to confirm membership.  On last decade, X-ray observations of young stars on star-forming regions were intensified thanks to the high spatial resolution and the improved broad-band ([0.2-12.0] and [0.5-10.0] keV) effective area of the \\xmm\\, and \\chandra\\, satellites. A first X-ray survey in the Carina region by \\cite{2003MNRAS.346..704C} was performed on the basis of two early \\xmm\\, observations (rev\\,\\#115 and \\#116) centered on \\etacar. Because of the spatial resolution of the EPIC\\footnote{European Photon Image Camera has about six times less spatial resolution than \\chandra ACIS-I camera.} camera and relatively short exposure time of the observations ($\\sim$35 ksec), they detected only 80 X-ray sources, most of them related to the massive OB-type stars with \\Lx$\\sim$10$^{32}$-10$^{34}$ \\ergs. Before the observation used here, three \\chandra\\, observations were obtained on this region, two (obsId\\,50 and 1249) in the timed exposure mode, and the third (obsId\\,51) in the continuous clocking mode, which produces no image. Using only observation obsId\\,1249, \\cite{2003ApJ...589..509E} presented luminosities and hardness ratios of the hot stars in Tr\\,16, and part of Tr\\,14. Low-resolution X-ray spectra of luminous sources were discussed by \\cite{2004ApJ...612.1065E}. However the short exposure time of such observation ($\\sim$ 9.5 ksec) was a serious limitation for the study of intermediate- and low mass stellar population of the region. This limitation exists even if obsId. 50 and obsId. 1249 are combined \\citep{2007ApJ...656..462S}, reaching completeness just at X-ray luminosity (\\Lx) of $\\sim$7$\\times$10$^{31}$ \\ergs, i.e. the X-ray  emission level typical of single O- and early B-type stars.  In this paper we present results of the analysis of the deepest X-ray observation ever done in this region ($\\sim$90 ksec). Section \\ref{obs} gives details on the observation and data reduction procedures. Section\\,3 explains the method used to detect the sources, photon extraction and the construction of the catalog. In section\\,4 we present results of the cross-correlation with existing near-IR and optical catalogs of objects and their characterization based on their color-color (CC) and color-magnitude (CM) diagrams. Section\\,5 presents a statistical study of variability in the X-ray domain. Section\\,6  is dealing with results of the analysis of extracted X-ray spectra.  In section\\,7 we discuss X-ray luminosities of stars and compare them statistically with the X-ray source population of ONC and Cygnus\\,OB2 star forming regions. In section\\,8 we discuss X-ray and stellar parameter of O- and early B- type stars. Finally, in section\\,9 we give a summary and draw conclusions of the paper. ",
        "conclusions": "\\label{concl}  We report here results of a deep {\\em Chandra} X-ray observation pointed toward the $\\sim$3\\,Myr old star forming region \\tr. Source detection was performed using the PWDdetect code, identifying 1035 X-ray sources in the 17'$\\times$17' ACIS-I FOV. Most of these seem to be outside the obscured V-shaped region of dust and gas. Star formation in this part of the masked region has probably been disrupted and/or diminished as the stellar winds are blocked inside the cloud, due to the efficiency of photo-evaporation processes caused by interactions of the nearby hot massive stars with the dense dust and gas structures.  Data extraction was performed using the semi-automated IDL-based {\\sc ACIS Extract} package, which is well suited to the analysis of observations of crowded fields such as ours.  The X-ray source list was cross-identified with optical and near-IR (2MASS) catalogs: 28 X-ray sources (of 44 within the FOV) were identified with optically characterized OB members of \\tr and 760 with 2MASS sources. Among these latter sources almost all are believed to be \\tr members. About 90 X-ray sources without optical/NIR counterparts are estimated to be of extragalactic nature (AGNs), while the remaining X-ray sources with no counterpart are likely associated with members that are fainter than the 2MASS completeness limit.  In order to characterize the previously unidentified likely cluster members with NIR counterparts, we placed them on NIR color-magnitude (K$_{\\rm s}$\\,vs.\\,H-K$_{\\rm s}$) and color-color (H-K$_{\\rm s}$\\,vs.\\,J-H) diagrams. A first estimate of interstellar extinction was obtained adopting a 3\\,Myr isochrone for the  low- and intermediate-mass stars and assuming that  O- and early B-type stars lie on the MS. We find a median visual absorption for OB stars of \\Av$\\sim$2.0 mag, while low mass likely members seem to be slightly more absorbed, \\Av$\\sim$3.6 mag. We also use the 3\\,Myr isochrone and the J magnitude to estimate masses of likely members assuming that they share the same distance and absorption. Our sample of X-ray selected members with near-IR counterparts reaches down to M=0.5-0.6\\,M$_\\odot$, and is likely complete down to $\\sim$1.5\\,M$_\\odot$. From the H-K$_{\\rm s}$\\,vs.\\,J-H diagram we estimate that $\\sim$15\\% (51/339) of low-mass stars have NIR excesses, finding it to be quite a high percentage with respect to the 2 Myr old Cyg\\,OB2. We believe that the disk fraction in young SFRs is more dependent on spatial morphology of gas and dust around massive stars, which may enhance photo-evaporation, and thus shorten disk lifetimes, than on the total number of massive stars in the region.  At least 77 sources, i.e. $\\sim$7.4\\%, were found to be variable within our observation with a confidence level greater than 99.9\\%. Only three of the 28 detected O- and early B-type stars were detected as variable during our 90-ksec observation, in spite of the high statistics of the OB stars' light curves. These exceptions are the known binary O3.5V+O8V star HD93205 (our source \\#242), showing a rather linear decay of the count rate during the observation plus a short-term variability, the B1.5\\,V star Tr16-11 (source \\#136) and the B1V star Tr16-5 (source\\#489). The latter two show a flare-like variability probably related to unresolved low-mass companions.  We modeled the ACIS X-ray spectra of sources with more than 20 photons and where f$_{\\rm cont}\\,<$1. We assumed an absorbed single-component thermal emission model. The median log(N$_{\\rm H}$) of the sources is 21.73 (cm$^{-2}$). This value agrees well with the median \\Av computed from the near-IR diagram. The median kT of low-mass stars is 2.6 keV. Sources associated with O- and early B-type stars are instead quite soft (median kT: 0.60 keV). Absorption corrected X-ray luminosities of OB stars were calculated from the best-fit spectral models. O and B-type stars are the most luminous, with L$_{\\rm x}$= 2.5$\\times10^{30}$-6.3$\\times10^{33}$  erg\\,s$^{-1}$. Their X-ray and bolometric luminosities are in rough agreement with the relation L$_{\\rm x}$/L$_{\\rm bol}=10^{-7}$, albeit with an order of magnitude dispersion. Low mass stars have L$_{\\rm x}$ ranging between $10^{30}$ and $10^{31}$ \\ergs (median \\Lx = $2.8\\times10^{30}$). Variable low mass stars are on average 0.5 dex brighter (log(\\Lx)$\\sim$31.0 \\ergs). These X-ray luminosities are consistent with those of similar mass stars in the slightly younger (2\\,Myr) Cyg~OB2 region. However, in the mass range 1.5-2.5 M$_\\odot$, the ONC\\,(1~Myr) shows higher X-ray activity level than observed in \\tr stars in the same mass range. We believe that the age-\\Lx activity connection is an acceptable explanation of this result."
    },
    "0807/0807.2647_arXiv.txt": {
        "abstract": "We report the results of the first self-consistent three-dimensional adaptive mesh refinement magnetohydrodynamical simulations of Population III star formation including the Biermann Battery effect.  We find that the Population III stars formed including this effect are both qualitatively and quantitatively similar to those from hydrodynamics-only (non-MHD) cosmological simulations.  We observe peak magnetic fields of $\\simeq 10^{-9}$~G in the center of our star-forming halo at $z\\simeq 17.55$. The magnetic fields created by the Biermann Battery effect are predominantly formed early in the evolution of the primordial halo at low density and large spatial scales, and then grow through compression and by shear flows.  The fields seen in this calculation are never large enough to be dynamically important (with $\\beta \\geq 10^{15}$ at all times), and should be considered the minimum possible fields in existence during Population III star formation, and may be seed fields for the stellar dynamo or the magnetorotational instability at higher densities and smaller spatial scales. ",
        "introduction": "The nature of the first generation of stars, as well as their influence on later structure formation, is a fundamental problem in cosmology.  A great deal of theoretical progress has been made (see recent reviews by \\citet{Bromm04},~\\citet{2005SSRv..117..445G}, and~\\citet{Ciardi05}.  In the past decade, cosmological hydrodynamic simulations of Population III star formation have achieved great success, and significantly different numerical methods have produced results that agree quite well \\citep{Abel02, 2002ApJ...564...23B,Yoshida03, O'Shea07}.  These calculations have given a reasonably clear picture of the formation process of Population III stars, and have provided some constraints on many of their important properties.  These calculations, while useful, largely ignore an important issue: the relevance of magnetic fields in Population III star formation.  Magnetic fields are widely observed in our galaxy, in other galaxies, and in galaxy clusters, and the origin of these fields is one of the most fundamental and challenging problems in astrophysics \\citep{Carilli02, Widrow02}. One possibility is that magnetic fields are created and amplified in the first generation of stars and are spread throughout the IGM when these stars explode, providing seed fields for later generations of stars and for further amplification by dynamo effects.  If some seed magnetic fields exist before Population III stars form, they may help to remove angular momentum from the star-forming clouds, significantly changing the ultimate mass range of these stars \\citep{Pudritz89, Davies00}.  Several groups have examined the importance of magnetic fields on the evolution of Population III protostellar disks using analytic or semi-analytic models.  These include \\citet{silk06}, \\citet{tan04a} and \\citet{tan04b}, who model (among other aspects of primordial star formation and evolution) dynamos in primordial accretion disks.  Other authors, including \\citet{flower03} and \\citet{maki07}, use one-zone calculations to examine the collapse of the primordial star-forming cloud and the assumptions of flux-freezing.  While useful, these models do not self-consistently include the effects of both magnetic fields and cosmological structure formation.  In this Letter, we report the results of the first magnetohydrodynamic simulations of Population III star formation including the Biermann Battery effect \\citep{Biermann50} within the context of cosmological structure formation.  We describe our numerical methods and the simulation used in Section~\\ref{sec:methods}, present our key results in Section~\\ref{sec:results}, and discuss and summarize our results in Section~\\ref{sec:discussion}. ",
        "conclusions": "\\label{sec:discussion}  In this Letter, we have shown results from the first cosmological magnetohydrodynamical simulation of Population III star formation including the Biermann Battery effect. This effect is one of the most robust methods of generating magnetic fields in the Universe \\citep{Biermann50}, and thus provides useful constraints on the minimum magnetic field expected in situations where Population III star formation will take place.  We find from our simulation that small magnetic fields are primarily generated via the Biermann Battery at relatively low overdensities, and are then amplified to values of nearly $10^{-9}$~G at the center of the cosmological halo via gravitational collapse.  While significant, this magnetic field is still quite small -- the plasma $\\beta$ is never smaller than $10^{15}$ at any point during the simulation.  This suggests strongly that the magnetic fields do not play a significant dynamical role up to n$_B \\sim 10^{10}$~cm$^{-3}$, when the simulation is terminated.  As a result, the evolution of the primordial gas within the cosmological halo is very close to the hydrodynamic results obtained by \\citet{O'Shea07}. At later times, however, the magnetic fields may be amplified to dynamically relevant values via a process such as the magnetorotational instability or a dynamo \\citep[][]{silk06,tan04a,tan04b}.  The simulation of further evolution of our evolving protostellar core is currently beyond the capability of our numerical tools.  However, there are multiple avenues that we can follow to more thoroughly explore the relevance of magnetic fields in Population III star formation.  We can begin our calculation with seed fields having strengths based on limits from observation and theory~\\citep{Widrow02}.  We can examine the effects of magnetic fields at higher densities by implementing more chemical and cooling processes \\citep[as discussed by ][]{turk08}. Both of these projects are underway, and we will report on the results in an upcoming paper.  To summarize, we have performed the first calculations that incorporate the Biermann Battery in cosmological magnetohydrodynamic simulations of Population III star formation.  Our key results are as follows:  1.  From an initial state with no magnetic fields, a combination of the Biermann Battery and compressional amplification can result in fields with strengths of $|B| \\simeq 10^{-9}$~G at n$_B \\simeq 10^{10}$~cm$^{-3}$ at the center of a cosmological halo where a Population III star will form.  2.  The Biermann Battery creates fields predominantly at low density (n$_B \\leq 10$~cm$^{-3}$) and large spatial scales in Population III star-forming halos.  3.  The magnetic fields created by the Biermann Battery are dynamically unimportant at all densities below n$_B \\simeq 10^{10}$~cm$^{-3}$ -- $\\beta \\equiv P_{th} / P_{B} \\geq 10^{15}$ at all times during the evolution of the halo."
    },
    "0807/0807.0532_arXiv.txt": {
        "abstract": "The present work aims at performing a comprehensive census and characterisation of the pre-main sequence (PMS) population in the cometary cloud L1615/L1616, in order to assess the significance of the triggered star formation scenario and investigate the impact of massive stars on its star formation history and mass spectrum. Our study is based on $UBVR_CI_C$ and $JHKs$ photometry, as well as optical multi-object spectroscopy. We performed a physical parametrisation of the young stellar population in L1615/L1616. We identified 25 new T~Tauri stars mainly projected on the dense head of the cometary cloud, almost doubling the current number of known members. We studied the spatial distribution of the cloud members as a function of the age and H$\\alpha$ emission. The star formation efficiency in the cloud is $\\sim$~7--8\\,\\%, as expected for molecular clouds in the vicinity of OB associations. The slope of the initial mass function (IMF), in the mass range $0.1{\\le}M{\\le}5.5$~$M_{\\odot}$, is consistent with that of other T and OB associations, providing further support of an universal IMF down to the hydrogen burning limit, regardless of environmental conditions. The cometary appearance, as well as the high star formation efficiency, can be explained in terms of triggered star formation induced by the strong UV radiation from OB stars or supernovae shockwaves. The age spread as well as both the spatial and age distribution of the PMS objects provide strong evidence of sequential, multiple events and possibly still ongoing star formation activity in the cloud. ",
        "introduction": "\\label{sec:Intro}  The Lynds clouds L1615 and L1616 \\citep{Lynds62} are both located at an angular distance of about $6\\degr$ West of the Orion OB1 associations. These clouds actually form a cometary-shaped single cloud with a ``head-tail'' distribution subtending about $40\\arcmin$ \\citep[5.2~pc assuming a distance of 450~pc;][]{Alcala04} roughly in the East-West direction. The dense head of the cloud complex (i.e. L1616 only), pointing toward East and facing the bright Orion-belt stars, harbours the IRAS Small Scale Structure X0504-034 \\citep{Helou88,Ramesh95} and a bright reflection nebula, NGC\\,1788 \\citep[= DG\\,51, Ced\\,40, vdB\\,33, RNO\\,35, LBN\\,916;][]{Lynds65}, which is illuminated by a small cluster of stars \\citep{Stanke02}. The two intermediate-mass stars HD\\,293815 and Kiso\\,A-0974~15 are the brightest visible members of this cluster.  L1615/L1616, like many other cometary clouds off the main Orion star forming regions, clearly shows evidence of ongoing star formation activity which might have been triggered by the strong impact of the UV radiation from the massive, luminous stars in the Orion complex \\citep{Maddalena86,Stanke02, Alcala04,Kun04,Lee05,Lee07}. The illumination of dense clumps in molecular clouds by OB stars could be responsible for their collapse and subsequent star formation. The UV radiation from the OB stars may sweep the molecular material of the cloud into a cometary shape with a dense core located at the head of the cometary cloud.  The radiation and wind of OB stars may also have an important impact on the mass accretion during the star formation process. While in a T~association a protostar may accumulate a significant fraction of mass, the mass accretion of a low-mass protostar in a region exposed to the wind of OB stars can be terminated earlier because of the photo-evaporation of the circumstellar matter \\citep{Kroupa01,Kroupa02}. Therefore, many low-mass protostars may not complete their accretion and hence can result as brown dwarfs (BDs). This mechanism might affect the low-mass end of the initial mass function (IMF). Recent studies have provided some information about the shape of the IMF in the very low-mass and sub-stellar regimes. While the IMF in the Orion Nebula Cluster appears to rise below 0.1 M$_{\\odot}$ \\citep{Hillenbrand00}, in T~association like Taurus-Auriga and Cha\\,II there is some indication of a deficit of sub-stellar objects \\citep{Luhman00,Briceno02,Spezzi08}. Other studies of the young cluster IC\\,348, which is devoid of very massive stars, have also revealed a deficit of BDs relative to the Orion Nebula Cluster \\citep{Preibisch03,Muench03,Lada06}.  Because of its vicinity to the Orion OB associations the L1615/L1616 cometary cloud constitutes an ideal laboratory to investigate the triggered star formation scenario. The most recent census of the pre-main sequence (PMS) stars in L1616 was provided by \\citet{Alcala04}, who presented a multi-wavelength study of the region, from X-ray to near-infrared wavelengths. They found 22 new low-mass PMS stars distributed mainly to the East of L1616, in about 1~square degree field. By adding the 22~new PMS stars to the previously confirmed members of the cloud \\citep{Cohen79,Sterzik95,Nakano95,Stanke02} and counting the millimeter radio-source, i.e. the Class-0 protostar MMS1A found by \\citet{Stanke02}, \\citet{Alcala04} ended up with a sample of 33~young stellar objects associated with L1616. However, the latter work could not investigate important aspects of the star formation in this cloud, like the IMF and the impact of environmental conditions on the mass spectrum, because their sample was rather incomplete.  Therefore, the aim of the present study is to perform a comprehensive census and characterisation of the PMS population in L1615/L1616, in order to investigate the star formation history, the relevance of the triggered scenario, and to study the IMF. To this aim, we report both optical and near-infrared observations, as well as multi-object optical spectroscopy in L1615/L1616. An important goal of our study is to compare the mass function in L1615/L1616 with that of other T and OB~associations.   The outline of the paper is as follows: we describe our observations and data reduction in Section~\\ref{sec:PhotObs} and \\ref{sec:Spec-Obs}; the results are reported in Section~\\ref{sec:PMS-Ident}, while the data analysis and the physical proprieties of the new PMS stars are presented in Section~\\ref{sec:Data analysis}. Our discussion and conclusions are developed in Section~\\ref{sec:Discussion}. Some details on the spectral-type classification as well as reddening, radius, and luminosity determination are given in the Appendix~\\ref{App-A} and \\ref{App-B}, respectively. ",
        "conclusions": "\\label{sec:Discussion}  Once the physical parameters of the PMS objects are known, we can study the star formation in L1615/L1616. In the next sub-sections we discuss the history, rate and efficiency of star formation, and some issues related to the mass-spectrum.  \\subsection{The star formation history} \\label{sec:Star-Form-Hist}  \\clearpage \\begin{figure*} \\centering \\resizebox{\\hsize}{!}{\\includegraphics[draft=false]{f4.eps}}% \\caption{IRAS 100$\\mu$m dust emission map  covering a sky-area of $5^\\circ\\times5^\\circ$ around L1615/L1616 (left panel). The Classical and Weak T Tauri stars are represented with filled dots and squares, respectively. The upper arrow shows the direction to the Orion OB1 associations, located at about 7.5$^\\circ$ ($\\sim60$ pc at a distance of 450 pc) to the North-East of L1615/L1616. The lower arrow points toward the Orion A Giant Molecular Cloud, also located at about 7.5$^\\circ$ to the South-East of L1615/L1616. The dashed square defines the area surveyed with WFI. This region is zoomed in the right panel, where the two intermediate-mass members HD\\,293815 and Kiso\\,A-0974~15 are marked with five-pointed star symbols.} \\label{Fig:IRAS-Image} \\end{figure*} \\clearpage  Based on the most complete census of the L1615/L1616's population performed in this work, we further investigated the triggered star formation scenario suggested by \\citet{Ramesh95}, \\citet{Stanke02}, and \\citet{Alcala04}.  In left panel of Figure~\\ref{Fig:IRAS-Image} the IRAS~100~$\\mu$m dust emission map of a $5^\\circ\\times5^\\circ$ sky-area around L1615/L1616 is shown. About $64$\\,\\% of the CTTSs are clustered, in the densest part of L1615/L1616, i.e.~within the boundaries of the NGC\\,1788 reflection nebula and to the West of the bright rim of the cloud (Figure~\\ref{Fig:IRAS-Image}, right panel). Such rim is located to the East of NGC\\,1788, at $\\sim6.5{\\arcmin}$ from the head of the cometary cloud (about 0.85 pc at a distance of 450 pc) and it is directly exposed to the UV radiation from the Orion OB stars. On the other hand, the WTTSs are more scattered (only $\\sim22$\\,\\% is projected on NGC\\,1788) and mainly occupy the side of the cloud facing the OB1 associations and to the East of the bright rim.  By applying the Kolmogorov-Smirnov test to the WTTS and CTTS age distributions we found that the two populations are similar at a confidence level of $\\sim 80$~\\%. Thus, it seems there are no age differences between CTTSs and WTTSs in our PMS sample. This result is consistent with previous PMS populations studies \\citep[see][and reference therein]{Feigelson99}; CTTSs are predicted to have ages between about 0.5 and 3 Myr, although some stars retain CTTS characteristics even at ages as old as 20~Myr. On the other hand, many WTTSs occupy the same region on the H-R diagram as CTTSs do, whereas some of them are approaching the zero-age main sequence.  The age distribution of the L1615/L1616 population peaks between 1 and 3 Myr, depending on the adopted evolutionary tracks (Figure~\\ref{Fig:Hist_Ages_Masses}, right panels). This is in agreement with the findings by \\citet{Alcala04}, but the age spread found by us is significantly higher than what these authors claim, exceeding the value expected on the basis of the uncertainties on luminosity and temperature. The L1615/L1616 members span a wide range in age, from less than 0.1~Myr up to about 30~Myr. This might suggest multiple events of star formation in the cloud, which would further support the hypothesis of triggered star formation.  In order to investigate a possible age difference between ``on-cloud'' and ``off-cloud'' PMS objects, we divided the sample in two groups, fixing as dividing line the bright rim of the cloud. Twenty of 56 objects in our sample are located within the boundaries of NGC\\,1788, i.e. to the West of the bright rim, while 36 are off-cloud members. The age distributions of the two groups are shown in Figure~\\ref{Compare_Age}. The on-cloud PMS stars are statistically younger than those located to the East of the bright rim, regardless of the adopted evolutionary tracks. Applying a Kolmogorov-Smirnov test we found that the probability that the two sets are sub-sample of the same statistical population is very low; in particular, we found confidence levels of 1.29, 0.18, and 0.19\\,\\% when using the PMS evolutionary tracks by \\citet{Baraffe98} \\& \\citet{Chabrier00}, \\citet{Dantona97}, and \\citet{Palla99}, respectively. We thus concluded that there is a clear age difference between the two groups of stars.  The above findings further support the scenario of triggered star formation in L1615/L1616, as proposed by \\citet{Stanke02}. In this context, the spatial dispersion and older age of the ``off-cloud'' members can be explained as a consequence of the ``rocket acceleration'' effect. As recently pointed out by \\citet{Kun04}, this acceleration continues after the onset of star formation and the parental cloud is further accelerated with respect to the newly formed objects. As a consequence, the cloud is soon swept off the newly formed stars. This hypothesis, together with the rapid dispersion typical of small clouds, causes the spatial displacement of the oldest cloud members. The conspicuous number of PMS stars found apparently isolated from classical star forming regions \\citep[e.g.][]{Alcala95,Covino97,Guillout98a,Guillout98b,Frasca03, Zickgraf05} might be a consequence of this mechanism as well.  \\clearpage \\begin{figure} \\centering \\resizebox{\\hsize}{!}{\\includegraphics[draft=false]{f5.eps}} \\caption{Mass (left panel) and age (right panel) distributions of the L1615/L1616 PMS population derived by using the evolutionary tracks by \\citet{Baraffe98} \\& \\citet{Chabrier00} (upper panels), \\citet{Dantona97} (central panels) and \\citet{Palla99} (lower panels).} \\label{Fig:Hist_Ages_Masses} \\end{figure}  \\clearpage \\thispagestyle{empty} \\setlength{\\voffset}{-15mm} \\begin{figure} \\centering \\resizebox{\\hsize}{!}{\\includegraphics[draft=false]{f6.eps}} \\caption{Age distributions of the on-cloud and off-cloud PMS stars in L1615/L1616 derived by using the evolutionary tracks by \\citet{Baraffe98} \\& \\citet{Chabrier00} (upper panel), \\citet{Dantona97} (central panel) and \\citet{Palla99} (lower panel).} \\label{Compare_Age} \\end{figure} \\clearpage \\setlength{\\voffset}{0mm}  \\subsection{The star formation efficiency} \\label{sec:SFE}  Based on $^{12}$CO and $^{13}$CO column density maps, \\citet{Ramesh95} estimated that the mass of L1616 alone is in the range $169-193$~$M_{\\sun}$, depending on the used tracer. \\citet{Yonekura99} mapped both L1615 and L1616 in the CO J=1-0 transition line. They inferred that the mass of the cloud system as a whole is $\\sim 530$~$M_{\\sun}$, $\\sim 350$~$M_{\\sun}$ and $\\sim 440 ~M_{\\sun}$ based on $^{12}$CO, $^{13}$CO and $^{18}$CO observations respectively. They also derived the mass of the $^{13}$CO and $^{18}$CO cores (i.e. L1616 only) finding a value of $\\sim 146$~$M_{\\sun}$ and $\\sim 161$~$M_{\\sun}$, in good agreement with the \\citet{Ramesh95}'s determination.  \\citet{Alcala04} estimated the SFE in L1616 to be $\\sim14$\\,\\%, i.e. higher than the average value measured in other low-mass star forming region ($\\la3$\\,\\%). The SFE reported by \\citet{Alcala04} for L1616 is based on the cloud mass reported by \\citet{Ramesh95} and the total stellar mass of the 32 members of the cloud investigated by the authors (i.e. $\\sim30~M_{\\odot}$).  Since the star formation history of L1615 and L1616 is intimately connected, we have re-calculated the SFE considering the system L1615 plus L1616 as a whole. By using the NICER color excess method by \\citet{Lombardi01}, we mapped the dust extinction of the cloud complex and derived its mass. In the NICER technique the $J-H$ and $H-Ks$ colours of the field stars are compared to the colours of stars in a nearby reference field. The colour-excesses of the field stars are then combined and transformed to the visual extinction $A_{\\mathrm V}$, fixing the form of the reddening curve. The normal interstellar extinction law derived by \\citep{Cardelli89} has been adopted by us for this purpose. We retrieved the 2MASS colours of all the point-like sources located in a $50{\\arcmin}\\times50{\\arcmin}$ region centered on NGC\\,1788 and in a reference field close to the L1615/L1616 cloud\\footnote{The SOFI survey is only centered on the densest part of the cloud and can not be used to map the cloud as a whole.}. The latter field was selected by using the nearby low-intensity regions of the IRAS~100$\\mu$m dust emission map, according to prescription by \\citet{Kainulainen06}. Finally, the visual extinction map resulting from the NICER method was converted to cloud mass by assuming the standard gas-to-dust ratio \\citep{Bohlin78}. We found a value of $\\sim550$~$M_{\\odot}$, in good agreement with the one inferred by \\citet{Yonekura99} on the basis of $^{12}$CO observations. Based on this value and on the total mass of the 56 present-day known members of the complex (i.e. 41--46~$M_{\\odot}$, depending on the adopted evolutionary tracks), we derived a SFE in L1615/L1616 of 7--8\\,\\%, i.e. significantly lower than the previous estimate by \\citet{Alcala04}, but still in good agreement with the one generally found in giant molecular clouds hosting OB associations (5--10\\,\\%) and predicted by theoretical calculations on the formation of OB associations \\citep[see][and reference therein]{Clark05}.  \\subsection{Density of star formation and star formation rate} \\label{sec:SFR}  An interesting question is whether L1615/L1615 can be considered as a cluster. According to the definition suggested by \\citet{Lada03}, a cluster is a group of some 35 members with a total mass density larger than 1.0~M$_{\\odot}$\\,pc$^{-3}$. To estimate the density of PMS objects in L1615/L1615, we considered the region spectroscopically surveyed by us (Figure~\\ref{Fig:SpecSurvey}). This region covers an area of approximately 0.25 square degrees and includes about 40 PMS objects. Assuming a distance of 450~pc, the resulting area is about 4~pc$^2$, which means a surface density of about 10-11 PMS objects per pc$^2$ and a volume density on the order of 7 PMS objects per pc$^3$. In the latter calculation we estimated the volume as $V = 0.752{\\times}Area^{1.5}$ \\citep[see][]{Jorgensen07}, assuming a locally spherical distribution of sources. The average mass of the 40 PMS objects, as determined from the results in Section~\\ref{sec:Masses-Ages}, is on the order of 0.7~M$_{\\odot}$, which implies a volume density of about 4-5~M$_{\\odot}$\\,pc$^{-3}$. Therefore, the group of 40 PMS objects confined in the spectroscopically surveyed area can be considered as a cluster according to the criterion by \\citet{Lada03}.  Now, it is interesting to estimate the rate at which the stars in this small cluster are formed. According to the results of Section~\\ref{sec:Masses-Ages}, we estimated that the total mass in PMS objects in that region is on the order of 41--46~M$_\\odot$. Therefore, considering the average age of 2~Myr for these objects (see Figure~\\ref{Fig:Hist_Ages_Masses}), we found that the cometary cloud is turning some 20-23~$M_{\\odot}$ into PMS objects every Myr, which is lower than the star formation rate in other clusters like those in Serpens \\citep{Harvey07}, but higher than in other T~associations like Chamaeleon~II \\citep{Alcala07} and Lupus \\citep{Merin07}. Thus we concluded that L1615/L1616 is a small cluster with a moderate star formation rate.  \\subsection{The Initial Mass Function}  Though based on a low-number statistics, the first clue on the Initial Mass Function in L1615/L1616 was given by \\citet{Alcala04}. Based on this study, the IMF in this region appears roughly consistent with that of the solar neighbourhood in the mass range $0.3<M<2.5~M_{\\odot}$ \\citep{Miller79}. The authors also found several candidates for young BDs and estimated the expected number of BDs relative to more massive PMS stars in L1615/L1616 to be intermediate between Taurus \\citep[$\\sim13$\\,\\%,][]{Briceno02} and the Trapezium cluster \\citep[$\\sim26$\\,\\%,][]{Luhman00}.  The depth of the spectroscopic survey conducted in this work goes down to $I_C\\approx19.0$~mag (i.e. $M\\approx0.03-0.05 ~M_{\\odot}$, see Section~\\ref{sec:BD-fraction}), leading to the discovery of young objects with a mass close to the Hydrogen burning limit. We have thus attempted a fair IMF determination down to the low-mass stellar regime (i.e. $M \\gtrsim 0.1~M_{\\odot}$).  All the present-day known members of L1615/L1616 have masses in the range between $\\sim$0.1 and 5.5~$M_{\\odot}$; their mean mass varies between 0.6 and 0.8~$M_{\\odot}$ depending on the adopted evolutionary tracks (see Table~\\ref{Tab:mass_age} and Fig~\\ref{Fig:Hist_Ages_Masses}, left panels). However, although our survey recovers most of the previously known young objects in L1615/L1616, the completeness of our PMS sample is different from the one of \\citet{Alcala04} sample. The latter includes bright and solar-like objects located in areas not covered by the deeper WFI survey. Therefore, we restriced our IMF analysis to the sample of objects that lie on the sky-area covered with WFI (Figure~\\ref{Fig:IRAS-Image}, right panel).  The shape of the IMF in the cloud may then be constrained using the usual approximation for the low-mass function ($\\frac{dN}{dM} \\propto M^{-\\alpha}$), as reported by \\citep{Moraux03}. We divided the mass range into mass bins of about $0.2~M_{\\odot}$; this value is larger than the accuracy on mass estimates as derived from the uncertainties on temperatures and luminosities and, at the same time, allows us to have a statistically significant number of objects in each bin.  \\clearpage \\begin{figure} \\resizebox{\\hsize}{!}{\\includegraphics[draft=false]{f7.eps}} \\caption{The L1615/L1616 IMF between 0.1 and 5.5~$M_{\\odot}$ as derived by using the evolutionary tracks by \\citet{Baraffe98} \\& \\citet{Chabrier00} (upper panel), \\citet{Dantona97} (central panel) and \\citet{Palla99} (lower panel).} \\label{Fig:IMF} \\end{figure} \\clearpage  Figure~\\ref{Fig:IMF} shows the IMF in L1615/L1616 as obtained by using the three sets of evolutionary tracks by \\citet{Baraffe98} \\& \\citet{Chabrier00}, \\citet{Dantona97} and \\citet{Palla99}. Values of {\\large $\\alpha$} in the three cases agree within the errors and we can derive a weighted mean value of: $$ \\alpha=0.84 \\pm 0.07 $$ as the best approximation of the IMF slope in L1615/L1616 in the mass range $0.1{\\le}M{\\le}5.5$~$M_{\\odot}$. We notice, however, that the evolutionary models by \\citet{Dantona97} tend to produce a steeper IMF with respect to those by \\citet{Baraffe98} \\& \\citet{Chabrier00} and \\citet{Palla99}, as already found by \\citet{Spezzi08}.  In Table~\\ref{Tab:IMF} the $\\alpha$ slope obtained for the IMF in L1615/L1616 is compared with those obtained in other star forming regions with different ages and environmental conditions. These values indicate a common shape of IMF in the mass range from $\\sim0.1$ to $\\sim1~M_{\\odot}$. This provides further evidence of a universal IMF also in the low-mass regime and is in agreement with the results already established for the intermediate and massive star domains \\citep{Kroupa02}.   \\subsection{No sub-stellar objects in L1615/L1616\\,?} \\label{sec:BD-fraction}  Prior to our survey, no confirmed BDs in L1615/L1616 were reported in the literature. However, \\citet{Alcala04} photometrically selected some 30--40 candidates for BDs in this region and estimated a fraction of sub-stellar objects relative to PMS stars ($R_{SS}$) in the range from 18\\,\\% to 25\\,\\%, i.e. similar to the value measured in the Trapezium cluster \\citep{Hillenbrand00,Briceno02,Muench02} and in other OB associations.  Our spectroscopic observations, which investigated about 50\\,\\% of the BD candidates with $I_C\\lesssim19.0$ mag reported by \\citet{Alcala04}, revealed 3--4 young objects with mass $M\\le0.1~M_{\\odot}$ (see Table~\\ref{Tab:mass_age}), close to the Hydrogen burning limit, and no objects with lower mass. However, taking into account the accuracy with which we estimated the masses, these objects could be BDs with a mass just below the Hydrogen burning limit. The $R_{SS}$ in the cloud would be then around 5--7\\,\\%. Though the very low-mass BDs ($M\\la0.03~M_{\\odot}$) and some deeply embedded low-mass objects might have escaped detection in our survey, this value should be slightly lower than the one estimated for other T~associations \\citep[12--15\\,\\%;][]{Briceno02,Lopez04}.  The photometric survey conducted by us in L1615/L1616 is spatially and photometrically complete down to $I_C\\approx21.5$~mag at $10\\,\\sigma$ level. However, only objects brighter than $I_C\\approx19.0$ mag could be observed spectroscopically with a S/N ratio sufficient for spectral type classification purpose. A 1--3~Myr old object (i.e. the typical age of the L1615/L1616 population) at the distance of 450~pc (i.e. a distance modulus of 8.27~mag) with a mass $M\\ga0.03~M_{\\odot}$ would have $I_C\\la19.0$~mag in the absence of extinction \\citep{Baraffe98}. Thus, in the off-cloud regions of L1615/L1616, where $A_{\\mathrm V}\\approx0$~mag, our survey is sensitive down to $\\sim0.03~M_{\\odot}$. The visual extinction occurring in the dense head of the cloud varies in the range 1--5~mag which displaces the limiting mass of the objects with $I_C\\la19.0$~mag to 0.05--0.50~$M_{\\odot}$ \\citep{Baraffe98}. We conclude that we are probably missing the very low-mass objects only in the small dense core of the cloud. Given the peculiar star formation history in this cloud (Section~\\ref{sec:Star-Form-Hist}), the knowledge of its sub-stellar mass spectrum could give crucial insights into the process forming BDs under particular environmental conditions.  \\clearpage \\begin{figure} \\centering \\resizebox{\\hsize}{!}{\\includegraphics[draft=false]{f8.eps}} \\caption{$I_C$ versus $(R_C-I_C)$ diagram for the point-like objects in L1615/L1616 detected in our WFI survey above $3\\,\\sigma$ level (grey little dots). The theoretical isochrones by \\citet{Baraffe98} \\& \\citet{Chabrier00} for the Cousins photometric system of Bessel (continuous line) and for the WFI-Cousins system \\citep[dashed lines,][]{Spezzi07} are overplotted. The PMS stars with both $R_C$- and $I_C$-band measurements (see Table~\\ref{Tab:phot}) are overplotted with black big dots. The $A_{\\mathrm V}=2$~mag reddening vector is also shown. The confirmed PMS object falling well below the 10~Myr isochrone is the veiled star TTS\\,050649.8$-$032104}. \\label{Fig:isoc_Bes} \\end{figure} \\clearpage  The BD candidates sample by \\citet{Alcala04} was selected by inspection of $I_C$ versus $(R_C-I_C)$ colour-magnitude diagram in comparison with the theoretical isochrones by \\citet{Baraffe98} \\& \\citet{Chabrier00}. This approach to search for PMS star and young BD candidates is the most direct way to tackle the problem using optical photometric data; however, in order to select cluster members on the basis of their magnitudes and colours, isochrones that adequately model the data must be used, since the shape of the isochrones may vary significantly from one photometric system to another. The theoretical isochrones for low-mass stars and BDs are provided by \\citet{Baraffe98} \\& \\citet{Chabrier00} in the Cousins photometric system \\citep{Cousins76}, while the observations by \\citet{Alcala04} were performed by using the WFI camera at the ESO 2.2m telescope. The WFI filters transmission curves are somewhat different from the original Cousins ones, in particular for the $I$-band. \\citet{Spezzi07} have computed \\emph{ad hoc} isochrones for the WFI-Cousins system. In Figure~\\ref{Fig:isoc_Bes} these isochrones are compared with those by \\citet{Baraffe98} \\& \\citet{Chabrier00}. The two sets of isochrones appear significantly different, in particular for the coolest objects, i.e. $(R_C-I_C)>1.7$, where the use of the isochrones in the Cousins system would produced many spurious member candidates. This consideration, together with the contamination by background/foreground stars in the field taken into account by the authors, could explain why \\citet{Alcala04} overestimated the number of BD candidates. The use of the adequate isochrones for the WFI-Cousins system would produce much less BD candidates, in agreement with the number of PMS objects with mass close to the Hydrogen burning limit confirmed by our spectroscopic follow-up (see Table~\\ref{Tab:mass_age}).  Though the $R_{SS}$ estimated by us for L1615/L1616 needs further refinements, the region seems to be poor of BDs, with a $R_{SS}$ ratio close to that observed in T~associations or, perhaps, even lower. Though in L1615/L1616 the radiation from the massive stars in the Orion OB1 associations is likely responsible for the star formation (Section~\\ref{sec:Star-Form-Hist}), this primordial process does not seems to favour the formation of BDs. As reported in \\citet{Alcala04}, the radial velocity (RV) distribution of the stars in L1615/L1616 shows a well defined peak at $22.3$~Km/s with the standard deviation of $4.6$~Km/s which is consistent with the average RV error (i.e. $\\sim5$~Km/s). This may indicate that the velocity dispersion of the stars must be less than $5$~Km/s. Considering a velocity dispersion of a few Km/s, $2$~Myr old objects would disperse over a distance of about $2$~pc, which is the approximate projected size of the head of the cometary cloud and has been completely covered by our photometric and spectroscopic observations.  Thus, the results of our study in L1615/L1616 does not play in favour of dynamical ejection or photoevaporation by ionising radiation from massive stars being triggering factors of the BDs formation mechanism."
    },
    "0807/0807.1702_arXiv.txt": {
        "abstract": "A joint spectral analysis of some \\chandra\\ ACIS X-ray data and Molonglo Observatory Synthesis Telescope radio data was performed for 13 small regions along the bright northeastern rim of the supernova remnant \\snr. These data were fitted with a synchrotron radiation model.  The nonthermal electron spectrum used to compute the photon emission spectra is the traditional exponentially cut off power law, with one notable difference: The power-law index is not a constant.  It is a linear function of the logarithm of the momentum.  This functional form enables us to show, for the first time, that the synchrotron spectrum of \\snr\\ seems to flatten with increasing energy.  The effective power-law index of the electron spectrum is 2.2 at 1~GeV (\\ie, radio synchrotron--emitting momenta) and 2.0 at about 10~TeV (\\ie, X-ray synchrotron--emitting momenta). This amount of change in the index is qualitatively consistent with theoretical models of the amount of curvature in the proton spectrum of the remnant.  The evidence of spectral curvature implies that cosmic rays are dynamically important instead of being ``test'' particles. The spectral analysis also provides a means of determining the critical frequency of the synchrotron spectrum associated with the highest-energy electrons.  The critical frequency seems to vary along the northeastern rim, with a maximum value of $1.1^{+1.0}_{-0.5} \\times 10^{17}$~Hz.  This value implies that the electron diffusion coefficient can be no larger than a factor of $\\sim$4.5--21 times the Bohm diffusion coefficient if the velocity of the forward shock is in the range 2300--5000~km~s$^{-1}$.  Since the coefficient is close to the Bohm limit, electrons are accelerated nearly as fast as possible in the regions where the critical frequency is about $10^{17}$~Hz. ",
        "introduction": "\\label{int}  Several observational and theoretical clues support the suggestion that Galactic cosmic rays, up to an energy of 100~TeV \\citep{lag83} or more \\citep{jok87,vol88,bel01}, are accelerated predominantly in the shocks of supernova remnants.  If Galactic supernovae occur at an average rate of one event every 30 years, then an average supernova remnant must transfer about 10\\% \\citep{dru89} of the initial $10^{51}$~ergs of kinetic energy of the ejecta to cosmic rays. In this case, the cosmic-ray energy density may be large enough to affect the structure of the shock (\\ie, cosmic rays may not be mere ``test'' particles).  Three consequences of a large cosmic-ray pressure are potentially observable.  One consequence is that the ambient material is slowed before it crosses the subshock.  This effect reduces the temperature of the shocked gas \\citep{che83,ell00,dec00}. An upper limit on the temperature is provided by the well-known relation between the shock speed $v_{s}$ and the postshock temperature of a gas whose ratio of specific heats $\\gamma = \\frac{5}{3}$: \\begin{equation} kT_{i} \\le \\frac{3}{16} m_{i} v_{s}^2, \\label{eqn1} \\end{equation} where $k$ is Boltzmann's constant and $m_{i}$ and $T_{i}$ are the mass and immediate postshock temperature, respectively, of particle species $i$. Since kinetic energy is transferred to cosmic rays at the expense of the thermalization of the shocked gas, the equality is appropriate only in the limit that the cosmic-ray energy density is negligible.  \\cite{hug00} analyzed the transverse motion of the X-ray--emitting material (mostly reverse-shocked oxygen and neon) in the supernova remnant 1E~0102.2$-$7219 and inferred a forward-shock velocity $v_{s} = 6200^{+1500}_{-1600}$~km~s$^{-1}$.  In this case, the temperature of the shocked protons must be lower than the right-hand side of equation~(\\ref{eqn1}) or else the fitted electron temperature will be too low to be explained by Coulomb heating \\citep{hug00}.  However, the proton temperature is sensitive to the shock velocity ($T \\propto v_{s}^{2}$), and a variety of speeds are reported for 1E~0102.2$-$7219. \\cite{fla04} report that the \\chandra\\ High Energy Transmission Grating \\ion{Ne}{10} line emission data are consistent with a model that includes radial \\ion{Ne}{10} velocities up to $1800 \\pm 450$~km~s$^{-1}$. \\citet{fin06} analyzed several features in \\ion{O}{3} images.  The mean transverse velocity of the features is $2000 \\pm 200$~km~s$^{-1}$. \\citet{eri01} report that the motion of some \\ion{O}{3}--emitting material is best described by a radial \\ion{O}{3} velocity of about 1800~km~s$^{-1}$. Since these velocity results differ and since none of the measurements provides a direct measure of the velocity of the forward shock, the claim that a significant fraction of the internal energy in 1E~0102.2$-$7219 has been transferred to cosmic rays requires additional support.  A second consequence is that the total compression ratio is larger than 4 \\citep{ell91,ber99}. Since cosmic rays slow the upstream material before it crosses the subshock, the velocity of the shocked material relative to the subshock is reduced.  As a result, the contact discontinuity is closer to the subshock than it would be in the absence of a large cosmic-ray pressure. \\citet{war05} report that in Tycho's supernova remnant the mean ratio of the radius of the contact discontinuity to the radius of the forward shock is $0.93 \\pm 0.02$. This separation corresponds to a total compression ratio $r = 5.1_{-1.0}^{+1.9}$.  Similarly, \\citet{cas08a} report that for the southeastern rim of \\snr\\ the ratio varies from a value of $0.91^{+0.03}_{-0.02}$ (\\ie, $r = 4.0^{+2.2}_{-0.5}$) between the bright X-ray synchrotron--emitting filaments to a value of 1 (\\ie, $r = \\infty$) along the filaments. The mean value between the filaments is $0.96 \\pm 0.03$ (\\ie, $r = 9.0^{+25}_{-3.6}$). These authors explore ways in which the location of the contact discontinuity can approach the location of the forward shock, but they cannot explain why these two features curiously appear to be coincident with one another along the filaments. \\citet{kse05a}, who also studied \\snr, infer a mean compression ratio $r = 5.2$.  A third consequence of a large cosmic-ray pressure is that the shock transition region is broadened or ``smeared out'' \\citep{ell91,ber99}.  Only the subshock has a short transition length.  In this case, low-energy cosmic rays, which have relatively small diffusion lengths, experience only a portion of the velocity gradient as they scatter back and forth across the subshock. Higher energy particles, which have relatively large diffusion lengths, experience a larger portion (or all) of the velocity jump. Since the rate of energy gain increases as the velocity difference increases, higher energy particles gain energy faster than lower energy particles. As a result, cosmic-ray spectra do not have power-law distributions.  The spectra flatten with increasing energy \\citep{bel87,ell91,ber99}. \\cite{jon03} and \\cite{vin06} report evidence of curvature in the synchrotron spectra of Cas~A and RCW~86, respectively.  One possible explanation for such curvature is that the cosmic-ray electrons producing the synchrotron emission have curved spectra.  Here we describe the results of an analysis of some radio and X-ray synchrotron data, which suggest that the synchrotron spectrum of \\snr\\ might be curved.  The data, assumptions, and analysis techniques are described in \\S\\ \\ref{dat}. The results of the analysis are discussed in \\S\\ \\ref{dis}. Our conclusions are summarized in \\S\\ \\ref{con}. ",
        "conclusions": "\\label{con}  We have performed a joint spectral analysis of some {\\sl Chandra} ACIS X-ray data and MOST radio data for 13 small regions along the bright northeastern rim of the supernova remnant \\snr.  The data were fitted with a model that includes a synchrotron emission component. This component is based on an electron spectrum that has a momentum-dependent spectral index. The rate of change in the index for each decade in momentum is a free parameter of the fit. If the assumptions described in \\S\\ \\ref{curv} are valid, then the results of the spectral analysis, which show that the synchrotron spectra of \\snr\\ are curved, can be interpreted as evidence of curvature in the GeV-to-TeV electron spectra. The mean amount of curvature in the electron spectra is qualitatively consistent with predictions of the amount of curvature in the proton spectrum of \\snr\\ \\citep{ell00}. The best-fit power-law index at 1~GeV (\\ie, at radio synchrotron--emitting momenta) is $2.221^{+0.013}_{-0.012}$. Including the effect of curvature, the effective spectral index at about 10~TeV (\\ie, at X-ray synchrotron--emitting momenta) is $2.005 \\pm 0.027$ (90\\% confidence level uncertainties).  This effective index is consistent with the predictions of \\citet{ber02}.  The evidence of curved electron spectra suggests that cosmic rays are not ``test'' particles.  The cosmic-ray pressure at the shock is large enough to modify the structure of the shock.  Since nonthermal electrons contain only about 0.1\\% (\\ie, $10^{48}$~ergs) or less of the total internal energy, the results provide indirect evidence of a much more energetic population of cosmic-ray protons.  Collectively, the evidence of (1) spectral curvature in Cas~A \\citep{jon03}, RCW~86 \\citep{vin06}, and \\snr, (2) an unusually low electron temperature in 1E~0102.2$-$7219 \\citep{hug00}, and (3) a compression ratio greater than 4 in Tycho \\citep{war05} and \\snr\\ \\citep{cas08a} suggests that efficient particle acceleration may be a common feature of young, shell-type supernova remnants.  The results of the spectral analysis also determine the ``cutoff critical frequency'' $\\nu_{m}$.  This frequency seems to vary from one region to another (see also Rothenflug \\etal\\ 2004), which implies that the exponential cutoff momentum of the electron spectrum and/or the strength of the magnetic field varies.  It is not possible to identify the cause of the variation using the synchrotron spectral data alone.  As described in the Appendix, the cutoff frequency can be used to set an upper limit on the mean diffusion coefficient $\\bar{\\kappa}$ of the highest-energy electrons.  Aside from the cutoff frequency, this limit depends (strongly) on the velocity of the forward shock $u_{1}$ and (weakly) on the compression ratio $r$ and the ratio of the upstream to downstream magnetic field strengths $B_{1} / B_{2}$.  If $\\nu_{m} = 1.1 \\times 10^{17}$~Hz, $u_{1} = 2300$--5000~km~s$^{-1}$, $r = 4$, and $B_{1} / B_{2} \\approx 0$, then $\\bar{\\kappa} < (4.5$--$21) \\bar{\\kappa}_{\\rm B}$, where $\\bar{\\kappa}_{\\rm B}$ is the mean Bohm diffusion coefficient.  This result implies that at least some of the highest-energy electrons in \\snr\\ diffuse close to the Bohm limit (\\ie, are accelerated about as fast as possible), which provides additional support for the idea that Galactic cosmic rays are predominantly accelerated by the shocks of supernova remnants."
    },
    "0807/0807.1969_arXiv.txt": {
        "abstract": " ",
        "introduction": "Since the first map of supernova remnant (SNR) RX J1713.7-3946 in very high-energy gamma-rays obtained by H.E.S.S. (Aharonian et al. \\cite{RX1713-2004nature}), there are few more SNRs with spatial distributions of VHE $\\gamma$-ray emission reported (see TeV Gamma-ray Source Catalog\\footnote{Mori M., TeV Gamma-ray Source Catalog [available at \\url{http://www.icrr.u-tokyo.ac.jp/~morim/TeV-catalog/}]} and reviews Rowell et al. \\cite{Rowell-2005}, Funk \\cite{Funk-2008}). Surface brightness distribution of $\\gamma$-ray emission of astrophysical objects is an important possibility to test models of kinetics of astroparticles as well as dynamics of magnetic field and turbulence in astrophysical plasma. Thus, an actual and important task is the modeling of the respective surface brightness distribution.  Inverse Compton (IC) electron-photon interactions is one of the most important processes of gamma-ray production in SNRs. Application of the exact formalism of IC emission to the numerical calculation of SNR images requires quite powerfull CPU resources. In such a situation, an accurate approximation may be of interest because it considerably reduces CPU time and also results in some new formulae in classical analysis (Sect. \\ref{ICdelta} and \\ref{ICThomson}).  A common approach is to deal with IC emissivity for a given energy of the initial (monochromatic) photons (e.g. Jones \\cite{jones-68}, Blumenthal \\& Gould \\cite{Blum-Gould-70}) with assumption of a given shape of electron spectrum (power-law as a common choice). The resulting IC photon spectrum is then given by the integration over the energy distribution of the field photons. Such an approach is essential for special cases of the field photon energy distributions.  It is known however, that the black-body radiation field may be used for IC emission in many astrophysical objects. In particular, for SNRs under typical conditions, one may consider just black-body photons (either with single or with few different temperatures representing CMB/IR/optical radiation). In SNRs with no clear IR emission assosiated, the contribution from CMB photons dominates the role of infrared and optical photons (see discussion in Appendix in Lazendic et al. \\cite{Lazendic-et-al-04}). The IR/optical photon fields may contribute typically 10\\%-15\\% of the IC flux in such SNRs (Gaisser et al. \\cite{Gaisser-et-al-98}, Baring et al. \\cite{Baring-et-al-99}).  In this note, we present an approximation for IC emissivity which may be applied to IC emission originating from the black-body photon field with some temperature $T$. Since $T$ is a parameter in our approach, the approximation may be used for calculation of IC radiation from different photon fields (CMB, IR, optical). The target radiation field around some SNRs (e.g. around Galatic center) may not be black-body and/or the contribution from IR/optical photons may dominate over CMB there (Porter et al. \\cite{Porter-et-al-06}, Hinton \\& Aharonian \\cite{Hinton-Aharonian-07}). In cases when different components of the target radiation field may be approximated by a superposition of multiple Planck distributions with different $T$, our approximation may be used in a similar fashion. Namely, the overall IC emission will be the weigthed sum of single approximations, each with different value of the temperature. In cases when the initial radiation field may not be approximately described by a sum of black-body distributions, our approximation is not applicable.  Another assumption in the present paper is the isotropy of the electron and photon fields. A thorough treatment of anisotropic IC scattering from cosmic-ray electrons is done by Moskalenko \\& Strong (\\cite{Moskalenko-Strong-00}).  Our approximation is given in terms of an energy of incident electrons rather than in commonly used terms of the field photon energy. Such an approach opens the possibility for accurate modeling of IC emission of electrons with energy spectra being different from power-law. For example, if we are interested in electron energies around maximum possible values. It is known that contribution from electrons accelerated by the shock to $E\\rs{max}\\sim 30-300$ TeV is important in interpretation of the H.E.S.S. observations of shell SNRs. ",
        "conclusions": "Numerical evaluation of the spatial distribution of the IC emission in SNRs requires essential computational resources because the IC volume emissivity (\\ref{PICdef}) consists in two enclosed integrations (over initial photon and electron energies). We developed the approximation (\\ref{calIappranyeta}) of the spectral distribution of the IC emission power $p(E\\rs{\\gamma})$ of electrons with Lorentz factor $\\gamma$ which are interacting with the isotropical black-body photon field. Namely, Eq.~(\\ref{calIappranyeta}) restores known results (Blumenthal \\& Gould \\cite{Blum-Gould-70}) with high enough accuracy in any regime, from Thomson to extreme Klein-Nishina limits. It may be used for different astrophysical objects. For $kTE\\rs{\\gamma}\\lesssim 100\\left(m\\rs{e}c^2\\right)^{2}$, i.e. in case of IC $\\gamma$-ray emission from electrons accelerated by the forward shock of SNRs, it is suitable to use a bit simple approximation (\\ref{calIappr}). In the Thomson limit, our approach results in a simple expression (\\ref{TomshICemis}).  Our approximation may be used in situation when the initial radiation field may be approximated with the Planck function with some temperature $T$ or when it may be represented by a superposition of the black-body distributions with different $T$. In addition, it assumes isotropy of electron and photon fields.  The approximation is given in terms of an energy of incident electrons rather than in commonly used terms of the field photon energy (there are known approximation for the latter approach, see e.g. Schlickeiser \\cite{Schlick-book}). Therefore, our approximation may be useful for analysis of the role of the electron spectrum with shapes different from power-law.  The main idea behind our approach is the possibility to split initial integral into two parts which, contrary to the original integral, may be scaled. This scaling is the reason of accuracy of the approximation over the wide range of parameters, from Thomson to extreme Klein-Nishina regime.  There is well known `monochromatic approximation' where electron is scattered by the monochromatic photons with energy $\\epsilon\\rs{o}$ (e.g. Schlickeiser \\cite{Schlick-book}). Fig.~\\ref{fig-d} shows that the spectral distribution of the radiation power of `single' electron with Lorentz factor $\\gamma$ scattered by photons distributed with the Planck function is peaked at some energy $E\\rs{\\gamma m}(\\gamma)$. This allows us to introduce -- similarly to the case of synchrotron emission -- the `delta-function' approximation for IC emission. In this approximation, Eq.~(\\ref{deltaapprox}), all radiated energy of electron is assumed to be at $E\\rs{\\gamma m}(\\gamma)$. In the Thomson limit, $E\\rs{\\gamma m}(\\gamma)\\approx 4\\epsilon\\rs{c}\\gamma^2$ where $\\epsilon\\rs{c}=kT$. In the classical `monochromatic approximation', the average $\\left\\langle E\\rs{\\gamma}\\right\\rangle=(4/3)\\epsilon\\rs{o}\\gamma^2$ is used as an estimator for the energy of emitted IC photons.  Our approach results in some new expressions which represent known results. Namely, Eq.~(\\ref{TomshICemis}) yeilds the spectral distribution of IC radiation power of a \"single\" electron and Eq.~(\\ref{Thomsemiss}) represents the spectrum of IC emission from the power-law spectrum of electrons, in the Thomson limit. These expression account for integration over all possible energies of the seed black-body photons. They can be derived thanks to the possibility of analytical integration of ${\\cal I}$ in the Thomson regime, Eq.~(\\ref{app3})."
    },
    "0807/0807.2067_arXiv.txt": {
        "abstract": "{ $\\theta$~Carinae belongs to a group of peculiar early-type stars (OBN) with enhanced nitrogen and carbon deficiency. It is also known as a binary system, but it is not clear yet whether the chemical anomalies can be explained by mass transfer between the two components. On the basis of the previously reported spectral variability of a few metal lines it may be expected that \\tcar{} possesses a weak magnetic field.  } {A study of the physical nature of this hot massive binary which is furthermore a well-known blue straggler lying $\\sim$2\\,mag above the turnoff of the young open cluster IC~2602 is important to understand the origin of its strong chemical anomalies.} { We acquired high resolution spectroscopic and low resolution spectropolarimetric observations to achieve the following goals: a) to improve the orbital parameters to allow a more in-depth discussion on the possibility of mass transfer in the binary system, b) to carry out a non-local thermodynamic equilibrium (NLTE) abundance analysis, and c) to search for the presence of a magnetic field.  } { The study of the radial velocities using CORALIE spectra allowed us to significantly improve the orbital parameters. A comparative NLTE abundance analysis was undertaken for \\tcar{} and two other early B-type stars with recently detected magnetic fields, \\tsco{} and \\xcma{}. The analysis revealed significantly different abundance patterns: a one-order-of-magnitude nitrogen overabundance and carbon depletion was found in \\tcar{}, while the oxygen abundance is roughly solar. For the stars \\xcma{} and \\tsco{} the carbon abundance is solar and, while an N excess is also detected, it is of much smaller amplitude (0.4--0.6\\,dex). Such an N overabundance is typical of the values already found for other slowly-rotating (magnetic) B-type dwarfs. For \\tcar{}, we attribute instead the chemical peculiarities to a past episode of mass transfer between the two binary components. The results of the search for a magnetic field using FORS\\,1 at the VLT consisting of 26 measurements over a time span of $\\sim$1.2\\,h are rather inconclusive: only few measurements have a significance level of 3$\\sigma$. Although we detect a periodicity of the order of $\\sim$8.8\\,min in the dataset involving  the measurements on all hydrogen Balmer lines with the exception of the H$\\alpha$ and H$\\beta$ lines, these results have to be confirmed by additional time-resolved magnetic field observations. } {} ",
        "introduction": "\\label{sect:intro}  It has long been assumed that massive stars do not have magnetic fields, as they lack the convective outer mantle prevalent in lower mass stars. However, indirect evidence is supporting that magnetic fields are indeed present in massive O and early B-type stars (e.g., Henrichs et al.\\ \\cite{henrichs}; Rauw et al.\\ \\cite{rauw}; Cohen et al.\\ \\cite{cohen}; Gagn{\\'e} et al.\\ \\cite{gagne}). Yet, only very few direct magnetic field detections have been reported in O-type stars so far (Donati et al.\\ \\cite{donati06a}; Wade et al.\\ \\cite{wade06}; Hubrig et al.\\ \\cite{hubrig07}). Among the hottest B-type stars, a magnetic field has been discovered in the B0.2V star \\tsco{} (Donati et al.\\ \\cite{donati06b}) and in the B0.7IV star \\xcma{}, which is one of the hottest $\\beta$\\,Cephei stars, with a rather large longitudinal magnetic field of up to 300\\,G (Hubrig et al.\\ \\cite{hubrig06}).  Walborn (\\cite{walborn06}) listed the hot B0.2V star $\\theta$~Carinae (HD\\,93030, HIP\\,52419, HR\\,4199; $m_V$ = 2.78)  among a few other massive stars with unexplained spectral peculiarities or variations for which a magnetic field could be expected. It has a peculiar, variable spectrum in both optical and UV. The spectral peculiarities in the blue-violet are an enhancement of nitrogen and deficiency of carbon, but also definite line-intensity variations as well as other line asymmetries (Walborn \\cite{walborn76,walborn79}).  The non-detection of a magnetic field was reported by Borra \\& Landstreet (\\cite{borra_land}), who used a photoelectric Balmer-line magnetograph to measure circular polarization in the wings of the H$\\beta$ line. Although \\tcar{} is a very bright target, easily observable with spectropolarimeters, no other magnetic field measurements have been reported so far in the literature, mainly due to the unavailability of instruments  equipped with polarization analyzing optics on telescopes located in the southern hemisphere.  \\tcar{} is an SB1 system with one of the shortest orbital periods known among massive stars ($P$ = 2.2\\,d; Lloyd et al.\\ \\cite{lloyd}). A discussion of the possibility of mass transfer in the binary system, which would be a natural explanation for the remarkable spectral peculiarities and for the singular location of this object in the H-R diagram of the 30\\,Myr old open cluster IC\\,2602, was presented by Walborn (\\cite{walborn76}). However, the spectral type of the companion remains unknown. Previously determined orbital elements were rather uncertain (Lloyd et al.\\ \\cite{lloyd}) and other periods of $\\sim$1\\,day and $\\sim$25~days have been suggested due to a systematic difference in the radial velocities from different observers.  Randich et al.\\ (\\cite{randich}) reported in their ROSAT/PSPC study of the cluster IC\\,2602 that  \\tcar{} is the brightest X-ray object among the studied cluster members, with a soft X-ray luminosity amounting to up to $\\log L_{\\rm X}$ = 30.99\\,ergs\\,s$^{-1}$. { Recently, Naz\\'e \\& Rauw (submitted), using XMM-{\\it Newton} observations, showed that the X-ray flux  of \\tcar{} is slightly lower than the flux typically observed in O and early B-type stars and confirmed the unusual softness of the X-ray emission. Further, they noted that X-ray lines appear narrow and unshifted, reminiscent of those of $\\beta$\\,Cru and the magnetic star \\tsco{} (Donati et al.\\ \\cite{donati06b}). }  Below, we present the results of our new spectroscopic and spectropolarimetric study of this peculiar massive star and discuss possible origins of its anomalies. ",
        "conclusions": "\\label{sect:disc}  \\subsection{Binarity and evolutionary history}  \\begin{figure} \\centering \\includegraphics[height=0.48\\textwidth,angle=270]{9972f8.eps} \\caption{Color-Magnitude diagram of IC\\,2602. Small crosses mark the position of main sequence stars with $T_{\\rm eff}$ = 30\\,000--32\\,000\\,K and $\\log g$ = 4.0--4.4, according to theoretical models. Solid lines represent isochrones for $\\log\\tau$ = 3.0, 7.5, and 8.0. The photometric position of \\tcar{} is marked with a filled square and other IC\\,2602 members  with filled circles. Photometric data and cluster parameters [$(m-M)$ = 6.1, $E(B-V)$ = 0.024] were taken from the WEBDA database (Mermilliod \\cite{mermilliod}). } \\label{fig:cmd} \\end{figure}  We have presented in Sect.~\\ref{sect:binarity} the first high quality radial velocity curve for \\tcar{}. It is well fitted by a Keplerian orbit with eccentricity $e$ = 0.127, leaving no doubt about the binary origin of these variations. The improved determination of the orbital parameters allows us to make a more sophisticated interpretation of the nature of this system and to review the possibility of a mass transfer between primary and secondary as the cause of the observed chemical anomalies. The physical parameters of \\tcar{} determined in  Sect.~\\ref{sect:abundances} suggest that it is located in the H-R diagram close to the ZAMS. According to the Geneva stellar models (Schaller et al.\\ \\cite{schaller}; Lejeune \\& Schaerer \\cite{lejeune_schaerer}) for solar abundances, the observed values of $T_{\\rm eff}$ = 31\\,000\\,K and $\\log g$ = 4.2 correspond to a star with mass $M_1$ = 15.25\\,M$_{\\sun}$, radius $R_1$ = 5.1\\,R$_{\\sun}$, and  age $\\log\\tau$ = 6.0. Considering a 1$\\sigma$ uncertainty in $\\log g$ and $T_{\\rm eff}$, we obtain $\\log\\tau \\le 6.7$. This age value is considerably lower than the accepted age for the cluster IC\\,2602 ($\\log\\tau = 7.83$, Kharchenko et al.\\ \\cite{kharchenko}; $\\log\\tau = 7.63$, Eggen \\cite{eggen}; $\\log\\tau = 7.32$, Whiteoak \\cite{whiteoak}; $\\log\\tau = 7.16$, Hill \\& Perry \\cite{hill_perry}), consistent with the classification of this star as a blue straggler. We note that the absolute magnitude and intrinsic colors interpolated in the same theoretical grid are in agreement with the position of \\tcar{} in the Color--Magnitude diagram of the cluster IC\\,2602 (see Fig.~\\ref{fig:cmd}).  \\begin{figure} \\centering \\includegraphics[width=0.35\\textwidth]{9972f9.eps} \\caption{Physical  parameters as a function of the orbital inclination: orbital semiaxis $a$, critical radius for the primary $R_1^{\\rm crit}$ and the secondary $R_2^{\\rm crit}$, primary radius derived from the rotational velocity $R_1^{\\rm vsini}$, and mass of the secondary star $M_2$. The full lines correspond to $M_1$ = 15.25\\,M$_{\\sun}$ and the dotted lines to $M_1$ = 13.9\\,M$_{\\sun}$ and $M_1$ = 18.2\\,M$_{\\sun}$.} \\label{fig:param} \\end{figure}  The uncertainty in the stellar mass and radius interpolated in the stellar model grid was estimated from the adopted errors for $T_{\\rm eff}$ and $\\log g$. The primary mass ranges from 13.9\\,M$_{\\sun}$ (corresponding to a ZAMS star with $T_{\\rm eff} = 30\\,000$\\,K) to 18.2\\,M$_{\\sun}$ (corresponding to $T_{\\rm eff} = 32\\,000$\\,K and $\\log g = 4.0$). From the estimated  primary mass, relevant information for the system can be derived from the radial velocity curve, even when the orbital inclination is in principle unknown. We show in Fig.~\\ref{fig:param}  the absolute value of the orbital semiaxis and the mass of the secondary star as a function of the orbital inclination. For each parameter plotted in this figure the solid line corresponds to the solution with $M_1 = 15.25$\\,M$_{\\sun}$, while solutions with 13.9\\,M$_{\\sun}$ and 18.2\\,M$_{\\sun}$ are plotted with dotted lines.  To evaluate the occurrence of a mass exchange in the system, we calculated the volume radius of the Roche lobe at periastron ($R_1^{\\rm crit}$ and $R_2^{\\rm crit}$), which are plotted in Fig.~\\ref{fig:param}. From this figure it is clear that these parameters are not strongly dependent on the orbital inclination unless the latter is low. However, a low inclination can be ruled out for the following reasons. On the one hand, as argued by Walborn (\\cite{walborn79}), the absence of a pronounced emission at H$\\alpha$ indicates that the equatorial velocity is lower than 300\\,km\\,s$^{-1}$ and consequently, from the  observed $v \\sin i$ value we obtain $i \\geq 22^\\circ$. This value is presented by the solid vertical line in Fig.~\\ref{fig:param} and corresponds to an upper limit of about 2.0\\,M$_{\\sun}$ for the secondary component. On the other hand, if we assume that the rotation is synchronized with the orbital motion at periastron, the relatively small radius deduced from the spectroscopic value of $\\log g$ is compatible with the observed $v \\sin i$ only for high inclinations. For the case of synchronous rotation we derive a rotational period of 1.692$\\pm$0.007\\,d and hence $R_1\\sin i$ = 3.78$\\pm$0.27\\,R$_{\\sun}$. The adopted range for $M_1$ (13.9--18.2\\,M$_{\\sun}$) corresponds to $R_1$ = 4.7--7.1\\,R$_{\\sun}$, and therefore we obtain $i = 32^\\circ - 54^\\circ$. These limits are marked with two  dashed vertical lines in Fig.~\\ref{fig:param}. Interestingly, this rotational period of 1.692$\\pm$0.007\\,d  is in good agreement with the 1.779\\,d period discovered by Walborn (\\cite{walborn79}) who studied intensity variations of \\ion{Si}{IV} $\\lambda$4089.  The radius derived from $v \\sin i$ assuming synchronization would coincide with the critical radius for $i \\approx 25^\\circ$. In this case $R_1 \\approx 9.1\\,$R$_{\\sun}$ and $\\log g \\approx 3.7$, which is lower than the spectroscopic value by  2.5$\\sigma$. Therefore, we expect that, at present, the primary star is probably not filling its Roche lobe, although it is close to that point.  A star of 15.25\\,M$_{\\sun}$ (13.9--18.5\\,M$_{\\sun}$) with a present radius of 5.1\\,R$_{\\sun}$ (4.7--7.1\\,R$_{\\sun}$) has an evolutionary age of 1.0\\,Myr (0--4.3\\,Myr) and, evolving as a single star, would have at the Terminal Age Main Sequence (TAMS) a radius of about 13.6\\,R$_{\\sun}$ (12.6--16.3\\,R$_{\\sun}$), which is significantly larger than the critical radius. Thus, we can expect mass-transfer to occur from the primary to the secondary during the main-sequence stage, at an evolutionary age of about 11\\,Myr.  Since  \\tcar{} is a bona-fide blue straggler of IC\\,2602, it is very likely that the system has suffered in the past a mass-transfer in the opposite direction. Eggen \\& Iben (\\cite{eggen_iben}) have discussed extensively this possibility as an explanation for its abnormal position in the Color-Magnitude diagram. The apparent evolutionary youth of \\tcar{} and its peculiar chemical abundances are best understood if the present primary star was originally the less-massive component and has accreted mass transfered from the companion. Its original mass could be less than half of the total mass of the system, similar to other stars near the turnoff point in the cluster. After considerable mass-transfer, it would appear as a non-evolved massive star above the cluster turnoff and near the ZAMS. As was noted by Eggen \\& Iben (\\cite{eggen_iben}), during the evolution of this binary system an important loss of angular momentum has taken place, making it possible to achieve a system with such a short period and low mass-ratio. The orbital angular momentum of a binary system can be written as $$ J_{\\rm orb} = \\frac{2 \\pi}{P} M a^2 \\frac{q}{(1+q)^2} \\sqrt{1-e^2} $$ while for the rotational momentum of each star we have $$ J_{rot} = \\frac{2 \\pi}{P_{rot}} \\beta^2 M R^2 $$ where $\\beta$ is the radius of gyration, which is typically 0.25 for a main-sequence star (Claret \\& Gim\\'enez \\cite{claret_gimenez}). In any case the rotational contribution to the total angular momentum is small. We can use these expressions to calculate the  angular momentum of the system  and compare the present value with that at the moment when the mass-ratio was close to unity. As illustration, let us assume that the original primary star (presently the secondary) with 9\\,$M_{\\sun}$ filled its Roche lobe when it was still in the main-sequence with a radius of the order of 7--9\\,$R_{\\sun}$. Therefore, at that time the separation between the stars was about 20--25\\,$R_{\\sun}$. Assuming a conservative mass transfer and adopting for the present configuration  $M_1$ = 15\\,$M_{\\sun}$, $M_2$ = 1.0\\,$M_{\\sun}$, and a = 19\\,$R_{\\sun}$,  we can estimate  that the present angular momentum is about 0.22--0.25 of the original value. In conclusion, most of the angular momentum has been lost.   \\subsection{Abundances and magnetic field}   The binary and blue straggler nature of \\tcar{} leads us to relate the observed chemical anomalies to the past history of the binary system. First, mass transfer of chemically processed material from the initially more massive component has dramatically altered the CNO surface abundances. Second, the spun up phase that ensued may have further enhanced nitrogen and depleted carbon in \\tcar{} as a result of rotational mixing (see e.g.\\ Langer et al.\\ \\cite{langer}). The tendency of magnetic B stars to show a higher incidence of an N excess compared to stars with no field detection (Morel et al.\\ \\cite{morel08}) also indicates that magnetic phenomena could play a role. The results of the previous studies, along with the fact that little is known about the magnetic properties of hot, massive stars, has motivated our search for a magnetic field in \\tcar{}.  There are only six massive stars with detected magnetic fields that are hotter than \\tcar{}: $\\theta^1$\\,Ori~C, HD\\,191612, HD\\,155806, HD\\,148937, 9\\,Sgr, and \\tsco{}. The star HD\\,191612 with a rotation period of 538\\,d was observed with ESPaDOnS only during four consecutive nights (Donati et al.\\ \\cite{donati06a}). Among the published magnetic field measurements of $\\theta^1$\\,Ori~C, only four measurements show significance at the 3$\\sigma$ level (Wade et al.\\ \\cite{wade06}). HD\\,155806 and HD\\,148937 were observed only once and 9\\,Sgr three times (Hubrig et al.\\ \\cite{hubrig07}; Hubrig et al., in preparation). The sixth star, \\tsco{}, with physical parameters very similar to those of \\tcar{}, was studied over the rotational period of 41\\,d. It possesses the weakest mean longitudinal magnetic field with a maximum value of 88\\,G and a very complicated geometry. The structure of its magnetic field topology features in particular a significantly warped torus of closed magnetic loops encircling the star and additional smaller networks of closed field lines (Donati et al.\\ \\cite{donati06b}).  In our magnetic field study of \\tcar{}, only very few measurements were achieved at a significance level of 3$\\sigma$ and their sporadic appearance is difficult to explain in the framework of the presence of a global large-scale organized magnetic field. On the other hand, the presence of a complex structure of the magnetic field in \\tcar{} may hypothetically be possible in an analogous manner to the magnetic field topology of \\tsco{}. The detected periodicity of the order of 8.8\\,min in the dataset of measurements carried out on hydrogen Balmer lines with the exclusion of H$\\beta$ is surprising, and if it is not spurious, its discovery would give rise to the important question whether the presence of pulsations could cause such a periodicity. No studies of short-time pulsations exist for \\tcar{} so far. B0 main-sequence stars are expected to pulsate in low radial order p- and g-modes with periods of the order of hours (typically 3 to 8 hours). For such massive stars, periods of the order of a few minutes would correspond to high radial order p-modes, but non-adiabatic codes do not predict their excitation. However, since current models do not take into account the presence of a magnetic field, one cannot exclude the possibility that magnetic fields would favour the excitation of these types of modes in analogy to roAp stars which do possess a magnetic field and pulsate in high radial order modes of the order of a few minutes. Our data with the CORALIE spectrograph were taken typically every 7~minutes. This leads to a sampling rate which is too low to allow us to detect variations of spectral line profiles on short time scales.  Since the mechanism of the generation and the maintenance of magnetic fields in massive stars is not well understood yet, we are cautious in drawing any conclusion on the presence and the behaviour of a magnetic field in \\tcar{}. Only additional time-resolved magnetic field observations will tell us about the presence and the structure of the magnetic field geometry of \\tcar{} and will help to discriminate among the hypotheses described above."
    },
    "0807/0807.2856_arXiv.txt": {
        "abstract": "Accurately understanding the interior structure of extra-solar planets is critical for inferring their formation and evolution. The internal density distribution of a planet has a direct effect on the star-planet orbit through the gravitational quadrupole field created by the rotational and tidal bulges. These quadrupoles induce apsidal precession that is proportional to the planetary Love number ($k_{2p}$, twice the apsidal motion constant), a bulk physical characteristic of the planet that depends on the internal density distribution, including the presence or absence of a massive solid core. We find that the quadrupole of the planetary tidal bulge is the dominant source of apsidal precession for very hot Jupiters ($a \\lesssim 0.025$ AU), exceeding the effects of general relativity and the stellar quadrupole by more than an order of magnitude. For the shortest-period planets, the planetary interior induces precession of a few degrees per year. By investigating the full photometric signal of apsidal precession, we find that changes in transit shapes are much more important than transit timing variations. With its long baseline of ultra-precise photometry, the space-based \\emph{Kepler} mission can realistically detect apsidal precession with the accuracy necessary to infer the presence or absence of a massive core in very hot Jupiters with orbital eccentricities as low as $e \\simeq 0.003$. The signal due to $k_{2p}$ creates unique transit light curve variations that are generally not degenerate with other parameters or phenomena. We discuss the plausibility of measuring $k_{2p}$ in an effort to directly constrain the interior properties of extra-solar planets. ",
        "introduction": "Whether studying planets within our solar system or planets orbiting other stars, understanding planetary interiors represents our best strategy for determining their bulk composition, internal dynamics, and formation histories. For our closest neighbors, we have had the luxury of sending spacecraft to accurately measure the higher-order gravity fields of these objects, yielding invaluable constraints on their interior density distributions. Using these observations, we have been able, for instance, to infer the presence of large cores, providing support for the core-accretion theory of planet formation \\citep{2005AREPS..33..493G}. Study of planets outside our solar system, however, has necessitated the development and usage of more indirect techniques. Nevertheless, as the number of well-characterized extra-solar planets grows, we gain more clues that help us answer the most fundamental questions about how planets form and evolve.  Guided by our current understanding of planetary physics, we have begun to study the interiors of extra-solar planets. This endeavor has been dominated by a model-based approach, in which the mass and radius of a planet are measured using radial velocity and transit photometry observations, and the interior properties are inferred by finding the model most consistent with those two observations. This strategy clearly requires a set of assumptions, not the least of which is that the physical processes at work in extra-solar planets are just like those that we understand for our own giant planets. While it does seem that this approach is adequate for explaining most of the known transiting planets, there does exist a group of planets for which the usual set of assumptions are not capable of reproducing the observations \\citep[e.g.,][]{2006A&A...453L..21G,2007ApJ...661..502B}. These are the planets with so-called positive ``radius anomalies'',  including the first-discovered transiting planet HD 209458b \\citep{2000ApJ...529L..45C}.  Though most of these planets can be explained by adjusting different pieces of the interior physics in the models (including opacities, equations of state, and heat deposition), it is currently impossible to discern which combination of these possible explanations is actually responsible for their observed sizes \\citep{2006A&A...453L..21G}.  Additional uncertainties also exist for planets at the other end of the size spectrum. For the group of under-sized extra-solar planets, such as HD 149026b, the canonical approach is to give the planet a massive highly condensed core of heavy elements in order to match the observed radius. This approach also provides a first order estimate of the planet's bulk composition, in terms of its fraction of heavy elements. There is also the added complication of how the assumed state of differentiation affects the inferred composition and predicted structure \\citep{Baraffe2008}.  Currently, the most promising approach to modeling the distinctive features of extra-solar planet interiors is to study the known transiting planets as an ensemble. The group can be used to develop either a single consistent model that reproduces all the observations \\citep[e.g.,][]{2006A&A...453L..21G} or to showcase the possible diversity in model parameters \\citep[e.g., opacities, as in][]{2007ApJ...661..502B}. Surely, a model-independent measure of interior structure would be valuable in order to begin disentangling otherwise unconstrained physics.  The idea of obtaining direct structural measurements for distant objects is by no means a new one. For decades, the interiors of eclipsing binary stars have been measured by observing ``apsidal motion,'' i.e. precession of the orbit due to the non-point-mass component of the gravitational field \\citep{1928MNRAS..88..641R,1938MNRAS..98..734C,1939MNRAS..99..451S,1939MNRAS..99..662S}. The signal of the changing orbit is encoded in the light curves of these systems by altering the timing of the primary and secondary eclipses. From these eclipse times, it is straightforward to determine the so-called apsidal motion constant which then constrains the allowed interior density distributions. Interior measurements inferred from apsidal precession were among the first indications that stars were highly centrally condensed. While it seems non-intuitive, we show in this paper that we can use a similar technique to measure the interior properties of very hot Jupiters. Most surprisingly, the interior structure signal for very hot Jupiters actually dominates over the signal from the star, yielding an unambiguous determination of planetary interior properties.  Our theoretical analysis is also extended to full simulated photometry in order to explore the observability of apsidal precession. We show that this precession is observable by measuring the subtle variations in transit light curves. The photometric analysis is focused on the data expected from NASA's \\emph{Kepler} mission, which successfully launched on March 6, 2009 \\citep{2003SPIE.4854..129B,2006Ap&SS.304..391K}. \\emph{Kepler} will obtain exquisite photometry on $\\sim$100,000 stars, of which about 30 are expected to host hot Jupiters with periods less than 3 days \\citep{2008arXiv0804.1150B}. \\emph{Kepler} has the potential to measure the gravitational quadrupoles of very hot Jupiters though the technique described below. If successful, this will constitute a major step towards an understanding of the diversity of planetary interiors.  In Section 2, we describe the background theory that connects interior structure and orbital dynamics and explore which effects are most important. Section 3 applies this theory to the observable changes in the transit photometry, including full \\emph{Kepler} simulated light curves. We show in Section 4 that the signal due to the planetary interior has a unique signature. Other methods for inferring planetary interior properties are discussed in Section 5. The final section discusses the important conclusions of our work. ",
        "conclusions": "The planetary mass and radius are the only bulk physical characteristics measured for extra-solar planets to date. In this paper, we find that the planetary Love number ($k_{2p}$, equivalent to $J_2$) can also have an observationally detectable signal (quadrupole-induced apsidal precession) which can provide a new and unique probe into the interiors of very hot Jupiters. In particular, $k_{2p}$ is influenced by the size of a solid core and other internal properties. Core sizes can be used to infer the formation and evolution of individual extra-solar planets \\citep[e.g.,][]{2009arXiv0901.0582D,2009arXiv0903.1997H}.  The presence of a nearby massive star creates a large tidal potential on these planets, raising significant tidal bulges which then induce non-Keplerian effects on the star-planet orbit itself. The resulting apsidal precession accounts for $\\sim$95\\% of the total apsidal precession in the best cases (Figure \\ref{odotfig}). Hence, we find that the internal density distribution, characterized by $k_{2p}$, has a large and clear signal, not to be confused with any other parameters or phenomena. We urge those modeling the interior structures of extra-solar planets to tabulate the values of $k_{2p}$ for their various models.  Encouraged by this result, we calculated full photometric light-curves like those expected from the \\emph{Kepler} mission to determine the realistic observability of the interior signal. We estimate that \\emph{Kepler} should be able to distinguish between interiors with and without massive cores ($\\Delta k_{2p} \\simeq 0.1$) for very hot Jupiters with eccentricities around $e \\sim 0.003$ (Figure \\ref{ek2planet}). Eccentricities this high may occur for some of the very hot Jupiters expected to be found by \\emph{Kepler}, though these planets usually have highly damped eccentricities. Much stronger constraints on apsidal precession can be obtained by combining \\emph{Kepler} photometry with precise secondary transits observed in the infrared. In cases where apsidal precession is not observed, the data can set strong upper limits on planetary eccentricities.  In analyzing \\emph{Kepler}'s photometric signal of apsidal precession, we find that transit timing variations are an almost negligible source of signal, though transit timing has been the focus of many observational and theoretical papers to date. The effect of ``transit shaping'' has $\\sim$30 times the photometric signal of transit timing for apsidal precession \\citep[see Figure \\ref{candywrapperpieces},][]{PK08,JB08}). At orientations where transit timing and shaping are weakest, the changing offset between primary and secondary transit times can be used to measure $k_{2p}$ (Figure \\ref{bowtiemain}). It may also be possible to measure $k_{2p}$ from high-precision multi-color photometry by directly detecting the planetary asphericity in transit. Such a measurement does not require a long baseline or an eccentric orbit.  Very hot Jupiters are also excellent candidates for detecting tidal semi-major axis decay, where we find that relatively small period changes of $\\dot{P} \\simeq 1$ ms/yr should be detectable. This could constitute the first measurements (or constraints) on tidal $Q_*$ for a variety of individual stars. We note that \\emph{Kepler} measurements of transit timing and shaping for eclipsing binaries should also provide powerful constraints on stellar interiors through apsidal motion and binary orbital decay (due to tides, if the components are asynchronous).  Accurately measuring the interior structure of distant extra-solar planets seems too good to be true. Nevertheless, the exquisite precision, constant monitoring, and 3.5-year baseline of the \\emph{Kepler} mission combined with the high sensitivity of transit light curves to small changes in the star-planet orbit make this measurement plausible.  Our focus on \\emph{Kepler} data should not be interpreted to mean that other observations will be incapable of measuring $k_{2p}$. In fact, the opposite is true since the size of the apsidal precession signal increases dramatically with a longer baseline. Combining \\emph{Kepler} measurements with future ground and space based observations can create a powerful tool for measuring $k_{2p}$. In the far future, many planets will have measured apsidal precession rates (like eclipsing binary systems have now) and inferred $k_{2p}$ values. Incorporating these measurements into interior models holds promise for greater understanding of all extra-solar planets."
    },
    "0807/0807.2251_arXiv.txt": {
        "abstract": "In the frame of the collapsar model for long gamma ray bursts (GRBs), we investigate the formation of a torus around a spinning BH and we check what rotational properties a progenitor star must have in order to sustain torus accretion over relatively long activity periods. We also study the time evolution of the BH spin parameter. We take into account the coupling between BH mass, its spin parameter and the critical specific angular momentum of accreting gas, needed for the torus to form. The large BH spin reduces the critical angular momentum which in turn can increase the GRB duration with respect to the Schwarzschild BH case. We quantify this effect and estimate the GRB durations in three cases: when a hyper accreting torus operates or a BH spins very fast or both. We show under what conditions a given progenitor star produces a burst that can last as short as several seconds and as long as several hundred of seconds. Our models indicate that it is possible for a single collapse to produce three kinds of  jets: (1) a very short, lasting between a fraction of a second and a few seconds, 'precursor' jet, powered only by a hyper accreting torus before the BH spins up, (2) an 'early' jet, lasting several tens of seconds and powered by both hyper accretion and BH rotation, and (3) a 'late' jet, powered only by the spinning BH. ",
        "introduction": "The commonly accepted  mechanism for a long gamma ray burst (GRB) production invokes a collapsar scenario (Woosley 1995; Paczy\\'nski 1998; MacFadyen \\& Woosley 1999). In this model the material from the collapsing star feeds the accretion disk, then the accretion energy is being transferred to the jet, which in turn produces gamma rays at some distance from the  central engine. Therefore the whole event cannot last much longer than the existence of a rotationally supported torus in the collapsar center. Within the collapsar model the jet can also be produced by a rotating black hole (BH) which can be spun up by the accreting torus material.  Among the most plausible mechanisms of the energy extraction from the accretion flow are the neutrino-antineutrino annihilation (Mochkovitch et al. 1993), or the magnetic fields (e.g. Blandford \\& Payne 1982; Contopoulos 1995; Proga et al. 2003). The neutrino cooling (e.g. Popham, Woosley \\& Fryer 1999; Di Matteo et al. 2002; Janiuk et al. 2004) is effective only if the accretion rate is  large ($\\dot m \\gtrsim 0.01 M_{\\odot}$ sec$^{-1})$).  Also, a large BH spin ($A_{\\rm Kerr} \\gtrsim 0.9$)  is thought to be a necessary condition for the jet launching: for $A_{\\rm Kerr}\\sim 0.9$, about 1\\% of the accreted rest-mass energy is emitted back as a Poynting jet (Blandford \\& Znajek 1977; McKinney 2005).  On the other hand, the rotationally supported torus may form only when the substantial amount of specific angular momentum is carried in the material.  In our recent article (Janiuk \\& Proga 2008; hereafter Paper I) we studied the problem of whether the collapsing star envelope contained enough specific angular momentum in order to support the formation of the torus.  This condition was parametrized by the so called critical specific angular momentum which in case of a non-rotating BH depends only on its mass. In the present work, we take into account also the BH rotation and the coupling between the specific angular momentum of the accreting material, the BH mass and its spin.  We show that as in Paper I, during the collapse the amount of the rotating material, which was initially available for the torus formation, may later become insufficient to support the torus. Moreover, the spin of the BH is changed by accretion (see e.g.  recent studies by Gammie et al. 2004; King \\& Pringle 2006; Belczy\\'nski et al. 2007). In our models, depending on the accretion scenario, both the spin-up and spin-down of BH are possible, because part of the infalling material has very small specific angular momentum.   The outline of this paper is as follows. In Section \\ref{sec:model}, we briefly describe the model of the evolution of the collapsing star together with the initial conditions. The results are presented in Section \\ref{sec:results}; the GRB durations are estimated in Section \\ref{sec:durations}.  In Section \\ref{sec:diss}, we discuss results in the context of a long GRB production mechanisms and conditions for the distribution of the specific angular momentum in a progenitor star.  In Appendix \\ref{sec:rafal} we provide some formulae for the description of the mass and spin evolution of the BH. ",
        "conclusions": "\\label{sec:diss}  In this article, we studied the collapsar model for long GRBs, powered by accretion onto a spinning BH, which formed from the core of a massive, rotating Wolf-Rayet star. To describe the rotation of the stellar interior, we adopted two different analytical functions, accounting for either a differential rotation (models {\\bf A1}, {\\bf A2}, {\\bf A3}), or a constant ratio between the gravitational and centrifugal forces (models {\\bf D1}, {\\bf D2}, {\\bf D3}). This study is an important test for the rotation models of the GRB progenitor stars (e.g. Heger et al. 2005; Yoon et al. 2006; Detmers et al. 2008).  To describe how the accretion proceeds during the collapse, we adopted three different scenarios: (1) uniform accretion, (2) two phase accretion, first from the poles and then from the torus and (3) only torus accretion. The accretion onto the BH is in our approach a homologous process, in which the subsequent shells of the envelope add their mass to the central object. The angular momentum is also accreted, but the limit for it is the critical angular momentum, to prevent the BH from spinning with $A_{\\rm Kerr} \\ge 1.0$. In this sense, we assume that the whole angular momentum with $l > l_{\\rm crit}$, i.e. in the torus, is transported outwards. We do not invoke any particular mechanism of transport (i.e. the viscosity), and the momentum is taken out by a negligibly small amount of mass (e.g., Pringle 1981). This simplified approach describes well a more realistic situation, in which the matter with small and large angular momentum can be mixed. Therefore some parts of the gas with large $l_{\\rm spec}$ might reach the BH, while some other parts might be blown out with the polar outflow.  We focused on the evolution of the BH spin during the collapse. The large BH spin is important for GRB production in two ways: first, to power the jet emission via the Blandford-Znajek (BZ) mechanism, and second, because it alters the condition for the torus formation, i.e. the critical specific angular momentum. We found that the spin of the BH strongly depends on both the model of the $l_{\\rm spec}$ distribution and on the accretion scenario.  In the torus accretion (i.e. either the second phase of the scenario 2, models {\\bf A2} and {\\bf D2}, or the scenario 3,  models {\\bf A3} and {\\bf D3}), the accreting material has  specific angular momentum always $l_{\\rm spec} \\ge l_{\\rm crit}$. This angular momentum must be transported outwards before reaching the BH, so that the gas which is changing the BH spin has the specific angular momentum equal to $l_{\\rm crit}$. Nevertheless, it is enough to spin up the BH to the maximal rotation, $A_{\\rm Kerr}=0.9999$, which happens in most cases at the very beginning of the collapse. The polar accretion, i.e. the first phase of scenario 2. (models {\\bf A2} and {\\bf D2}), leads only to the BH spin-down in all the models. The uniform accretion scenario is the most complex. In the model {\\bf A1} it leads only to a temporary increase of the BH spin, while during the accretion of the outer shells, the BH is spinning down. In the model {\\bf D1}, the BH spin first decreases, while later during the collapse it may increase, provided the stellar envelope contains enough $l_{\\rm spec}$.  We found that in model {\\bf A1}, the final BH spin after the collapse is always about $A_{\\rm end} \\sim 0.85$, and it does not depend on the normalization the specific angular momentum contained in the stellar envelope, i.e. on $x$. However, the pattern of the BH spin evolution is very sensitive to this parameter. Therefore for small values of $x$ it may happen that even for a short time during the collapse, the BH never reaches a spin $A_{\\rm Kerr}>0.9$, which we consider necessary to power the jet with the BZ mechanism. However, in the same models, the torus does exist and the accretion rate in the torus is large enough to power the jet via the neutrino annihilation. This might lead to a relatively short living (less than $\\sim$ 7-8 s) GRB central engine without a very rapidly spinning BH. On the other hand, for $A_{\\rm Kerr}>0.9$, the stage of a rapidly spinning BH begins very shortly after the collapse has started, and lasts much longer after the accretion rate in the torus has dropped below $\\dot m=0.01$ M$_{\\odot}$s$^{-1}$. For instance, a GRB powered by the BZ mechanism may last almost $\\sim 120$ s, while that powered by the neutrino annihilation (concurrent with the spinning BH) lasted only $\\sim 40$ s. A very short time required for the BH to spin up, while the collapse proceeds, is of the order of $\\sim 1.5$ s.  In model {\\bf D1} the situation is different. Here we do not find any models with only the neutrino-powered bursts, i.e. with a large accretion rate but not accompanied with a rapidly spinning BH. In other words, whenever there exists a torus with a large accretion rate, the BH is spun up to $A_{\\rm Kerr}>0.9$, and the timescale for this spin up is a fraction of a second ($\\sim 0.15$ s). Similarly to model {\\bf A1}, the stage of a large BH spin can last much longer, after the accretion rate has dropped below  $\\dot m=0.01$ M$_{\\odot}$s$^{-1}$. For instance, the BZ-powered burst lasting $\\sim 430$ s is accompanied by a $\\sim 100$ s burst powered by both BZ and neutrino mechanisms.  Observationally, this behavior may have led to three kinds of jets. The first is a very short, lasting between a fraction of a second and few seconds, 'precursor' jet, powered by only the neutrino annihilation, before the BH spins up. The second is an 'early' jet, lasting several tens of seconds and powered by both neutrino and BZ mechanisms. The third is a 'late' jet, powered by only the spinning BH via the BZ mechanism. In our models, we can have the GRB jets with all the three components, as well as the 'orphan precursor' jets, when the BH failed to spin up.  The precursors have been detected by Ginga, BeppoSAX, BATSE, INTEGRAL and Swift in some GRBs (e.g. Murakami et al. 1991; Piro et al. 2005; Lazzati 2005; Romano et al. 2006; McBreen et al. 2006). These GRB precursors are an important observational test for their theoretical models (e.g. Ramirez-Ruiz et al. 2002; Umeda et al. 2005; Morsony et al. 2007;  Wang \\& Meszaros 2008). For instance, in the sample of BATSE bursts, studied by Lazzati (2005), about 20\\% of the bursts had a precursor, which was characterized by a non-thermal spectrum and contained less than 1 per cent of the total counts. The main GRB in these events was delayed with respect to the precursor by 10-200 seconds. As argued by Morsony et al. (2007), who in the 2-D numerical MHD simulations identified three distinct phases during the jet propagation, this large gap in the emission might be a selection effect. Because of different opening angles of these three jets, some observers located at large viewing angles may see a 'dead' phase, i.e. the break in the emission, related to the second jet. Another explanation of the gap  between the precursor and the main jet could be the development of the instabilities in the hyper-accreting disk (Perna et al. 2006; Janiuk et al. 2007), possibly combined with the viewing angle effects.  We therefore conclude that in the present model, the 'dead' phase would  refer to an 'early' jet, which is powered by both neutrino and BZ mechanisms and can be collimated to a much narrower angle than the 'late' jet. For the viewing angle larger than  the 'early' jet but smaller than the 'late' jet opening angle, the observer should see the precursor, followed by a gap in the emission on the order of 40-150 seconds, and then see the 'main' GRB. We also notice that recently, the observation of the bright, long GRB 080319B (Racusin et al. 2008), seems to have confirmed that the jet's opening angle may vary, indicating for the two types of jets.  Finally, comparing our current models with the results presented for a non-rotating BH (Paper I), we notice that the GRB durations are similar in case of model {\\bf A1}, i.e. $\\sim 40$ s vs. $\\sim 50$ for the Schwarzschild and Kerr BH main jet, respectively. In model {\\bf D1}, the discrepancy is more pronounced, namely $\\sim 50$ s vs. $\\sim 100$ s, respectively. On the other hand, in the current work, the model  {\\bf D1} produces GRBs powered by neutrino annihilation only for a very narrow range of parameters (i.e. $x$), while in Paper I for this model we found no limitations for $x$."
    },
    "0807/0807.2910_arXiv.txt": {
        "abstract": "{Some recent observational results impose significant constraints on all the models that have been proposed to explain the Galactic black-hole X-ray sources in the hard state.  In particular, it has been found that during the hard state of Cyg X-1 the power-law photon number spectral index, $\\Gamma$, is correlated with the average time lag, $<t_{\\rm lag}>$, between hard and soft X-rays. Furthermore, the peak frequencies of the four Lorentzians that fit the observed power spectra are correlated with both $\\Gamma$ and $<t_{\\rm lag}>$. } {We have investigated whether our jet model can reproduce these correlations.} {We performed Monte Carlo simulations of Compton upscattering of soft, accretion-disk photons in the jet and computed the time lag between hard and soft photons and the power-law index $\\Gamma$ of the resulting photon number spectra.} {We demonstrate that our jet model naturally explains the above correlations, with no additional requirements and no additional parameters. } {} ",
        "introduction": "The states of Galactic black-hole X-ray sources (GBHs) are determined from the following quantities:  i) their disk fraction, which is the ratio of the disk flux to the total flux, both unabsorbed, at $2-20$ keV; ii) the photon number spectral index, $\\Gamma$, of the  hard band,  power-law component at energies below any break or cutoff;  iii) the root-mean-square (rms) power in the power spectral density (PSD) integrated from 0.1--10 Hz;  iv) the integrated rms amplitude of any quasi-periodic oscillation (QPO) detected in the range 0.1--10 Hz \\citep{remi06}.  In particular, the so-called hard state is characterized by a power-law X-ray spectrum with photon  number spectral index $1.5 \\simless \\Gamma \\simless 2.1$, which suffers an exponential cutoff at $\\sim$ 100  or few 100 keV, and contributes more than 80\\% of the 2--20 keV flux.  Several models have succeeded in reproducing the observed X-ray spectrum in this state \\citep[see e.g.][]{tita94,esin97,pout99,reig03,gian04}. When a GBH is in the hard state, radio emission  is also detected and a jet is either seen or inferred \\citep{fend01}. Again, several models are successful in reproducing the energy spectrum from the radio domain to the hard X-rays \\citep[see e.g.][]{mark01,vada01,corb02,mark03,gian05}.  The multiplicity of models that can fit the time average spectrum of GBHs well indicates that this alone is  not enough to distinguish the most realistic among them.  Timing data of the observed intensity in various energy bands provide a totally different ``dimension\" than the spectral one, and can in principle provide information on the location, size, physical conditions and kinematics of the plasma responsible for the X-ray emission in these objects. For example, the X-ray light curves of GBHs  in different energy bands show similar variations, but those in the higher energy band generally lag the variations detected in the lower energy band  \\citep{nowa99,ford99}. The time delay, $t_{\\rm lag}$, usually depends  on frequency, $\\nu$.  For Cyg X-1  the time delays decrease with increasing frequency roughly as  $t_{\\rm lag} \\propto \\nu^{-\\beta}$, where $\\beta \\approx 0.7$. As constraining as it is, this relation has also been explained in more than one way  \\citep[see e.g.][]{pout99,reig03,kord04}.  Even more constraining to models is the fact that the high-frequency power spectrum flattens as the photon energy increases \\citep{nowa99}. Equivalent to this is the fact that the width of the  temporal autocorrelation function of the light curves of Cyg X-1  decreases with increasing photon energy \\citep{macc00}. These observational facts are contrary to what one expects from simple Comptonization models, yet possible  explanations for them  have also been offered \\citep{koto01,gian04}.  Significant advances in our understanding of how the X-ray mechanism in GBHs (and compact accreting objects in general) works  can be achieved when we combine variability and spectral information. This has been possible in the last few years, due mainly to \\xte.  For example, in the case of Cyg X-1, \\citet{pott03}  used 130 \\xte\\ observations between 1998 and 2001 to study  the long-term evolution of the source's power and energy spectrum in detail. They used 512 s long light curves to estimate the  power spectral density (PSD), which they fitted with a model that consisted of the sum of multiple Lorentzian profiles. They generally achieved good fits using four broad Lorentzian components with typical peak frequencies of $\\nu_1\\sim 0.2$ Hz, $\\nu_2\\sim 2$ Hz, $\\nu_3\\sim 6$ Hz, and $\\nu_4\\sim 40$ Hz, which vary in unison with time: their ratios remain constant despite each frequency varies by up to a factor of $\\sim 5$. Up to now, no physical explanation of these Lorentzians has been offered. As for the energy spectrum, they find that a simple model consisting of a power-law component of index $\\Gamma$, a multi-temperature disk black body and a reflection from neutral material component could fit the energy spectrum of the source below 20 keV.  They present convincing evidence that the timing and spectral properties of the source are closely linked. In fact, a number of the correlations they present impose stringent constraints on all the  models proposed so far. These correlations are:  a) The observed photon number spectral index, $\\Gamma$, correlates with the observed average time lag, $<t_{\\rm lag}>$, between 2--4 keV and 8--13 keV, averaged over the 3.2--10  Hz frequency band, in the sense that  $<t_{\\rm lag}>$ increases as the spectrum steepens. Therefore, a model that can successfully explain the spectral slope $\\Gamma$ should also predict the right amount of time delay between these specific energy bands and vice versa: a model that can fit well  the time lag vs. Fourier frequency relation, should also predict the appropriate value of the index $\\Gamma$  for a given time delay.  b) The observed photon number spectral index, $\\Gamma$, also  correlates with the Lorentzian peak frequencies: the frequencies increase as the spectrum steepens \\citep[see also][]{shap06,shap07}.  In summary, the \\citet{pott03} results show that a specific value of  $\\Gamma$, say $\\Gamma = 2$, which is relatively easy to obtain with a simple Comptonization model, must also be associated with {\\it specific} values of the Lorentzian peak frequencies and a {\\it specific} value of  $<t_{\\rm lag}>$! These are challenging relations that  any model must be able to explain. In this work we investigate whether the jet model  that we have proposed can meet these constraints under reasonable assumptions.  We think that the jets in GBHs  play a central role in all the observed phenomena,  not only in the radio emission. We have shown that Compton upscattering in the jet  of soft photons from the accretion disk can explain: 1) the X-ray spectrum and the time-lag dependence on Fourier frequency \\citep[][hereafter Paper I]{reig03}; 2) the narrowing of the autocorrelation function and the increase of the  rms amplitude of variability with increasing photon energy  \\citep[][hereafter Paper II]{gian04}; 3) the energy spectrum from radio to hard X-rays of XTE J1118+480, the only source for which we have simultaneous observations for all energy bands \\citep[][hereafter Paper III]{gian05}.  In this paper we show that the jet model proposed and developed in the above work can explain the $\\Gamma$ -- $<t_{\\rm lag}>$ relation of \\citet{pott03}, with no additional parameters or requirements,  if we let just two model parameters, namely the radius at the base of the jet and the optical depth along its axis, vary over a reasonable range of values. At the same time, the model can also explain the \"$\\Gamma$ -- Lorentzian peak frequencies $\\nu_i$\"  relation  qualitatively if we identify $1/\\nu_i$, $i=1, 2, 3, 4$, with quantities proportional to the characteristic timescale $M_{\\rm out}/{\\dot M}_{\\rm out}$,  where ${\\dot M}_{\\rm out}$ is the rate of mass outflow in the jet and $M_{\\rm out}$  the available mass to power the jet. Finally, the model parameter variations that are needed to explain the  \\citet{pott03} relations also predict that, on long time scales,  the radio variations should also correlate with $\\Gamma$,  which, as we show, is indeed the case in Cyg X-1.  In Sect. 2 we describe briefly our model, in Sect. 3 we present and discuss our results, and we close with our conclusions in Sect. 4. ",
        "conclusions": "In a series of three papers (Papers I, II, and III) it has been shown that our jet model can explain a number of observations regarding black-hole X-ray binaries in the hard state.  These are: a) the energy spectrum from radio to hard X-rays;  b) the time-lag of the hard photons relative  to the softer ones as a function of Fourier frequency; c) the flattening of the power spectra  at high frequencies with increasing photon energy and the narrowing  of the autocorrelation function.  In this paper we investigate whether our model can explain the long term variations in the spectral and timing properties of Cyg X-1, which are closely linked. The main result of our work is that we can reproduce the $\\Gamma-<t_{\\rm lag}>$ correlation of \\citet{pott03}  if we assume reasonable variations in just two  model parameters around their mean values: the  width of the jet at its base and the optical depth along the jet's axis (or equivalently the electron density at the base of the jet). We find that, if $\\tau_{\\parallel}\\approx 1.5-10$, $R_0\\approx 5-50 r_s$ and $\\tau_{\\parallel}\\propto R_0^{-1}$, then the model predicts a $\\Gamma-<t_{\\rm lag}>$ relation that is almost  identical to the one that \\citet{pott03} observed. Furthermore, if the flux variability operates on a time scale which is proportional to $M_{\\rm out}/\\dot{M}_{\\rm out}$ (i.e. the characteristic time needed to eject all the available matter to the jet), then the same $\\tau$ and $R_0$ variations can also explain the {\\it shape} of the ``spectral index - Lorentzian peak frequencies\" correlation of \\citet{pott03}. Finally, given the model parameter variations that are needed to explain the observed \"spectral - timing\" correlations, the model predicts a \"radio flux -- $\\Gamma$\" correlation, which, as we show, does indeed hold in Cyg X-1.  The $\\tau_{\\parallel}\\propto R_0^{-1}$ relation, which is  necessary to explain  the observed  $\\Gamma-<t_{\\rm lag}>$ relation, is a  natural outcome of the model if simply $\\dot{M}_{\\rm out}\\propto R_0$. We propose that this is the underlying ``fundamental\" relation that governs the  {\\it long-term} evolution of the physical characteristics of the jet,  hence of the observational spectral and timing properties of the source. The physical reason for this relation is unclear at the moment; but whatever its origin, the relation itself does make sense because as the rate of the mass ejected in the jet increases, it seems reasonable to assume that the jet will become ``bigger\", i.e. $R_0$ will also increase. In fact, in this way, one can explain reasonably well all the long-term observed changes with variations in one fundamental physical parameter of the system, namely the accretion rate. As the accretion rate increases, $\\dot{M}_{\\rm out}$ should also increase. If the jet size increases proportionally to $\\dot{M}_{\\rm out}$, then $n_0 \\propto R_0^{-1}$, and  hence the $\\Gamma$ -- $<t_{\\rm lag}>$ relation.  We hope that this ``accretion rate -- ejection rate -- jet size\"  relation will provide an important and useful hint for all models that try to explain the connection between accretion flows and outflows, hence the formation of jets in regions close to the central engines in accreting compact objects.  We find that the $\\dot{M}_{\\rm out}\\propto R_0$  relation predicts that the radio flux should increase linearly with the photon number spectral index. To the best of our knowledge, such a relation has never been suggested or even detected either in Cyg X-1 or in any other GBH in the past.  We show that a linear  ``radio flux -- spectral index\" relation does exist  on long time scales in Cyg X-1. We believe that this result strongly supports our view that a) X-rays in Cyg X-1 are produced by Compton upscattering in the jet of soft photons from the accretion disk and b) the jet's size and optical depth evolve with time because $\\dot{M}_{\\rm out}\\propto R_0$. Furthermore,  that the ``radio flux --  spectral index\" relation breaks at index values higher than $\\sim 2.2$,  provides interesting hints into what may be happening as the source moves to the intermediate and thermal state. Most probably the magnetic field  decreases and/or there is weaker particle acceleration in the jet (i.e. the electron distribution index, $\\alpha$, becomes larger)   as the accretion rate increases (as inferred from the increase in the disk thermal emission).  The explanation of the short time-scale variability properties of Cyg X-1 is significantly more challenging. Our model cannot explain why  four Lorentzian components exist in the source's power spectrum (assuming of course that the Lorentzians are the fundamental building blocks of the power spectrum in Cyg X-1 and other GBHs). However, it can explain, to some extent, the way they vary with the source's photon  number spectral index. If indeed the long-term evolution of the source is governed by variations in $R_0$ and $\\dot{M}_{\\rm out}$, it is natural to assume that the short-term variations are caused by short-term, random variations of $\\dot{M}_{\\rm out}$ (and hence of $n_0$ or $\\tau_{\\parallel}$, and finally X-ray flux), around its mean value. We find that, if these variations operate on a time scale that is proportional to the characteristic  time needed to eject all the available matter to the jet, then the model can explain the $\\it shape$ of the ``$\\Gamma$ - Lorentzian peak frequencies\" relations of \\citet{pott03}. There is no reason why the model  ``$\\Gamma -$ inverse time needed to eject the  mass to the jet\" relation should have the same shape as the observed ``$\\Gamma -$ Lorentzian peak frequencies\". Thus we do not find this agreement to be coincidental. Instead, it is one more evidence in favor of the jet model and the parameters' evolution we present in this work, and we find it  operates on long time scales in Cyg X-1.  Our results suggest that the characteristic  frequencies seen in the power spectrum should increase with increasing size of the source, i.e. $R_0$, a result that goes against current views. However, the characteristic frequencies may not scale directly with the source size, but with accretion rate, hence  $\\dot{M}_{\\rm out}$.  As argued above, it seems reasonable to think that, as the accretion rate increases, so too the mass ejection rate to the jet. Perhaps, then, it is the accretion rate (through $\\dot{M}_{\\rm out}$)  that drives the variations in both $R_0$ and on the characteristic time scales. A  physical explanation as to how this may happen and as to the value (and number) of characteristic time scales, requires knowledge of the details of the acceleration/ejection process of the jet, and is beyond the scope of the present work. However, as we mentioned above, we hope that the results we present here will help the investigation of the proper mechanism responsible for the launch of the jets and winds in accreting compact objects.  We note that our conclusion, presented in Sect. 3.2, that the Lorentzian peak frequencies $\\nu_i$ {\\it increase} with the size $R_0$ of the base of the jet, holds for a given source (in our case Cyg X-1).  If one wanted to extend our model to other sources, then the mass of the black hole comes in.  The argument goes as follows:  Making the reasonable assumption that $M_{\\rm out} \\propto R_0^2$ and using $n_0 \\propto R_0^{-1}$, we find that the $\\nu_i$ scale as $1/R_0$.  Since $R_0$ is typically a few times $r_g$, we conclude that {\\it the Lorentzian peak frequencies $\\nu_i$ are inversely proportional to the mass of the black hole}.  In closing, we want to point out that the above results and conclusions are consistent with the fact that, in the thermal state of Cyg X-1, the time lags versus Fourier frequency are {\\it identical} to those in the hard state \\citep{pott00}.  In the thermal state, the radio emission is very weak and possibly undetectable,  {\\it but we think that the jet is there} and that it produces, by Comptonization,  both the time lags versus Fourier frequency and the steep power-law energy spectrum above $\\sim 10$ keV. According to our model, the suppression of both the radio and the X-ray emission in this state should be caused by a magnetic field suppression and a weaker particle acceleration, respectively."
    },
    "0807/0807.1122_arXiv.txt": {
        "abstract": "{Numerical simulations of forced turbulence in elongated shearing boxes are carried out to demonstrate that a nonhelical turbulence in conjunction with a linear shear can give rise to a mean-field dynamo. Exponential growth of magnetic field at scales larger than the outer (forcing) scale of the turbulence is found. Over a range of values of the shearing rate $S$ spanning approximately two orders of magnitude, the growth rate of the magnetic field is proportional to the imposed shear, $\\gamma\\propto S$, while the characteristic spatial scale of the field is $\\lb\\propto S^{-1/2}$. The effect is quite general: earlier results for the nonrotating case by \\cite{}Yousef et~al. (2008) are extended to shearing boxes with Keplerian rotation; it is also shown that the shear dynamo mechanism operates both below and above the threshold for the fluctuation dynamo. The apparently generic nature of the shear dynamo effect makes it an attractive object of study for the purpose of understanding the generation of magnetic fields in astrophysical systems.} ",
        "introduction": "Turbulence is generally considered to play a fundamental role in the generation and maintenance of magnetic fields found in a wide range of astrophysical systems. It is a numerically well established property of turbulence to amplify magnetic fluctuations at the same or smaller scales than the scales of the turbulent motions \\citep{meneguzzi81,sch04,haugen04,iskakov07,sch07}. This type of dynamo is known as small-scale, or fluctuation, dynamo. It is believed to be a universal property of turbulent systems and, at least in the case of large magnetic Prandtl numbers, a simple theoretical picture exists of the field amplification via random stretching by turbulent motions (\\citealt{moffatt64,zeldovich84}; see also \\citealt{sch04,SC07}). It often turns out to be more difficult to either establish or explain the generation of magnetic fields at much larger scales than the turbulence scale. Such fields are observed or believed to exist in many astrophysical bodies (stars, galaxies, disks), so the question of their origin is an important theoretical challenge.  Mechanisms for the generation of such large-scale fields are known as large-scale, or mean-field, dynamos. A motley of such mean-field dynamos has been studied in the literature \\citep[e.g.,][]{moffatt78,krause80,brand05review,raedler07}. The precise way in which they operate often seems to be system dependent and introducing ever more realistic features into one's theoretical model produces ever more complicated behaviour. While such modelling is necessary for quantitative understanding, it is quite interesting to ask what are the generic ingredients required to produce large-scale fields. One such ingredient appears to be the presence of net kinetic helicity in the system: in many mean-field dynamos, the key generation mechanism is the so-called $\\alpha$ effect \\citep{steenbeck66}, whereby an assembly of non-mirror-symmetric velocity fluctuations having a nonzero net helicity are responsible for magnetic-field generation. The existence of a mean-field dynamo in helically forced turbulence is well established numerically \\citep{brand01,maron02,brand08}. However, the requirement of net helicity may somewhat limit the applicability of the $\\alpha$ effect. It is also far from certain that the direct link between kinetic helicity and mean-field generation via the $\\alpha$ effect found in the model case where the helicity is injected by the random forcing, carries over to the cases where the helicity arises naturally (e.g., in a rotating convective layer; see \\citealt{cattaneo06}). Therefore, there is a strong motivation for seeking alternative mean-field dynamo processes driven by nonhelical turbulence.  It is clear that a nonhelical turbulence by itself cannot make large-scale fields. In recent years, many authors have argued that large-scale magnetic fields can be generated by nonhelical velocity fluctuations when acted upon by a large-scale shear: theoretical paradigms put forward to support such a dynamo, which we refer to as the shear dynamo, have included the shear-current effect \\citep{RK03,RK04}, the stochastic $\\alpha$ effect (\\citealt{vishniac97,silantev00,fedotov03,fedotov06,proctor07,brand07}; see, however, \\citealt{kleeorin08}), negative-diffusivity type theories \\citep{urpin99a,urpin99b,ruediger01,urpin02,urpin06}, shear amplification of the fluctuation-dynamo-generated  small-scale fields \\citep{blackman98}. While there is not yet agreement between theoreticians about the validity or areas of practical applicability of these models, it is clear that they are addressing a fundamental issue. Indeed, shear is an extremely generic property of astrophysical systems, so the idea of a { shear dynamo} gives us a particularly attractive scenario for ubiquitous generation of large-scale magnetic fields.  Until recently, a numerical demonstration of this type of dynamo remained elusive. Originally motivated by the predictions of \\citet{RK03,RK04}, we have previously performed numerical simulations of nonhelical turbulence with a superimposed linear shear and demonstrated the existence of the shear dynamo \\citep{yousef08}. Theoretical understanding of these numerical results is still poor. More work and, we believe, more information gathered from numerical experiments are needed in order to make progress towards understanding the properties of this dynamo and the underlying physical processes that produce it. This paper, which is an extension of the work by \\citet{yousef08}, aims at presenting a collection of new numerical results regarding the existence and behaviour of the shear dynamo in various regimes. We focus on three different cases of astrophysical interest, namely shear dynamo in the presence of forced nonhelical turbulence and a linear velocity shear, shear dynamo in the presence of forced nonhelical turbulence and Keplerian differential rotation, and finally shear dynamo in the presence of forced nonhelical turbulence, Keplerian differential rotation and a small-scale fluctuation dynamo. In order to study the effects of shear and rotation, we adopt a local rotating shearing sheet model. This model and the corresponding numerical set-up are presented in \\secref{sec:model}. In \\secref{sec:linear}, we consider the case of linear shear without rotation. Results for the Keplerian regime are presented in \\secref{sec:kepler}. \\Secref{sec:fd} describes some preliminary results on the shear dynamo in the presence of small scale magnetic fluctuations generated by the fluctuation dynamo. A short discussion concludes the paper (\\secref{sec:conc}). ",
        "conclusions": "\\label{sec:conc}  Can nonhelical turbulence in combination with large-scale velocity shear act as a mean-field dynamo and generate magnetic fields with length scales much larger than the outer scale of the turbulence? This has been the subject of considerable recent debate \\citep{vishniac97,urpin99a,urpin99b,silantev00,urpin02,fedotov03,RK03,RK04,brand05,fedotov06,raedler06,ruediger06,raedler07,proctor07,brand07,kleeorin08}. To our knowledge, the paper by \\citet{yousef08} and this paper present the first set of dedicated numerical experiments that demonstrates that such a generation mechanism is feasible. However, in retrospect, one might conjecture that the shear-dynamo might have already been seen in several earlier numerical studies that combined large-scale flows, and consequently large-scale shear, with nonhelical forcing at smaller (or, in some cases, similar) scales and reported generation of magnetic fields at scales larger than the forcing scale \\citep{brand05,ponty05,mininni05,shapovalov06}.  We have carried out a suite of numerical experiments on the shear dynamo effect in vertically elongated shearing boxes and for magnetic Reynolds numbers subcritical with respect to the fluctuation dynamo. For the values of the imposed shear $S$ between $1/8$ and $8$ (corresponding to $S\\tau\\sim 0.04 \\dots 3$, where $\\tau$ is the turnover rate of the randomly forced velocity fluctuations), we have found that the dynamo growth rate is $\\gamma\\propto S$ and the characteristic length scale of the generated mean magnetic field is $\\lb\\propto 1/\\sqrt{S}$.  The first key result of this paper, compared with the earlier study by \\citet{yousef08}, is that the shear dynamo works both in the nonrotating case and for the case of Keplerian rotation ($\\Omega=-2S/3$). There does not appear to be much difference, qualitative or quantitative, between the rotating and nonrotating cases, although perhaps it would be interesting to look at non-Keplerian cases and try to identify the role of rotation via a parameter scan in $\\Omega$ independent of the one in $S$.  The second key result, claimed on the basis of only a preliminary study, is that the shear dynamo works both for situations that are sub- and supercritical with respect to the fluctuation (small-scale) dynamo. In the latter case, the overall magnetic energy grows very quickly due to the fluctuation dynamo effect independent of the presence of the shear. Imposing the shear on the magnetohydrodynamic turbulence resulting from the saturation of the fluctuation dynamo, leads to the emergence of magnetic fields that have spatial scales larger than the outer scale of the turbulence and that fluctuate on very long time scales compared to the turbulent turnover time.  Thus, the shear dynamo effect appears to be quite general and robust. As the combination of a shear flow and turbulence is a very common feature in astrophysical systems, the shear dynamo potentially represents a generic mechanism for making large-scale fields. While much needs to be understood about its properties before its relevance to real astrophysical systems can be more than an appealing speculation, the simplicity of the idea of the shear dynamo certainly makes it a worthwhile object of study. It is also important to determine how generic the shear dynamo is and how it combines with other large-scale features present in real astrophysical systems: various differential rotation laws, temperature and density gradients, linear instabilities, etc..  Studies in this vein are already being undertaken. For example, recent numerical experiments by \\citet{kapyla08} have shown that large-scale dynamo action is also possible in local simulations of magnetoconvection with imposed horizontal shear. They also report a growth rate $\\gamma \\sim S$ and find large-scale dynamo action for magnetic Reynolds numbers above the critical threshold for the fluctuation dynamo. Another topical recent study is by \\citet{gressel08}, who simulated the supernova-driven galactic turbulence and found that they needed to impose a linear velocity shear to obtain the amplification of a large-scale field. These studies clearly demonstrate the key role of shear in producing a mean-field dynamo. However, in comparing their results to ours, one has to keep in mind that their simulations had rotation and vertical stratification, so the turbulence in these simulations is likely to be helical and may also host an $\\alpha$ effect.  In motivating our choice of Keplerian rotation law, we mentioned the possible relevance to accretion-disk turbulence, which is believed to be driven by the magnetorotational instability (MRI) \\citep{balbus03}. Given a (weak) large-scale field, the MRI will generate velocity and magnetic-field fluctuations at small scales. These fluctuations, in conjunction with Keplerian rotation and shear, must then amplify the large-scale field to close the loop. The mechanisms for such an ``MRI dynamo'' have been discussed and simulated for some time \\citep[e.g.,][]{brand95,hawley96,fromang07,rincon07b,rincon08,lesur08}. It is tempting to observe in the context of the results reported above that a combination of small-scale turbulence (magnetohydrodynamic turbulence in the case of the MRI) and large-scale shear does indeed appear to work as a dynamo giving rise to a large-scale azimuthal magnetic field ($B_y$ in the shearing sheet model). We note, however, that \\citet{lesur08}, who have analysed this process in detail, find some important differences between what happens in MRI-driven shearing sheet simulations and the forced case studied by us.  To conclude, we believe that the discovery of shear dynamo has opened a number of new and exciting avenues of research and produced some promising leads towards unravelling the ways in which cosmic magnetic fields emerge. Further investigations will help assess the range of applicability and relevance of the shear dynamo effect and the physical mechanisms that are responsible for it."
    },
    "0807/0807.1913_arXiv.txt": {
        "abstract": "Rotation plays a major role in the evolution of massive stars. A revised grid of stellar evolutionary tracks accounting for rotation has recently been released by the Geneva group and implemented into the Starburst99 evolutionary synthesis code. Massive stars are predicted to be hotter and more luminous than previously thought, and the spectral energy distributions of young populations mirror this trend. The hydrogen ionizing continuum in particular increases by a factor of up to 3 in the presence of rotating massive stars. The effects of rotation generally increase towards shorter wavelengths and with decreasing metallicity. Revised relations between star-formation rates and monochromatic luminosities for the new stellar models are presented. ",
        "introduction": "The star-formation rates of galaxies are commonly determined from the integrated light emitted at a certain wavelengths, such as the $V$ band. Comparison with theoretically predicted  mass-to-light ($M/L$) ratios can in principle provide the total stellar mass and, in combination with an appropriate timescale, the star-formation rate. This fundamental methodology goes back to Tinsley's (1980) pioneering work. Although the reasoning is immediately intuitive, two major challenges need to be tackled.  The first challenge is the a priori unknown stellar initial mass function (IMF). Observed values of $M/L_{\\rm V}$ in galaxy centers are around 5~--~100. At the same time we know that the main contributors to the galaxy light are upper main-sequence (MS) and evolved low-mass stars. These stars have $M/L_{\\rm V} \\approx 1$. To put these numbers into perspective, the average $M/L_{\\rm V}$ of all known stars within 20~pc of the Sun is about 1~--~2 (Faber \\& Gallagher 1979), and an early-type MS star has $M/L_{\\rm V} \\approx 10^{-2}$. Since the mass-to-light ratio in galaxy centers is not too sensitive to dark matter (at least for disk galaxies; cf. E. Brinks' talk at this conference), the apparent discrepancy suggests that most of the stellar mass is hidden from view because most stars have lower luminosity, and therefore have lower mass than indicated by the spectrum. For the purpose of this paper I will assume we can correct for this effect by assuming a known, universal IMF (Kroupa 2007 and this conference).  The second challenge involves the $M/L$ of individual stars of all masses. At the high-mass end, this quantity is not accessible to direct measurements and can only be predicted by stellar evolution models: masses are poorly known because of the scarcity of very massive binaries with mass determinations, and luminosities are elusive because most of the stellar light is emitted in the ionizing ultraviolet. The purpose of this paper is to discuss how the latest generation of stellar evolution models including rotation differs from its predecessor, and how these new models affect the predictions of the evolutionary synthesis code Starburst99 (Leitherer et al. 1999; V\\'azquez \\& Leitherer 2005). Some of the results presented here can be found in V\\'azquez et al. (2007). ",
        "conclusions": ""
    },
    "0807/0807.4477_arXiv.txt": {
        "abstract": "Within both dynamical and thermodynamical approaches using the equation of state for neutron-rich nuclear matter constrained by the recent isospin diffusion data from heavy-ion reactions in the same sub-saturation density range as the neutron star crust, the density and pressure at the inner edge separating the liquid core from the solid crust of neutron stars are determined to be $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 together with the observed minimum crustal fraction of the total moment of inertia allow us to set a new limit for the radius of the Vela pulsar significantly different from the previous estimate. It is further shown that the widely used parabolic approximation to the equation of state of asymmetric nuclear matter leads systematically to significantly higher core-crust transition densities and pressures, especially with stiffer symmetry energy functionals. ",
        "introduction": "Having been the major testing grounds of our knowledge on the nature of matter under extreme conditions, neutron stars are among the most mysterious objects in the universe. To understand their structures and properties has long been a very challenging task for both the astrophysics and the nuclear physics community~\\cite{Lat04}. Theoretically, neutron stars are expected to have a solid inner crust surrounding a liquid core. Knowledge on properties of the crust plays an important role in understanding many astrophysical observations~\\cite% {BPS71,BBP71,Pet95a,Pet95b,Lat00,Lat07,Ste05,Lin99,Hor04,Bur06,Owe05}. The inner crust spans the region from the neutron drip-out point to the inner edge separating the solid crust from the homogeneous liquid core. While the neutron drip-out density $\\rho _{out}$ is relatively well determined to be about $4\\times 10^{11}$ g/cm$^{3}$ \\cite{Rus06}, the transition density $% \\rho _{t}$ at the inner edge is still largely uncertain mainly because of our very limited knowledge on the equation of state (EOS), especially the density dependence of the symmetry energy, of neutron-rich nucleonic matter~\\cite{Lat00,Lat07}. These uncertainties have hampered our accurate understanding of many important properties of neutron stars~\\cite{Lat04,Lat00,Lat07}.  Recently, significant progress has been made in constraining the EOS of neutron-rich nuclear matter using terrestrial laboratory experiments (See Ref.~\\cite{LCK08} for the most recent review). In particular, the analysis of isospin-diffusion data \\cite{Tsa04,Che05a,LiBA05c} in heavy-ion collisions has constrained tightly the density dependence of the symmetry energy in exactly the same sub-saturation density region around the expected inner edge of neutron star crust. Moreover, the obtained constraint on the symmetry energy was found to agree with isoscaling analyses in heavy-ion collisions \\cite{She07}, the isotopic dependence of the giant monopole resonance in even-A Sn isotopes \\cite{LiT07}, and the neutron-skin thickness of $^{208}$Pb \\cite{Ste05b,LiBA05c,Che05b}. In this paper, using the equation of state for neutron-rich nuclear matter constrained by the recent isospin diffusion data from heavy-ion reactions in the same sub-saturation density range as the neutron star crust, we determine the inner edge of neutron star crusts. Consequently, the limit on the radius of the Vela pulsar is significantly different from the previous estimate. In addition, we find that the widely used parabolic approximation (PA) to the EOS of asymmetric nuclear matter enhances significantly the transition densities and pressures, especially with stiffer symmetry energy functionals. ",
        "conclusions": "In summary, the density and pressure at the inner edge separating the liquid core from the solid crust of neutron stars are determined to be $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, using the MDI EOS of neutron-rich nuclear matter constrained by the recent isospin diffusion data from heavy-ion reactions in the same sub-saturation density range as the neutron star crust. These constraints allow us to set a new limit on the radius of the Vela pulsar. It is significantly different from the previous estimate and thus puts a new constraint for the mass-radius relation of neutron stars. Furthermore, we find that the widely used parabolic approximation to the EOS of asymmetric nuclear matter leads systematically to significantly higher core-crust transition densities and pressures, especially for the energy density functional with stiffer symmetry energies. Our results thus indicate that one may introduce a huge error by assuming {\\it a priori} that the EOS is parabolic with respect to isospin asymmetry for a given interaction in locating the inner edge of neutron star crust."
    },
    "0807/0807.3844_arXiv.txt": {
        "abstract": "Cosmic rays with energies exceeding $10^{17}$~eV are frequently registered by measurements of the fluorescence light emitted by extensive air showers.  The main uncertainty for the absolute energy scale of the measured air showers is coming from the fluorescence light yield of electrons in air.  The fluorescence light yield has been studied in laboratory experiments.  Pioneering measurements between 1954 and 2000 are reviewed. ",
        "introduction": "\\Label{intro}  Cosmic rays with energies exceeding $10^{17}$~eV are frequently observed by measurements of the fluorescence light emitted in air showers. The latter are induced by high energy cosmic rays interacting in the Earth's atmosphere initiating a cascade of secondary particles. Relativistic electrons (and positrons) are the most numerous charged particles in air showers. On their way through the atmosphere they excite nitrogen molecules. The nitrogen molecules release their excitation energy partly through isotropic emission of fluorescence light.  The fluorescence light is measured with imaging telescopes in air shower experiments, allowing for a three-dimensional reconstruction of the shower profile in the atmosphere. The main uncertainty of the absolute energy scale of the measured showers is coming from the fluorescence light yield of electrons in air.  The bulk of the fluorescence light is induced by electrons with MeV energies. Thus, the emission mechanism can be studied in laboratory experiments using particles, mainly electrons, from radioactive sources or particle accelerators. Particles are injected into volumes of nitrogen and air under well controlled circumstances and the fluorescence light yield is measured as a function of various parameters, like electron energy, gas pressure, temperature, and humidity.  Recent developments have been discussed during the 5th Fluorescence Workshop in El Escorial, Spain from September 16th to 20th, 2007.  The results are summarized in an accompanying article \\Cite{summary}.  The present article summarizes measurements of the fluorescence light yield conducted between 1954 and 2000.  The objective is to compile information relevant for air shower observations from the (sometimes hardly accessible) historical papers.  Their implications on the contemporary understanding of the subject are discussed in \\Cite{summary}.  \\Label{sec:fl_history} ",
        "conclusions": ""
    },
    "0807/0807.4180_arXiv.txt": {
        "abstract": "Extrasolar terrestrial planets with the potential to host life might have large obliquities or be subject to strong obliquity variations. We revisit the habitability of oblique planets with an energy balance climate model (EBM) allowing for dynamical transitions to ice-covered snowball states as a result of ice-albedo feedback. Despite the great simplicity of our EBM, it captures reasonably well the seasonal cycle of global energetic fluxes at Earth's surface.  It also performs satisfactorily against a full-physics climate model of a highly oblique Earth--like planet, in an unusual regime of circulation dominated by heat transport from the poles to the equator.  Climates on oblique terrestrial planets can violate global radiative balance through much of their seasonal cycle, which limits the usefulness of simple radiative equilibrium arguments. High obliquity planets have severe climates, with large amplitude seasonal variations, but they are not necessarily more prone to global snowball transitions than low obliquity planets. We find that terrestrial planets with massive $\\rm CO_2$ atmospheres, typically expected in the outer regions of habitable zones, can also be subject to such dynamical snowball transitions. Some of the snowball climates investigated for $\\rm CO_2$--rich atmospheres experience partial atmospheric collapse. Since long-term $\\rm CO_2$ atmospheric build-up acts as a climatic thermostat for habitable planets, partial $\\rm CO_2$ collapse could limit the habitability of such planets. A terrestrial planet's habitability may thus depend sensitively on its short-term climatic stability. ",
        "introduction": "\\label{obl_sec:intro}  The Earth's obliquity is remarkably stable: the angle between the spin--axis and the normal to the orbital plane varies by no more than a few degrees from its present value of $\\sim 23.5\\degr$.  This stability is maintained by torque from the Moon \\citep{laskar_et_al1993,nerondesurgy+laskar1997}.  Even within our own Solar System, though, the obliquity of other terrestrial planets has varied significantly more; the analysis of \\citet{laskar+robutel1993} indicates that Mars' obliquity exhibits chaotic variations between $\\sim 0\\degr$ and $\\sim 60\\degr$.   How does climate depend on obliquity and its possible variations in time?  How does the range of orbital radii around a star at which a planet could support water-based life depend on the planet's obliquity?  Has the stability of Earth's obliquity made it a more climatically hospitable home?  The answers to these questions will be important to evaluate the fraction of stars that have potentially habitable planets.  There are now more than 300 extrasolar planets known,\\footnote{See http://exoplanet.eu/.} several of which are close to the terrestrial regime with masses less than 10 times that of the Earth (e.g., \\citealt{beaulieu_et_al2006}, \\citealt{udry_et_al2007}, \\citealt{bennett_et_al2008}, \\citealt{mayor_et_al2008}).  {\\it COROT}, which has already launched, and {\\it Kepler}, scheduled to launch in less than one year, are dedicated space-based transit-detecting observatories that will monitor a large number of stars to detect the small decreases in stellar flux that occur when terrestrial planets cross in front of their host stars \\citep{baglin2003,borucki_et_al2003,borucki_et_al2007}.  These missions are expected to multiply by perhaps several hundredfold or more the number of known terrestrial planets, depending on the distribution of such planets around solar-type stars (\\citealt{borucki_et_al2007, borucki_et_al2003, basri_et_al2005}; although see revised predictions in \\citealt{beatty+gaudi2008}).  NASA and ESA have plans for ambitious future missions to obtain spectra of nearby Earth-like planets in the hope that they would reveal the first unambiguous signatures of life on a remote world: NASA's {\\it Terrestrial Planet Finder} and ESA's {\\it Darwin} \\citep{leger+herbst2007}.  The design of such observatories, and the urgency with which they will be built and deployed, will depend on the habitability potential of terrestrial planets that will be found in the next 5-10 years.  Over the last 50 years, various authors have addressed how to predict the way in which terrestrial planet habitability depends on star-planet distance (see \\citealt{kasting+catling2003} for a recent review).  Several of the important initial calculations predated the first discoveries of extrasolar planets, including \\citet{dole1964}, \\citet{hart1979}, and the seminal work of \\citet{kasting_et_al1993}. \\citet{selsis_et_al2007}, \\citet{vonbloh_et_al2008}, and \\citet{barnes_et_al2008} have reconsidered habitability in light of recent exoplanetary detections.  \\citet{williams_et_al1996}, \\citet[hereafter WK97]{williams+kasting1997} and \\citet{williams+pollard2003} have tackled precisely the questions relating to obliquity posed above, and have concluded that variations in obliquity do not necessarily render a planet non-habitable (see also \\citealt{hunt1982}, \\citealt{williams1988b,williams1988c}, \\citealt{oglesby+ogg1998}, \\citealt{chandler+sohl2000} and \\citealt{jenkins2000,jenkins2001,jenkins2003} in the context of Earth's paleoclimate studies).  Here we seek to generalize these analyses to model planets that are less close analogs to Earth than have been considered previously.  In \\citet[hereafter SMS08]{spiegel_et_al2008}, we examined how regionally and temporally habitable climates are affected by variations in the efficiency of latitudinal heat transport on a planet, and by variations in the ocean fraction. Importantly, we found that otherwise habitable Earth-like terrestrial planets can be subject to dynamical climate transitions into globally-frozen snowball states. Since it is not trivial to escape a snowball state (e.g., \\citealt{pierrehumbert2005}) and such globally-frozen climates may have profound influences on the development or existence of life (e.g., \\citealt{hoffman+schrag2002}), identifying the likelihood of such transitions on terrestrial exoplanets should be central to any habitability assessment.  Following in the footsteps of our first analysis, here we focus on obliquity and consider the influence on habitability of several planetary attributes a priori unknown for exoplanets, such as the efficiency of latitudinal heat transport and the land-ocean distribution.  The remainder of this paper is structured as follows: In \\S~\\ref{obl_sec:model} we describe the energy balance climate model we use. In \\S~\\ref{obl_sec:valid} discuss several validation tests in which our model performs well enough to give us some confidence in its behavior for conditions that differ from those found on Earth. In \\S~\\ref{obl_sec:results} we examine the influence on regional and seasonal habitability of various excursions from Earth-like conditions.  Finally, we conclude in \\S~\\ref{obl_sec:disc+conc}. ",
        "conclusions": "\\label{obl_sec:disc+conc}  We have presented a series of energy balance models to address the variety of climatic conditions that might exist on oblique terrestrial planets with circular orbits.  We considered dynamic climate forcings and responses determined by several planetary attributes a priori unknown for extrasolar planets, including obliquity, rotation rate, distribution of land/ocean coverage, and the detailed nature of the radiative cooling and heating functions.  We find that planets with small ocean fractions or polar continents can experience very severe seasonal climatic variations, but that these planets also might maintain seasonally and regionally habitable conditions over a larger range of orbital radii than more Earth-like planets.  Climates on high obliquity planets with nonuniform distributions of land and ocean can be far from global radiative balance, as compared to the Earth. Our results provide indications that the modeled climates are somewhat less prone to dynamical snowball transitions at high obliquity.  Fast rotating Earth-like planets may fall victim to global glaciation events at closer orbital radii than slower rotating planets. This is also the case for planets with massive CO$_2$ atmospheres, which are expected to be found in the outer orbital range of habitable zones. Snowball transitions could be particularly significant for such planets since partial collapse of their CO$_2$-rich atmospheres may occur and possibly interfere with the thermostatic effect of their carbonate-silicate weathering cycle, thus affecting their long-term habitability."
    },
    "0807/0807.3769_arXiv.txt": {
        "abstract": "We investigate the effects of weakly-interacting massive particle (WIMP) dark matter annihilation on the formation of Population III.1 stars, which are theorized to form from the collapse of gas cores at the centers of dark matter minihalos. We consider the relative importance of cooling due to baryonic radiative processes and heating due to WIMP annihilation. We analyze the dark matter and gas profiles of several halos formed in cosmological-scale numerical simulations. The heating rate depends sensitively on the dark matter density profile, which we approximate with a power law $\\rho_{\\chi}\\propto r^{-\\alpha_\\chi}$, in the numerically unresolved inner regions of the halo. If we assume a self-similar structure so that $\\alpha_{\\chi}\\simeq 1.5$ as measured on the resolved scales $\\sim 1$~pc, then for a fiducial WIMP mass of 100~GeV, the heating rate is typically much smaller ($<10^{-3}$) than the cooling rate for densities up to $n_{\\rm H}=10^{17}\\:{\\rm cm^{-3}}$. In one case, where $\\alpha_{\\chi}=1.65$, the heating rate becomes similar to the cooling rate by a density of $n_{\\rm H}=10^{15}\\:{\\rm cm^{-3}}$. The dark matter density profile is expected to steepen in the central baryon-dominated region $\\lesssim 1$~pc due to adiabatic contraction, and we observe this effect (though with relatively low resolution) in our numerical models. From these we estimate $\\alpha_{\\chi}\\simeq 2.0$. The heating now dominates cooling above $n_{\\rm H}\\simeq 10^{14}\\:{\\rm cm^{-3}}$, in agreement with the previous study of Spolyar, Freese \\& Gondolo. We expect this leads to the formation of an equilibrium structure with a baryonic and dark matter density distribution exhibiting a flattened central core. Examining such equilibria, we find total luminosities due to WIMP annihilation are relatively constant and $\\sim 10^3\\:L_\\odot$, set by the radiative luminosity of the baryonic core. We discuss the implications for Pop III.1 star formation, particularly the subsequent growth and evolution of the protostar. Even if the initial protostar fails to accumulate any additional dark matter, its contraction to the main sequence could be significantly delayed by WIMP annihilation heating, potentially raising the mass scale of Pop III.1 stars to masses $\\gg100\\:M_\\odot$. ",
        "introduction": "Population III stars are defined to be those whose formation and evolution are independent of metallicity (McKee \\& Tan 2008; O'Shea et al. 2008), since their metallicity is extremely low: close to or equal to that arising from big bang nucleosynthesis. Population III.1 stars are defined as having their formation be independent of other stars or other astrophysical objects, so that their initial conditions are determined solely by cosmology. These stars will be the first objects to form in a given region of the universe and they are likely to play an important role in cosmic reionization and in laying the foundations for the formation and build-up of galaxies. It is possible that they are the direct or indirect progenitors of supermassive black holes.  Within the commonly accepted $\\Lambda$CDM framework, Pop III.1 stars form within dark matter halos. Indeed, in those halos that do form stars, only one star appears to form in the center of each halo (Abel et al. 2002; Bromm, Coppi, \\& Larson 2002).  One of the most theoretically well-motivated cold dark matter candidates is a Weakly-Interacting Massive Particle (WIMP). Supersymmetric theories with R-parity naturally provide a stable dark matter candidate which could compose all or part of the dark matter in the Universe.  It has been pointed out that, if dark matter consists of a WIMP such as the supersymmetric neutralino, the energy released by the annihilation of these particles could influence early structure formation, star formation and protostellar evolution (Ripamonti, Mapelli \\& Ferrara 2007; Ascasibar 2007; Spolyar, Freese, \\& Gondolo 2008; Iocco 2008; Freese, Spolyar \\& Aguirre 2008; Freese et al. 2008b,c). The effects on stellar evolution at fixed mass have also been investigated (Taoso et al. 2008; Yoon, Iocco \\& Akiyama 2008). Spolyar et al. (2008) show that, for the adiabatically-contracted Navarro, Frenk \\& White (1996) (NFW) dark matter density profiles they considered, dark matter heating can overwhelm gas cooling in the innermost region of a star-forming minihalo, and they propose that this can then lead to a dark matter powered star.  In this paper we revisit the scenario investigated by Spolyar et al. (2008). In \\S\\ref{S:analytic} we derive an analytic expression for the dark matter heating rate, including a simplified treatment of radiative transport. In \\S\\ref{S:results} we present our results of the assessment of the importance of WIMP annihilation heating for several halos formed in numerical simulations of cosmic structure formation (O'Shea \\& Norman 2007). We describe the dark matter density structure in \\S\\ref{S:dm} and the properties of the baryons in \\S\\ref{S:baryon}. We compare the WIMP annihilation heating rates and the baryonic cooling rates and discuss the equilibrium structure of dark matter powered protostars in \\S\\ref{S:equilibrium}. We discuss the implications for subsequent protostellar evolution in \\S\\ref{S:protostar}. We conclude in \\S\\ref{S:conclusions}. ",
        "conclusions": "\\label{S:conclusions}  We have investigated the effects of WIMP dark matter annihilation on the formation of Population III.1 stars by analyzing the results of cosmological simulations that follow the gravitational collapse of baryons and dark matter. While these simulations (O'Shea \\& Norman 2007; Yoshida et al. 2006) have followed the baryons to very high densities at scales $\\lesssim 1$~AU, the dark matter is only well-resolved down to scales $\\sim 1$~pc. Thus we have considered various power law extrapolations of the dark matter density profile towards the center.  If one assumes the dark matter profile is self-similar, extending inwards from the dark matter dominated regime with $\\rho_{\\chi}\\propto r^{-\\alpha_{\\chi}}$ and $\\alpha_{\\chi}\\simeq 1.5$, then, for a fiducial WIMP mass of 100~GeV, the dark matter annihilation heating is typically negligible ($\\sim10^{-4}$ of the cooling rate). This conclusion would be unchanged for a reduction in the WIMP mass by a factor of 10 or more. One of the simulated minihalos (C) exhibits a slightly steeper density profile in the well-resolved region ($\\alpha_\\chi=1.65$, and in this case WIMP annihilation heating does become important at $n_{\\rm H}\\gtrsim 10^{15}\\:{\\rm cm^{-3}}$.  However, there are theoretical reasons to expect that the dark matter density profile will steepen because of adiabatic contraction in the baryon-dominated core. Indeed, this process appears to be occurring in the simulations of O'Shea \\& Norman (2007), although it is not well-resolved. A value of $\\alpha_{\\chi}\\simeq2.0$ appears to be a better description of the dark matter density profile in this region. For such a profile, and again for a 100~GeV particle, the dark matter annihilation heating now exceeds baryonic cooling for densities $n_{\\rm H}>10^{14}\\:{\\rm cm^{-3}}$, in agreement with the previous study of Spolyar et al. (2008).  We considered the properties of equilibrium halos in which the density distributions of the baryons and dark matter exhibit a constant density central core. The luminosity that is generated is $\\sim 10^3\\:L_\\odot$ and is relatively invariant, being set by baryonic cooling processes. The sizes of the central cores range from $\\sim 1$ to 40~AU.  The detailed effects of this extra heating on the protostellar structure remain to be determined. We expect that subsequent baryonic growth of the protostar will occur more rapidly than its accumulation of dark matter, because the baryons are undergoing rapid free fall collapse followed by disk accretion. However, even if the initial protostar does not gain any additional dark matter, its initial dark matter content in the subsolar mass core could be sufficient to prevent contraction to the zero age main sequence for masses of 100~$M_\\odot$ or greater. Such circumstances could have dramatic implications for the masses of Pop III.1 stars, conceivably raising the mass scale to a regime important for the formation of supermassive black holes. These conclusions depend sensitively on the initial protostellar core, which now needs to be studied with self-consistent cosmological simulations that include the influence of WIMP annihilation heating on the baryons."
    },
    "0807/0807.3133_arXiv.txt": {
        "abstract": "In this paper, we study the photon-photon pair production optical depth for 10 GeV--1 TeV gamma rays from 3C 279 due to the diffuse radiation of broad-line region (BLR). Assuming a power-law spectrum of $E_{\\gamma}^{-a_2}$ for the photon intensity of very high energy (VHE) gamma rays, $a_1 \\gtrsim 405$ and $a_2\\gtrsim 6.4$ are inferred by the integrated photon fluxes measured by MAGIC and HESS. Based on this power-law spectrum, the pre-absorbed spectra are inferred by correcting the photon-photon absorption on the diffuse photons of the BLR (internal absorption) and the extragalactic background light (external absorption). Position of gamma-ray emitting region $R_{\\rm{\\gamma}}$ determines the relative contributions of this two diffuse radiation to the total absorption for 10 GeV--1 TeV gamma rays. The internal absorption could make spectral shape of gamma rays more complex than only corrected for the external absorption, and could lead to the formation of arbitrary softening and hardening gamma-ray spectra. It should be necessary for the internal absorption to be considered in studying 10 GeV--1 TeV gamma rays from powerful blazars. The energies of annihilated gamma-ray photons due to the internal absorption are likely to be mainly reradiated around GeV. Our results indicate that $R_{\\rm{\\gamma}}$ may be between the inner and outer radii of the BLR for 3C 279. This implies for powerful blazars that $R_{\\rm{\\gamma}}$ might be neither inside the BLR cavity nor outside the BLR, but be within the BLR shell. Observations of $\\it GLAST$, MAGIC, HESS, and VERITAS in the near future could give more constraints on the position of the gamma-ray emitting region relative to the BLR. ",
        "introduction": "The classical flat spectrum radio quasar (FSRQ) 3C 279 is one of the brightest extragalactic objects in the gamma-ray sky. It was detected by the EGRET, and its spectrum does not show any signature of gamma-ray absorption by pair production up to $\\sim$ 10 GeV (Fichtel et al. 1994; von Montigny et al. 1995). With the detection of high energy gamma rays in 66 blazars, containing 51 FSRQs and 15 BL Lac objects, in the GeV energy range by the EGRET experiment aboard the Compton Gamma Ray Observatory (Catanese et al. 1997; Fichtel et al. 1994; Lin et al. 1997; Mukherjee et al. 1997; Thompson et al. 1995, 1996; Villata et al. 1997; Hartman et al. 1999; Nolan et al. 2003), an exceptional opportunity is presented for the understanding of the central engine operating in blazars. Some of blazars have been also firmly detected by atmospheric Cerenkov telescopes at energies above 1 TeV, such as Mrk 421 (Punch et al. 1992), and Mrk 501 (Quinn et al. 1996). The High Energy Stereoscopic System (HESS) is an imaging atmospheric Cerenkov detector with the energy threshold above 100 GeV (Funk et al. 2004; Hinton 2004; Hofmann 2003). The Very Energetic Radiation Imaging Telescope Array System (VERITAS) provides unprecedented sensitivity to photon energies between 50 GeV and 50 TeV (see e.g., Holder et al. 2006). The Major Atmospheric Gamma Imaging Cerenkov telescope (MAGIC) is currently the largest single-dish Imaging Air Cerenkov Telescope in operation with the lowest energy threshold, $\\sim$ 30 GeV, among the new Cerenkov telescopes (see e.g., Baixeras et al. 2004). At present, 23 active galactic nuclei (AGNs) have been detected in very high energy (VHE) gamma rays, containing 21 BL Lac objects, one radio source M87, and the first FSRQ 3C 279 with the highest redshift $z=0.536$ in these VHE AGNs\\footnote{http://www.mppmu.mpg.de/$\\sim$rwagner/sources/}. The MAGIC telescope detected VHE gamma rays from 3C 279 (Teshima et al. 2007). HESS observations measured an upper limit of integrated photon flux (Aharonian et al. 2008). The Large Area Telescope instrument on the Gamma-Ray Large Area Space Telescope ($\\it GLAST$), the new-generation high energy gamma-ray telescope with sufficient angular resolution to allow identification of a large fraction of their optical counterparts, will observe gamma rays with energies from 20 MeV to greater than 300 GeV, and have the unique capability to detect thousands of gamma-ray blazars to redshifts of at least $z=4$ (see e.g., Chen et al. 2004). $\\it GLAST$ was launched on June 11 2008. $\\it GLAST$, combined with new-generation TeV instruments such as MAGIC, HESS, and VERITAS, will tremendously improve blazar spectra studies, filling in the band from 20 MeV to 10 TeV with high significance data for hundreds of AGNs (see Gehrels \\& Michelson 1999). Future measurements of the gamma-ray spectrum shape and its variability of blazars would tremendously improve our understanding to blazars.  These gamma rays from blazars are generally believed to be attributed by emission from a relativistic jet oriented at a small angle to the line of sight (Blandford \\& Rees 1978). These gamma-ray components are contributed by inverse Compton emissions, including synchrotron self-Compton (SSC) scattering synchrotron seed photons, and external Compton (EC) scattering seed photons from sources outside the jet (see e.g., B\\\"ottcher 1999). The diffuse radiation fields of broad-line region (BLR) could have a strong impact on the expected EC spectra of powerful blazars, FSRQs (Liu \\& Bai 2006; Reimer 2007; Tavecchio \\& Ghisellini 2008). The external soft photon fields not only provide target photons for the EC processes to produce these gamma-ray components, but also absorb gamma rays from the EC processes, because gamma rays between 10 GeV and 1 TeV interact with infrared-ultraviolet photons to be attenuated by photon-photon pair production. Many efforts to study the absorption of gamma rays focus on the photon-photon annihilation by the diffuse extragalactic background radiation at the infrared (IR), optical, and ultraviolet (UV) bands (see e.g., Stecker et al. 1992, 2006, 2007; Stecker \\& de Jager 1998; Oh 2001; Renault et al. 2001; Chen et al. 2004; Dwek \\& Krennrich 2005; Schroedter 2005). This external absorption of gamma rays on the diffuse extragalactic background light (EBL) is also proposed and used to probe the EBL (Renault et al. 2001; Chen et al. 2004; Dwek \\& Krennrich 2005; Schroedter 2005). Indeed, the internal absorption of gamma rays inside FSRQs could result in serious problem for the possibility to use the external absorption of gamma rays to probe the IR--optical--UV extragalactic background, because the intrinsic spectra of gamma rays are masked by the internal absorption inside blazars (Donea \\& Protheroe 2003; Liu \\& Bai 2006; Reimer 2007). The intrinsic spectra of gamma rays are complicated by the complex spectra of the diffuse radiation fields of the BLRs in FSRQs (Tavecchio \\& Ghisellini 2008). The mean of intrinsic spectral index is around $2.3$ for 17 BL Lac objects detected in the VHE regime (Wagner 2008).  The positions of gamma-ray emitting regions are still an open and controversial issue in the researches on blazars. It is suggested that gamma rays are produced inside the BLRs and the gamma-ray emitting radius $R_{\\gamma}$ ranges roughly between 0.03 and 0.3 pc (Ghisellini \\& Madau 1996). It is argued by Georganopoulos et al. (2001) that the radiative plasma in relativistic jets of powerful blazars are within cavities formed by the BLRs. However, other researchers argued that the gamma-ray emitting regions are outside the BLRs (Lindfors et al. 2005; Sokolov \\& Marscher 2005). In our previous research (Liu \\& Bai 2006, hereafter Paper I), the position of gamma-ray emitting region is a key parameter to determine whether high energy gamma rays could escape the diffuse radiation fields of the BLRs for FSRQs. It is unknown whether these gamma rays could be detected by $\\it GLAST$, MAGIC, HESS, and VERITAS even if blazars intrinsically produce 10 GeV--1 TeV gamma rays, because the gamma-ray emitting radii are unknown. These gamma rays around 200 GeV is optically thick for 3C 279 if the emitting region is within the BLR cavity (see Paper I). However, the MAGIC observations on 2006 February 23 have shown a clear gamma-ray signal in the VHE regime (Teshima et al. 2007). This indicates that the emitting region of these VHE gamma rays should not be inside the BLR cavity, otherwise, the intrinsic flux of these VHE gamma rays is likely to be extremely high. In Paper I, we addressed an important topic in gamma-ray astrophysics, namely the absorption of high energy gamma rays inside of FSRQs by photons of the BLR. In this paper, we attempt to address the particular topic of absorption in the gamma-ray quasar 3C 279 using the available observational data, and its potential effect on the spectra of gamma rays. In order to constrain the position of gamma-ray emitting region in 3C 279, we study the internal absorption of gamma rays by the diffuse radiation from the BLR and its potential effect on the spectra of gamma rays from 10 GeV to 1 TeV.  The structure of this paper is as follows. $\\S$ 2 presents intensity of VHE gamma rays. $\\S$ 3 presents theoretical calculation, and consists of two subsections. $\\S$ 3.1 presents calculations of temperature profiles, and $\\S$ 3.2 photon-photon optical depth for 3C 279. $\\S$ 4 presents external absorption on IR--optical--UV extragalactic background. $\\S$ 5 is spectral shape of VHE gamma rays. $\\S$ 6 presents pair spectrum due to photon-photon annihilation and their radiation. $\\S$ 7 is for discussions and conclusions. Throughout this paper, we use a flat cosmology with a deceleration factor $q_0=0.5$ and a Hubble constant $H_0=75 \\/\\ \\rm{km \\/\\ s^{-1} \\/\\ Mpc^{-1}}$. ",
        "conclusions": "HESS observations in 2007 January measured an upper limit of integrated photon flux $F(E_{\\gamma}>300\\/\\ \\rm{GeV})<3.98\\times 10^{-12}\\/\\ \\rm{cm^{-2}s^{-1}}$ (Aharonian et al. 2008). In calculations, we assumed $F(E_{\\gamma}>300\\/\\ \\rm{GeV})=3.9\\times 10^{-12}\\/\\ \\rm{cm^{-2}s^{-1}}$, allowed by HESS observations. Combining this flux with integrated photon flux $F(E_{\\gamma}>200\\/\\ \\rm{GeV})=3.5\\times 10^{-11}\\/\\ \\rm{cm^{-2}s^{-1}}$ measured by MAGIC (Teshima et al. 2007), a power-law spectrum of $\\frac{dI}{dE_{\\gamma}}=a_1 E_{\\gamma}^{-a_2}$ is inferred with $a_1 =505$ and $a_2 =6.4$. If $F(E_{\\gamma}>300\\/\\ \\rm{GeV})=3.6\\times 10^{-12}\\/\\ \\rm{cm^{-2}s^{-1}}$ is adopted for HESS observations, $a_1 =1510$ and $a_2 =6.6$. If $F(E_{\\gamma}>300\\/\\ \\rm{GeV})=3.98\\times 10^{-12}\\/\\ \\rm{cm^{-2}s^{-1}}$ is adopted, $a_1 =405$ and $a_2 =6.36$. Thus $a_1 \\gtrsim 405$ and $a_2 \\gtrsim 6.4$ are likely to be limited by MAGIC and HESS observations for a VHE gamma-ray spectrum of $\\frac{dI}{dE_{\\gamma}}=a_1 E_{\\gamma}^{-a_2}$. Based on the gamma-ray spectrum of $\\frac{dI}{dE_{\\gamma}}=505 E_{\\gamma}^{-6.4}$ but corrected for the internal and external absorption (see Figs. 2$b$, 3$b$, and 4$b$), photon indices of these corrected VHE gamma rays are around 5 as $R_{\\gamma}=0.1 \\/\\ \\rm{pc}$, 2.0 as $R_{\\gamma}=0.25 \\/\\ \\rm{pc}$, and 1.7 as $R_{\\gamma}=0.4 \\/\\ \\rm{pc}$. These inferred photon indices as $R_{\\gamma}=$ 0.25 and 0.4 $\\rm{pc}$ are inside of those known intrinsic photon indices. If $a_2 =6.6$, photon indices of $1.9\\pm0.03$ are inferred for $R_{\\gamma}=0.4 \\/\\ \\rm{pc}$, $2.2\\pm0.01$ ($a_{\\ast}=0.5$) and $2.1\\pm 0.003$ ($a_{\\ast}=0.998$) for $R_{\\gamma}=0.25 \\/\\ \\rm{pc}$, and $6.0\\pm0.2$ ($a_{\\ast}=0.5$) and $5.1\\pm0.2$ ($a_{\\ast}=0.998$) for $R_{\\gamma}=0.1 \\/\\ \\rm{pc}$. As $R_{\\gamma}$ varies between 0.1 and 0.4 $\\rm{pc}$, photon indices of pre-absorbed VHE gamma-ray spectra are likely to be inside of intrinsic photon index range from 1.3 to 3.6 (Wagner 2008). If one smaller $F(E_{\\gamma}>300\\/\\ \\rm{GeV})=2.0\\times 10^{-12}\\/\\ \\rm{cm^{-2}s^{-1}}$ is adopted, one larger $a_2=8.0$ is obtained. As $a_2 =8.0$, photon indices of $3.3\\pm0.03$ are inferred for $R_{\\gamma}=0.4 \\/\\ \\rm{pc}$, $3.6\\pm0.01$ ($a_{\\ast}=0.5$) and $3.5\\pm 0.003$ ($a_{\\ast}=0.998$) for $R_{\\gamma}=0.25 \\/\\ \\rm{pc}$, and $7.3\\pm0.2$ ($a_{\\ast}=0.5$) and $6.5\\pm0.2$ ($a_{\\ast}=0.998$) for $R_{\\gamma}=0.1 \\/\\ \\rm{pc}$. Thus, larger photon indices of these pre-absorbed spectra are obtained for a fixed $R_{\\gamma}$ as larger $a_2$ is adopted. As $R_{\\gamma}=r_{\\rm{BLR,in}}$, photon indices of these pre-absorbed gamma-ray spectra are always larger than the typical value $\\Gamma_{\\rm{in}}=2.3$ and the maximum $\\Gamma_{\\rm{in}}=3.6$ of those known intrinsic photon indices. Thus, it is unlikely for 3C 279 that $R_{\\gamma}$ be inside the BLR cavity, i.e. it is likely $R_{\\gamma}> r_{\\rm{BLR,in}}$. If $F(E_{\\gamma}>300\\/\\ \\rm{GeV})$ adopted is closer to the upper limit of HESS observations, $a_2$ is not too large. Too large $a_2$, such as $a_2\\simeq 10$ corresponding to $F(E_{\\gamma}>300\\/\\ \\rm{GeV})\\simeq 10^{-12}\\/\\ \\rm{cm^{-2}s^{-1}}$, seems to be impossible for VHE gamma-ray spectra because these spectra with $a_2 \\gtrsim 10$ are much steeper (softer) than the steepest (softest) spectrum measured for PG 1553+113 (Albert et al. 2007). When gamma-ray emitting region is already beyond the BLR, the EC mechanisms, where external photons originate from the BLR, are made inefficient to produce the observed gamma rays. For 3C 279, soft photon energy density around $r_{\\rm{BLR,out}}$ is lower by a factor of 6 to 7 than that around $r_{\\rm{BLR,in}}$ for the BLR diffuse radiation produced by the spherical shell of clouds (see Fig. 1 in Paper I). Thus it is unlikely $R_{\\gamma}>r_{\\rm{BLR,out}}$, i.e. it is likely $R_{\\gamma}\\lesssim r_{\\rm{BLR,out}}$.  The external absorption softens the observed spectra relative to the emission ones in the interval from 10 GeV to 1 TeV. Whether the internal absorption softens or hardens the observed spectra relies on gamma-ray emitting radius $R_{\\gamma}$ and energy $E_{\\gamma}$. Photon index variations of gamma-ray spectra relative to $a_2=6.4$ are calculated for the internal absorption, and the internal and external absorption (see Figure 7). As $R_{\\gamma}=r_{\\rm{BLR,in}}$, the local photon indices of gamma-ray spectra decrease with $E_{\\gamma}$ from 10 to 110 GeV, and increase with $E_{\\gamma}$ from 110 GeV to 1 TeV (see Figure 7$a$). The local photon indices of gamma-ray spectra corrected for the internal absorption recover around 270 GeV ($\\Delta \\Gamma(\\rm{in})=0$). The internal absorption make gamma-ray spectra softer and softer from 10 to 110 GeV and from 270 down to 110 GeV, and make gamma-ray spectra harder and harder from 270 GeV to 1 TeV. After corrected for the internal and external absorption, the local photon indices of gamma rays recover around 500 GeV ($\\Delta \\Gamma(\\rm{in+ext})=0$). The external absorption can not change the trend of the local photon indices of gamma-ray spectra corrected for the internal absorption. As $R_{\\gamma}=(r_{\\rm{BLR,in}}+r_{\\rm{BLR,out}})/2$, the local photon indices of gamma-ray spectra corrected for the internal absorption behave as in the case of $R_{\\gamma}=r_{\\rm{BLR,in}}$ (see Figure 7$b$). These local photon indices recover around 400 GeV ($\\Delta \\Gamma(\\rm{in})=0$). The internal absorption make gamma-ray spectra softer and softer from 10 to 110 GeV and from 400 down to 110 GeV, and make gamma-ray spectra harder and harder from 400 GeV to 1 TeV. After corrected for the internal and external absorption, the local photon indices of pre-absorbed gamma-ray spectra basically decrease with $E_{\\gamma}$. The external absorption can not change the trend of the local photon indices of gamma-ray spectra corrected for the internal absorption below 100 GeV, but change that above 100 GeV. Gamma-ray spectra become softer and softer relative to pre-absorbed ones from 10 GeV to 1 TeV ($\\Delta \\Gamma(\\rm{in+ext})<0$). As $R_{\\gamma}=r_{\\rm{BLR,out}}$, the local photon indices have similar trend to that in the case of $R_{\\gamma}=(r_{\\rm{BLR,in}}+r_{\\rm{BLR,out}})/2$ (see Figure 7$c$). The internal absorption make gamma-ray spectra softer and softer from 10 to $\\sim$ 420 GeV and from 1 TeV down to $\\sim$ 420 GeV. After corrected for the internal and external absorption, the local photon indices monotonically decrease with $E_{\\gamma}$, and gamma-ray spectra become softer and softer relative to pre-absorbed ones from 10 GeV to 1 TeV ($\\Delta \\Gamma(\\rm{in+ext})<0$). The external absorption change the trend of the local photon indices of gamma-ray spectra corrected for the internal absorption. The calculated results presented in Figure 7 basically are independent on the particular value of $a_2$. For example, the trends of the local photon indices for $a_2=5.4$ are basically identical to those in the case of $a_2=6.4$. In summary, the internal absorption could make spectral shape more complex than only considering the external absorption, and could lead to the formation of arbitrary softening and hardening gamma-ray spectra (see Figs. 2--4 and 7). Thus, it should be necessary for the internal absorption to be considered in studying gamma rays of 10 GeV--1 TeV from FSRQs.  Assuming $\\Gamma_{\\rm{jet}}=15$ for 3C 279, most of gamma rays are contained within a radiation cone with a half open angle of $\\Delta \\varphi\\sim 1/ \\Gamma_{\\rm{jet}} \\sim 3.8^{\\circ}$, because of the relativistic beaming effect. If the central IR--optical--UV photons coming directly from accretion disks travel through the radiation cone, the IR--optical--UV photons can have photon-photon pair creation processes with gamma rays within the radiation cone. The gamma-ray emitting region $R_{\\rm{\\gamma}}\\sim 0.1 \\/\\ \\rm{pc}$ is assumed, and the radii of the UV radiation regions are $R_{D}\\lesssim 30 r_{\\rm{g}}\\sim 0.0004 \\/\\ \\rm{pc}$ (see Fig. 1), the angle between the jet direction and the travelling direction of UV photons at $R_{\\rm{\\gamma}}$ is $\\sim \\arcsin R_{\\rm{D}}/R_{\\rm{\\gamma}}\\lesssim 0.2^{\\circ}$. Then the photons within the radiation cone have the collision angles of $\\theta \\lesssim 4.0^{\\circ}$. For the UV photons at the frequency $\\nu \\simeq10^{16.5} \\/\\ \\rm{Hz}$ with energies of $\\varepsilon_2\\simeq 2.56\\times 10^{-4}$ and the gamma rays of $\\varepsilon_1\\simeq 2.0\\times 10^{6}$ corresponding to energies around 1 TeV, the left of threshold condition (eq. [3] in Paper I) has a upper limit of $\\lesssim 0.6$, which is less than unity, and thus the two kinds of photons cannot be absorbed by the photon-photon pair creation processes. Therefore, the central UV radiation have negligible contributions to the absorption for gamma rays relative to the diffuse radiation from the BLR. For optical photons, the radii of the optical radiation regions are $R_{D}< 200 r_{\\rm{g}}\\sim 0.003 \\/\\ \\rm{pc}$, the angle between the jet direction and the travelling direction of optical photons at $R_{\\rm{\\gamma}}$ is $\\sim \\arcsin R_{\\rm{D}}/R_{\\rm{\\gamma}}\\lesssim 1.7^{\\circ}$. Then the photons within the radiation cone have the collision angles of $\\theta \\lesssim 5.5^{\\circ}$. For the optical photons at the frequency $\\nu \\simeq10^{15} \\/\\ \\rm{Hz}$ with energies of $\\varepsilon_2\\simeq 8.1\\times 10^{-6}$ and the gamma rays of $\\varepsilon_1\\simeq 2.0\\times 10^{6}$, the left of threshold condition has a upper limit of $\\lesssim 0.04$, which is less than unity, and thus the two kinds of photons cannot be absorbed by the photon-photon pair creation processes. Therefore, the central optical radiation have negligible contributions to the absorption for gamma rays relative to the diffuse radiation from the BLR. For IR photons, the central radiation also have negligible contributions to the absorption for gamma rays. Thus, the central IR--optical--UV radiation have negligible contributions to the absorption for gamma rays relative to the diffuse radiation from the BLR.  Absorption for gamma rays by photon-photon annihilation and where the energies carried by the annihilated gamma rays reradiate are important to gamma-ray research. The intense creation of pairs would produce a strong radiation at low energy X-rays or at GeV energy (Protheroe \\& Stanev 1993; Saug\\'e \\& Henri 2004; Zdziarski \\& Coppi 1991). The electron-positron pair cascade could cause the soft X-ray excesses (Zdziarski \\& Coppi 1991). The produced electron-positron pairs could make a difference around 1 GeV, where the pileup of photons below the absorption threshold results in a significant flattening in the observed spectrum relative to the emission spectrum (Protheroe \\& Stanev 1993). In $\\S$ 6, we studied the pair spectrum due to photon-photon absorption, and the synchrotron and EC spectra emitted by the equilibrium pair distribution of steady state. The synchrotron radiation peaks around keV X-rays, and the EC radiation peaks around GeV gamma rays (see Fig. 6$b$). Thus, pairs due to annihilation absorption for gamma rays by the diffuse radiation fields of the BLR are likely to make a difference around 1 GeV for 3C 279.  In this paper, in order to limit the gamma-ray emitting radius $R_{\\gamma}$, we used a BLR model to study the photon-photon absorption by the diffuse radiation of the BLR in 3C 279 for gamma rays of 10 GeV to 1 TeV in the observed spectrum. We calculated the internal absorption of gamma rays from 10 GeV to 1 TeV for $R_{\\gamma}=r_{\\rm{BLR,in}}$, $r_{\\rm{BLR,out}}$, and $(r_{\\rm{BLR,in}}+r_{\\rm{BLR,out}})/2$ (see Figs. 2$a$, 3$a$, and 4$a$). For a fixed $R_{\\gamma}$, dependence of photon-photon absorption optical depth $\\tau_{\\gamma\\gamma}$ on energies of gamma rays $E_{\\gamma}$ relies on $R_{\\gamma}$. Dependence of $\\tau_{\\gamma\\gamma}$ on $R_{\\gamma}$ was also studied for a fixed $E_{\\gamma}$ (see Fig. 5). For a fixed $E_{\\gamma}$, $\\tau_{\\gamma\\gamma}$ decreases with increasing $R_{\\gamma}$. The external absorption on the IR--optical--UV EBL was also estimated for gamma rays of 10 GeV--1 TeV, and it monotonically increases as $E_{\\gamma}$ increases. Comparing the internal absorption with the external one shows that $R_{\\gamma}$ determines the relative contributions of the internal and external absorption to the total photon-photon annihilation absorption of 10 GeV--1 TeV gamma rays (see Figs. 2$a$, 3$a$, and 4$a$). Based on MAGIC and HESS observations, a power-law spectrum as in equation (1) was adopted for the photon intensity of VHE gamma rays. The pre-absorbed gamma-ray spectra are inferred by this power-law corrected for the internal and external absorption. The internal absorption could make spectral shape of gamma rays more complex than that only corrected for the external absorption, and could lead to the formation of arbitrary softening and hardening gamma-ray spectra (see Figs. 2$a$, 3$a$, 4$a$, and 7). Thus, it should be necessary for the internal absorption to be considered in studying 10 GeV--1 TeV gamma rays from FSRQs. $R_{\\gamma}$ significantly influences the variations of spectral shape due to the internal absorption. Calculations imply that the energies of annihilated gamma rays due to the internal absorption are mainly reradiated around GeV regime (see Fig. 6$b$). Considering the possible variations of photon index $a_2$, the photon indices of the pre-absorbed VHE gamma-ray spectra were compared with those known intrinsic photon indices. As $R_{\\gamma}=r_{\\rm{BLR,in}}$ and $a_2\\gtrsim 6.4$, the photon indices of the pre-absorbed gamma-ray spectra are always larger than the typical value $\\Gamma_{\\rm{in}}=2.3$ and the maximum $\\Gamma_{\\rm{in}}=3.6$ of those known intrinsic photon indices, and $\\tau_{\\gamma\\gamma}$ is larger than unity. Thus, it is likely $R_{\\gamma}> r_{\\rm{BLR,in}}$ for 3C 279. As $R_{\\gamma}=r_{\\rm{BLR,out}}$ and $6.4 \\lesssim  a_2 \\lesssim 8.3$, photon indices of pre-absorbed gamma-ray spectra are not larger than $\\Gamma_{\\rm{in}}=3.6$. For a fixed $a_2$, photon indices of pre-absorbed gamma-ray spectra decrease as $R_{\\gamma}$ increases. Too large $a_2$ seems to be impossible for VHE gamma-ray spectra. In addition, the EC processes may be inefficient to produce the observed gamma rays as gamma-ray emitting region is already beyond the BLR. Thus, it is likely $R_{\\gamma}\\lesssim r_{\\rm{BLR,out}}$ for 3C 279. Our results suggest that $R_{\\gamma}$ for powerful blazars might be neither inside the BLR cavity nor outside the BLR, but be within the BLR shell. This is neither consistent with suggestions of Ghisellini \\& Madau (1996) and Georganopoulos et al. (2001) nor consistent with suggestions of Lindfors et al. (2005) and Sokolov \\& Marscher (2005).  Our results are model dependent, especially dependent on the assumed power-law spectrum for the VHE gamma rays. If Teshima et al. published the spectral indices of gamma-ray spectra measured by MAGIC (Teshima et al. 2007), the power-law spectrum assumed in this paper could be tested. Tavecchio \\& Ghisellini (2008) used the photoionization code CLOUDY, described by Ferland et al. (1998), to calculate the detailed spectra from the BLRs for powerful blazars, and then used these spectra to calculate the EC spectra. Approximate spectra of the BLRs are used in this paper and Paper I, and another approximate spectra are used by Reimer (2007). Difference between the detailed and approximate spectra should exist. It should be useful in the future researches to study the effects of this difference on the results of previous researches using approximate spectra (e.g., Liu \\& Bai 2006; Reimer 2007). Observations of $\\it GLAST$, MAGIC, HESS, and VERITAS in the near future could give more observational constraints on the gamma-ray emitting regions and the BLRs for the powerful blazars. Publications of intrinsic photon indices predicted by theoretical researches and photon indices measured by observations in the VHE regime could give stronger constraints on $R_{\\gamma}$."
    },
    "0807/0807.4844_arXiv.txt": {
        "abstract": "The Tautenburg Exoplanet Search Telescope (TEST) is a robotic telescope system. The telescope uses a folded Schmidt Camera with a 300mm main mirror. The focal length is 940mm and it gives a $2.2^{\\circ}$ x $2.2^{\\circ}$ field of view. Dome, mount, and CCD cameras are controlled by a software bundle made by Software Bisque. The automation of the telescope includes selection  of the night observing program from a given framework, taking darks and skyflats, field identification, guiding, data taking, and archiving. For the search for transiting exoplanets and variable stars an automated psf photometry based on IRAF and a lightcurve analysis based on ESO-Midas are conducted. The images and the results are managed using a PostgreSQL database. ",
        "introduction": "In 1999 the first transiting extrasolar planet was discovered (\\cite[Charbonneau et al. 2000]{Charbonneau_etal00}). As the knowledge of transit parameters allows an unique insight into the nature of the planet, many efforts have been made to increase the number of known transiting planets. Until now about 50 transiting extrasolar planets are discovered. Inspired by the Berlin Exoplanet Search Telescope (BEST) (\\cite[Rauer et al. 2004]{Rauer_etal04}), which was operated by the Deutsches Zentrum f{\\\"u}r Luft- und Raumfahrt at the Th{\\\"u}ringer Landessternwarte Tautenburg in the years 2001~-~2003, a new small aperture telescope has now been installed at this site to search for transiting exoplanets. ",
        "conclusions": ""
    },
    "0807/0807.1185_arXiv.txt": {
        "abstract": "The Extreme ultraviolet Imaging Spectrometer (EIS) on board \\hinode\\ is the first solar telescope to obtain wide slit spectral images that can be used for detecting Doppler flows in transition region and coronal lines on the Sun and to relate them to their surrounding small scale dynamics. We select EIS lines covering the temperature range $6\\times10^4$~K to $2\\times10^6$~K that give spectrally pure images of the Sun with the 40\\arcsec\\ slit. In these images Doppler shifts are seen as horizontal brightenings. Inside the image it is difficult to distinguish shifts from horizontal structures but emission beyond the image edge can be unambiguously identified as a line shift in several lines separated from others on their blue or red side by more than the width of the spectrometer slit (40 pixels). In the blue wing of \\HeII, we find a large number of events with properties (size and lifetime) similar to the well-studied explosive events seen in the ultraviolet spectral range. Comparison with X-Ray Telescope (XRT) images shows many Doppler shift events at the footpoints of small X-ray loops. The most spectacular event observed showed a strong blue shift in transition region and lower corona lines from a small X-ray spot that lasted less than 7~min. The emission appears to be near a cool coronal loop connecting an X-ray bright point to an adjacent region of quiet Sun. The width of the emission implies a line-of-sight velocity of 220~\\kms. In addition, we show an example of an \\FeXV\\ shift with a velocity about 120~\\kms, coming from what looks like a narrow loop leg connecting a small X-ray brightening to a larger region of X-ray emission. ",
        "introduction": "The Extreme ultraviolet Imaging Spectrometer \\citep[EIS;][]{Culhane07} on \\hinode\\ obtains images and spectra of many transition region and coronal lines in the wavelength ranges $170-211$~\\AA\\ and $246-292$~\\AA. The EIS wide slits are an interesting compromise between narrow slit spectra where the time to raster across a particular structure is often longer than the lifetime of the structure and filter images which give no direct measurement of flow velocities. The 40\\arcsec\\ wide slit provides overlapping spectra from the observed  40\\arcsec\\ wide region of the Sun. Images in different lines overlap if the lines are separated by less than the 40 pixel width of the slit image times the wavelength plate scale, 0.0223 \\AA/pixel. Figure~\\ref{slots}, shows a 40\\arcsec\\ wide slit spectrum of a small active region for the wavelength region $268-292$~\\AA. Isolated lines produce well-defined 40\\arcsec\\ wide images.  Within the EIS spectrum there are several relatively strong isolated lines that can be used to study the dynamics of the outer atmosphere. \\citet{Hansteen07} report on rapid temporal variations in quiet Sun \\HeII\\ 256~\\AA\\ and \\FeXII\\ 195~\\AA\\ wide slit emission features observed with a cadence of 30~s. They also mention the detection of many \\HeII\\ blue shifts in quiet Sun narrow slit rasters observed after the wide slit sequence.  In this paper we show how the detection of \\HeII\\ blue shifts and large-scale temporal variations, can be made simultaneously with wide slit observations. The key is that the emission beyond the edge of the 40\\arcsec\\ main line image is either line broadening/Doppler shifts from the main line or emission  in a neighbouring line. Figure~\\ref{slots} demonstrates both effects. \\FeXV\\ and \\FeXIV\\ (labelled) are both well separated lines and the images are basically straight along both edges. On the edge that cuts the active region, the emission from both lines bulges slightly because the bright active region lines are broader. A closer inspection of the emission shows that both lines have the same structure along the edge. On the other hand, the strong unmarked line on the left is mainly \\FeXIV\\ 270.5~\\AA, but there is clearly a blend producing extensions on the left hand side. Positions of the projections coincide with bright \\MgVII\\ and \\SiVII\\ because the blend is the lower coronal line \\MgVI\\ 270.4~\\AA.  \\begin{figure} \\centering \\includegraphics[width=12 cm]{slots.eps} \\caption{A 40\\arcsec\\ wide slit spectrum of a small active region, showing well separated and blended lines. Several of the lines discussed in this paper are marked.} \\label{slots} \\end{figure}  When the Sun's structure is known from images in isolated lines, it is possible to work out from the spectrum whether neighbouring lines are expected to produce emission. If there is no spectral line that can cause the emission, it is very likely due to Doppler shifts/broadening in the image line. Thus the EIS 40\\arcsec\\ wide slit can be used to obtain both images of the dynamics over a wide range of temperatures, and Doppler shifts from structures on the edge of the images.  Inside the image, Doppler broadened lines appear as horizontal streaks which are difficult to distinguish from a jet-like structure. Comparison with filter images is very important in order to positively identify shifts inside the image. Simultaneous images obtained by other instruments on \\hinode\\ can, in principle, be used to help identify shifts but there will always be some uncertainty because the filters to not detect exactly the same plasma.  In this letter we show examples of Doppler shift events seen in the quiet Sun with the EIS 40\\arcsec\\ wide slit. Comparisons are made with X-ray images obtained by the X-Ray Telescope \\citep[XRT;][]{Golub07} in order to show the relationship to hot loop structures. Further work will report on detailed analyses with data from the Solar Optical Telescope \\citep[SOT;][]{Tsuneta08} on \\hinode\\ and the Solar Ultraviolet Measurements of Emitted Radiation \\citep[SUMER;][]{Wetal95} spectrograph on SoHO. ",
        "conclusions": "There is a wealth of information in the EIS 40\\arcsec\\ wide slit images. The primary use of wide slit images is to see rapid time variations of structures in specific spectral lines \\citep{Hansteen07}, typically with a cadence 20 times (the slit width) faster than rastering with the 2~\\arcsec\\ slit. After passing through the slit, the light is dispersed producing rows of overlapping spectral images. If a Doppler shift occurs in the center of the image, this is indistinguishable from a horizontal structure. It is therefore beneficial to have simultaneous filtergrams to show the underlying structures.  Here we show how the wide slit can be used to detect Doppler shifts at the edge of the image. We have highlighted three types of Doppler shift event seen in EIS wide slit spectra: explosive events in \\HeII\\ at the image edge; a fast jet related to an X-ray spot brightening; and \\FeXV\\ flow connecting  an X-ray spot brightening to a larger X-ray region. The latter observations were done in sit-and-stare at two positions spaced 40\\arcsec\\ apart in order to pick up first the blue and then the red wings of the lines. At present our understanding of the line profile at the edge of the slit is not sufficiently well understood to determine accurate line shifts but flows producing emission more than two pixels away from the average line profile at the edge are clearly detectable.  This suggests new ways of investigating, for example, explosive event dynamics because one can see the loop motion and heating beyond the position of the explosive event. It could also be useful when investigating Doppler shift fluctuations from hot loops because narrow slit observations in sit-and-stare are not able to distinguish real intensity fluctuations at a point from movement of loops into and out of the slit field-of-view \\citep[cf][]{Wetal03b, Mariska07}.  The aim of this short paper is to show the possibilities for future studies, not to present event details and their scientific implications. Discussions on the events themselves, including detailed analyses from SOT on \\hinode\\, and SUMER on SOHO will be given in more extensive papers."
    },
    "0807/0807.3239_arXiv.txt": {
        "abstract": "We discuss a new criterion to estimate the mass in the outer, non-equilibrium region of galaxy clusters, where the galaxy dynamics is dominated by an overall infall motion towards the cluster centre. In the framework of the spherical infall model the local mean velocity of the infalling galaxies at every radius provides information about the integrated matter overdensity $\\delta$. Thus, a well-defined value of the  overdensity $\\delta_t$ is expected at the turnaround radius $r_t$, i.e. the radius where the Hubble flow balances the infall motion. Within this scenario, we analysed the kinematical properties of a large catalogue of simulated clusters, using both dark matter particles and member galaxies as tracer of the infall motion. We also compared the simulation with analytical calculation performed in the spherical infall approximation, to analyze the dependence of the results on cosmology in spatially flat universe. If we normalize cluster mass profiles by means of the turnaround mass $M_t$ (i.e. the mass within $r_t$), they are consistent with an exponential profile in the whole non-equilibrium region ($0.5\\la r/r_t\\la 2$). Turnaround radii are proportional to virialization radii ($r_t\\simeq 3.5 r_v$), while turnaround masses are proportional to virialization masses, i.e. $M_t\\simeq 1.7 M_v$, where $M_v$ is the mass within $r_v$. Actually, the mass evaluated within the turnaround radius is a more exhaustive evaluation of the total mass of the cluster. These results can be applied to the analysis of observed clusters. ",
        "introduction": "The gravitational collapse of galaxies towards the centre of clusters is usually described within the framework of the spherical infall model, as the motion of a set of concentric, spherically symmetrical mass shells \\citep{GG,Si,Sc}. Actually the spherical infall model is widely accepted in literature, since it describes fairly well the dynamics of the non-equilibrium region of galaxy clusters, defined as the region where the effects of virialization and the crossing of the above-mentioned shells are negligible and some overall infall motion of member galaxies is recognizable. Under the spherical symmetry assumption, the infall motion produces a pattern of caustic surfaces in the galaxy redshift-space distribution (which is obtained representing the line-of-sight velocities of galaxies \\emph{versus} their projected position on the sky plane). These caustics envelop all galaxies whose infall motion overwhelms the Hubble flow \\citep{K}. Caustics with a characteristical ``trumpet'' shape were actually observed in the redshift-space distribution of clusters \\citep{Oal}. Diaferio \\& Geller \\citep{DG} and Diaferio \\citep{D} showed that the caustic amplitude provides a direct measure of the escape velocity of galaxies, and therefore allows to estimate the mass profile of the cluster in the innermost part of the non-equilibrium region, up to the turnaround radius $r_t$ (i.e. the radial distance where the velocity of the infall motion is equal to the Hubble flow velocity.). The caustic technique was applied to the observation of many local clusters \\citep{GDK,Ral1,Ral2,Ral3}. These mass estimates are consistent with those based on virial theorem \\citep{Giral,BivGir} and weak lensing observations (\\citealt{DGR}, and references therein). In fact, up to now the sampled volumes were always restricted within the turnaround radius, and this is due to the definition of caustics surfaces \\citep{RG}.  In this paper, we discuss an approach to the issue of mass estimation, which can be applied to larger sampled volumes, well beyond the turnaround radius. We use the radial velocity of galaxies as the key quantity, instead of the escape velocity, as in the caustic technique. According to Silk \\citep{Si}, Peebles \\citep{P1,P2}, and Gunn \\citep{G}, within the spherical symmetry hypothesis, the velocity of the matter infall motion at a certain distance from the centre depends on the encompassed mass. Our purpose is to use this dependence to constrain the value of the overdensity at the turnaround radius. In fact, we see that the turnaround radius is far outside the virialization core of clusters, and is therefore a suitable normalization scale for the cluster mass profile in the non-equilibrium region \\citep{VH}. To test our assumptions and to verify the results, we will analys a large galaxy population extracted from a simulated cluster catalogue \\citep{Borgal,Bivial}. We will study all clusters both as a whole and one by one. We will prove that the actual turnaround overdensity of clusters is in good agreement with the predictions of the spherical infall model. Moreover, we will show that the normalized mass profiles are generally consistent with a power-law profile, which extends the standard Navarro--Frenk--White profile \\citep[hereafter NFW]{NFW1,NFW2,NFW3} to the non-equilibrium region.  In Section \\ref{sec:model} we present the details of our model, concerning the theoretical framework (\\ref{sec:theory}) and the simulated data sample (\\ref{sec:data}). In Section \\ref{sec:results} we discuss the results of our analysis, focusing on the mass estimation at the turnaround radius (\\ref{sec:turn}) and in the whole non-equilibrium region, up to 8 virialization radii (\\ref{sec:prof}). Finally, in Section \\ref{sec:concl} we draw the conclusions of our work. ",
        "conclusions": "\\label{sec:concl}  We analized a large sample of simulated galaxy clusters in order to reconstruct the mass profile in the non-equilibrium region, where the galaxy dynamics is dominated by an overall infall motion towards the cluster centre. Within the assumptions of the spherical infall model, the turnaround overdensity $\\delta_t$ can be theoretically computed as a function of only the matter density parameter $\\Omega_0$, assuming a spatially flat universe. We obtained the overdensity $\\delta_t\\simeq 6-15$, depending on the infall velocity profile we adopted.  We interpolated the infall velocity profile of member galaxies extracted from the simulated clusters of our catalogue, and we showed that: \\begin{enumerate} \\p The turnaround radius $r_t$ can be quite well approximated by a multiple of the virialization radius $r_v$: $r_t\\simeq 3.5r_v$; \\p The turnaround overdensity $\\delta_t$ is consistent with the prediction of the spherical infall model, as long as the infall velocity profile is described by the Meiksin approximation \\citep{VD}. \\end{enumerate} Points (i) and (ii) are in agreement with \\citet{VH} and \\citet{RG} and imply a proportionality between the turnaround mass $M_t$ and the virialization mass $M_v$: $M_t\\simeq 1.7 M_v$. Moreover, $M_t$ turns out to depend on the 3-d DM velocity dispersion within the virialization core $\\sigma_{v,\\rmn{DM}}$ approximately in the form of a cubic relation.  The turnaround values can be assumed as a suitable normalization scale for the mass profiles in the non-equilibrium region of clusters. We showed that the normalized mass profiles are generally consistent with a cosmic profile, which can be described (for $0.5\\la r/r_t\\la 2$) by: \\be\\label{eq:M_concl} M(r)\\simeq M_t\\exp\\left[\\f{0.6}{\\Omega_0^{1/4}}\\left(\\f{r}{r_t}-1\\right)\\right]. \\ee While in the inner, relaxed or almost-relaxed regions the mass can be considered independent on cosmological parameters, in the outer regions a dependence on $\\Omega_0$ (even if small) has to be taken, at least in principle, into account.  We used a synthetic cluster, obtained by summing all catalogue clusters, to determine a robust estimate of $r_t$ and $\\delta_t$. If we assume this values, our model is able to predict the mass profile in the non-equilibrium region for about $80\\%$ of clusters. So, it is possible to speak about a mass profile even in the region where mass accretion takes places along isolated radial filaments rather than in a spherically symmetric way.  Our model may be useful in observational analysis in order to estimate the total mass of clusters using the redshift-space distribution of galaxies. The method is the following: \\begin{enumerate} \\item one estimates the virialization radius and the virialization mass from the galaxy velocity dispersion in the cluster core; \\item using equation (\\ref{eq:st}) and equation (\\ref{eq:Mt}) one computes the turnaround radius $r_t$ and the turnaround mass of the cluster $M_t$; \\item once the turnaround radius and the turnaround mass are known, one can estimate the mass profile in the non-equilibrium region using the exponential law in equation (\\ref{eq:M_concl}). \\end{enumerate} The advantage of this approach lies in the possibility to estimate the mass profiles up to the far outskirts of clusters, where the caustic pattern is not generally recognizable \\citep{DG,D}; up to now, these cluster outer volumes have been usually neglected in the evaluation of cluster total masses.  Actually, the turnaround mass is a more exhaustive evaluation of the total mass of the cluster. The steps leading to it consist in the abovementioned points (i), (ii), and (iii), and are expected to be applied to observed clusters."
    },
    "0807/0807.2523_arXiv.txt": {
        "abstract": "A dynamical analysis of an effective homogeneous and irrotational Weyssenhoff fluid in general relativity is performed using the $1+3$ covariant approach that enables the dynamics of the fluid to be determined without assuming any particular form for the space-time metric. The spin contributions to the field equations produce a bounce that averts an initial singularity, provided that the spin density exceeds the rate of shear. At later times, when the spin contribution can be neglected, a Weyssenhoff fluid reduces to a standard cosmological fluid in general relativity. Numerical solutions for the time evolution of the generalised scale factor $R(t)$ in spatially-curved models are presented, some of which exhibit eternal oscillatory behaviour without any singularities. In spatially-flat models, analytical solutions for particular values of the equation-of-state parameter are derived. Although the scale factor of a Weyssenhoff fluid generically has a positive temporal curvature near a bounce (i.e. $\\ddot{R}(t)>0$), it requires unreasonable fine tuning of the equation-of-state parameter to produce a sufficiently extended period of inflation to fit the current observational data. ",
        "introduction": "The Einstein-Cartan (EC) theory of gravity is an extension of Einstein's theory of general relativity (GR) that includes the spin properties of matter and their influence on the geometrical structure of space-time ($\\cite{Cartan:1922}$; see also $\\cite{Kleinert:1989}$, $\\cite{Kleinert:2008}$). In GR, the energy-momentum of the matter content is assumed to be the source of curvature of a Riemannian space-time manifold $V_4$. In the EC theory, the spin of the matter has been postulated, in addition, to be the source of torsion of a Riemann-Cartan space-time manifold $U_4$ $\\cite{Hehl:1973}$. Weyssenhoff and Raabe $\\cite{Weyssenhoff:1947}$ were the first to study the behaviour of perfect fluids with spin. Obukhov and Korotky extended their work in order to build cosmological models based on the EC theory $\\cite{Obukhov:1987}$ and showed that by assuming the Frenkel condition$\\footnote{Note that the Frenkel condition arises naturally when performing a rigorous variation of the action. It simply means that the spin pseudovector is spacelike in the fluid rest frame.}$ the theory may be described by an effective fluid in GR where the effective stress-energy momentum tensor contains some additional spin terms.  The aim of this publication is two-fold. First, we wish to investigate the possibility that the spin contributions for a Weyssenhoff fluid may avert an initial singularity, as first suggested by Trautman $\\cite{Trautman:1973}$. Second, since any realistic cosmological model has to include an inflation phase to fit the current observational data, it is also of particular interest to see if the spin contributions are able to generate a dynamical model endowed with an early inflationary era, as first suggested by Gasperini $\\cite{Gasperini:1986}$. Scalars fields can generate inflation, but they have not yet been observed. Therefore, it is of interest to examine possible alternatives, such as a Weyssenhoff fluid. In contrast to the approaches of Trautman $\\cite{Trautman:1973}$ and Gasperini $\\cite{Gasperini:1986}$, our use of the $1+3$ covariant formalism enables us to determine the dynamics of a Weyssenhoff fluid without assuming any particular form for the space-time metric.  The study of the dynamics of a Weyssenhoff fluid in a $1+3$ covariant approach was initiated by Palle $\\cite{Palle:1998}$. His work has been revised and extended in our previous publication $\\cite{Brechet:2007a}$. The present paper builds on $\\cite{Brechet:2007a}$ to extend the work carried out first by Trautman $\\cite{Trautman:1973}$ in an isotropic space-time, and Kopczynski $\\cite{Kopczynski:1973}$ and Stewart $\\cite{Stewart:1973}$ in an anisotropic space-time. It also generalises the analysis of the inflationary behaviour of Weyssenhoff fluid models made by Gasperini $\\cite{Gasperini:1986}$ to anisotropic space-times.  In our dynamical analysis, we choose to restrict our study to a spatially homogeneous and irrotational Weyssenhoff fluid. This particular choice, which implies a vanishing vorticity and peculiar acceleration, has been motivated by underlying fundamental physical reasons. For a vanishing vorticity, the fluid flow is hypersurface-orthogonal, which means that the instantaneous rest spaces defined at each space-time point should mesh together to form a set of 3-surfaces in space-time $\\cite{Ellis:1971}$. These hypersurfaces, which are surfaces of simultaneity for all the fluid observers, define a global cosmic time coordinate determined by the fluid flow. Moreover, by assuming that any peculiar acceleration vanishes, the cosmic time is then uniquely defined. It is worth mentioning that the absence of vorticity is an involutive property, which means that if it is true initially then it will remain so at later times as shown by Ellis et al $\\cite{Ellis:1998}$. Finally, the assumption that there is no vorticity on all scales implies that the fluid has no global rotation. This is in line with recent Bayesian MCMC analysis of WMAP data performed by Bridges et {\\it al.} $\\cite{Bridges:2006}$. Their work confirms that a physical Bianchi $\\mathrm{VII_h}$ model, which has a non-vanishing vorticity, is statistically disfavored by the data.  It is worth pointing out that Szydlowski and Krawiec $\\cite{Szydlowski:2004}$ have considered an isotropic and homogeneous cosmological model in which a Weyssenhoff fluid is proposed as a potential candidate to describe dark energy at late times. In a subsequent publication $\\cite{Krawiec:2005}$, the authors showed that it is not disfavoured by SNIa data, but it may be in conflict with CMB and BBN observational constraints. By contrast, in this paper, we consider the full evolutionary history of an, in general, anisotropic universe with a Weyssenhoff fluid as its matter source, concentrating in particular on the `early universe' behaviour when the spin terms are significant. Indeed, at late times, when the spin contributions can be neglected, the Weyssenhoff fluid reduces to a standard cosmological fluid. We thus allow for the presence of a non-zero cosmological constant, in accord with current observational constraints.  In $\\Sref{Weyssenhoff fluid description}$, we give a concise description of a Weyssenhoff fluid using a 1+3 covariant approach outlined in {\\it Appendix} A. The spatial symmetries and macroscopic spin averaging procedure are discussed in $\\Sref{Spatial symmetries and macroscopic spin averaging}$. In $\\Sref{Dymamics of a homogeneous and irrotational Weyssenhoff fluid}$, we establish the relevant dynamical relations for a homogeneous and irrotational Weyssenhoff fluid. In $\\Sref{Geodesic singularity analysis}$, we perform a geodesic singularity analysis for such a fluid. In $\\Sref{Dynamical evolution: general considerations}$, we analyse the fluid dynamics. The behaviour of the generalised scale factor $R(t)$ of such a fluid in a spatially-curved models is discussed in $\\Sref{Quantitative dynamical evolution of spatially-curved models}$ and explicit analytical solutions in spatially-flat models are given in $\\Sref{homogeneous solutions}$ . For the reader's convenience, certain main results obtained in our earlier work $\\cite{Brechet:2007a}$ will be repeated in the case of a homogeneous and irrotational Weyssenhoff fluid in $\\Sref{Weyssenhoff fluid description}$ and $\\Sref{Dymamics of a homogeneous and irrotational Weyssenhoff fluid}$. In this paper, we use the $(+,-,-,-)$ signature. To express our results in the opposite signature used by Ellis $\\cite{Ellis:1998}$, the correspondence between physical variables can be found in $\\cite{Brechet:2007a}$. ",
        "conclusions": "We have used the $1+3$ covariant approach to perform a dynamical analysis of an effective homogeneous and irrotational Weyssenhoff fluid. Contrary to the case of a perfect fluid in GR, the effective spin contributions to the fluid dynamics act like centrifugal forces preventing the formation of singularities for isotropic and anisotropic models satisfying the spin-shear constraint $\\eref{cond spin shear}$. The temporal evolution of the models is symmetric with respect to $t=0$.  In a cosmological context, the energy density at the bounce state $\\rho_0$ has to be sufficiently dense in order to seed large scale structures from primordial quantum fluctuations. For cosmological parameters which are consistent with current cosmological data $\\eref{lambda condition}$ $\\eref{curvature condition}$, the temporal curvature of scale factor of a Weyssenhoff fluid is positively defined near the bounce $\\eref{a}$. However such a fluid is not a suitable candidate for inflation given that the only way to include an inflation phase of about $50-70$ e-folds, is by considering a fluid with a very fine-tuned equation-of-state $\\eref{State tunning}$, which does not reduce to the standard cosmological fluid at later times.  It is worth emphasizing that the time evolution of the scale factor of a homogeneous and irrotational Weyssenhoff fluid exhibits eternal oscillations, without any singularities. By contrast, the corresponding solutions obtained for a perfect fluid in GR are cycloids, which do exhibit singularities. Hence, the absence of singularities for a specific range of parameters is a genuinely new feature of cosmological models based on a Weyssenhoff fluid.  \\ack  S~D~B thanks the Isaac Newton Studentship and the Sunburst Fund for their support. The authors also thank Reece Heineke for giving an insightful talk entitled ``Inflation via Einstein-Cartan theory\", and John M. Stewart for useful discussions.  \\appendix"
    },
    "0807/0807.2653_arXiv.txt": {
        "abstract": "{ We employ a long \\xmmn observation of the core of the Perseus cluster to validate claims of a non-thermal component discovered with \\chandran. From a meticulous analysis of our dataset, which includes a detailed treatment of systematic errors, we find the 2-10 keV surface brightness of the non-thermal component to be smaller than about 5$\\times10^{-16}$erg$~$cm$^{-2}$s$^{-1}$arcsec$^{-2}$. The most likely explanation for the discrepancy between the \\xmmn and \\chandra estimates is a problem in the  effective area calibration of the latter. Our \\epic based magnetic field lower limits are not in disagreement with Faraday rotation measure estimates on a few cool cores and with a minimum energy estimate on Perseus. In the not too distant future \\emph{Simbol-X} may allow detection of non-thermal components with intensities more than 10 times smaller than those that can be measured with \\epicn; nonetheless even the exquisite sensitivity within reach for \\emph{Simbol-X} might be insufficient to detect the IC emission from Perseus. ",
        "introduction": "\\label{sec: intro}   Although the bulk of the energy radiated from clusters is of thermal nature, non-thermal mechanisms play an important role. Indeed the characterization of non-thermal components provides much needed clues on the physical process presiding over the formation and evolution of clusters. In some of the more disturbed objects, evidence of non-thermal processes has been known for quite some time. Radio observations indicate that merging clusters are often the site of cluster-wide synchrotron emission, the so called radio halos and radio relics \\citep[and references therein]{cassano07}. Radio data, polarimetric or not, has been used to provide estimates of magnetic fields in clusters. Unfortunately estimates based on Faraday rotation measures (hereafter RM) or minimum-energy arguments are affected by large uncertainties. In the case of RM estimates the unknown field topology \\citep{ensslin03} and the accessability of only a few, not necessarily representative, lines of sight \\citep{rudnick03}, are a major source of concern, while for minimum energy arguments the proton to electron ratio, and  the applicability of the argument itself play an equally critical role.  It has been recognized for quite some time that detection of inverse Compton (hereafter IC) emission at X-ray wavelengths can provide an alternative method to estimate cluster magnetic fields \\citep{rephaeli87}. Detection and characterization of the so called hard tails is by no means a trivial task: the signal, if there is one, is caught between the hammer of the thermal emission and the anvil of the instrumental background. So far, only detections at a few sigma level have been  reported \\citep[for a recent review see][]{rephaeli08}, and at least in one case, Coma \\citep{fusco99,fusco04,fusco07}, they have been challenged \\citep{rossetti04,rossetti07}.  Aim of the present work is to employ a long \\xmmn observation of the Perseus core to validate claims about a non-thermal emission component discovered with \\chandra \\citep{sand04,sand05,sand07}. The major advantage of the \\xmmn observation is the better sensitivity at high energies ($\\sim$ 6 keV) of the \\epic detectors with respect to the \\stn ; this is illustrated in Fig.~\\ref{fig: cxo_epc} where we show the ratio of background to source intensity in a representative region, a ring with bounding radii 1$^{\\prime}$-2$^{\\prime}$ centered on the emission peak. The \\chandra ratio rises very rapidly from 1\\% at 5keV to unity at 8 keV; at 10 keV the \\st background is about 20 times larger than the source. \\mos and \\pn ratios rise above 1\\% around 7 keV and remain below 30\\% and 10\\% respectively at all energies.  \\begin{figure} \\centering \\resizebox{75mm}{!}{\\includegraphics[angle=-90]{fig_cxo_pn_m2.ps}} \\\\ \\caption{Ratio of background to source spectra for the \\st (blue), \\md (green) and \\pn (red) instruments from the 1$^{\\prime}$-2$^{\\prime}$ ring. The \\st ratio rises rapidly from 1\\% at 5keV to unity at 8 keV; at 10 keV the \\st background is about 20 times larger than the source. \\mos and \\pn ratios rise above 1\\% around 7 keV and remain below 30\\% and 10\\% respectively at all energies.} \\label{fig: cxo_epc} \\end{figure}  The outline of the paper is as follows: in Sect.~\\ref{sec: ev_prep} we provide details on our data preparation; in Sect.~\\ref{sec: spec_anal} we describe results from the spectra analysis; in Sect.~\\ref{sec: chandra} we compare results from the analysis of \\epic data with those from \\chandran; in Sect.~\\ref{sec: discussion} we discuss our results and in Sect.~\\ref{sec: summary} we summarize them. We assume $H_0 = 70$~km s$^{-1}$Mpc$^{-1}$ so that 1 arcsec$=$0.35 kpc at the redshift of NGC1275 \\citep[0.0176,][]{huchra99}; all errors are 1$\\sigma$ unless otherwise stated. ",
        "conclusions": "\\label{sec: discussion}    From a general perspective  our analysis may be viewed as an attempt to characterize the spectrum of an X-ray source beyond a simple one/two component model. Given the modest intensity of the  component we are specifically interested in and the lack of spectral signatures, such as lines or edges, the task is by no means a trivial one. Clearly it can only be attempted on sources where adequate statistics is available. It is often the case that, when the statistics is abundant, systematics become the dominant source of error. At the present time X-ray astronomers do not have a standard method of analyzing data under these predicaments and mostly resort, has we have done, to  a trial and error approach. The development of a standard strategy, possibly descending from first principles, is highly desirable particularly if we consider that in the not too distant future X-ray missions such as \\emph{XEUS} or \\emph{Con-X} will have, as one of their primary goals, the detailed characterization of relatively bright X-ray spectra. In some instances the advent of high resolution spectrometry will provide a valid solution. In others, where features that need to be characterized are broad, either because they are of  non-thermal nature, such as the putative hard tails in clusters, or because they are smeared by other processes (i.e. the broad iron line observed in nearby AGN), solutions will have to be found elsewhere. Recent work by \\citet{drake07} provides an interesting starting point.  \\subsection{\\epic estimate} \\label{sec: dis_epic}  In this paper we have characterized  the surface brightness of the non-thermal component in Perseus  by  resorting to a heuristic approach which can be divided into 2 steps: 1) inclusion of 2\\% systematic errors on spectra; 2) modification of spectral models. As far as the first step is concerned, we find a significant reduction of the surface brightness measured from the \\pn spectra when the 2\\% systematic error is included, no significant changes are found for the \\mos measurements ( note however that a reduction is observed starting from a 3\\% systematic error, see Sect.~\\ref{sec: ring_1_2_mos} ). As far as the second step is concerned, the stability of our results with respect to the inclusion of corrections for some imperfections in redistribution matrix, effective area, energy scale and  choice of  astrophysical model indicate that our measures are relatively robust with respect to the residual systematic effects we have been able to identify. Comparison of \\mos with \\pn results shows that the former favour a somewhat larger value for the surface brightness of the non-thermal component than the latter, roughly speaking 2 against 1.  The difference between the two measures is most likely  caused by minor cross-calibration issues, possibly in the high energy response of either \\mos or \\pnn. Since we do not have firm evidence as to which of the two experiments is better calibrated, rather than privileging measures from one over the other, we  take the common envelope of \\mos and \\pn measurements (i.e. 0- 5$\\times10^{-16}$erg$~$cm$^{-2}$s$^{-1}$arcsec$^{-2}$) as our best estimate for the interval constraining the surface brightness of the non-thermal component. Some of our readers might find this range rather broad, particularly in light of the high statistical quality of our data, we reiterate that the major source of indetermination here as elsewhere are systematic and not statistical errors.  \\subsection{\\epic vs. \\acis } \\label{sec: dis_acis}  Results from the analysis of \\epic spectra, either \\pn or \\mosn, cannot be reconciled with those obtained with \\stn. As discussed in Sect.~\\ref{sec: chandra} the difference is to be ascribed to a cross-calibration issue between \\epic and \\acisn. Similar problems have been identified by a group of calibrators who have compared \\chandra and \\xmmn observations of hot clusters \\citep{david07}. Recent efforts by \\chandra calibrators \\citep{david07,marshall08} have shown that the above difference follows  from problems in the calibration of the \\chandra high energy effective areas.  \\subsection{Magnetic field estimates} \\label{sec: dis_b}  Under the assumption that the non-thermal emission originates from inverse Compton (IC) scattering of seed microwave and infrared photons by relativistic electrons, responsible of the synchrotron emission in the mini radio-halo, the Perseus core magnetic field can be estimated from the energy density of the seed photon field, $U_{rad}$, the synchrotron and inverse-Compton luminosities, $L_R$ and $L_X$, i.e. $ U_B = U_{rad} \\cdot L_R / L_X $, where $ U_B$ is the energy density of the magnetic field. \\citet{sand05}, from their measures of the non-thermal component, find magnetic field intensities ranging from several $\\mu$G at the very center, to $\\sim$ 1.0~$\\mu$G at 8 kpc (0.4$^\\prime$) and $\\sim$ 0.1~$\\mu$G at 40 kpc (2$^\\prime$) \\citep[see Fig.~9 of][]{sand05}. Our measurements can be used to provide a revised estimate of the B field. As discussed above we estimate the surface brightness of the non-thermal component to be somewhere between 0 and 5$\\times 10^{-16}$erg$~$cm$^{-2}$s$^{-1}$arcsec$^{-2}$.  We therefore start by assuming an upper limit of 5$\\times 10^{-16}$erg$~$cm$^{-2}$s$^{-1}$arcsec$^{-2}$, from this we subtract a value of 2.4$\\times 10^{-16}$erg$~$cm$^{-2}$s$^{-1}$arcsec$^{-2}$, as in \\citet{sand05}, to account for projection  effects. The resulting surface brightness is used to derive lower limits for the magnetic field using essentially the same method and radio measures described in \\citet{sand05}. In the 0.5$^{\\prime}$-1$^{\\prime}$ and 1$^{\\prime}$-2$^{\\prime}$ annuli our X-ray estimate convert into lower limits of roughly 0.4~$\\mu$G and 0.3~$\\mu$G respectively. We note that these numbers have been derived through a series of rather drastic approximations, (i.e. we have a rather poor determination of the X-ray upper limit, the radio emission is measured at a frequency which is substantially larger than that at which non-thermal electrons responsible for the inverse Compton emission emit via synchrotron, the B field is assumed to be constant, etc.) what really matter is that they are lower limits and not detections.  Faraday rotation measures (RM) in cool cores on scales of tens of kpc are in the order of several hundreds to thousands \\citep{taylor02}. Estimates of magnetic fields from rotation measures have undergone some revision in the last few years with more recent estimates typically in the order of a few $\\mu$G \\citep{clarke04,ensslin06}. In the case of the Perseus mini radio-halo  Faraday rotation measures are available only on very small scales \\citep{taylor06}, i.e. few tens of pc. RM estimates are in the order of $\\sim$7000 rad m$^{-2}$ leading to B field values of $\\sim 25~\\mu$G under the assumption the screen is localized in the ICM. This, however, appears to be unlikely as variations  of 10\\% in the RM are observed on scales of $\\sim$ 1 pc \\citep{taylor06}, while ICM magnetic fields are expected to be ordered on significantly larger scales \\citep[few kpc:][]{taylor02,vogt05,ensslin06}.  Application of the classical minimum-energy argument to the Perseus mini radio-halo data, leads to estimates for the central (i.e. $r=0$) magnetic field strength of $\\sim 7~\\mu$G \\citep{pfrommer04}.  IC estimates of the magnetic field based on \\chandra measurements \\citep{sand05}  are about an order of magnitude below the RM and minimum-energy estimates detailed above. A similar discrepancy has been found when comparing magnetic field estimates on cluster wide scales. There IC measures are in the 0.2-1~$\\mu$G range \\citep[and references therein]{rephaeli08}, while RM estimates are about an order of magnitude larger \\citep[e.g.][]{carilli02,rephaeli08}. While the intricate nature of the astrophysical scenario lends itself to various possible explanations for the discrepancy \\citep[e.g.][]{carilli02,govoni04}, a more trivial alternative, namely that the IC estimates might be incorrect, should also be considered. There is at least one object, Coma, where IC magnetic field measures \\citep{fusco99,fusco04,fusco07} have been challenged \\citep{rossetti04,rossetti07}. Moreover first measures from the \\emph{Suzaku} mission on a few objects, (i.e. A3667, Coma) are turning out to be  lower limits on IC measures \\citep[and references therein]{fuka07}. \\xmmn measures presented in this paper seem to play a similar role in the determination of  the Perseus core IC magnetic field estimates.        \\subsection{Future prospects} \\label{sec: dis_future}  It is unlikely that either \\chandra or \\xmmn  will provide significantly better estimates of the intensity of the non-thermal component in Perseus than those reported here. Clearly experiments with sensitivities extending into the hard X-ray band are the most appropriate for these kind of studies. Recently the \\emph{BAT} experiment on board the \\emph{Swift} satellite has been used to determine that an extrapolation of the non-thermal flux measured with  \\chandra to the 50-100 keV band overshoots the flux measured with \\emph{BAT} by a factor of about 4 \\citep{ajello08}. The \\emph{Suzaku} high energy experiment \\citep{taka07}, limited as it is by modest spatial resolution and by a non optimal background treatment \\citep{kokubun07}, may provide some useful indications but will most likely not allow substantial advancement. Amongst missions that are currently under development \\emph{Simbol-X} \\citep{ferrando05} is arguably one of the most promising. The combination of large throughput in the 1-60 keV range, low instrumental background and \\epicn-like spatial resolution, will allow a sensitive measurement of non-thermal components extending beyond the thermal cutoff. We have performed simulations of \\emph{Simbol-X} spectra based on  our best fits for the 1$^{\\prime}$-2$^{\\prime}$ annulus trying out different values of the normalization of the non-thermal component to determine how far down we might go. We started by taking a value of the surface brightness of 1$\\times10^{-16}$erg$~$cm$^{-2}$s$^{-1}$arcsec$^{-2}$. We find that the non-thermal component dominates thermal emission above $\\simeq$ 20 keV and remains above background emission up to $\\simeq$ 30 keV. Thus, if the non-thermal emission has an intensity of this magnitude or larger it will be detected rather easily. Repeating the same exercise with a surface brightness that is ten times smaller, i.e. 1$\\times10^{-17}$erg$~$cm$^{-2}$s$^{-1}$arcsec$^{-2}$, we find that formal fitting of the simulated spectrum still allows a detection of the non-thermal component at the $\\sim 3-4 \\sigma$ level, provided the observation is longer than $3\\times 10^5$s and current estimates of the \\emph{Simbol-X} instrumental background are within  10-20\\% of the real value. We note that in this case the relative intensity of the non-thermal component in the 20-30 keV band is about 10\\%.  As we have learned from the analysis of our \\epic data, at these intensity levels, systematics, which are at this time unknown for \\emph{Simbol-X}, will play an important, possibly dominant role. In the specific case of \\emph{Simbol-X} measurements, the critical elements which need to be kept under control at the few percent level are: the effective area calibration; the cross-calibration between the high and low energy detectors and the background. These requirements, albeit challenging, are not beyond reach, particularly if we consider that the formation flight strategy adopted by \\emph{Simbol-X} will entail some important advantages over previous missions. Extensive calibration of the telescope effective areas can be performed in flight by comparing observations of calibration sources performed with and without the telescope in the optical path. Moreover direct illumination of the \\emph{Simbol-X} focal plane, with the telescope removed, will allow an in-flight verification of the detector quantum efficiency.  Assuming that \\emph{Simbol-X} can indeed reach a sensitivity of 1$\\times10^{-17}$erg$~$cm$^{-2}$s$^{-1}$arcsec$^{-2}$ this will allow to detect IC components with associated B fields of roughly 1~$\\mu$G, thereby providing important constraints on the magnetic field in Perseus. Given the current magnetic field estimates from Faraday rotation measures on a few  cool cores \\citep[few $\\mu$G:][]{ensslin06} and the  classical minimum-energy argument estimate on Perseus \\citep[$\\sim~7\\mu$G:][]{pfrommer04}, even the exquisite sensitivity within reach for \\emph{Simbol-X} might be insufficient to detect the IC emission.      We have carried out a detailed analysis of a long \\epic observation of the Perseus core in an attempt to detect and characterize a non-thermal component, our main findings may be summarized as follows:  \\begin{itemize} \\item  systematic uncertainties play an important role in the characterization of the non-thermal component; in the absence of a strategy descending from first principles we have developed a heuristic approach to include them in our analysis; \\item  at variance with our preliminary estimates, we find that the non-thermal component is not detected; the surface brightness is determined to be smaller than $\\sim$ 5$\\times10^{-16}$erg$~$cm$^{-2}$s$^{-1}$arcsec$^{-2}$; \\item  our \\xmmn estimates are at variance with \\chandra estimates from \\citet{sand05,sand07} and from our own analysis; the most likely explanation for the discrepancy between \\chandra and \\xmmn is a problem in the \\chandra effective area calibration; \\item  our \\epic based upper-limit on the surface brightness converts into IC magnetic field lower limits of $\\sim$0.4~$\\mu$G for the 0.5$^{\\prime}$-1$^{\\prime}$ annulus and $\\sim$0.3~$\\mu$G for the 1$^{\\prime}$-2$^{\\prime}$ annulus; these measures are  not in disagreement with RM estimates on a few cool cores \\citep[few $\\mu$G:][]{ensslin06} and the minimum energy estimate on Perseus \\citep[10~$\\mu$G:][]{pfrommer04}; \\item in the not too distant future \\emph{Simbol-X} may allow detection of non-thermal components with intensities more than 10 times smaller than those that can be measured with \\epicn; nonetheless even the exquisite sensitivity within reach for \\emph{Simbol-X} might be insufficient to detect the IC emission from Perseus. \\end{itemize}"
    },
    "0807/0807.1424.txt": {
        "abstract": "We have conducted aperture polarimetry of $\\sim$500 stars of the Orion Nebula Cluster (ONC) in M42 based on our wide-field ({$\\sim$}8$'${$\\times$}8$'$) $JHKs$ band polarimetry. Most of the near-infrared (NIR) polarizations are dichroic, with position angles of polarization agreeing, both globally and locally, with previous far-infrared (FIR) and submillimeter observations, having taken into account the 90$^\\circ $ difference in angles between dichroic absorption and emission. % thermal emission polarizations, %\\cup except for the 90$^{\\circ}$ difference of the position angles. This is consistent with the idea that both NIR dichroic polarizations and FIR/submillimeter thermal polarizations trace the magnetic fields in the OMC-1 region. The magnetic fields inferred from these observations show a pinch at scales less than 0.5 pc with a centroid near IRc2. The hourglass-shaped magnetic field pattern is explained by the models in which the magnetic field lines are dragged along with the contracting gas and then wound up by rotation in a disk. The highly polarized region to the northwest of IRc2 and the low-polarized region near the bright bar are also common among NIR and FIR/submillimeter data, although a few regions of discrepancy exist.  We have also discerned {$\\sim$}50 possible highly polarized sources whose polarizations are more likely to be intrinsic rather than dichroic. Their polarization efficiencies ($P(H)/A(H)$) are too large to be explained by the interstellar polarization. These include 10 young brown dwarfs that suggest a higher polarization efficiency, which may present geometrical evidence for (unresolved) circumstellar structures around young brown dwarfs. ",
        "introduction": "Magnetic fields are believed to play an important role in the evolution of molecular clouds, from their large-scale structures to dense cores, protostar envelopes, and protoplanetary disks. Magnetic fields can be measured by observing polarizations of electromagnetic waves. Dichroic polarizations are believed to be caused by spinning non-spherical dust grains becoming aligned with their short axis precessing around the direction of the local magnetic field.  In absorption this produces polarization parallel to the magnetic field and perpendicular to the magnetic field in emission (e.g., Weintraub et al. 2000). %Dichroic polarizations are believed to be caused by magnetically aligned, nonspherical grains, %which thus follow the magnetic fields (e.g., Weintraub et al. 2000). Optical stellar polarimetry is useful for revealing magnetic field structures in the periphery of molecular clouds (e.g., Appenzeller 1974; Vrba et al. 1976; Moneti et al. 1984; Goodman et al. 1990).   However, because of heavy dust extinction, optical polarimetry is not always a good tracer of the magnetic fields inside molecular clouds. Since the pioneering work by Vrba et al. (1976) and Wilking et al. (1979), near-infrared (NIR) stellar polarimetry in star-forming regions has been a powerful tool for tracing magnetic field structures in molecular clouds (Tamura et al. 1987; Sato et al. 1988; Tamura et al. 1988; Klebe \\& Jones 1990). The stars could be either background stars or even embedded sources if they are not associated with circumstellar structures (i.e., have no intrinsic polarization). Although arguments have been made against the usefulness of NIR polarimetry in tracing magnetic fields inside molecular clouds and despite a suggestion that the dust grains are not well aligned in the dense regions (Goodman et al. 1995), recent submillimeter polarimetry has clearly shown aligned dust grains, even in cold regions such as starless globules and cores (Kirk, Ward-Thompson, \\& Crutcher 2006). With the advent of wide-field NIR polarimetric capability (Tamura et al. 1996), reevaluating the utility of NIR stellar polarimetry is warranted because the wide field allows simultaneous polarimetry of many stars in the field of view. The Orion Nebula Cluster (ONC), or Trapezium cluster, is an ideal site for such a study. The ONC is one of the nearest (450 pc), massive, star-forming regions to the Sun and the most populous young cluster within 2 kpc, composed of some 3500 young, low-mass stars (O'Dell 2001).\\\\ %  The magnetic field structure of the Orion region was most extensively studied by observing linearly polarized thermal emission from aligned dust grains. Houde et al. (2004) and Schleuning (1998) showed at 350 $\\mu$m and 100$\\mu$m, respectively, that the magnetic field in OMC-1 is generally oriented northwest--southeast, with the field pinched along the northeast--southwest direction on a scale of several arcmin. The polarization percentage is low at the location of the Becklin--Neugebauer object and Kleinmann--Low nebula (BN/KL) compared to elsewhere, but that may be due to a small-scale variation that is undetectable in their 12$''$ - 35$''$ resolution maps. See Cudlip et al. 1982, Hildebrand et al. 1984, Dragovan 1986, Barvainis et al. 1988, Gonatas et al. 1990, Leach et al. 1991, and Rao et al. 1998 for earlier or other millimeter and submillimeter studies in the Orion region.  Imaging polarimetry of M42 has been conducted either only at optical wavelengths (Pallister et al. 1977) or toward a small region near IRc2 or BN at near-infrared wavelengths (Minchin et al. 1991; Jiang et al. 2005; Simpson et al. 2006). % Hough et al. (1986), Burton et al. (1991), and Chrysostomou et al. (1994) measured the polarization of the H$_2$~v~=~1-0~S(1) line at $2.12\\mu$m around IRc2. They attributed the polarization in the vicinity of IRc2 to dichroic absorption, and in particular, the twist in the polarization vectors 5$''$ east of IRc2 to a twist in the magnetic field direction. % At NIR wavelengths, the other IRc sources (IRc2 is not readily visible at 2~$\\mu$m, probably because it is behind the edge of the hot core) are all strongly polarized, with centrosymmetric polarization vectors indicating scattered light; that is, the IRc sources are illuminated by either BN (IRc1) or some star in the vicinity of IRc2 (Werner et al. 1983; Minchin et al. 1991; Chrysostomou et al. 1994).  The high degree of polarization of BN, which is proportional to the extinction with the FIR-inferred magnetic field directions, has been used to deduce that BN is polarized by dichroic absorption at all NIR and MIR wavelengths where it has been measured (e.g., Lee \\& Draine 1985; Hough et al. 1996; Aitken et al. 1997; Smith et al. 2000). However, Jiang et al. (2005) used NIR adaptive optics-aided, very high-resolution imaging polarimetry to successfully reveal an outflow/disk system around BN. % %    In this paper, we present the aperture polarimetry in the $J, H,$ and $Ks$ bands of $\\sim 500$ stars of the ONC with the polarimetry mode of the IRSF/SIRIUS instrument (SIRPOL). We also discuss possible intrinsically highly polarized sources and the polarization of young brown dwarfs. ",
        "conclusions": "We conducted deep $JHKs$ imaging polarimetry of a 7$\\farcm$7 $\\times$ 7$\\farcm$7 area of M42 and here present the results of the aperture polarimetry. Our main conclusions are summarized as follows:  \\begin{enumerate}  \\item Most of the NIR polarizations are dichroic. Their global and local vector patterns are in good agreement with previous FIR and submillimeter polarization patterns, except for the 90$^{\\circ}$ difference of position angles.  \\item A positive correlation exists between $P(H)$ and $H-Ks$; most of $P(H)$ is below $P_{max}$, except for possible highly polarized sources. This supports an intracloud origin of the NIR polarizations, excluding (a) a non-intracloud magnetic origin of polarization, and (b) a weaker alignment of dust grains in the cloud.  \\item We argue that both NIR dichroic polarizations and FIR/submillimeter thermal polarizations trace the magnetic fields in the OMC-1 region.  \\item The magnetic fields are pinched at scales less than 0.5 pc with a centroid near IRc2. The hourglass-shaped magnetic field pattern is explained by models in which the magnetic field lines are dragged along with the contracting gas and then wound up by rotation in a disk.  \\item The highly polarized region to the northwest of IRc2 and the low-polarized region near the bright bar, commonly seen in both NIR and FIR/submillimeter data, are explained by the field geometry; the latter is the field along and the former perpendicular to the line of sight.   \\item We also discriminated $\\sim$50 possible highly polarized sources whose polarizations are more likely to be intrinsic rather than dichroic. Their polarization efficiencies ($P(H)/A(H)$) are too large to be explained by normal interstellar polarization.  \\item For 9 young brown dwarf candidates, we also suggest the existence of higher polarization efficiency, which may present geometrical evidence for (unresolved) circumstellar structures around young brown dwarfs.  \\end{enumerate}  We are grateful to Shuji Sato for helpful suggestions. Thanks are due to the staff in SAAO for their kind help during the observations. We thank D. C. Lisfor kindly providing the 350 \u0083\u00cam map in FITS format. We also thank Noboru Ebizuka Tetsuo Nishino, and Toshihide Kawai for their technical support in the development of SIRPOL. The IRSF/SIRIUS project was initiated and supported by Nagoya University, National Astronomical Observatory of Japan, and The University of Tokyo in collaboration with South African Astronomical Observatory under a financial support of Grants-in-Aid for Scientific Research on Priority Area (A) No. 10147207 and No. 10147214, and Grants-in-Aid No. 13573001 and No. 16340061 of the Ministry of Education, Culture, Sports, Science, and Technology of Japan. This study was partly supported by a MEXT Grant-in-Aid for Scientific Research in Priority Areas, \u0081gDevelopment of Extrasolar Planetary Science,\u0081h and by grants-in-aid fromMEXT(Nos. 16077204, 16077171, and 19204018)."
    },
    "0807/0807.2605_arXiv.txt": {
        "abstract": "The mass reinserted by young stars of an emerging massive compact cluster shows a bimodal hydrodynamic behaviour. In the inner part of the cluster, it is thermally unstable, while in its outer parts it forms an out-blowing wind. The chemical homogeneity/inhomogeneity of low/high mass clusters demonstrates the relevance of this solution to the presence of single/multiple stellar populations. We show the consequences that the thermal instability of the reinserted mass has to the galactic super-winds. ",
        "introduction": "The open star clusters and stellar moving groups have internally homogeneous chemical composition. Clusters like the Hyades, Collinder 261, the Herculis stream or the moving group HR 1614 are chemically unique, distinguishable one from the other, showing no pollution from secondary star formation \\cite[(De Silva et al., 2008)]{Silva3}. The chemical homogeneity of open star clusters like the Hyades \\cite[(De Silva et al., 2006)]{Silva1} and Collinder 261 \\cite[(De Silva et al., 2007)]{Silva2} proves that they have been formed out of a well-mixed cloud and that any self-enrichment of stars did not take place there.  Young and massive stellar clusters, frequently called super star clusters, are preferentially observed in interacting galaxies. Their stellar mass amounts to several million M$_{\\odot }$ within a region less than a few parsecs in diameter. They represent the dominant mode of star formation in starburst galaxies. Their high stellar densities resemble those of globular clusters, where several stellar populations have been observed \\cite[(Piotto, 2008)]{GP}.  To explain the presence of multiple stellar generations in globular clusters, the slow wind emerging from a first generation of fast rotating massive stars is invoked by \\cite[Decressin et al. (2007)]{TDetal}. (See also the review by \\cite[Meynet (2008)]{Meynet} in this volume.) The authors argue that the fast rotating massive stars function as a filter separating the H-burning products from later products of He-burning. However, it is not clear why all the massive stars rotate fast, or why the slow wind produced by stellar rotation is just retained inside the potential well of the stellar cluster.  An alternative  solution, how to form the  second generation of  stars in massive star clusters, is proposed in models of star cluster winds described by \\cite[Tenorio-Tagle et al. (2007)]{GTTetal}, \\cite[W\\\" unsch et al. (2007)]{RWetal1}, and \\cite[W\\\" unsch et al. (2008)]{RWetal2}. There, we argue that a critical mass of a cluster exists, below which the single-mode hydrodynamical solution to the cluster winds applies. Such clusters should have one stellar generation only, and show strong winds corresponding to the momentum and energy feedback of all their stars. The clusters above the critical mass should follow the bi-modal solution to their winds, where only the outer skin of the cluster participates in the wind. Their inner parts are thermally unstable, and hence being the potential places of secondary star formation. ",
        "conclusions": "We propose a bimodal solution, where in the central part of a massive star cluster, a thermally unstable region forms (see Fig. 2). In this region, the thermal instability creates cold regions surrounded by hot medium imploding into them. The high-velocity wing of broad spectral lines \\cite[Gilbert, \\& Graham (2007)]{AMG&JRG} observed in SSC may be created by imploding shock in the vicinity of thermally unstable parcels of gas.  The second generation of stars may be formed out of cold clumps produced by thermal instability. During the early evolution of a massive cluster, the first Myrs, the mechanical energy input is dominated by stellar winds \\cite[(Leitherer et al. 1999)]{Leithereretal}. The efficiency of thermalization $\\eta $ may be low in this case, and a massive cluster may be above $M_{SC, crit}$, since its value is low. Later, the importance of winds fades out, and the mechanical input is dominated by supernovae. This may increase the thermalization efficiency $\\eta $, increasing at the same time the value $M_{SC, crit}$. Thus the same massive cluster, which was initially in the bi-modal situation moves to single-mode situation.  The cold parcels of gas form, in the cluster central part during the early bi-modal situation, from the winds that are enriched by products of H-burning. The later He-burning products are inserted into the cluster volume when the mechanical energy input is dominated by supernovae, which may mean that the cluster is in the single-mode situation, and the wind clears its volume from He burning products.  Thus, the thermal instability, which operates during a few initial Myr in the central part of the cluster, may produce a second generation of stars enriched by H-burning products. Later, the cluster moves into the single-mode situation, which means that it is able to expel the He-burning products.  The feedback of massive stars in super star clusters creates galactic winds, or super winds, reaching to large distances from the parent galaxies, transporting the products of stellar burning into intergalactic space.  The bimodal solution, providing a possible explanation of multiple stellar populations in globular clusters, limits the super winds. During the initial period of cluster evolution, when the  stellar winds dominate the mechanical energy input, the super wind is restricted only to the outer skin of the cluster, which means that it is rather weak. Only later, when supernovae become dominant  in mechanical energy input, strong super winds blowing out from all the cluster volume may reach  large distances from their parent galaxies. How effective the super winds of super star clusters can be in transporting the products of stellar evolution into intergalactic spaces should be discussed in future."
    },
    "0807/0807.2433_arXiv.txt": {
        "abstract": "Much progress has been made in measuring black hole (BH) masses in (non-active) galactic nuclei using the tight correlation between stellar velocity dispersions $\\sigma$ in galaxies and the mass of their central BH. The use of this correlation in quasars, however, is hampered by the difficulty in measuring sigma in host galaxies that tend to be overpowered by their very bright nuclei.  We discuss results from a project that focuses on $z\\sim0.3$ quasars suffering from heavy extinction at shorter wavelengths. This makes it possible to obtain clean spectra of the hosts in the spectral regions of interest, while broad lines (like H$\\alpha$) are still visible at longer wavelengths. We compare BH masses obtained from velocity dispersions to those obtained from the BLR and thus probe the evolution of this relation and BH growth with redshift and luminosity.  Our preliminary results show an offset between the position of our objects and the local relation, in the sense that red quasars have, on average, lower velocity dispersions than local galaxies.   We discuss possible biases and systematic errors that may affect our results. ",
        "introduction": "Black hole (BH) mass is believed to be one of the fundamental parameters that characterize quasar activity and much effort has been devoted to obtaining accurate BH masses for quasars and other AGN \\citep[\\eg][]{ho99}. In recent years, much progress has been made in measuring BH masses in galactic nuclei, particularly with the remarkable discovery by \\citet{geb00a} and \\citet{fer00} of a tight correlation between stellar velocity dispersion in galaxies and the mass of their central BH (M$_{\\rm BH}\\propto\\sigma^{n}$). The use of this correlation to derive BH masses in AGN, however, is hampered by the difficulty in measuring velocity dispersions in host galaxies that tend to be overpowered by their very bright nuclei.  Nevertheless, the correlation has been shown to be present at low redshift ($z<0.1$) in low luminosity AGN (\\eg\\ BL Lac objects: \\cite{barth03}; or Seyfert galaxies: \\cite{geb00b}).  Seyfert galaxies at higher redshift ($z\\sim0.36$ and $z\\sim0.57$), however, appear to show an offset from the local relation \\citep[][and references therein]{Woo2008}.  It is not yet known whether the M$_{\\rm BH}-\\sigma$ correlation holds for the highest luminosity AGN. A loose correlation has been found by using the width of [O\\,III] lines in active nuclei \\citep{nel00,shi03}, but the width of these lines is dependent upon other parameters (outflows, radio luminosity, etc.) and therefore lead to a correlation with a large scatter. BH masses derived from [O\\,III] emission line widths can only be accurate to within a factor of five at best \\citep{bor03}.  More accurate determinations are necessary if we hope to use them to disentangle some of the other fundamental relationships among quasar parameters. ",
        "conclusions": ""
    },
    "0807/0807.0436_arXiv.txt": {
        "abstract": " ",
        "introduction": "Cataclysmic variable stars (CVs) consist of a Roche-lobe filling main sequence star, called the secondary star, orbiting a white dwarf - the primary star. Mass flows from the secondary star through the inner Lagrangian point towards the primary star, ultimately leading to the formation of an accretion disc (AD). After the AD's formation, a bright spot (BS) is formed at the contact point where the gas stream hits the AD. Dwarf novae (DNe) are a subclass of cataclysmic variables with two different brightness states, quiescence and outburst. In the latter, the system usually brightens by 2--5\\,mag with a time-scale that can vary from tens to hundreds of days. DNe can also show a third brightness state, the superoutburst, where the system brightens by typically an additional magnitude or two compared to normal outburst.  IP Peg is the brightest eclipsing ($i \\approx81^{\\circ}$) DN above the period gap, and was discovered by \\citet{lip81}. It has a 3.8-h orbital period and a $V$ magnitude in the range of 14--12\\,mag, the extreme values corresponding to quiescence and outburst, respectively.  From the first photometric light curves during quiescence and return to quiescence  \\citep{woo86,woo89}, it was evident that the BS produces a prominent orbital hump, when it rotates into view from behind the disc with a single-stepped ingress and double-stepped egress superimposed on a more gradual disc-eclipse. Infrared light curves showed 0.2\\,mag ellipsoidal variations from the secondary star superposed on a deep eclipse of the WD and the BS \\citep{szk86}. \\citet{woo89}, based on the analysis of 39 eclipses, found out that IP Peg has a rapidly increasing orbital period. However, a new study performed by \\citet{wol93} and combining 49 eclipses along with previously reported ones,  reached the conclusion that the orbital period varies sinusoidally with a period of 4.7 years. This variation they attributed to a third body, a late M dwarf star with a most probable mass of 0.10\\,$\\rm M_\\odot$. They also showed that the rate of shrinkage in disc radius during decline from outburst is constant during different outburst cycles, following an exponential decay. A study of the flickering in IP Peg \\citep{bru00} showed that the BS is the dominant source of flickering, with the AD and WD contributing practically nothing.  Spectroscopy of the red star \\citep{mar89} showed TiO bands and the NaI doublet (red dwarf features), several broad double-peaked emission lines (disc features) and a narrow third peak in the emission lines moving in anti-phase with the double peaks (a chromospheric emission feature). \\citet{bee00} performed a TiO study of IP Peg and concluded that TiO instead of the NaI doublet should be used for radial velocity studies of the secondary star. In this way they found that the contamination from disc emission, which in other studies returned an elliptical orbit, could be avoided and higher signal-to-noise ratio could be achieved. They gave a final projected orbital velocity of the companion star $K_2=331.3\\,\\rm km\\,s^{-1}$ and refined values of the primary and secondary star masses $M_1=1.05\\,\\rm M_\\odot$ and $M_2=0.33\\,\\rm M_\\odot$, respectively. Covering a complete cycle from quiescence to outburst and quiescence again, the spectral characteristics of IP Peg will now be described.  {\\it Quiescence spectra} \\citep{mar88} show broad double-peaked emission lines on a blue continuum, a strong orbital hump in the continuum (but much weaker in the lines) and a flat Balmer decrement indicative of optically thick emission. The trailed spectra show the signature of a prograde rotation of the AD, where the blue-shifted emission is eclipsed prior to the red-shifted one, a weak S-wave component in the red-shifted wing and a 10\\% blue/red asymmetry before eclipse, most probably arising from line emission from the BS.  {\\it Early rise-to-outburst spectra} \\citep{wol98} show Balmer emission, TiO absorption bands, KI and NaI absorption doublets, a weak S-wave and two bright spots in both Doppler and eclipse maps. The study of \\citet{wol98}, as well as a re-analysis of their results using an extension of the classical eclipse mapping method \\citep{bob99}, concluded that the first spot should be identified as the BS while two possibilities existed for the second spot. It either was the beginning of formation of a spiral arm or a second inner BS due to gas-stream overflow as predicted by, e.g., simulations performed by \\citet{arm96}.  {\\it Rise-to-outburst spectra} \\citep{har94} show a decrease in the equivalent width of H$\\alpha$ emission line, down to a level similar to that observed immediately after outburst and, simultaneously, weak H$\\alpha$ emission from the secondary due to irradiation of the secondary star by the boundary layer.  {\\it Near-outburst spectra} \\citep{pic89} show unusually strong high-excitation lines, Balmer lines in emission, indications of a very disturbed AD with a non-uniform He emission distribution, an outflowing wind from the WD vicinity and chromospheric emission from the secondary.  {\\it Outburst spectra} as presented by \\citet{mar90} show: i) strong Balmer and HeII $\\lambda$4686 double-peaked emission lines rising from the AD, ii) non-Keplerian emission filling in the centre of HeII $\\lambda$4686, probably due to a compact outflowing wind or inflowing magnetic accretion columns from close to the WD, iii) Balmer but not HeII emission from the secondary, located in the polar region on the side facing the AD, iv) no evidence for an enhanced gas-stream compared to quiescence, v) large azimuthal asymmetry in the outer disc, with the blue-shifted side of the disc being brighter and possibly more extended at phase 0 and vi) a sharp transient component in the HeII line possibly generated from the interaction of the gas stream and disc. \\citet{mar90} suggested that He emission can be explained as photoionisation of the disc by the boundary layer and found that Balmer emission in outburst can be attributed by 3/4 to the disc, 1/4 to irradiation of the red star and also to boundary layer irradiation. Another study performed by \\citet{ste96} showed, from trailed spectra and Doppler maps, asymmetric disc emission, strong secondary star emission and a low velocity emission in the H$\\alpha$ line only. This stationary emission, appearing as an emission blob in the centre of the Doppler maps, was interpreted as slingshot prominences, material trapped in magnetic loops and co-rotating with the secondary star.  During another outburst maximum, \\citet{har99a}  found metal line emission such as MgII. Doppler tomography of the He and metal lines locates those lines' emission on the inner Roche lobe of the secondary star.  {\\it Early decline-from-outburst spectra} \\citep{hes89} reveal for the first time in such strength a blend of high-excitation lines of HeII, CIII and NIII along with a very strong chromospheric emission-line component from the secondary.  {\\it Decline-from-outburst spectra} \\citep{mor00} show irradiation-induced emission from the companion star in Balmer, He, MgII and CII lines. This emission, located only near the poles of the secondary, suggests shielding of the equatorial regions of the secondary star by the vertically extended AD.  {\\it Early-quiescence spectra} in the infrared \\citep{lit01} reveal the existence of mirror eclipses in IP Peg. A mirror eclipse is an eclipse of the secondary star by an optically thin AD and appears in the trailed spectra as a reduction in the equivalent widths of the lines, at phase 0.5. A mirror eclipse also has a mirror-symmetry with the classical rotational disturbance, i.e. the red-shifted part of the line is eclipsed prior to the blue-shifted one.  Even though IP Peg had shown many peculiarities and novelties in the studies cited previously, it had to reveal another breakthrough. \\citet{ste97} were the first to find convincing evidence for the existence of spiral structure in the AD of a close binary. Spiral waves had long been predicted in studies of tidal interactions and numerical simulations (see \\citet{bof01} and references therein) and may play a crucial role in our understanding of the angular momentum transport. \\citet{ste97} obtained spectrophotometric observations of IP Peg during the rise to an outburst, covering H$\\alpha$ and HeI $\\lambda$6678. Disc emission was centred on the WD and had a strong azimuthal asymmetry in the forms of two spiral arms. The repeatability of the two-armed spiral structure was confirmed during another outburst maximum. \\citet{har99a} applied Doppler tomography of the HeII $\\lambda$4686 line and clarified the shock nature of the spiral arms which showed a jump of a factor of more than 2 in intensity. Additionally they calculated an azimuthal extent of 90$^{\\circ}$, a 30$^{\\circ}$ upper limit in opening angle and a 15\\% contribution of the shocks to the total disc emission. \\citet{mor00} showed that the spiral pattern persists even 5 and 6 days after the outburst maximum, without showing any diminution in strength. Time-resolved optical spectra obtained during late outburst by \\citet{bap00} reveal, through the eclipse mapping technique, that spirals are present even 8 days after the onset of an outburst. Their study showed a similar contribution to the continuum emission, as that measured during the peak of the outburst but also a retrograde azimuthal rotation of 58$^{\\circ}$ compared to outburst maximum.  \\citet{bap05} performed a re-analysis of previous spectroscopic data of IP Peg during outburst \\citep{mor00}, based on the eclipse mapping technique. The difference in their new study was based on the fact that they separated the inner and outer regions of the emission line light curves, obtaining in this way the ``core'' and ``spiral arms'' lights curves, respectively. In this way they improved the previously noisy view of the spiral arms and were able to trace them in detail. The two-armed spiral structures were clearly visible in all their eclipse maps, albeit the ``blue'' arm appearing farther out from the disc's centre than the ``red'' one. However, their crucial findings were the sub-Keplerian velocities along the spiral structures, as well as the clear correlation between the opening angle of the spirals and the outburst stage (the later in the outburst, the smaller the opening angle of the spirals). Both these findings greatly favour the interpretation of the spiral asymmetric structures as tidally-induced spiral shocks instead of mere radiation patterns on an otherwise unperturbed AD.  A study performed by \\citet{ste01} presents a review of the spiral structure that far, and plots in detail the azimuthal angles and extent of the spirals for the HeII $\\lambda$4686 emission line during an outburst presented by \\citet{har99}. In more detail, it is shown that the two spirals can be traced for almost 180$^{\\circ}$, while their velocity varies from 495 to 780\\,km\\,s$^{-1}$. Under the assumption of Keplerian velocities, they show that the arms cover a substantial part of the disc, 0.3--0.9 times the distance to the inner Langragian point $L_1$. It is pointed out that tracing the spiral arms and plotting their position and intensity as a function of azimuth throughout an outburst, or even better applying this procedure to a range of emission lines simultaneously, can yield important results on the evolution of the spirals. A future paper, as stated in the Discussion section, is intended to address this issue. A first approach on tracing the spiral arms of two different emission lines (H$\\alpha$, HeI $\\lambda$6678), as well as plotting their emissivity as a function of azimuth, is presented in \\citet{ste01}. It is shown that, even though there is a similar position for the spirals in terms of velocity as a function of azimuth, the intensity modulation is significantly different.  \\begin{figure} \\centering \\includegraphics[width=13cm]{fig01.eps} \\caption{AAVSO long-term light curve of visual validated data. The arrow shows the timing of our observations, just 1 day after the peak of the outburst.} \\label{f_aavso} \\end{figure}  Our work is part of a project of studying CVs through high resolution time-resolved echelle spectra and applying the indirect imaging technique of Doppler tomography. The simultaneous study of several emission lines can yield important results on the AD structure, as presented for another CV system in \\citet{pap08}. Following the introduction on IP Peg one can conclude on the benefits of applying such an analysis to this object, especially since it was observed during outburst. In this way we can study the behaviour of the spiral arms and map the structure of the AD in detail. The application of Modulation Doppler tomography, an extension of the classic Doppler tomography method \\citep{ste03}, gives us the opportunity to map emission sources which can vary harmonically with the orbital period. Such an analysis is a first for IP Peg.  This paper is organised as follows: in Sect.~2 we present our observations, in Sec.~3 the data reduction, in Sect.~4 we perform the spectral analysis and Sect.~5 deals with the application and results of Modulation Doppler tomography. Last, Sect.~6 is devoted to the discussion and summary of our results. ",
        "conclusions": ""
    },
    "0807/0807.1430_arXiv.txt": {
        "abstract": "Resonant relaxation (RR) is a rapid relaxation process that operates in the nearly-Keplerian potential near a massive black hole (MBH). RR dominates the dynamics of compact remnants that inspiral into a MBH and emit gravitational waves (extreme mass ratio inspiral events, EMRIs). RR can either increase the EMRI rate, or strongly suppress it, depending on its still poorly-determined efficiency. We use small-scale Newtonian $N$-body simulations to measure the RR efficiency and to explore its possible dependence on the stellar number density profile around the MBH, and the mass-ratio between the MBH and a star (a single-mass stellar population is assumed). We develop an efficient and robust procedure for detecting and measuring RR in $N$-body simulations. We present a suite of simulations with a range of stellar density profiles and mass-ratios, and measure the mean RR efficiency in the near-Keplerian limit. We do not find a statistically significant dependence on the density profile or the mass-ratio. Our numerical determination of the RR efficiency in the Newtonian, single-mass population approximations, suggests that RR will likely \\emph{enhance} the EMRI rate by a factor of a few over the rates predicted assuming only slow stochastic two-body relaxation. ",
        "introduction": "Dynamical relaxation processes near massive black holes (MBH) in galactic centers affect the rates of strong stellar interactions with the MBH, such as tidal disruption, tidal dissipation, or gravitational wave (GW) emission \\citep[e.g.][]{ale05}. These relaxation processes may also be reflected by the dynamical properties of the different stellar populations there \\citep{hop+06a}, as observed in the Galactic Center \\citep{gen+00,pau+06}. Of particular importance, in anticipation of the planned Laser Interferometer Space Antenna (LISA) GW detector, is to understand the role of relaxation in regulating the rate of GW emission events from compact remnants undergoing quasi-periodic extreme mass ratio inspiral (EMRI) into MBHs.  Two-body relaxation, or non-coherent relaxation (NR), is inherent to any discrete large-N system, due to the cumulative effect of uncorrelated two-body encounters. These cause the orbital energy $E$ and the angular momentum $J$ to change in a random-walk fashion ($\\propto\\!\\sqrt{t}$) on the typically long NR timescale $T_{\\mathrm{NR}}$. In contrast, when the gravitational potential has approximate symmetries that restrict orbital evolution (e.g. fixed ellipses in a Keplerian potential; fixed orbital planes in a spherical potential), the perturbations on a test star are no longer random, but correlated, leading to coherent ($\\propto\\! t$) torquing of $J$ on short timescales, while the symmetries hold. Over longer times, this results in resonant relaxation (RR) (\\citealt{rau+96,rau+98}; \\S \\ref{ss:RR}), a rapid random walk of $J$ on the typically short RR timescale $T_{\\mathrm{RR}}\\!\\ll\\! T_{\\mathrm{NR}}$. RR in a near-Keplerian potential can change both the direction and magnitude of $\\mathbf{J}$ ({}``scalar RR''), thereby driving stars to near-radial orbits that interact strongly with the MBH. RR in a near-spherical potential can only change the direction of $\\mathbf{J}$ ({}``vector RR'').  RR is particularly relevant in the potential near a MBH, where compact EMRI candidates originate. \\citet{hop+06a} show that RR dominates EMRI source dynamics. Depending on its still poorly-determined efficiency, RR can either increase the EMRI rate over that predicted assuming NR only, or if too efficient, it can strongly suppress the EMRI rate by throwing the compact remnants into infall (plunge) orbits (cf Fig. \\ref{f:EMRI} below) that emit a single, non-periodic and hard to detect GW burst. A prime motivation for the systematic numerical investigation of RR efficiency presented here, are the still open questions about the implications of RR for EMRI rates and orbital properties.  This paper is organized as follows. In \\S \\ref{s:theory} we briefly review the theory of NR and RR relaxation and derive a new relation between scalar and vector RR. In \\S \\ref{s:ACF} we describe our method of analyzing and quantifying the effects of RR in $N$-body simulations, which are described in \\S \\ref{s:simulations}. We present our results in \\S \\ref{s:results} and discuss and summarize them in \\S \\ref{s:conclusions}. ",
        "conclusions": "\\label{s:conclusions}  We characterized and measured the mean efficiency coefficients of NR ($\\alpha_{\\Lambda},\\eta_{s,\\Lambda},\\eta_{v,\\Lambda}$) and RR ($\\beta_{s},\\beta_{v}$) in Newtonian $N$-body simulations of isotropic, thermal, near-Keplerian stellar cusps around a MBH. We derived a simple analytical form for the rms RR auto-correlation curves of $\\mathbf{J}$ and $J$, and showed that the two are simply proportional to each other. We then measured these coefficients in a large suite of small scale $N$-body simulations with different stellar density distributions and MBH/star mass ratios. We found no statistically significant trends in the values of these coefficients as function of the system properties. This may require better statistics.  Our measured RR efficiency suggests that RR increases the EMRI rate by a factor of $\\sim\\!5$ above what is predicted for NR only. This estimate of RR efficiency is consistent with that suggested by the analysis of the dynamical properties of the different stellar populations in the Galactic Center \\citep{hop+06a}. However, this conclusion is still preliminary, since several important open issues remain, which should be addressed by larger scale simulations. These include (1) The dependence of the RR coefficients on the orbit of the test star, for example its eccentricity \\citep{gur+07}. We find that $N\\!\\sim\\!200$ is not enough for reliable statistics on sub-samples within given \\emph{E} or \\emph{J}-bins. The eccentricity dependence of RR is particularly relevant for the supply rate of stars to the MBH from $J\\!\\rightarrow\\!0$ orbits (the loss-cone refilling problem). (2) The effects of a stellar mass spectrum. This will likely affect RR by decreasing the RR timescale, and changing the stellar density distribution through strong mass segregation (\\citealt{ale07}; Alexander \\& Hopman 2008, in prep.). (3) The robustness of RR against perturbations from the larger non-Keplerian stellar system in which the inner near-Keplerian region of interest is embedded. (4) The role of post-Newtonian effects in RR, such as General Relativistic precession and GW emission. These are expected to play a key role in enabling inspiral by quenching RR just as the compact remnant enters the EMRI phase, and in regulating the GW inspiral rate \\citep{hop+06a}."
    },
    "0807/0807.3329_arXiv.txt": {
        "abstract": "{  We review the results of the first multi-scale, hydrodynamical simulations of mergers between galaxies with central supermassive black holes (SMBHs) to investigate the formation of SMBH binaries in galactic nuclei. We demonstrate that strong gas inflows due to tidal torques produce nuclear disks at the centers of merger remnants whose properties depend sensitively on the details of gas thermodynamics. In numerical simulations with parsec-scale spatial resolution in the gas component and an effective equation of state appropriate for a starburst galaxy, we show that a SMBH binary forms very rapidly, less than a million years after the merger of the two galaxies, owing to the drag exerted by the surrounding gaseous nuclear disk. Binary formation is significantly suppressed in the presence of a strong heating source such as radiative feedback by the accreting SMBHs. We also present preliminary results of numerical simulations with ultra-high spatial resolution of $0.1$~pc in the gas component. These simulations resolve the internal structure of the resulting nuclear disk down to parsec scales and demonstrate the formation of a central massive object ($\\sim 10^8 \\Mo$) by efficient angular momentum transport due to the disk's extended spiral arms. This is the first time that a radial gas inflow is shown to extend to parsec scales as a result of the dynamics and hydrodynamics involved in a galaxy merger, and has important implications for the fueling of SMBHs. Due to the rapid formation of the central clump, the density of the nuclear disk decreases significantly in its outer region, reducing dramatically the effect of dynamical friction and leading to the stalling of the two SMBHs at a separation of $\\sim 1$~pc. We discuss how the orbital decay of the black holes might continue in a more realistic model which incorporates star formation and the multi-phase nature of the ISM. ",
        "introduction": "In recent years, compelling dynamical evidence has indicated that supermassive black holes (SMBHs) are ubiquitous in galactic nuclei (e.g., Ferrarese \\& Ford 2005). According to the  standard modern theory of cosmological structure formation, the Cold Dark Matter (CDM) paradigm (e.g., Blumenthal et al. 1984), galaxies in the Universe grow through a complex process of continuous mergers and agglomeration of smaller systems. Thus, if more than one of the protogalactic fragments contained a SMBH, the formation of SMBH binaries during galaxy assembly will be almost inevitable (e.g., Begelman et al. 1980).  In a purely stellar background, as the binary separation decays, the effectiveness of dynamical friction slowly declines, and the pair can become tightly bound via three-body interactions, namely by capturing stars that pass close to the black holes and ejecting them at much higher velocities (e.g., Milosavljevi{\\' c} \\& Merritt 2001). If the hardening continues to sufficiently small relative distances, gravitational wave emission becomes the dominant source of orbital energy loss and the two SMBHs may coalesce in less than a Hubble time. However, the binary orbit may stop shrinking before gravitational radiation becomes relevant as there is a finite supply of stars on intersecting orbits (e.g., Berczik et al. 2005).  During the assembly of galaxies, especially at high $z$, their SMBHs likely evolve within gas-rich environments. Merging systems such as the Ultraluminous Infrared Galaxies (ULIRGs) NGC 6240 and Arp 220 harbor large concentrations of gas, in excess of $10^9 \\Mo$, at their center, in the form of either a turbulent irregular structure or of a kinematically coherent, rotating disk (e.g., Downes \\& Solomon 1998). Massive rotating nuclear disks of molecular gas are also ubiquitous in galaxies that appear to have just undergone a major merger, such as Markarian 231 (Davies et al. 2004). Gas dynamics may thus profoundly affect the pairing of SMBHs both during and after their host galaxies merge (e.g., Escala et al. 2004; Kazantzidis et al. 2005).  Recent simulations of the orbital evolution of SMBHs within an equilibrium, rotationally-supported, gaseous disk have shown that dynamical friction against the gaseous background leads to the formation of a tightly bound SMBH binary with a final separation of $<1$~pc in about $10^7$~yr (Escala et al. 2005; Dotti et al. 2006; Dotti et al., these proceedings). Here we review the results of high-resolution $N$-body + smoothed particle hydrodynamics (SPH) simulations of mergers between galaxies with central SMBHs having enough dynamic range to follow the black holes from hundreds of kiloparsecs down to sub-parsec scales, bridging more than ten orders of magnitude in density. ",
        "conclusions": ""
    },
    "0807/0807.3745_arXiv.txt": {
        "abstract": "We revisit arguments that simple models of inflation with a small red tilt in the scalar power spectrum generically yield an observable tensor spectrum. We show that criteria for fine-tuning based upon the algebraic simplicity of the potential depend strongly upon the explicit assumptions they incorporate, particularly regarding the end of inflation. In addition, some models with algebraically simple potentials require carefully tuned initial field configurations, and not all types of fine-tuning are identifiable via the algebraic simplicity of the potential. Conversely, in the absence of a strong prior on the mechanism that ends inflation, we demonstrate the existence of potentials with vanishingly small tensor amplitudes which are natural in terms of both their algebraic form and initial conditions. We thus argue that proposed experiments (CMBPol or BBO) which make highly sensitive measurements of the tensor amplitude cannot definitively rule out the inflationary paradigm. ",
        "introduction": "In the simplest scenarios, inflation is driven by a single scalar field, the ``inflaton'', rolling in a slowly varying potential  \\cite{Guth:1981,Sato:1981,Mukhanov:1981,Linde:1982,Albrecht:1982}. Sufficient inflation ($\\approx 60$ $e$-folds) smooths the universe and solves the flatness, horizon, and monopole problems. Initial inhomogeneities originate as quantum fluctuations in the inflaton field and the metric, sourcing scalar \\cite{Hawking:1982,Guth:1982, Starobinsky:1982, Bardeen:1983} and tensor \\cite{Grishchuk:1975, Starobinsky:1979} perturbations which are almost Gaussian and scale-invariant \\cite{Bardeen:1983}. These broad-brush predictions are consistent with observations  (e.g. \\cite{Komatsu:2008hk,Peiris:2006ug,Peiris:2006sj,Peiris:2008be}).  The shape of the inflaton potential is not predicted by theory; it is only required to satisfy the slow-roll conditions. However, the shape of the potential is constrained by the scalar spectral index, $n_s$, which measures the deviation of the scalar power spectrum from scale invariance, and the tensor-scalar ratio, $r$, which measures the amplitude of the tensor modes relative to that of the scalar modes. Data from the {\\sl Wilkinson Microwave Anisotropy Probe} (WMAP) combined with a compilation of supernovae and baryon acoustic oscillations data gives $n_s =0.960^{+0.014}_{-0.013}$ at 68\\% CL if the tensor-scalar ratio $r=0$, and $n_s=0.968 \\pm 0.015$ when we marginalize over $r$  \\cite{Dunkley:2008ie, Komatsu:2008hk}. There is currently no evidence for the existence of a primordial tensor background: the 95\\% CL upper limit from the same data compilation is $r<0.20$.  A detection of primordial tensors would have tremendous implications for cosmology as their existence is a key prediction of inflation \\cite{Starobinsky:1979}. Moreover, measuring a non-zero value of $r$ would eliminate the ekpyrotic/cyclic scenarios, which predict an unobservably small primordial tensor amplitude \\cite{Khoury:2001bz,Boyle:2003km}. Finally, the observed amplitude of the tensor spectrum fixes the energy scale of inflation, removing a major uncertainty in inflationary model building. On the other hand, an observable tensor spectrum would imply the need for super-Planckian field variations during inflation in simple single field models \\cite{Lyth:1996im}, which raises specific theoretical challenges.  Cosmic variance, polarized foregrounds, and weak lensing of the $E$-mode may ensure that $r \\sim 10^{-3}$ is the effective lower limit for a detection of tensor modes in the cosmic microwave background (CMB) (e.g. \\cite{Knox:2002pe,Kesden:2002,Verde:2005ff,Amarie:2005in}). Consequently, it is important to know whether inflation makes a {\\em generic\\/} prediction for the value of $r$ and, if so, whether it falls within the detectable range. A plethora of explicit inflationary models with varying degrees of physical motivation have been proposed. Many models do have observable values of $r$ (e.g. chaotic inflation \\cite{Linde:1983}); conversely, a detection of tensors would rule out many \\cite{Baumann:2006cd}, but not all \\cite{Dimopoulos:2005ac,Easther:2005zr,Silverstein:2008}, explicit string-theory constructions. Consequently, to show that inflation has a ``natural'' range of $r$ values, we must establish a ``weight'' that will give preference to certain classes of models.  One weighting proposal is that of \\BST\\ (BST), who offer a fine-tuning statistic based on the algebraic simplicity of the potential, along with five conditions which a successful inflationary potential should satisfy. These criteria are then applied to quartic polynomial potentials, with the conclusion that a high degree of fine-tuning is required for $n_s>0.98$, or for $r<10^{-2}$ with $n_s>0.95$. If future observations confirm that $0.95 < n_s< 0.98$, the BST criterion suggests that $r$ is within the observable range.  In addition to tuning the {\\em parameters\\/} of the potential, some potentials need a very special initial field configuration if inflation is to begin. This tuning can involve both the suppression of the kinetic energy relative to the potential energy  (to ensure that the inflaton does not ``overshoot'' the inflationary portion of the potential) or the absence of inhomogeneities on the initial inflationary patch. This topic has been investigated for many of the common models of inflation \\cite{Belinsky:1985, Goldwirth:1990, Mendes:2000,Underwood:2008dh, Chiba:2008}.  In particle physics, models are regarded as fine-tuned if they require correlations between parameters beyond those imposed by the symmetries of the underlying theory \\cite{tHooft:1979bh}, or when certain parameters are arbitrarily set to zero; models which are not fine-tuned in this sense are known as ``technically natural'' \\cite{Hotchkiss:2008sa}.  There are other definitions of naturalness based on the sensitivity of observables to small changes in their parameters \\cite{Barbieri:1988, Anderson:1994dz}. Ref. \\cite{Hotchkiss:2008sa} examined technically natural potentials which did not obey all of the BST viability conditions, and found red-tilted examples with $r<0.01$ which were not fine-tuned under their definition. Other studies came to similar conclusions \\cite{Efstathiou:2006, Smith:2008pf}. Inflationary potentials are often written down without reference to an underlying theory, so we do not necessarily know the relevant symmetries, nor the couplings between the inflaton and other fields which contribute loop corrections to the inflaton potential. Consequently, while technical naturalness is a powerful tool, it is less useful when applied to generic inflationary potentials. Moreover, inflationary potentials typically require some tuning in order that inflation happens at all.  We tackle two questions. Firstly, we show that even a mild relaxation of the rules imposed by BST significantly weakens their conclusions for the likely range of $r$. In addition, we ask whether the inflaton's initial conditions must be tuned in order for viable inflation to occur. We consider only the ratio of the kinetic to potential energy, treating the universe as a homogeneous system and ignoring the ability of primordial inhomogeneities to suppress the onset of inflation. However, this requirement places a strong constraint on the inflationary model space, which does not overlap with the BST proposal, demonstrating the interplay between the {\\em definition\\/} of tuning, and the resulting conclusions for the likely values of $r$ and $n_s$. Our aim is not to assess the intrinsic merit of different naturalness criteria. Rather, our purpose is to investigate whether conclusions about the natural value of $r$ are sensitive to assumptions encoded in the tuning criteria. This question is of considerable importance when assessing and justifying CMB $B$-mode or gravitational wave experiments sensitive to $r$. Likewise, any conclusions we reach regarding the ``natural'' value of $r$ in inflation will determine the significance of a null result from these experiments. ",
        "conclusions": "\\label{sec:conclusions}  We have investigated whether, given the observed limits on the scalar tilt $n_s$, it is possible to claim that certain ranges of values for the tensor-scalar ratio $r$ are ``fine-tuned'' in a way that is independent of the definition of fine-tuning. We find that, on the contrary, different criteria are sensitive to different types of fine-tuning, and hence can give differing answers about the naturalness of a given potential. The BST criterion, based on the number of unnecessary ``features\" in the potential, tells us that models with $r<10^{-2}$ are fine-tuned for $n_s > 0.95$, while an alternative criterion based on sensitivity to initial conditions shows that models with $n_s<0.94$ and $r<10^{-3}$ are fine-tuned, but with $n_s>0.94$ there are ``natural'' models with $r \\sim 10^{-6}$.  Moreover, BST's conclusions are not robust against relatively minor modifications to their criteria.  In particular, the claim that models with $n_s>0.98$ are fine-tuned is softened by a small technical change in the fine-tuning criterion $Z_\\eta$: our upper limit on natural values of the spectral index is $n_s <0.99$.  This may seem like a fine distinction, but recall that the breaking of scale invariance is measured by $n_s-1$, and thus differs by a factor of 2.  Moreover, Planck is expected to measure  $n_s$ with a 1-$\\sigma$ error of $0.0045$ \\cite{Planck} and  $n_s = 0.99$ cannot be easily distinguished from $n_s = 1$, but $n_s=0.98$ could be detected with considerable confidence.  Likewise,  BST's  conclusion that $n_s>0.95$ implies $r>10^{-2}$ in the absence of fine-tuning only holds if the potential evolves smoothly to an analytic minimum with $V(0) \\approx 0$. This is equivalent to requiring that the parameters fixing the shape of the potential as cosmological perturbations are laid down also govern the shape of the potential as inflation ends.   While it is not inherently unreasonable, this restriction immediately excludes both hybrid inflation and many stringy constructions, even though their underlying potentials may be technically natural from a particle physics perspective. Pending a better understanding of physics at the energy scale of inflation, this condition therefore seems overly restrictive.  Our alternative criterion is based on the sensitivity of a particular potential to the inflaton's initial conditions. It shows that potentials with $n_s<0.94$ and $r<10^{-3}$ are very sensitive to their initial conditions, as are all models  with $r<10^{-6}$. Since inflation was originally designed to solve the initial condition problems of a hot big bang, one should certainly check that inflation itself does not require the universe to be in a special initial state.   For our ``scoring'' of the initial conditions problem, there is a small preference for models with higher $r$, but the {\\em specific\\/} results we find depend on the explicit construction of our fine-tuning criterion, so we do not see this as a robust conclusion. Rather, we emphasize that this criterion singles out a different set of potentials from the BST criterion, and neither set of potentials is a strict subset of the other.   The above conclusions apply to potentials that are well-described by a quartic potential. In Appendix \\ref{potentialparams}, we explore the regions of parameter space  which yield untuned potentials. \\begin{figure*}[!htp] \\includegraphics[scale=0.175]{ri-05_10-zn-1.png}  \\hfill \\includegraphics[scale=0.175]{cubic-both-05.png} \\\\ \\caption{(Top) The set of models which are deemed natural by both the BST and initial conditions criteria, i.e. $Z_\\eta<2$ and $R_i>0.5$. The color bar shows $R_i$. (Left) Imposing all five BST conditions. (Right) Neglecting selection conditions (\\ref{cond:min}) and (\\ref{cond:meta}). Potentials which have low $r \\sim 10^{-4}$ and $n_s \\sim 0.95$ generically contain a flat plateau terminating in a drop, similar to Case 2 in Figure~\\ref{fig:cases}.} \\label{combinedft} \\end{figure*}  We have shown that conclusions drawn from   BST-like criteria depend on the explicit assumptions they encode.   We also discussed a second type of fine-tuning -- sensitivity to the initial field configuration -- which is distinct from that described of BST.   One may restrict attention to models deemed acceptable  by  {\\em both\\/} criteria. The left panel of Figure~\\ref{combinedft} indicates that, at face value, this would lead one to conclude that the ``natural'' lower limit on the tensor amplitude is   $r \\gtrsim 0.04$, and thus well within the range that is experimentally accessible.   However, the right panel shows that this conclusion only holds if condition (\\ref{cond:min}) is imposed.   Moreover, there are types of fine-tuning missed by our initial conditions criterion (e.g. initial inhomogeneity or spatial curvature); including these would further reduce the number of models for which inflation can be regarded as generic.   The BST criterion aims to characterize the fine-tuning of the {\\it shape} of the potential.  The initial conditions criterion might, at face value, appear to measure only the sensitivity to an initially large kinetic term. However, it is important to realize that it also encodes information about the functional form of the potential: for instance, a potential with a steep hill, a small plateau, and another steep downhill stretch will have difficulty supporting inflation without carefully tuned initial conditions. Thus, $R_i$ penalizes the functional form of the potential, but in a different way to $Z_\\eta$. Hence, Occam's razor cannot  characterize fine-tuning (even of a single property, such as the shape of the potential) in an unequivocal way. These ambiguities will remain after the arrival of data with sufficient precision to select a small region of the $(r, n_s)$ plane. If this region is disfavored according to a specific fine-tuning criterion, should we interpret this as evidence against inflation, or even (in a more limited sense) as evidence against a class of inflationary models? Absent a detailed understanding of the physics of inflation, the answer is unclear, since we would not know whether the fine-tuning identified by that criterion is relevant to the underlying theory.    As is well known, simple single field models of inflation have a direct  correlation between the total excursion of the inflaton field  and the value of $r$.  If values of $r>10^{-3}$  are ruled out by experiment, this would also eliminate all  single field models with $\\Delta \\phi>M_P$, with dramatic consequences for inflationary model building,  as well as for alternative theories of the early universe \\cite{Khoury:2001bz,Boyle:2003km}. Many proposed inflationary models predict signals in this range but only a handful of models live in the lower part, the so-called ``tensor desert'' at  $10^{-3}<r<10^{-2}$ \\cite{Smith:2008pf,Alabidi:2006fu, Efstathiou:2006}. Conversely, many other models have $r \\ll 10^{-3}$.   Consequently, it is clear that the inflationary paradigm would not be ruled out by a failure to detect primordial tensors, although  a large and important class of models would be conclusively eliminated.  Our experience with different fine-tuning criteria suggests that it is likely to be impossible to construct a set of criteria which defines a ``generic inflationary model'' that is simultaneously (a) robust against small changes, (b) does not eliminate apparently reasonable models from consideration, and (c)  leads to a definitive conclusion for the likely value of the tensor-scalar ratio $r$. Rather, the value of analyses like that of BST lies in the extent to which they identify a category of inflationary models which would be falsified if future experiments fail to detect tensors. However, determining that $r\\lesssim 0.001$ -- or any other value likely to be experimentally accessible -- cannot be construed as a definitive test of inflation.  \\appendix"
    },
    "0807/0807.4353_arXiv.txt": {
        "abstract": "The shear viscosity of the crust might have a damping effect on the amplitude of r-modes of rotating neutron stars.  This damping has implications for the emission of gravitational waves.  We calculate the contribution to the shear viscosity coming from the ions using both semi-analytical methods, that consider binary collisions, and Molecular Dynamics simulations.  We compare these results with the contribution coming from electrons.  We study how the shear viscosity depends on density for conditions of interest in neutron star envelopes and outer crusts.  In the low density limit, we find good agreement between results of our molecular dynamics simulations and classical semi-analytic calculations. ",
        "introduction": "Neutron stars are good resonators where oscillation modes can be excited. In particular, the r-modes of rotating neutron stars involve currents associated with very small density variations. These modes are unstable at all rates of rotation in a perfect fluid star \\cite{AnderssonGW}. The instability is due to the emission of gravitational radiation, suggesting the possibility of gravitational wave detection with Laser Interferometer Gravitational-Wave Observatory (LIGO) \\cite{Owen-LIGO}. This instability is expected to spin down newly born hot neutron stars \\cite{LindblomGW}. However, the observation of colder rapidly rotating neutron stars, suggests the existence of a damping mechanism of the r-mode instability. Several works have been done to explain this mechanism. For example, in Ref. \\cite{Bildsten2000} the damping mechanism is suggested to be the result of a viscous layer at the interface between the solid crust and the fluid core. Other works discuss the r-mode dynamics of superfluid neutron stars \\cite{Andersson2001,Lindblom2000} finding that a core filled with neutron and proton superfluids limits the amplitude growth of the modes \\cite{Lee_r-modes}.  In addition, high multipolarity p-mode oscillations may impact the pulse shape of some radio pulsars \\cite{pulseshape}.  For p-modes the primary restoring force is the pressure and the modes may be damped by the shear viscosity of the neutron star crust \\cite{Chugunov-shear}.  In the outer crust, the total shear viscosity has contributions from electrons and ions both of which transport momentum. Previous works have calculated the electron contribution to the shear viscosity in the Born approximation \\cite{Flowers1979} and including non-Born corrections \\cite{Chugunov-shear}.  Recently Horowitz and Berry calculated the electron contribution to the shear viscosity of non-spherical nuclear pasta phases \\cite{viscos_pasta}.  In this work we study the dependency of the shear viscosity with the density. We calculate the contribution of the ions to the shear viscosity in the neutron stars crust by two different methods: via Molecular Dynamics (MD) simulations and calculating momentum transport cross sections. The first method follows the Kubo formalism \\cite{Kubo}, and calculates the autocorrelation function of the pressure tensor. The second method considers binary collisions and allows one to consider both the classical and quantum systems. We focus our study to the case in which the ions form a dilute Hydrogen One Component Plasma (OCP). The conditions of temperature and density are chosen to reproduce those of the envelope and outer crust. However, the formalism can be applied to calculate other transport properties, such as diffusion coefficients. In particular the numerical procedure can be extended to study Multi-Component Plasmas (MCP).  This paper is organized as follows:  in Sec. \\ref{sec:formalism} we describe our ion-ion interaction model, the Kubo formalism, and the semi-analytical procedure to calculate transport coefficients from transport cross sections. In Sec. \\ref{dilute} we present our numerical and semi-analytical results for the dependency of the shear viscosity with density. In Sec. \\ref{quantum} we discuss quantum results vs classical ones. In Sec. \\ref{electron-shear} we compare our results with the contribution coming from the electrons, and finally in Sec. \\ref{conclusions} we conclude. The appendix contains the description of the method employed to calculate the phase shifts. ",
        "conclusions": "\\label{conclusions} We have calculated the ion contribution to the shear viscosity in the outer crust of neutron stars. We use two different methods for this purpose. One of these is based on MD simulations while the other is semi-analytical and calculates the transport cross sections for classical and quantum systems.  We find good agreement between the two methods in the low density, classical limit.  Using the MD trajectories, we have used autocorrelation functions of the pressure tensor in a Hydrogen OCP to calculate its viscosity. We have studied the case in which the plasma was dilute. In this case the coupling between the ions is weak. We have found that the viscosity is constant at low densities and then becomes an increasing function of density, contrary to the prediction including only binary collisions, where the viscosity remains constant with density.  This fact is due to the correlations between ions. The dependence of the plasma viscosity with temperature, at a fixed density, is the one expected for a gas. However, as the temperature decreases changes in the viscosity with density, are larger due to stronger ion correlations.  In the dilute gas limit and for light ion systems, calculations based on the Chapman-Enskog formalism, indicate that viscosities with quantum transport cross sections are nearly a factor of ten larger than the classical results for a fixed screening length of 10 fm. However, for larger more realistic screening lengths, the difference between quantum and classical calculations is smaller.   Using the structure factor of the ions we calculate the contribution to the shear viscosity due to electron-ion scattering. We find that the contribution to the shear viscosity from the ions is negligible compared to the former one for the conditions of interest in the outer crust of neutron stars. Therefore, we do not expect major damping of the r-modes from the ions in the crust.  Our method of calculating the electron contribution via the structure factor of the ions also allows us to use multi-component plasmas, and the contribution of electron-impurity scattering will be automatically included. On the other hand, we have used ion autocorrelation functions of the pressure tensor, and the Kubo formalism to calculate the viscosity of the ions. This method could be extended to calculate other mechanical properties of the ions, like the shear modulus, of great importance in the understanding of starquakes."
    },
    "0807/0807.2449_arXiv.txt": {
        "abstract": "We introduce the method and the implementation of a cosmological simulation of a class of metric-variation $f(R)$ models that accelerate the cosmological expansion without a cosmological constant and evade solar-system bounds of small-field deviations to general relativity. Such simulations are shown to reduce to solving a non-linear Poisson equation for the scalar degree of freedom introduced by the $f(R)$ modifications. We detail the method to efficiently solve the non-linear Poisson equation by using a Newton-Gauss-Seidel relaxation scheme coupled with multigrid method to accelerate the convergence. The simulations are shown to satisfy tests comparing the simulated outcome to analytical solutions for simple situations, and the dynamics of the simulations are tested with orbital and Zeldovich collapse tests. Finally, we present several static and dynamical simulations using realistic cosmological parameters to highlight the differences between standard physics and $f(R)$ physics. In general, we find that the $f(R)$ modifications result in stronger gravitational attraction that enhances the dark matter power spectrum by $\\sim 20\\%$ for large but observationally allowed $f(R)$ modifications. More detailed study of the non-linear $f(R)$ effects on the power spectrum are presented in a companion paper. ",
        "introduction": "\\label{sec:intro}  Explaining the accelerated cosmological expansion has become one of the greatest challenges in modern day physics \\citep[see e.g.,][]{frieman08a}. Of the many potential explanations, two classes are popular today: ones that propose new and unseen form of energy with negative pressure, and ones that modify the Einstein-Hilbert action. Of particular interest in recent literature are a class of modifications to the Einstein action by an addition of non-linear functions of the Ricci scalar, $R$, which have been demonstrated to cause accelerated expansion for a range of $f(R)$ functions \\citep{carroll04a,nojiri03a,capozziello03a}.  In addition to explaining the cosmic acceleration, such $f(R)$ modifications must satisfy stringent constraints imposed by solar system tests \\citep{williams96a,anderson98a,bertotti03a,williams04a} as well as cosmological constraints imposed by recent dark energy measurements \\citep{knop03a,riess04a,spergel07a}. Recently, a new class of $f(R)$ models was proposed by Hu \\& Sawicki \\citep[][hereafter HS model]{hu07a} that has the flexibility to reproduce the observed cosmological acceleration as well as to satisfy solar system bounds on deviations from the Newtonian limits of general relativity. In the proposed model, the scalar degree of freedom introduced by the $f(R)$ modification becomes massive in regions of high curvature and is hidden by the so-called chameleon mechanism \\citep{khoury04a,khoury04b}.  In such models that evade both the cosmological background test and solar system tests, the next strongest constraint comes from the effects of the $f(R)$ modifications to the intermediate lengths scales, in the regime of cosmological large scale structures. Outside of dense dark matter halos, the light scalar degree of freedom produces a long range fifth force that enhances gravitational attraction and thereby potentially change clustering of matter. Such effects were studied in the linear regime \\citep{song07a,hu07a,pogosian08a}, in which large deviations of dark matter power spectrum were obtained even for small $f(R)$ modifications. Furthermore, an attempt to model some of the non-linear chameleon physics using a parametrized framework on top of a halo model are presented in \\cite{hu07b}. However, those studies remain inconclusive because the effects of non-linear growth of structure and the chameleon mechanism (which requires non-linear growth) are not included in the analyses. For such non-linear study, a realistic N-body simulation of structure formation in $f(R)$ cosmologies is needed.  In this paper, we introduce a formalism and an implementation of a cosmological N-body simulation with fully non-linear treatment of $f(R)$ modifications. Although somewhat limited in the attainable dynamic range, we show that our implementation can capture the essential non-linear features of the HS model that have large impact on the cosmological large scale structure formation. In a companion paper, we apply our methodology to realistic cosmological simulations and discuss the implications of the $f(R)$ model on the large scale structure of the Universe.  This paper is organized as follows.  In \\S \\ref{sec:fr}, we review the general properties of $f(R)$ cosmology and give details on the particular model of $f(R)$ that we choose. The numerical algorithms to solve for the $f(R)$ cosmological large scale structure formation are developed in \\S \\ref{sec:numerical}, and are tested in \\S \\ref{sec:tests}. Several test cases under the HS model, highlighting the different aspects of the $f(R)$ model, are presented in \\S \\ref{sec:results}. In addition, a brief discussion of full cosmological simulations with $f(R)$ modifications are presented in \\S \\ref{sec:results}; a more complete discussion will be presented in a companion paper. Finally, conclusions and discussions of possible future extension of our formalism, are given in \\S \\ref{sec:conclusion}. ",
        "conclusions": "\\label{sec:conclusion} In this paper, we have introduced a method to construct a cosmological N-body simulation with a class of viable $f(R)$ modifications to general relativity. Our simulation code, implemented in standard C++, solves the HS model of $f(R)$ modification using particle mesh method, particularly suited because the scalar degree of freedom in the model is well represented as a scalar field (i.e., a mesh) coupled to the Newtonian potential. The non-linear field equation for $f_R$ is solved on the mesh using an efficient multigrid scheme that is shown to converge robustly to the desired solution. The code is also shown to reproduce the chameleon behavior of the HS model.  The code should allow us to gain insight into the quasi-nonliear physics of $f(R)$ modified gravity models. In general, gravitational forces are enhanced under the HS model, and we show that the resulting power spectra are also enhanced compared to the $\\Lambda$CDM power spectrum. We further show hints that the power spectrum enhancement is smaller than the linear predictions due to the chameleon mechanism that has not so far been modelled. In a companion paper \\citep{oyaizu08b}, we further study the effects of the HS model on the dark matter power spectrum and its cosmological implications. In addition, effects on the cluster mass function and weak lensing statistics are potentially interesting and should be investigated in the future.  Lastly, the work presented here can be used as a general framework for non-linear simulation of any minimally coupled scalar field model of gravity, provided that the resulting field equations are sufficiently stable under relaxation and multigrid schemes. In particular, an interesting exercise might be to apply this framework to the Dvali, Gabadadze, and Porrati (DGP) model \\citep{dvali00a}."
    },
    "0807/0807.2814_arXiv.txt": {
        "abstract": "s{HiRes has made the first statistically significant observation of the Griessen-Zatsepin-Kuzmin suppression, by fitting the ultra high energy cosmic ray spectrum observed in monocular mode by the two HiRes detectors to a broken power law model.  We find the break to be at an energy of $5.6\\times10^{19}$ eV, with a significance of $5.3\\sigma$.  The significance is determined by Poisson statistics where we expect 43.2 events above the break point, but observe only 13.  We have also looked for correlations between HiRes stereo events and active galactic nuclei.  We observe no statistically significant correlation in a number of different tests.  We have performed a search for upward going showers, which would be indicative of neutrinos propagating through the Earth and interacting just below the surface.  Observing no such events, we set a limit on the electron neutrino flux.  This limit is significantly lower in this topology than limits on the other types of neutrinos due to the LPM effect greatly increasing the available target mass for neutrino interactions.  Finally we look forward to improving these and other results in our work in the Telescope Array.}  The High Resoltuion Fly's Eye Experiment (HiRes) was operated for nine years, from June 1997 to April 2006, on Dugway Proving Grounds in Utah, USA.  It was designed to observe ultra high energy cosmic rays (UHECR's) using the fluorescence technique, with two sets of telescopes placed on dessert hills separated by 12.6 km.  Details of the fluorescence method and the analysis methods used by HiRes can be found elsewhere\\cite{Baltrusaitis-1985-NIMA-240-410,AbuZayyad-1999-ICRC-26-5-349,Boyer-2002-NIMA-482-457}. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0807/0807.0811_arXiv.txt": {
        "abstract": "We present a radio continuum study of the pulsar wind nebula (PWN) DA\\,495 (G65.7+1.2), including images of total intensity and linear polarization from 408 to 10550\\,MHz based on the Canadian Galactic Plane Survey and observations with the Effelsberg 100-m Radio Telescope. Removal of flux density contributions from a superimposed \\ion{H}{2} region and from compact extragalactic sources reveals a break in the spectrum of DA\\,495 at 1.3\\,GHz, with a spectral index ${\\alpha}={-0.45~\\pm~0.20}$ below the break and ${\\alpha}={-0.87~\\pm~0.10}$ above it (${S}_\\nu \\propto{\\nu^{\\alpha}}$).  The spectral break is more than three times lower in frequency than the lowest break detected in any other PWN. The break in the spectrum is likely the result of synchrotron cooling, and DA\\,495, at an age of $\\sim$20,000 yr, may have evolved from an object similar to the Vela~X nebula, with a similarly energetic pulsar. We find a magnetic field of $\\sim$1.3~mG inside the nebula. After correcting for the resulting high internal rotation measure, the magnetic field structure is quite simple, resembling the inner part of a dipole field projected onto the plane of the sky, although a toroidal component is likely also present. The dipole field axis, which should be parallel to the spin axis of the putative pulsar, lies at an angle of ${\\sim}50\\degr$ east of the North Celestial Pole and is pointing away from us towards the south-west. The upper limit for the radio surface brightness of any shell-type supernova remnant emission around DA\\,495 is $\\Sigma_{1\\,GHz} \\sim 5.4 \\times 10^{-23}$~Watt~m$^{-2}$~Hz$^{-1}$~sr$^{-1}$ (assuming a radio spectral index of $\\alpha = -0.5$), lower than the faintest shell-type remnant known to date. ",
        "introduction": "The Crab Nebula, the remnant of the supernova explosion of 1054, is an astrophysical laboratory of immense value where we can study the development of a young pulsar and the synchrotron nebula generated by the particles that it injects. The Crab Nebula has become the archetype of a class of objects earlier known as filled-center, plerionic, or Crab-like supernova remnants (SNRs), but now called pulsar wind nebulae (PWNe) in recognition of their energy source. PWNe are a minority in SNR catalogs, less than 10\\% in number, but if we include the composite SNRs, synchrotron nebulae centrally placed within the more usual SNR shell, their ranks swell to 15\\% of the Galactic SNR catalog (Green 2004).  PWNe are characterized by a flat radio spectral index, $\\alpha$, in the range ${0} > {\\alpha} > {-0.3}$ (where the flux density $S$ at frequency $\\nu$ varies as ${S} \\propto {{\\nu}^{\\alpha}}$) and a central concentration of emission gradually declining to the outer edge. Their appearance reflects the central energy source and their radio spectrum reflects the energy spectrum of the injected particles. The more common shell SNRs exhibit steep outer edges and steep radio spectra (generally ${\\alpha} < {-0.3}$).  We have historical records of fewer than ten supernova events (Stephenson \\& Green 2002).  By chance two of those have left PWNe, the Crab Nebula and 3C 58, with no detectable trace of accompanying shells, and one a composite remnant, G11.2$-$0.3. Thus we have three examples in which we can closely follow the evolution of young PWNe, but the relatively small number of known PWNe has provided little observational basis for our understanding of the evolution of these objects through later stages. This paper presents and analyzes new observations of DA\\,495 (G65.7+1.2), which seems likely to be a PWN of advanced age.  DA\\,495, while an outlier among PWNe, is without doubt a member of that class. Its radio appearance, first seen in the high-resolution images of \\citet{land83}, has the signature of a PWN, a central concentration of emission smoothly declining with increasing radius, with no evidence of a surrounding shell (although there is a depression in intensity near its center). The only indication that it might not be a PWN is its measured radio spectral index, ${\\alpha} \\approx {-0.5}$ \\citep{koth06b}, much steeper than is usual for this class of objects, and more closely resembling that of a shell-type SNR. Nevertheless, shell remnants have steep outer edges where kinetic energy is being converted to radio emission near the shock front, while in DA\\,495 the emissivity distribution is centrally concentrated \\citep[Fig.~3]{land83} indicating that the energy source is central. Nevertheless, on the basis of the central depression in intensity, \\citet{velu89} argue that DA\\,495 might be a shell remnant with an exceptionally thick shell formed by a reverse shock interacting with an unusual amount of slowly moving ejecta.  Recent analysis of X-ray data for DA\\,495, obtained from the ROSAT and ASCA archives, dramatically confirms the PWN interpretation of this object by revealing a compact central object surrounded by an extended non-thermal X-ray region \\citep{arzo04}. Analysis of Chandra data \\citep{arzo08} shows a non-thermal X-ray nebula of extent $\\sim$30\\arcsec\\ surrounding the central compact object. The compact source is presumably the pulsar (required to power the synchrotron nebula), and we shall refer to it as such (even though pulsed emission has not been detected). In this paper we present new radio images of DA\\,495 which include linear polarization information. Analysis of high-resolution radio images allows us to remove superimposed sources. After this the integrated spectrum of the source exhibits a spectral break characteristic of a synchrotron nebula, and the observed parameters allow an estimate of the strength of the magnetic field within the nebula. Polarization images reveal the structure of the field within this PWN. ",
        "conclusions": "We have imaged the radio emission from the pulsar wind nebula DA\\,495 in total intensity and polarized emission across a range of frequencies from 408 to 10550~MHz.  After correcting for the flux contribution of extraneous superimposed sources we have demonstrated that there is a break in the spectrum at about 1.3\\,GHz which we can convincingly attribute to synchrotron cooling. The spectral break at such a low frequency, the lowest known for any PWN, can only be produced if the pulsar originally was very energetic, with a high energy loss rate.  We consider that DA\\,495 is an aging pulsar wind nebula which has a pulsar that in its earlier stages resembled the Crab pulsar or the pulsar in Vela~X. The magnetic field inside the nebula is remarkably well-organized and has the topology of a dipole field centered on the pulsar; the field probably also has a toroidal component, wrapped around the dipole axis. The rotation measure is high, as would be expected in a nebula where the magnetic field has the high value of $\\sim$1.3~mG, and the nebula becomes Faraday thick at a very high frequency.  We have set a very low upper limit for the surface brightness for any shell of SNR emission surrounding DA\\,495. The absence of a shell implies that the SN explosion occurred in a low-density environment, and this also explains how the magnetic field structure has remained so regular to an age of $\\sim$20,000 yr."
    },
    "0807/0807.2025_arXiv.txt": {
        "abstract": "{} {Flares on dM stars contain plasmas at very different temperatures and thus affect a wide wavelength range in the electromagnetic spectrum. While the coronal properties of flares are studied best in X-rays, the chromosphere of the star is observed best in the optical and ultraviolet ranges. Therefore, multiwavelength observations are essential to study flare properties throughout the atmosphere of a star.  } {We analysed simultaneous observations with UVES/VLT and \\emph{XMM-Newton} of the active M5.5 dwarf CN~Leo (Gl 406) exhibiting a major flare. The optical data cover the wavelength range from 3000 to 10\\,000 \\AA.} {From our optical data, we find an enormous wealth of chromospheric emission lines occurring throughout the spectrum. We identify a total of 1143 emission lines, out of which 154 are located in the red arm, increasing the number of observed emission lines in this red wavelength range by about a factor of 10. Here we present an emission line list and a spectral atlas. We also find line asymmetries for \\ion{H}{i}, \\ion{He}{i}, and \\ion{Ca}{ii} lines. For the last, this is the first observation of asymmetries due to a stellar flare. During the flare onset, there is additional flux found in the blue wing, while in the decay phase, additional flux is found in the red wing. We interpret both features as caused by mass motions. In addition to the lines, the flare manifests itself in the enhancement of the continuum throughout the whole spectrum, inverting the normal slope for the net flare spectrum. } {} ",
        "introduction": "Flaring is a commonly observed phenomenon on late-type stars. During a flare event, large amounts of energy are released probably from magnetic field reconfigurations and emitted over a wide range of the electromagnetic spectrum. Although white light flares are relatively rare for the Sun, they are quite common among M dwarf stars, because of their lower photospheric background emission in the optical. Chromospheric emission lines react sensitively to flares in amplitude, line width, line shape, and wavelength shifts. In addition, optical flares often have X-ray counterparts, although there is no one-to-one relationship between X-ray and optical flares. During large flares, dM stars can show an increase in X-ray luminosity by factors up to 100 \\citep[e.\\,g.][]{Guedel_ProxCen_1,Guedel_ProxCen_2}, and even greater magnitude increases have occasionally been reported in the optical \\citep[e.\\,g.][]{Eason, Hawley_Pettersen}.  The thermal properties of the coronal flare plasma can be diagnosed from the emitted soft X-ray emission, while gyrosynchrotron radio and microwave emission typically traces nonthermal particles in the corona. Additionally, emission lines in the ultraviolet mainly diagnose the transition region. Different wavelength ranges thus provide complementary information, so that multiwavelength observations of flares allow to study the contemporaneous reaction of different layers of the stellar atmosphere. A prototypical multiwavelength campaign was carried out by \\citet{deJager} for UV~Cet, which was simultaneously observed in X-ray, optical, and radio. As a more recent example, \\citet{UVCet_Xray_UV} observed a correlation between ultraviolet and X-ray flux for stellar flares. Another recent multiwavelength campaign was aimed at the flare star EV Lac and provided data in the radio, the optical, the ultraviolet, and in X-rays \\citep{Osten_EV_Lac_flare, Osten_EV_Lac_quiescence}. These data cover a huge radio flare and a number of optical and X-ray flares. \\citet{Osten_EV_Lac_flare} note that, for a flare in one wavelength region, counterparts in other wavelength regions are often missing. An example of a giant flare on AD~Leo observed in the ultraviolet and optical is described by \\citet{Hawley_Pettersen}. AD~Leo was again observed in multiwavelength by \\citet{Hawley_Allred}, where a number of medium-sized flares are covered and discussed. \\citet{Smith} compare radio and X-ray flares on five active M dwarfs and provide possible mechanisms for uncorrelated flares. Multiwavelength observations also show that quiescent X-ray, optical, and radio emission strengths are no longer correlated for late M and early L dwarfs \\citep{Berger1,Berger3,Berger2}.  Our target star \\object{CN Leo} (Gliese~406) is a well-known nearby flare star with a spectral type of M5.5 \\citep{Reid} or M6.0 \\citep{Kirkpatrick91}. \\citet{Reiners} find a mean magnetic field of  $Bf \\sim 2.2$ kG, exhibiting night-to-night and, with less certainty, intra-night flux variations.  CN~Leo has previously been the target of a multiwavelength campaign in December 2005, covering the optical (VLT/UVES) and X-ray regimes (\\emph{XMM-Newton}; \\citet{fuhrmeister_liefke}). During these observations, CN~Leo once again proved to be quite active, showing two medium-sized and a number of smaller flares. Unfortunately, both larger flares occurred at the end of the optical observations and were therefore not properly covered. Now, CN~Leo was again observed in the framework of a similar multiwavelength campaign, during which a spectacular flare occurred. This paper is the first in a series analysing the flare data, and it is structured as follows: In Sect. 2 we describe  our observations obtained with UVES and \\emph{XMM-Newton}. In Sect. 3 we give an overview over the timing behaviour during the flare. The wide variety of chromospheric emission lines is reported in Sect. 4, while asymmetries found in the line shape of some emission lines are described in Sect. 5. A discussion and our conclusions are presented in Sect. 6. ",
        "conclusions": "We presented  the optical part of a multiwavelength campaign that caught CN~Leo in a spectacular flare. There are only few observations of such giant flares, but even these few examples show very different timescales and lightcurve characteristics. While the CN~Leo flare lasted about 50 minutes, a 5-magnitude flare in the U-Band on UV~Cet lasted only 12 minutes \\citep{Eason}, while the large flare on AD~Leo from 1985 lasted more than 4 hours \\citep{Hawley_Pettersen}. Compared to somewhat smaller (but nevertheless very large) flares, these can be of the same duration or even longer than the CN~Leo flare. For example, a 2.1 magnitude flare in the U-band was observed on AD~Leo to last about 50 minutes \\citep{AD-Leo1984}, while a large flare on AT Mic lasted  about 1.5 hours \\citep{AT-Mic}.  Most of the large flares show a double (or triple) peaked lightcurve in the U-band and also for Balmer emission lines like the CN~Leo flare \\citep{Hawley_Pettersen, AD-Leo1984, AT-Mic}. The UV~Cet flare had an even more complicated substructure with at least 4 peaks in the major outburst \\citep{Eason}. This multiple-peak flare pattern is usually interpreted in terms of different reconnection events, which may trigger each other in succession. The CN~Leo flare shows two major peaks that can be easily noticed in the exposuremeter data  in Fig.~\\ref{lightcurve}, with a minor peak between the two. The time between the two major peaks is about 200 seconds. For the 2.1 mag flare on AD~Leo, the two major peaks are separated by about 9 minutes \\citep{AD-Leo1984}. For all other mentioned flares (besides UV~Cet), the separation is even longer.  The flux peaks of chromospheric lines are typically delayed with respect to the U-band flux peak, with the \\ion{Ca}{ii} H and K lines even more delayed than the Balmer lines \\citep[e.\\,g.][]{AT-Mic, Hawley_Pettersen}. The time resolution of our spectra is unfortunately too low to observe any delay in the peak, but as is apparent in Fig.~\\ref{line_flux}, the flux of the \\ion{Ca}{ii} triplet line decays more slowly than the other lines. This behaviour is compatible with what is normally observed for the \\ion{Ca}{ii} H and K lines, which decay more gradually than Balmer and Helium lines \\citep[e.\\,g.][]{AD-Leo1984}, consistent with an interpretation of a faster evolution of high-temperature lines in the chromosphere \\citep{Houdebine3}.  The optical spectral data allowed us to trace the chromospheric reaction to the flare (with one \\ion{He}{ii} transition line covered). We found a wealth of chromospheric emission lines that has not been observed before, especially in the red part of the spectrum; also, \\ion{V}{ii} and \\ion{Sc}{ii} lines show up in the blue arm. Especially in the flare onset spectrum no.~2, the higher Balmer lines develop very broad damping wings that can only  be fitted involving a Lorentzian profile. With respect to line shapes and profiles, we  find asymmetries varying from blue excess flux during flare onset to red excess flux during the decay phase. In the red excess for H$\\alpha$, not only a general decay is found, but also variations from spectrum to spectrum, providing strong evidence for a scenario involving multiple flare kernels, each triggering chromospheric downward condensations. We ascribe the blue excess flux during flare  The optical quasi-continuum is also quite massively affected by the flare. The blue spectrum is clearly blackbody-dominated during flare onset. The same is true for the red spectrum with the quiescent spectrum removed. The derived temperatures for the gas producing the blackbody radiation in the first few seconds of the flare onset exceed of 20\\,000 K and then rapidly cool.  This set of excellent optical data, which characterises the photosphere and chromosphere of CN~Leo during a giant flare is suited well to classical semi-empirical chromospheric flare modelling with a stellar atmosphere code. Modelling using a multitude of emission lines will be presented in a forthcoming paper in this series \\citep{paper3}."
    },
    "0807/0807.1817_arXiv.txt": {
        "abstract": "The presence of exponential bulges and anti-truncated disks has been noticed in many lenticular galaxies. In fact, it could be expected because the very formation of S0 galaxies includes various processes of secular evolution. We discuss how to distinguish between a pseudobulge and an anti-truncated disk, and also what particular mechanisms may be responsible for the formation of anti-truncated disks. Some bright examples of lenticular galaxies with the multi-tiers exponential stellar structures are presented, among them -- two central group giant S0s seen face-on and perfectly axisymmetric. \\\\ Keywords: galaxies: elliptical and lenticular, cD; galaxies: bulges; galaxies: disks; galaxies: evolution ",
        "introduction": "Bulges have been traditionally thought to have de Vaucouleurs' brightness profiles just as elliptical galaxies (e.g. \\cite[Freeman 1970]{freeman}). However John Kormendy (\\cite[Kormendy 1982a]{korm82a}, \\cite[Kormendy 1982b]{korm82b}, \\cite[Kormendy 1993]{korm93}) has proved that there exists a type of bulges named `pseudobulges' which resemble disks from the dynamical point of view: they are rather cold and demonstrate fast rotation. Despite their dynamical properties, they are bulges from the geometrical point of view: they are `fat' and have rather large scaleheights. Pseudobulges are thought to be products of secular evolution; but if it is so, then dynamical simulations (e.g. \\cite[Pfenniger \\& Friedli 1991]{pf_fri}) predict that they must have exponential brightness profiles. And indeed, small bulges of late-type spirals which could be easily made by secular evolution, practically all have exponential brightness profiles (e.g. \\cite[Andredakis et al. 1995]{phot_apb}, \\cite[Graham 2001]{graham}). But it is not only small bulges in late-type spirals that have exponential brightness profiles. Lenticular galaxies whose bulges are not usually small at all demonstrate typically the mean Sersic coefficient lesser than Sa galaxies, and often their $n  <2$: the Sersic parameter $n$ peaks at the Sa--Sb morphological type and falls further toward S0s (\\cite[Graham 2001]{graham}, \\cite[Laurikainen et al. 2005]{lauri05}, \\cite[Laurikainen et al. 2007]{lauri07}). In fact, it is quite natural because the very event of a S0 galaxy transformation from a spiral must include various processes of secular evolution resulting in matter re-distribution over the radius and bulge reshaping.  But galactic disks outside the bulges can also consist of several exponential segments. Now it becomes clear that so called anti-truncated disks consisting of two exponential segments with the outer scalelengths larger than the inner ones, may be the dominant type of galactic disks among some types of galaxies (\\cite[Erwin et al. 2008a]{erwinstat}). We reported finding a few large nearby spiral galaxies with such two-tiers disks during the last 10 years. A significant population of such disks has been found since the brightness profiles begin to reach 27th and 28th magnitudes from one square arcsecond (\\cite[Pohlen \\& Trujillo 2006]{pohtru}). Does it mean that the outer segments of anti-truncated disks represent always low surface-brightness (LSB) disks? The statistics by Erwin et al. (\\cite[Erwin et al. 2008b]{erwinbar}) for barred galaxies says so. Among our findings in unbarred spiral galaxies, only one galaxy, NGC~5533, has a LSB outer disk (\\cite[Sil'chenko et al. 1998]{n5533}). Other galaxies (NGC~7217, \\cite[Sil'chenko \\& Afanasiev 2000]{n7217}; NGC~615, \\cite[Sil'chenko et al. 2001]{n615}; NGC~4138, \\cite[Afanasiev \\& Sil'chenko 2002]{we2002}; NGC~7742, \\cite[Sil'chenko \\& Moiseev 2006]{n7742}) have on the contrary quite normal outer disks, if to compare with the Freeman's (\\cite[Freemen 1970]{freeman}) reference value of the central $B$ surface brightness of $21.7$; and simultaneously they have compact bright inner disks. Can we distinguish the compact inner disk from a pseudobulge? Yes, by using multi-variant approach. For example, in NGC~7742 all spirals and current star formation are confined to the inner exponential disk, and the outer one is smooth and featureless (\\cite[Sil'chenko \\& Moiseev 2006]{n7742}), so we can conclude that the inner exponential stellar component is dynamically cold and cannot be a bulge. But the key property is a visible geometry of the stellar component with the exponential profile. To be a disk, it must be thin. For disks inclined to the line of sight, a good check is isophotal analysis. The outer disk is always assumed to be thin: then its isophote axis ratio characterizes the cosine of the inclination. If the inner component has the same visible axis ratio as the outer one and if two disks are coplanar, the inner structure cannot be a spheroid, it must be a disk. For face-on galaxies, this approach does not work: their isophotes are always round independently of the scaleheights. But for face-on galaxies we can use kinematical data and estimate their thickness by measuring vertical stellar velocity dispersions.  We have now started a program of studying systematically multi-tiers (anti-truncated) exponential structures in early-type, presumably lenticular, galaxies. The study will include photometric as well as spectral observations. Some first results are presented in this talk. ",
        "conclusions": "among barred galaxies, the anti-truncated disks contribute one third of all, among unbarred -- more than 50\\%. This inconsistency is a complete puzzle yet."
    },
    "0807/0807.1210_arXiv.txt": {
        "abstract": "The propagation of photons, electrons and positrons at ultra-high energies above $\\sim10^{19}\\,$eV can be changed considerably if the dispersion relations of these particles are modified by terms suppressed by powers of the Planck scale. We recently pointed out that the current non-observation of photons in the ultra-high energy cosmic ray flux at such energies can put strong constraints on such modified dispersion relations. In the present work we generalize these constraints to all three Lorentz invariance breaking parameters that can occur in the dispersion relations for photons, electrons and positrons at first and second order suppression with the Planck scale. We also show how the excluded regions in these three-dimensional parameter ranges would be extended if ultra-high energy photons were detected in the future. ",
        "introduction": "Many modern extensions of the Standard Model of particle physics, including string theory and various other approaches aiming at unification of quantum mechanics with general relativity, suggest that the Lorentz symmetry may be broken or modified at energy and length scales approaching the Planck scale. While such effects necessarily have to be tiny at energies up to the electroweak scale in order to satisfy laboratory constraints~\\cite{Kostelecky:2008ts}, they can be magnified at higher energies as they can occur in astrophysics. The observation of standard radiation processes as they are predicted in the absence of Lorentz invariance violation then often allows to derive strong constraints on Lorentz invariance violating (LIV) effects~\\cite{Jacobson:2005bg,Bietenholz:2008ni}.  LIV effects could, for example, change the propagation and thus spectra and composition of the highest energy particles observed in Nature \\cite{Coleman:1998ti,Aloisio:2000cm}. In particular, photons are produced as secondaries of cosmic rays but are quickly reabsorbed in the cosmic microwave background (CMB) if Lorentz symmetry holds. If photon interactions with the microwave background were inhibited by new physics near the Planck scale $M_{\\rm Pl}$, a significant fraction of ultra-high energy (UHE) cosmic rays above $10^{19}$ eV would be photons, in contradiction to current experimental upper limits~\\cite{Risse:2007sd}. Based on this argument, in Ref.~\\cite{Galaverni:2007tq} we put strong constraints on the coefficients of LIV terms of the form $(E/M_{\\rm Pl})^n$, $n=1,2$, in the dispersion relation for photons of energy $E$, assuming for simplicity, that the corresponding parameters for electrons and positrons are significantly smaller. These latter terms can, however, play an equally important role because they can influence the kinematics of pair production, $\\gamma\\,\\gamma_{\\rm b}\\rightarrow e^-\\,e^+$, the most important process for absorption of high energy photons in the background of low energy photons $\\gamma_{\\rm b}$, such as the CMB. In the present work we, therefore, extend those studies to all three LIV terms that are allowed at first and second order suppression with the Planck scale in the dispersion relations of photons, electrons and positrons from general principles of effective field theory, without making any assumptions on the relative size of these parameters. We obtain three-dimensional exclusion plots for the coefficients of these terms, following from the non-observation of photons above $\\simeq10^{19}\\,$eV. We also show how the excluded regions in these three-dimensional parameter ranges would be extended if ultra-high energy photons were detected in the future. Throughout this paper we assume that primary cosmic rays are dominated by protons.  For the coefficients of LIV terms linearly suppressed with the Planck scale values larger than $\\sim 10^{-5}$ for electrons and positrons are currently ruled out by the observation of synchrotron radiation from the Crab Nebula \\cite{Maccione:2007yc}. We find that the observation of photons above $\\simeq10^{19}\\,$eV would rule out that any one of the three LIV coefficients for electrons, positrons and photons has absolute value $\\gtrsim 10^{-14}$. This is in agreement with Ref.~\\cite{Maccione:2008iw} which considers a two dimensional subset of the general three-dimensional parameter range.  In contrast, we will find that for LIV terms quadratically suppressed with the Planck scale, arbitrarily large values of one of the LIV terms for electrons and positrons can not be ruled out by UHE photon observations if the coefficients of the two other LIV terms have absolute value $\\lesssim10^{-6}$.  The remainder of this paper is structured as follows: In Sect.~II we introduce the modified dispersion relations, in Sect.~III we determine the thresholds for pair production and photon decay, and in Sect.~IV we derive the resulting constraints on the LIV parameters. Finally, we summarize and conclude in Sect.~V. We use natural units, $c=1$, throughout. ",
        "conclusions": "In the present paper we derived general constraints on LIV dispersion relations in the QED sector from the propagation of UHE photons. The kinematics of pair production and photon decay is discussed in terms of all three LIV parameters $\\xi_n$, $\\eta_n^+$ and $\\eta_n^-$ entering in the dispersion relations of photons, electrons and positrons, both for linear ($n=1$) and for second order ($n=2$) suppression with the Planck scale.  The upper limits on the flux of UHE photons require that combinations of the LIV parameters for which the UHE photons are stable and cannot pair produce on low-energy (e.g. CMB) photons are excluded. Similarly, the detection of photons of $\\sim10^{19}\\,$eV would exclude those combinations of the LIV parameters for which both photon polarizations are unstable.  For terms in the dispersion relation linearly suppressed by the Planck scale the resulting constraints are very strong: Using the non-detection of an UHE photon flux and anticipating the detection of a $\\sim10^{19}\\,$eV photon, it will be possible to exclude all LIV parameters with absolute value $\\gtrsim 10^{-14}$.  In contrast, in the $n=2$ case, the LIV parameter region excluded using these arguments based on UHE photon propagation does not completely surround the origin. If UHE photons are eventually detected, the maximum absolute value allowed for LIV parameters in the photon sector will be $\\sim 10^{-6}$, whereas currently only values smaller than $\\sim -10^{-6}$ are ruled out. However, even if UHE photons were detected, constraints on the electron parameters can be evaded for some particular combinations: If, for example, the modulus of one of the two parameters $\\eta_2^+$ or $\\eta_2^-$ is smaller than $\\sim 10^{-6}$, then the modulus of the other parameter is constrained neither by the upper limit on the UHE photon flux nor by a putative future detection of a $\\sim10^{19}\\,$eV photon. However, the case where the moduli of both $\\eta_2^+$ and $\\eta_2^-$ are $\\gtrsim 10^{-6}$ can still be excluded once UHE photons are detected."
    },
    "0807/0807.4573_arXiv.txt": {
        "abstract": "We present the serendipitous discovery of a young stellar cluster in the Galactic disk at $l$=12$^\\circ$. Using Keck/NIRSPEC, we obtained high-- and low--resolution spectroscopy of several stars in the cluster, and we  identified one red supergiant and two blue supergiants. The radial velocity of the red supergiant provides a kinematic cluster  distance of 4.7$\\pm$0.4 kpc, implying luminosities of the stars consistent with their spectral types. Together with the known Wolf--Rayet star located 2\\farcm4 from the cluster center, the presence of the red supergiant and the blue supergiants suggests a cluster age of $6-8$ Myr, and an initial mass of $\\ga 2000$ \\Msun. Several stars in the cluster are coincident with X--ray sources, including the blue supergiants and the Wolf--Rayet star. This is indicative of a high binary fraction, and is reminiscent of the massive young cluster Westerlund 1. The cluster is coincident with two supernova remnants, SNR G12.72$-$0.0 and G12.82$-$0.02, and the highly magnetized pulsar associated with the TeV $\\gamma$--ray source HESS J1813$-$178.  The mixture of spectral types suggests that the progenitors of these objects had initial masses of $20-30$ \\Msun. ",
        "introduction": "Young stellar clusters are important tools for investigating the structure, the chemical enrichment, and the current star formation of a galaxy.  Furthermore, they are natural laboratories to study the evolution of massive stars.  Massive stars explode as supernovae creating neutron stars and black-holes, and are believed to be sources of $\\gamma$--ray bursts, which are the most energetic explosions in the Universe. Because of their short lifespans, massive stars are rare and predominantly observed in young (few Myrs) massive ($\\sim 10^{4}$ \\Msun) clusters.  Only a handful of massive stellar clusters are known in the Galaxy: important examples include Arches, Quintuplet, Westerlund 1, and Westerlund 2 \\citep{figer08}.  As these clusters lie in the plane, their detection is hampered both by interstellar extinction and by the relatively low numbers of member stars. Indeed, the census is severely incompletely, as demonstrated by the recent discovery of two other massive clusters \\citep[RSGC1 and RSGC2,][]{figer03,davies07}.  We have serendipitously discovered a new stellar cluster at RA= 18\\fha 13\\fma 24\\fs31 and DEC=$-$17\\fdga 53\\farcma 30\\farcs83 (J2000). Previously undetected, this cluster clearly appears as a stellar over--density at infrared wavelengths (Fig.\\ \\ref{fig1}), and is also visible at optical wavelength. It is located adjacent to the star forming region W33 in an extremely rich and complex region. Two supernovae remnants (SNRs), G12.82$-$0.02 and G12.72$-$0.0 \\citep{brogan06}; a $\\gamma$--ray source, HESS J1813$-$178 \\citep{hess06}; and numerous X--ray sources \\citep{funk07,helfand07} have been detected in its surrounding area.  \\begin{figure}[!] \\begin{center} \\plotone{f1.eps} \\caption{\\label{fig1} \\Ks--band image of the cluster from 2MASS, with overlayed contours of 90~cm data \\citep{white05}. Contour levels, in mJy beam$^{-1}$, are: 40, 60, 80, 100, 120, 140, 160, 200. The beam size is 24\\arcsec $\\times$ 18\\arcsec, FWHM, and the position angle of the major axis is along the North--South direction.  Stellar identification numbers are from Table~1.} \\end{center} \\end{figure}   In this report, we present near--infrared spectra of four cluster members, of which three are identified as early B--type stars, and one as a red supergiant (RSG).  Also, we analyze existing photometric infrared data and X--ray data, and find that several bright cluster members are associated with an X--ray source, one of which is a known Wolf--Rayet (WR) star \\citep{hadfield07}.  From these results, we estimate the age and distance of the cluster, which appears associated with the star forming region W33 and the high--energy source HESS J1813$-$178. ",
        "conclusions": ""
    },
    "0807/0807.3954_arXiv.txt": {
        "abstract": "Inverse Compton (IC) scattering is one of  two viable mechanisms that can produce the prompt non-thermal soft gamma-ray emission in Gamma-Ray Bursts. IC requires low energy seed photons and a population of relativistic electrons that upscatter them. The same electrons will  upscatter the gamma-ray photons to even higher energies in the TeV range. Using the current upper limits on the prompt optical emission we show  that under general conservative assumption the IC mechanism suffers from an ``energy crisis\". Namely, IC will over-produce a very high energy component that would carry much more energy than the observed prompt gamma-rays, or alternatively it will require a low energy seed that is more energetic  than the prompt $\\gamma$-rays. Our analysis is general and it makes no assumptions on the specific mechanism that produces the relativistic electrons population. ",
        "introduction": "The mechanism that produces the prompt gamma-ray emission in Gamma Ray Burst (GRBs) is still uncertain. The non-thermal character of the spectrum points out towards Inverse Compton (IC) and Synchrotron as the two natural candidates. The latter become, somehow, the ``standard\" process but the former remained always a serious alternative \\citep[][and others]{s94,sd95,snp96,sp97,w97,glcr00,sp04,k07}. The observations of numerous bursts with low energy spectral slopes that are inconsistent with synchrotron  \\citep{c97,p98,glcr00,p02} provided additional motivation to consider IC.  Recently  \\cite{km08} have shown further inconsistency with the overall synchrotron model and suggested that  Synchrotron Self-Compton (SSC) can resolve some of these problems.  The recent observations of a naked eye optical flash from GRB080319b \\citep{r08,b08,d08} that coincided in time with the prompt $\\gamma-$ray emission provided further motivation to consider IC as the source of the prompt $\\gamma$-rays. Among the different models that appeared so far \\citep{kp08,fp08,zps08,ywd08}, several favor models in which the prompt $\\gamma$-ray emission is IC of the optical flash and there have been suggestions that this is generic to many GRBs.  Motivated by these ideas we examine, here,  the possibility that IC is the source of the prompt soft $\\gamma$-ray emission in GRBs. This requires  a soft component at the IR-UV range that serves as the seed for the IC process. The flux of these seed photons is constrained by observations (or upper limits) of the prompt optical emission. GRB 990123 \\citep{ab99} and GRB 080319B \\citep{r08} are rare exceptions with very strong optical emission,  $\\sim 9$ and $\\sim 5.3$ mag respectively. However most bursts are much dimer optically with observations or upper limits around 14 mag \\citep{ya07}. This should be compared with fluxes of $mJy$ in soft gamma rays  for a modest burst.  What is important, in this work is the flux ratio $F_\\g /F_{opt}$ which is typically larger than 0.1 during the peak soft gamma emission \\citep{ya07}.   The basic problem of the IC  model can be explained simply. If the low energy seed  emission is in the optical, while the observed soft $\\gamma$-ray spectrum is the first IC component, then second IC scatterings would create a TeV component. Upper limits or observations of the prompt optical signal show that the $Y$ parameter, i.e. the ratio between the energy in the first IC component to that in the low energy seed photons is very large, typically greater than thousands. Theory would then show that the second IC component in the TeV range would carry an even larger amount of energy, again by a factor of $Y \\gg 1$, producing an ``energy crisis\" for this model, and possibly violating upper limits from EGRET (Energetic Gamma-Ray Experiment Telescope) \\footnote{Deeper upper limits on a wider energy range, may soon come up from Fermi, making our argument stronger.} \\citep{gs05,ans08}. This problem is generic and it does not depend on the specific details of the overall model.   The above analysis is oversimplified and two factors may alleviate the energy catastrophe. First, the frequency of the  seed photons may differ from that where upper limits exist, allowing larger seed flux and reducing the lower limits on $Y$. Second, the Klein-Nishina (KN) suppression, which does not affect the first scattering, may affect the second, resulting in a lower $Y$ parameter for the second scattering than the first one. In this article, we explore the parameter space to see weather there exist a regime where a combination of these two factors allows for less energy in the second IC component (Typically in the TeV range) than in the $\\gamma$-rays. We find that possible solutions are limited to a very small region in the parameters space in which the seed photons are in the IR, the bulk Lorentz factor is very low $(\\le 200$) and the electrons' Lorentz factor is very large ($\\ge 2000)$. However, this solution implies a healthy emission in the IR, while self absorption limits it. Therefore, when taking self-absorption into account, this solution is  ruled out as well. A second possible solution exists if the seed photons are in the UV. This solution requires a very low electrons' Lorentz factor $\\le 100$, and a seed photon flux that carries comparable energy to the observed prompt $\\gamma$-rays. Furthermore, prompt X-ray observations limit the high energy tail of the UV component and practically rule out this model.  \\begin{figure} \\centerline{\\includegraphics[width=0.9\\textwidth]{FigSchematic4.pdf}} \\caption{A schematic description of the IC process. Low energy photons at the IR (marked in dotted lines), optical or UV (marked in solid thin lines) are IC scattered to produce the observed soft gamma ray emission (marked in bold lines). A second IC scattering brings the soft gamma photons to the TeV region. If the initial seed photons are softer the higher energy component is harder. If the initial seed is in the IR then the second IC process might be in the KN regime, in which case this component is suppressed (dashed-dotted line). The seed low energy emission is constraint by upper limits on the optical prompt observations (bold solid arrow).} \\label{fig:schematic} \\end{figure}     We take the Lorentz factor of the electrons and the bulk Lorentz factor as free parameters and we estimate what is the second IC fluence (at TeV or multi GeV) given the observed prompt gamma-ray flux and the limits on the prompt optical band.  Most of our analysis is insensitive to the size of the source, which appears only in the final section when we estimate the self absorption flux. In our numerical examples we use very conservative parameters. For example we use R magnitude of 11.2 as an upper limit on the optical flux, while many limits are much stronger and the $\\gamma$-ray flux we take,  $10^{-26} {\\rm erg\\,cm^{-2}\\,s^{-1}\\,Hz^{-1}}$, is quite modest. Similarly we use conservative rather rather than ``canonical\" values for the spectral slopes. ",
        "conclusions": "For a typical GRB, IC has to  amplify the total energy  of a low energy seed photon flux by a factor of $\\approx 1000$ to produce the observed prompt gamma-ray flux. The same relativistic electrons will, however, continue and upscatter the gamma-ray flux to very high energies in the TeV range. In many cases this second generation IC will be in the Klein-Nishina regime (that is the photon's energy will be larger than the electrons rest mass, in the electron's rest frame). This will suppress somewhat the efficiency of conversion of gamma-rays to very high energy gamma-rays, however it won't stop it altogether.   Our analysis focused on the case that the low energy seed photons are produced within the moving region that includes the  IC scattering relativistic electrons. Such will be the case, for example, in Synchrotron self-Compton. Related considerations, that will be published elsewhere, apply when the seed photons are external and constrain IC processes in this case as well. The analysis is also limited to the important implicit assumption that the emitting region is homogenous. It is possible that very strong inhomogeneities could change this picture.   We have shown that, under quite general conservative assumptions, if IC produces the prompt MeV photons, then a second scattering will over produce a very high (GeV-TeV) prompt component that will carry significantly more energy than the prompt gamma-rays themselves. On the  theoretical front such a component will cause an ``energy crisis\" for most current progenitor models. From an observational point  of view, this component is possibly already ruled out by EGRET upper limits\\citep{gs05,ans08}. Fermi should very soon put much stronger limit to (or verify) this possibility. For example, a burst with isotropic energy $E_{\\gamma,iso}=10^{53}$erg, locating at $z=1$, would produce $\\sim 100 Y_H (E_H/10{\\rm GeV})$ photons detected by Fermi.  One may not over produce a high energy component if the seed photons are in the UV regime. However, in this case, the needed seed photon energy should be equal or larger than the observed prompt gamma-ray energy.  Downwards extrapolation of the X-ray observations put strong limits on this solution and probably rule it out as well.  Moreover this UV solution requires pair loading to be efficient.  We thank R.  Mochkovitch, E. Nakar, J. Poutanen, P. Kumar and X. F. Wu for helpful discussions.  The research was partially supported by the ISF center for excellence in High Energy Astrophysics, by an Israel-France collaboration grant, by an IRG grant, by an ATP grant and by the Schwartzman chair."
    },
    "0807/0807.1744.txt": {
        "abstract": "We present the results of NICMOS imaging of a sample of 19  high mass passively evolving galaxies with $1.2<z<2$, taken primarily from the Gemini Deep Deep Survey (GDDS). Around $80\\%$ of galaxies in our GDDS sample have spectra dominated by stars with ages $\\gtrsim1$~Gyr.  Our rest-frame $R$-band images show that most of these objects have compact regular morphologies which follow the classical R$^{1\\slash 4}$~law.  These galaxies scatter along a tight sequence in the size vs. surface brightness parameter space which defines the Kormendy relation. Around one-third ($3\\slash 10$) of the massive red objects in the GDDS sample are extraordinarily compact, with effective radii under one kiloparsec. Our NICMOS observations allow the detection of such systems more robustly than is possible with optical (rest-frame UV) data, and while similar systems have been seen at $z\\ga2$, this is the first time such systems have been detected in a rest-frame optical survey at $1.2<z<2$. We refer to these compact galaxies as `red nuggets', and note that similarly compact massive galaxies are completely absent in the nearby Universe. We introduce a new `stellar mass Kormendy relation' (stellar mass density vs size) which we use to single out the effects of size evolution from those of luminosity and color evolution in stellar populations. The $1< z < 2$ passive galaxies have mass densities that are an order of magnitude larger then early type galaxies today and are comparable to the compact distant red galaxies at $2 < z < 3$. We briefly consider mechanisms for size evolution in contemporary models focusing on equal-mass mergers and adiabatic expansion driven by stellar mass loss. Neither of these mechanisms appears able to transform the high-redshift Kormendy relation into its local counterpart, leaving the origin and fate of these compact `red nuggets' unresolved.  \\bigskip ",
        "introduction": "The formation mechanism of elliptical galaxies has long been controversial and remains a key test of more general galaxy formation models. The original `nature' (\\citet{elbs62} monolithic collapse) vs. `nurture' (formation through mergers, \\citep[e.g.,][]{schweizer87,searle78,tt72} debate is still with us, but is now set in a  $\\Lambda$CDM cosmological context which attempts to connect the stellar component of galaxies to an underlying evolutionary picture for the clustering of dark matter halos. Testing this model requires studying the evolution of galaxies over a large redshift range.  A wide range of selection techniques have been effective in selecting galaxies in various redshift ranges on the basis of their current star formation rates (e.g. Lyman break galaxies, sub-mm sources etc), or from the spectral signatures of passively evolving old stellar populations (e.g., extremely red objects (EROs) and other color selections). The most massive local elliptical galaxies have the oldest stellar populations \\citep{ght84}, so identifying the progenitors of local early-type galaxies within the high-redshift galaxy population is of particular interest. There is a consensus that the mass density in the red sequence is evolving strongly in the $1<z<2$ range (\\citealp[GDDS Paper VIII,][]{abr07}; \\citealp[GDDS Paper III,][]{gla04}; \\citealp{fon04,ru03}), a process that continues at redshifts below unity as well \\citep{fab07,bell04}, although the magnitude of the evolution is uncertain \\citep{brown07,chen03}. Massive morphologically-confirmed elliptical galaxies have been found up to $z=2$ (\\citealp[GDDS Paper IV,][] {mcc04}; \\citealp{cim04}) with spectra consistent with formation epochs up to $z>5$. These observations were in in direct contradiction with early  $\\Lambda$CDM models where stellar mass assembly traced the build up of cold dark matter haloes, although additional feedback mechanisms on the baryons have more recently been able to better account for this \\citep[e.g.,][]{KJS06}. A complication recently added to this picture is the observation that the space density of ellipticals is found to evolve strongly over $1<z<2$ \\citepalias{abr07} even while their stellar populations evolve weakly, suggesting that one must be careful to decouple morphological evolution from evolution of the underlying stellar populations. This is seen at higher redshifts also, where the paucity of passively evolving galaxies at $z>2$ in deep $J-K$ and 3.5~$\\mu$m selected samples \\citep{kr06,la05,cim02} shows that the $assembly$ epoch for the red sequence may be decoupled from the epoch of the earliest star formation.  Studies of star formation history and morphology can only go so far in unraveling the puzzle of galaxy formation; dynamical and chemical probes are needed to connect progenitors to descendants.  Clustering signatures offer one dynamical approach to connecting progenitors to descendants and the strong clustering of the passive red galaxies \\citep{dad05a,dad04,brown03,mcc01} strongly suggest that they are linked to today's massive ellipitical galaxies.  Theoretical attempts to explain these observations have resulted in greatly improved  $\\Lambda$CDM models which decouple mass assembly from this \u00d4stellar population downsizing\u00d5. An example is the semi-analytic model of \\citet{dlb07}. Here the small ellipticals and their stars form early by disc mergers. Massive ellipticals can then grow bigger and more numerous at late times through dissipationless or \u00d4dry\u00d5 merging. This may even have been observed \\citep{bell06a} though there is some disagreement as to whether the $\\Lambda$CDM  merger rate is high enough \\citep{bte07}. At this stage it is perhaps fair to say that dry merging simulations show that it does not disrupt elliptical scaling relations  \\citep{bk06,bk05,gva03} as one might naively expect. However only a limited number of simulations of this process have been done and they have not yet been incorporated into cosmological models in a detailed way such that they can be compared with data (e.g., numbers, sizes and masses of galaxies). Further it is not clear that a dry merging hierarchy consistent with cosmological downsizing can also be made consistent with the evolving mass-metallicity relation \\citep{pip08}. A contrasting picture is painted by \\citet{na07} using a SPH model of individual systems. They argue for a formation mode dominated by something very close to early monolithic collapse, but in a  $\\Lambda$CDM  cosmological context, with mergers (along with accretion) playing only a minor role in stellar mass growth at late times.  High spatial resolution studies of the morphologies and structures of passive galaxies offer one approach to gauging the importance of recent major merger events. A number of studies with the {\\em Hubble Space Telescope} (HST) have shown that half or more of red galaxies in color-selected samples have simple early type morphologies. Most of these studies are confined to redshifts of $\\sim1.5$ and less, and the early-type fraction varies from $\\sim50\\%$ to $70\\%$ \\citep{mou04,yan03}. At higher redshifts a significant fraction of the red galaxies appear to be discs (e.g., \\citetalias{abr07}, \\citealp{fon04}). Understanding the connection between these two classes of objects naturally focuses on the importance of mergers, since nearly equal-mass mergers are thought to transform discs into spheroids. Mergers, both gas-rich and dissipationless, are also thought to be important in the growth of the red sequence  and evidence, both direct and indirect, supports that this is occuring at intermediate and low redshifts \\citep[e.g.,][and the references therein]{bell06b}. It appears that much of the high-redshift merging activity may be of the dissipationless variety where the main effect of merging is to reorganize existing stellar population {\\em without} triggering new star formation. It is difficult to envision how this might operate unless the merging systems are themselves gas-poor, which is not generally expected \\citep{van05}. In any case, the signatures of such `dry' mergers are difficult to detect at high redshifts.  Recently, several imaging studies have shown that red galaxies at $z>1$ appear smaller than their likely present-day descendants with the same stellar mass \\citep{lon07,cim08}. The implications of these observations are seen most clearly in the structural and dynamical scaling relations, the Fundamental Plane and its projections (the Faber-Jackson (1976) and Kormendy (1977) relations). In the present paper we explore the nature of the Kormendy relation, (mean surface brightness within the effective radius, $\\langle\\mu\\rangle_e$, versus effective radius, R$_e$). This is the most observationally accessible projection of the fundamental plane at high-redshift. Our analysis spans the redshift range $1.2 < z < 2$ using HST NICMOS observations of a sample of quiescent high-redshift galaxies taken mainly from the  Gemini Deep Deep Survey \\citep[GDDS Paper I,][]{abr04}. We present NICMOS F160W images for ten of the twenty $z>1.3$ passive red galaxies from \\citetalias{mcc04}. These systems all have spectra dominated by old stellar populations. This extends to higher redshifts ($z>1.7$) than the earlier NICMOS work of \\citet{lon07} from the Munich Near-IR Cluster Survey \\citep[MUNICS,][]{dr01}. We also independently analyze the archival NICMOS data of \\citet{lon07} in the redshift range $1.2 < z < 1.7$ to supplement our sample and confirm their findings. At the higher redshifts previous findings of compact galaxies were based on optical data obtained with the Advanced Camera for Surveys (ACS) onboard HST \\citep{cim08}. Our use of NICMOS allows us to more robustly show that the old components in the galaxies are truly compact. Finally, we are able to unify the optical and infrared work by introducing a new `stellar mass Kormendy relation' which we use to  better quantify evolution in the sizes of early-type galaxies as a function of stellar mass over the redshift range $1<z<2$. We briefly examine the likelihood that dry mergers explain such size evolution, and examine whether an alternative process, adiabatic expansion, might be more important. We describe the observations in section~\\ref{obs}, our analysis in section~\\ref{analysis}, and present our results in section~\\ref{res}. In section~\\ref{disc} we discuss the implications of our observations for simple models for galaxy size growth. Throughout we use standard cosmological parameters; H$_0=70$~km~s$^{-1}$~Mpc$^{-1}$, $\\Omega_\\textrm{m}=0.3$, $\\Omega_\\Lambda=0.7$. Unless stated otherwise, all magnitudes are based on the AB system. ",
        "conclusions": ""
    },
    "0807/0807.3171_arXiv.txt": {
        "abstract": "{Analysis of bulges to redshifts of up to $z \\sim 1$ have provided ambiguous results as to whether bulges as a class are old structures akin to elliptical galaxies or younger products of the evolution of their host disks. } {We aim to define a sample of intermediate-$z$ disk galaxies harbouring central bulges, and a complementary sample of disk galaxies without measurable bulges. We intend to provide colour profiles for both samples, as well as measurements of nuclear, disk, and global colours, which may be used to constrain the relative ages of bulges and disks. } {We select a diameter-limited sample of galaxies in images from the HST/WFPC2 (Wide-Field Planetary Camera 2 at the Hubble Space Telescope) Groth Strip survey, which is divided into two subsamples of higher and lower inclination to assess the role of dust in the measures quantities. Mergers are visually identified and excluded. We take special care to control the pollution by ellipticals. The bulge sample is defined with a criterion based on nuclear surface brightness excess over the inward extrapolation of the exponential law fitted to the outer regions of the galaxies. We extract colour profiles on the semi-minor axis least affected by dust in the disk, and measure nuclear colours at 0.85 kpc from the centre over those profiles. Disk colours are measured on major axis profiles; global colours are obtained from 2.6\\arcsec\\-diameter apertures. Colour transformations and K-corrections are calculated using SEDs covering bands $U B V I J K$, from the GOYA photometric survey. } {We obtain a parent sample containing 248 galaxies with known redshifts, spectroscopic or photometric, spanning $0.1 < z < 1.2$. The bulge subsample comprises 54 galaxies ($21.8\\%$ of the total), while the subsample with no measureable bulges is $55.2\\%$ of the total (137 galaxies).  The remainder ($23\\%$) is composed of mergers.  We list nuclear, disk, and global colours (observed and rest-frame) and magnitudes (apparent and absolute), as well as galaxy colour gradients for the samples with and without bulges, and make them available in electronic form{\\thanks{Tables~\\ref{tab:bul_datos} to~\\ref{tab:nobul_gradient} are available in electronic form at the CDS via anonymous ftp to \\textsf{cdsarc.u-strasbg.fr (130.79.128.5)} {\\rm or via} \\textsf{http://cdsweb.u-strasbg.fr/cgi-bin/qcat?J/A+A/}}}. We also provide images, colour maps, plots of spectral energy distributions, major-axis surface brightness profiles, and minor-axis colour profiles for both samples.} {} ",
        "introduction": "\\label{sec:introduction}  Bulges of spiral galaxies are central ingredients in the study of galaxy formation.  Traditionally seen as elliptical galaxies that happen to live inside disks \\citep{Renzini99}, bulges are still often used by modellers as yardsticks of the importance of violent processes vs gentle accretion in the mass assembly of galaxies \\citep[e.g.,][]{Kauffmann96bul, Cole00, Saiz01}. In the now termed 'classical' hypothesis,  bulges form before disks, at high redshifts, via major mergers \\citep{Baugh98} or primordial collapse \\citep{Eggen62}. Subsequently, a new disk forms out of the remaining gas. However, other hypotheses have been presented for bulge formation \\citep[see][for two recent reviews]{Wyse97,Kormendy04}.  In a second scenario, bulges form after disks through bar instabilities \\citep{Pfenniger90}, clumpy fragmentation of the existing disk \\citep{Noguchi00,Immeli04}, or other secular processes; see \\citet{Athanassoula05} for a discussion of the relation of these processes to disky and peanut-shaped bulges.  A related model posits that bulges may be the product of initial star formation in a galaxy, which proceeds until the bulge becomes massive enough to allow subsequent incoming gas to form a stable disk \\citep{vandenBosch98}. Minor mergers probably have a role in the growth of bulges and disks \\citep{Aguerri01,Abadi03,Bournaud05,elichebulges06}, leading perhaps to a coeval growth of both components.  Knowing the relative ages of bulges and disks is key to being able to falsify one or more of the hypothesis for bulge formation.  A prevalence of blue bulges would argue against an old formation age for bulges.  In the disk instability model, blue, starbursting bulges would need to be found in sufficient quantities to account for the mass growth of today's massive, red bulges. The third scenario would be supported by the finding of similar colour for bulges and disks.  Usually, two approaches are followed for obtaining bulge ages and exploring  mechanisms of bulge formation and evolution. The first is the study of kinematic, luminosities, colours, and chemical abundances of local bulges to recover the tracks of their evolution. The second is the analysis of global properties of galaxies at high redshifts with the aim of studying their evolution nearer the epoch at which bulges may have formed.  In the local Universe, \\citet{Peletier96} found that the age differences between bulges and disks was, in most cases, consistent with zero.  Similarly, \\citet{Bouwens99} concluded that dating uncertainties prevent us from distinguishing between the various bulge formation models. Indeed, age dating of stellar populations has large uncertainties. In the present series of papers, we address the \\citet{Peletier96} study at redshift up to $z \\sim 1$.  To study bulges at high redshift, we need large samples from deep images at high spatial resolution. Deep surveys with HST provide enough resolution to resolve big bulges at cosmological distances; e.g. at $z=1$, a bulge with a size of 2 kpc subtends 0.26\\arcsec\\, about twice the point spread function (PSF) of the HST/WFPC2. But defining and characterising bulge components faces several practical difficulties derived from the properties of the imaging survey and the adopted selection criteria. First, pixelation and convolution with the PSF make the detection of small bulges difficult. Second, cosmological surface brightness dimming prevents detection of extended disks. Hence, samples selected by magnitude will favor compact, spheroidal galaxies, while diameter-limited selection will include only those bulges embedded in moderate- or high-surface brightness disks. Bulge-disk decompositions of the surface brightness profiles, already difficult in the local Universe, are made harder at high redshifts because bulges are not mapped with enough resolution and also sometimes because of the irregular morphologies of many disks. Moreover, the bulge colours based on bulge-disk decomposition have the additional problem of assigning a single representative colour for the bulges and for the disks. As the colour of the disk is usually not uniform, one can substract from bulge region disk light that is too blue or to red, which biasses the bulge colour.  Many recent studies of high-redshift galaxies are relevant to the formation of bulges \\citep[e.g.,][to name just a few]{Marleau98, Menanteau01, Ellis01, Hammer01, Simard02, Conselice04, Koo05, Koo05HDF, Elmegreen07,MacArthur07}.  The range of agreement and discrepancies between the conclusions of these studies is strongly affected by differences in sample selection and in the methodology to infer age information from the data. In the present paper, we approach the definition of bulge samples at high redshift in a complementary way to previous studies. The aim is to work with samples of disk galaxies, i.e., we focus on bulges residing inside disks.  We then isolate samples with and without measurable bulges, obtain colour profiles, and measure nuclear, disk and global colours. Specifically, we work with a diameter-limited sample in the Groth Strip survey, divided in  two subsamples by inclination. From those samples we exclude mergers (visually identified by their morphology), and take care to control sample pollution by bonafide ellipticals, which should be minimised in the high-inclination sample. We define a prominence index $\\eta$ as the surface brightness excess over the inward extrapolation of the outer disk profile, and use it to identify the sample of galaxies with measurable bulges. As a result, we get the sample of intermediate-$z$ bulge galaxies, but also we obtain a comparison sample of disk galaxies with no measurable bulges. We provide nuclear, disk, and global colours for galaxies in both samples in tabular form.  Details on the approach followed here are given in Sect.~\\ref{sec:approach}.  Section~\\ref{sec:data} describes the HST and ground-based data. In Sect.~\\ref{sec:sample}, we describe the characteristics of the high- and low-inclination samples, compare different methods to select bulge sample and describe our bulge sample selection method, discuss completeness issues as well as comparison samples at intermediate redshfits in the same survey. Bulge and disk colour measurements and K-corrections are detailed in Sect.~\\ref{sec:methodology}. The characteristics and colours of different galaxy samples with and without bulges are tabulated in Appendix~\\ref{sec:datatables}; in Appendix~\\ref{sec:samplesbulges}, we show the HST/WFPC2 $F814W$ images, also colour maps, semi-major axis surface brightness profiles and semi-minor axis colour profiles of the bulge sample.  In a second paper \\citep[hereafter Paper~II]{Dominguez08II}, we will analyse the nuclear, disk, and global colours of the galaxies presented here, establish relations between different parameters, and discuss implications of the results for galaxy formation and evolution models. A cosmology with  $\\Omega_M = 0.3$, $\\Omega_\\Lambda = 0.7$, $H_0 = 70$ \\kmsMpc\\ is assumed throughout.  Magnitudes are expressed in the Vega system. ",
        "conclusions": "\\label{sec:conclusions}  We have presented a sample of 248 disk galaxies from the Groth strip with known redshifts in which we have identified candidate bulge components. We provide global photometry as well as  colour measurements for bulges, disks, and for each galaxy as a whole.  For galaxies without bulges, we provide colour measurements at the same radial distances and following the same method as done for the galaxies with bulges.  These colours are analysed in Paper~II.  The tables in the Appendices are available in electronic format from the authors."
    },
    "0807/0807.1945_arXiv.txt": {
        "abstract": "I study the dynamics of a superfluid vortex in a random potential, as in the inner crust of a neutron star. Below a critical flow velocity of the ambient superfluid, a vortex is effectively immobilized by lattice forces even in the limit of zero dissipation. Low-velocity, translatory motion is not dynamically possible, a result with important implications for understanding neutron star precession and the dynamical properties of superfluid nuclear matter. \\bigskip\\bigskip ",
        "introduction": " ",
        "conclusions": ""
    },
    "0807/0807.4421_arXiv.txt": {
        "abstract": "The trigger system for the Auger fluorescence telescopes is implemented in hard- and software for an efficient selection of fluorescence light tracks induced by high-energy extensive air showers. The algorithm of the third stage uses the multiplicity signal of the hardware for fast rejection of lightning events with above 99\\,\\% efficiency. In a second step direct muon hits in the camera and random triggers are rejected by analyzing the space-time correlation of the pixels. The trigger algorithm was tested with measured and simulated showers and implemented in the electronics of the fluorescence telescopes. A comparison to a prototype trigger without multiplicity shows the superiority of this approach, e.g. the false rejection rate is a factor 10 lower. ",
        "introduction": "The Pierre Auger Southern Observatory in the Province of Mendoza, Argentina in the vicinity of Malarg\\\"ue measures ultra-high energy cosmic rays by an array of 1600 Water-Cherenkov tanks (SD) and 24 fluorescence telescopes (FD). It combines the surface and the fluorescence detection techniques in a hybrid design for minimal systematic errors and optimum energy and angular resolution.  Already during construction the Auger collaboration measured the energy spectrum, the arrival direction, and the mass composition of cosmic rays above $10^{18}$\\,eV with high statistical significance. Recently published results \\cite{auger_collaboration1} shed first light on the astrophysical important questions of possible sources, the energy spectrum and GZK-cutoff.  In the following sections the trigger system of the fluorescence telescopes is described in general with a particular focus on the software-implemented Third Level Trigger (TLT) as part of a multi-level trigger system. We specify the selection criteria (cuts) for efficient background rejection and measure the figure of merit of this new implementation.  \\subsection{FD system}  The 24 FD telescopes are located in four FD stations around the perimeter of the 3000\\,km$^{2}$ large SD array. They are designed to measure the fluorescence produced by secondary cosmic ray particles due to their interaction with $\\rm N_2$ molecules in the air. Each FD station comprises 6 telescopes with a Schmidt optical system that covers a $30^{\\circ}$ azimuth times $28.6^{\\circ}$ elevation field of view. Fluorescence light falls through a 2.2\\,m wide aperture with an UV transmitting filter and a corrector lens annulus on a spherical mirror and is focused on a photomultiplier camera \\cite{FDPaper}. The 440 camera pixels consist of hexagonal PMTs (type XP3062) arranged in a matrix of 22 rows and 20 columns.  The required high sensitivity of the PMT to single photons constrains the telescope operation to clear nights with dim moonlight. For safety reasons the stations are therefore equipped with a sophisticated slow control system which allows remote operation of all FD buildings from the central campus (CDAS) in Malarg\\\"ue via a telecommunication network .  \\subsection{FD electronics} The PMT signals are fed to 20 front-end boards hosted in a 19'' sub-rack. Each board processes the signals of 22 pixels of a single camera column. After analogue filtering and gain control at the  First Level Trigger (FLT) board the signals are digitized with 12-bit resolution and 10 MHz sampling rate.  The digital data are continuously stored in a $\\rm 64K\\,\\times$ 16-bit memory per pixel, which is organized in 64 ring-buffers of 1000 words each. If the following trigger stage finds a fluorescence light track in the camera image, the data recording continues after a post-trigger delay of $70\\,\\mu{\\rm s}$ with the next available ring-buffer. Otherwise, the current ring-buffer data is overwritten after $100\\,\\mu{\\rm s}$.  \\subsection{First Level Trigger (FLT)}  While ADC values are stored in memory, the running sum of the last 10 samples is calculated to smooth the random fluctuations of the sky background. If the running sum exceeds an adjustable threshold the corresponding pixel is marked as triggered for a $\\rm 20\\,\\mu s$ coincidence time.  The rate of triggered pixels is measured and kept constant around 100\\,Hz by regulating the individual thresholds. With this robust regulation scheme we prevent increasing random trigger rates despite slowly changing background light intensities during the night.   \\subsection{Second Level Trigger (SLT)}  While the FLT trigger selects pixels seeing light levels above the common night sky background, we have implemented a further stage in hardware identifying tracks of triggered pixels in the camera image. For that purpose and for the readout of the data the sub-rack is equipped with one Second Level Trigger (SLT) board. It reads the triggered pixels from the FLT to construct an image of the camera in its memory. Within this image the SLT logic scans for small track segments of 5 adjacent pixels showing the patterns in fig.~\\ref{fig:Pattern} by using highly parallel, pipelined pattern recognition logic \\cite{SLTpaper}. In total, 37\\,163 different combinations of pixels are checked every microsecond. If at least 4 out of 5 pixels forming a pattern have triggered within a $20\\,\\mu{\\rm s}$ long coincidence time, the recorded data are kept as event.  \\begin{figure}[ht] \\centering \\includegraphics[width = 0.40\\textwidth]{fig1.eps} \\caption{Basic topological patterns of triggered pixels as used by the SLT to identify shower track segments} \\label{fig:Pattern} \\end{figure}    \\subsection{Third Level Trigger (TLT)} \\label{sec:old_TLT} Different event types pass the first two hardware trigger levels. On the one hand there are extensive air showers we want to measure (s.~fig.~\\ref{fig:example_shower}a), on the other hand different types of background events.  A small fraction of the background consists of random triggers of few noisy pixels without temporal sequence (s.~fig.~\\ref{fig:example_shower}b). Another type of background are direct hits of muons in the PMTs, that result in tracks of synchronously triggered pixels (s.~fig.~\\ref{fig:example_shower}c). But the main problem is distant lightning causing large areas of the camera to trigger in fast bursts (s.~fig.~\\ref{fig:example_shower}d).  \\begin{figure*}[p] \\hfill \\begin{minipage}[t]{0.37\\textwidth} \\includegraphics[width = \\textwidth]{fig2a.eps} \\includegraphics[width = \\textwidth]{fig2b.eps} \\includegraphics[width = \\textwidth]{fig2c.eps} \\includegraphics[width = \\textwidth]{fig2d.eps} \\end{minipage}\\hfill \\begin{minipage}[t]{0.40\\textwidth} \\includegraphics[width = \\textwidth]{fig2e.eps} \\includegraphics[width = \\textwidth]{fig2f.eps} \\includegraphics[width = \\textwidth]{fig2g.eps} \\includegraphics[width = \\textwidth]{fig2h.eps} \\end{minipage} \\hfill ~ \\caption{Examples of different event types handled  by the third level trigger. \\quad \\emph{Left panels} show the camera view with 440 hexagonal pixels, signal times are coded by color lightness: a dark pixel corresponds to an early signal. \\emph{The right panels} give the corresponding multiplicity signal with the multiplicity parameters like explained in section~\\ref{sec:multiplicity_parameters}. \\quad \\emph{From top to bottom:} \\textbf{a)} Extensive air shower event showing a temporal development in vertical direction, \\textbf{b)} random trigger of few pixels, \\textbf{c)} camera hit by a cosmic muon resulting in equal pixel times and \\textbf{d)} an event caused by strong lightning.} \\label{fig:example_shower} \\end{figure*}  The software implemented third level trigger has to distinguish between background and extensive air showers, varying from dim and short tracks for distant showers up to an illumination of the whole camera for very close events. The key to the trigger decision lies in the timing information of the pixels: Real shower events always show a temporal development along their track.  The prototype TLT identifies background by analyzing the ADC data and checking for a temporal sequence inside the 5-pixel SLT-pattern. The major disadvantage of this approach was the very time consuming readout of ADC traces, needed to determine the pixel time. Especially during lightning, bursts of events with many pixels were rejected too slowly, all FLT buffers were filled and dead time was introduced.  In this paper we will describe a new two-stage TLT. In a first step  lightning events are fast and efficiently rejected by cuts on the hardware multiplicity data. In the second step the time sequence of the pixels data is analyzed using their ADC data.  In the DAQ chain the TLT is followed by the Event Builder, which merges concurrent events in neighboring telescopes together. A subsequent algorithm tries to calculate the preliminary shower direction to form a hybrid trigger (T3) for the suface detector. Finally the event is stored on hard disk and transmitted later on to the central campus via microwave uplink.   \\subsection{Data sample} \\label{sec:data_sample} For the developement of the new trigger we need a set of background events and extensive air showers. Therefore we randomly selected 8839 events from one year of data taking and manually separated them into 2425 shower and 6414 background events. This classification was done very conservative to prevent the misinterpretation of potential showers as background. The separation was successfully cross-checked by a reconstruction with the detector simulation and analysis framework Offline\\cite{offlinepaper}. None of our events classified as background, but 25\\,\\% of our shower candidates can be reconstructed. This confirms the conservative character of our manual classification.  After using this data sample for the trigger developement, we determine the background rejection efficiency with a second batch of background events independentely from the training data. The acceptance of shower events, especially at high primary energies, is checked with Monte-Carlo-simulated showers. ",
        "conclusions": "We have developed a new TLT software on the basis of the existing fluorescence telescope hardware. The main goals of the development were an improved background rejection and fast trigger decision in order to handle bursts of background events during bad weather conditions.  These aims are achieved by using a two stage algorithm: In a first step we use the temporal developement of the number of triggered pixels in an event. This multiplicity signal is calculated in hardware. Five independent cuts on multiplicity parameters are used to reject background events.  Although the multiplicity trigger works fine for the main data load of events with many illuminated pixels, it fails in distinguishing short shower traces from background events with less than 25 pixels. For these events we need a more detailed analysis of the temporal appearance of the triggered pixels along the shower track.  The acceptance of extensive air showers was determined to be better than 99\\,\\% using simulated events in the energy range from $10^{18}\\,\\rm eV$  to $10^{21}\\,\\rm eV$. At the same time 95\\,\\% of the background is rejected, during bad weather conditions the rejection efficiency of events with more than 25 pixels rises above 99\\,\\%. As the algorithm abstains from single pixel analysis for these events, the decision is done fast enough to reject background in almost real-time.  The new TLT was installed into the trigger chain in September 2007. This enhancement improves the data quality of the fluorescence detector and noticeably simplifies data taking especially during the rainy season."
    },
    "0807/0807.3037_arXiv.txt": {
        "abstract": "The advent of explicit Dirac--Born--Infeld (DBI) inflationary models within string theory has drawn renewed interest to the cosmological role of unusual scalar field dynamics, usually referred to as k-inflation. In this situation, the standard method used to determine the behavior of cosmological perturbations breaks down. We present a generic method, based on the uniform approximation, to analytically derive the power spectra of scalar and tensor perturbations. For this purpose, a simple hierarchy of parameters, related to the sound speed of the cosmological fluctuations and its successive derivatives, is introduced in a k-inflation analogue of the Hubble flow functions. The scalar spectral index and its running are obtained up to next to next to leading order for all k-inflationary models. This result relies on the existence of a well-motivated initial state, which is not trivial in the present context: having the wavelength of the Fourier mode smaller than the sonic horizon is indeed not enough and some conditions on the dynamics of the sound speed are also required. Our method is then applied to various models encountered in the literature.  After deriving a generic slow-roll trajectory valid for any DBI model, simple formulae for the cosmological observables are obtained. In particular, the running, as the spectral index, for the so-called UV and IR brane inflationary models is found to be uniquely determined by the 't~Hooft coupling. Finally, the accuracy of these cosmological predictions is assessed by comparing the analytical approximations with exact numerical integrations. ",
        "introduction": "\\label{sec:intro}  The primordial matter perturbations held responsible for growth and formation of large scale structure are commonly traced back to quantum fluctuations of a scalar field $\\varphi$ which should have dominated the energy density in the Universe at early times. The primordial power spectrum of both the scalar and tensor perturbations is a calculational output of the inflationary scenario, and the recent Cosmic Microwave Background (CMB) experiments have gathered considerable evidence in favor of it~\\cite{Gold:2008kp, Hill:2008hx, Hinshaw:2008kr, Nolta:2008ih, Dunkley:2008ie, Komatsu:2008hk}.  \\par  A compelling virtue of the inflationary paradigm is the fact that it can be sustained by a whole class of scalar field potentials $V(\\varphi)$, provided these exhibit characteristics (\\ie in their slope and curvature) in agreement with the ``slow-roll'' conditions: while inflation is under way, the potential must dominate over the kinetic energy $\\dot{\\varphi}^{2}/2$, that is, $\\ln V$ should be flat enough not to accelerate the field quickly. The slow-roll regime for inflation is described by a hierarchy of parameters $\\epsilon_{i}$, assumed to be small, in which the primordial power spectra can be analytically expressed as a Taylor expansion~\\cite{Stewart:1993bc, Martin:1999wa, Martin:2000ak, Schwarz:2001vv, Leach:2002ar, Schwarz:2004tz}.  \\par  A large body of  literature is devoted to the search for inflationary scenarios, in the sense that they should be naturally motivated by a high energy physics theory~\\cite{Lyth:1998xn}. For example, the toolkit of ten-dimensional super-string theory and its various compactifications to an effective four-dimensional field theory has been used to design candidate potentials~\\cite{Cline:2006hu, Kallosh:2007ig,McAllister:2007bg}.  \\par  An interesting approach studies the effect of modifications to the kinetic term of the inflaton (k-inflation)~\\cite{ArmendarizPicon:1999rj}; in the perturbation treatment, these modifications manifest themselves as a (possibly time-dependent) ``speed of sound'' $\\cs\\neq1$ (the speed of light being $c=1$) for the scalar Fourier modes~\\cite{Garriga:1999vw}. It turns out that string inflation models where the inflaton field is an open string mode are typically of this kind, with the speed of sound being a function of the background geometry~\\cite{Kachru:2003sx, Alishahiha:2004eh}. Hence inflationary model building in string theory combines both candidate potentials and non-canonical field evolution.  \\par  The prime example of these scenarios are the so-called brane inflation models, where the inflaton $\\varphi$ corresponds to the position of a $D$-brane within in a higher-dimensional manifold~\\cite{Dvali:1998pa, Dvali:2001fw, Alexander:2001ks}. Being an open string mode, $\\varphi$ has a Dirac-Born-Infeld (DBI) action, which is essentially the square root of the induced metric on the brane. This metric, in turn, contains information about the chosen background compactification of the extra-dimensions through the position-dependent brane tension $T(\\varphi)$. The inflaton potential $V(\\varphi)$ can be of various shape, and its exact calculation remains a matter of active research~\\cite{Baumann:2007ah, Baumann:2007np}. Typically, it receives contributions involving $T(\\varphi)$, but may also be affected by finer geometric detail such as the presence of other branes in the extra-dimensional background geometry. As a natural consequence of the non-canonical interplay between potential and dynamics in the k-inflation case, one can no longer trust the intuition that flat potentials support inflation. Through $T(\\varphi)$, the warping of the background acts as a break on the field, allowing potential energy domination and accelerated expansion of the Universe even when $\\ln V$ is steep.  \\par  As  it is clear from the above description, the standard methods of slow-roll inflation break down in the case of k-inflation. In particular, the formula expressing the background trajectory is modified since the shape of $T(\\varphi)$ (and not only the shape of the potential as in the ordinary case) affects the motion of the mobile brane. Another new feature of k-inflation, which is of prime concern for this article, is that one can no longer calculate the cosmological perturbations' power spectra using the standard techniques. Indeed, as already mentioned above, the scalar perturbations have now a time-dependent speed of propagation which prevents us to integrate the equations of motion in terms of Bessel functions. Moreover, concerning the perturbations' evolution, the ``usual'' Hubble flow functions $\\epsilon_{i}$ do no longer provide a sufficient description. Therefore, the main goal of this paper is to put forward a general formalism for k-inflation, resembling as far as possible the usual slow-roll formalism, where all these issues can be addressed in a consistent and unified way.  \\par  The paper is organized as follows. In Sec.~\\ref{sec:dbisr}, we use a combined hierarchy $\\left(\\epsilon_{i},\\delta_{i}\\right)$ of Hubble and ``sound'' flow functions such that $\\epsilon_{i},\\delta_{i}\\ll1,i\\geq1$ is the analogue of the standard slow-roll approximation~\\cite{Alishahiha:2004eh, Chen:2006nt, Kinney:2007ag, Peiris:2007gz}. Then, using the background equations of motion, the slow-roll trajectory is expressed as a quadrature. Our new formula can be applied to any DBI model (but not to k-inflationary models in general) characterized by the functions $T(\\varphi)$ and $V(\\varphi)$ and is valid under a single assumption, namely $\\epsilon _1\\ll 1$. In a next step, we carry on through the calculation of the k-inflationary perturbation spectra in full generality (but assuming, as usual, the smallness of the Hubble and sound flow parameters), including the non-trivial effects induced by a varying sound speed, in the so-called uniform approximation~\\cite{Habib:2002yi, Habib:2004kc}. In particular, we derive the scalar spectral index and the running at the next to next to leading order for a general model of k-inflation. Moreover, we show that a time-dependent sound horizon may lead to sub-sonic Fourier modes starting their evolution out of the Wentzel-Kramers-Brillouin (WKB) regime. This generically results in oscillations in the primordial scalar power spectrum. At the end of this section, we also discuss how the new parameters can be used to generalize to k-inflation the classification of (single field) inflationary models described in Ref.~\\cite{Schwarz:2004tz}. Then, in Sec.~\\ref{sec:examples}, our approach is applied to various example models encountered in the literature. In particular, we recover or extend previous results concerning power-law DBI inflation~\\cite{Spalinski:2007dv,Spalinski:2007qy}, the Kachru--Kallosh--Linde--Maldacena--McAllister--Trivedi (KKLMMT) model~\\cite{Kachru:2003sx}, chaotic Klebanov--Strassler (CKS) inflation~\\cite{Klebanov:2000hb, Bean:2007hc, Peiris:2007gz}, and derive new results for models having a CKS potential plus constant~\\cite{Shandera:2006ax}. At various stages, we assess the accuracy of our approximation scheme from exact numerical integrations. In Sec.~\\ref{sec:conclusion}, we recap and discuss our main findings. Finally, in three short appendices, we briefly compare our new hierarchy of parameters to the parameters already considered before in the literature and we recall how the master equation for the Mukhanov-Sasaki variable can be derived in k-inflation. ",
        "conclusions": "\\label{sec:conclusion}  In this section, we briefly recap and discuss the results obtained before. The main goal of our article was to present a generic formalism, as similar as possible to the standard approach, to investigate the features and the observational predictions of k-inflation. In particular, we have offered a new formula (but valid for DBI inflation only) for computing the classical trajectory for any warp function $T(\\varphi)$ and potential $V(\\varphi)$. Then, using the uniform approximation, we have computed the scalar and tensor power spectra, see Eqs.~(\\ref{eq:scalarspectrum}) and~(\\ref{eq:tensorspectrum}). These formulas are based on a double expansion. There is first an expansion around a pivot scale and, then, each coefficient of this expansion is in turn evaluated by means of another expansion, this time in terms of the Hubble and sound flow parameters. Finally, we have illustrated our results  on various example models encountered in the literature.  In particular, we have derived new results concerning the CKS models with constant term and obtained simple expressions for the cosmological observables.  \\par  Having generalized the standard formalism to k-inflation, the next step is of course to compare the predictions to the observations. This is done in Ref.~\\cite{Lorenz:2008je} in the case of the CMB WMAP5 data."
    },
    "0807/0807.0728_arXiv.txt": {
        "abstract": "From the radial velocities of the N~{\\sc iv} $\\lambda$4058 and He~{\\sc ii} $\\lambda$4686 emission lines, and the N~{\\sc v} $\\lambda$4604-20 absorption lines, determined in digital spectra, we report the discovery that the X-ray bright emission line star Wack~2134 (= WR 21a) is a spectroscopic binary system with an orbital period of 31.673$\\pm$0.002 days. With this period, the N~{\\sc iv} and  He~{\\sc ii} emission and N~{\\sc v} absorption lines, which originate in the atmosphere of the primary component, define a rather eccentric binary orbit (e=0.64$\\pm$0.03). The radial velocity variations of the N~{\\sc v} absorptions have a lower amplitude than those of the He~{\\sc ii} emission. Such a behaviour of the emission line radial velocities could be due to distortions produced by a superimposed absorption component from the companion. High resolution echelle spectra observed during the quadrature phases of the binary show H and He~{\\sc ii} absorptions of both components with a radial velocity difference of about 541~km s$^{-1}$. From this difference, we infer quite high values of the minimum masses, of about 87M$_\\odot$ and 53M$_\\odot$ for the primary and secondary components, respectively, if the radial velocity variations of the He~{\\sc ii} emission represent the true orbit of the primary. No He~{\\sc i} absorption lines are observed in our spectra. Thus, the secondary component in the Wack~2134 binary system appears to be an early O type star. From the presence of H, He~{\\sc ii} and N~{\\sc v} absorptions, and N~{\\sc iv} and C~{\\sc iv} emissions, in the spectrum of the primary component, it most clearly resembles those of Of/WNLha type stars. ",
        "introduction": "The X-ray source 1E 1024.0-5732, detected with {\\it EINSTEIN}, was identified by Caraveo et al. (1989) with the emission line star 2134 in the Wackerling (1970) catalogue, and suggested to be a binary system composed of an O star with a compact companion. Further X-ray data of this source obtained with {\\it ROSAT} were analyzed by Mereghetti et al. (1994), who described the optical spectrum to be of Wolf-Rayet type, and proposed that the X-rays could arise in the colliding winds of a WN+OB binary system. Thus, the star was added to the Catalog of Galactic Wolf-Rayet stars (Van der Hucht 2001) as WR~21a. Recent interferometric radio observations detected a weak non-thermal source at the position of Wack~2134, which was also interpreted as due to a colliding wind region in a WN+OB binary (Benaglia et al. 2005).  However, indications of orbital binary motion had not been found thus far. Here we present a radial velocity study of Wack~2134, showing it to be a binary system with an orbital period of 31.673$\\pm$0.002 days and a rather high eccentricity, e=0.64$\\pm$0.03. ",
        "conclusions": "From the radial-velocity variations of the He~{\\sc ii} $\\lambda$4686 emission line in the spectrum of the WN-type star Wack~2134, we found a most probable period P=31.673$\\pm$0.002 days, thus revealing Wack~2134 as an eccentric binary system (e=0.64$\\pm$0.03).  Our higher-resolution spectra taken during quadratures show H and He~{\\sc ii} resolved into two components. We could obtain an orbital solution for the secondary component measuring the He~{\\sc ii} $\\lambda$5411 absorption line in those spectra, but this orbit has only one measurement at one of the quadratures (see Fig.~\\ref{vrs5411}). Using the He~{\\sc ii} $\\lambda$4686 emission line and the He~{\\sc ii} $\\lambda$5411 absorption line to follow the motion of the WN and O component, respectively, we obtained very high minimum masses $M_{\\rm WN}$$\\sim$87M$_\\odot$ and $M_{\\rm O}$$\\sim$53M$_\\odot$. In spite of the high minimum masses obtained, no light variations indicative of eclipses are found in the photometry performed by ASAS. These minimum masses have to be considered as preliminary, as more high resolution and S/N spectra are needed to confirm these values. Also, more precise photometry could shed light on this new massive binary system. Although we prefer to be cautious, the WN component in the Wack~2134 system could be amongst the most massive stars ever measured from binary motion in the Galaxy, as both WN6ha components in the binary system WR~20a (83+82$M_\\odot$; Rauw et al., 2004; Bonanos et al., 2004), located about 16\\arcmin\\ from Wack~2134, and those in the new Wolf-Rayet binary system NGC3603-A1, which has masses of 116+89$M_\\odot$ (Schnurr et al., 2008; Moffat et al., 2004)."
    },
    "0807/0807.2946_arXiv.txt": {
        "abstract": "{} {I present the Monte Carlo radiative transfer code TRADING (Transfer of RAdiation through Dust In Galaxies). The code computes self-consistently the extinction of radiation in a dusty medium (including absorption and scattering) and the dust emission.} {A binary-tree adaptive grid is used for the description of the dust distribution. Dust radiation is computed at thermal equilibrium or under stochastic heating condition, for a distribution of grains of different radii and materials. The code is applied to the case of a clumpy galactic disk, including both diffuse dust and a distribution of spherical clouds modelled on the GMCs of local galaxies. Diffuse and localised sources of starlight are used, with independent spectra.} {A model is provided for the edge-on galaxy NGC~891. The SED of the galaxy from the UV to the submm/mm range can be well reproduced by: a bulge/disk configuration of old stars together with an extended dust disk, as suggested by the analysis of optical/near-infrared images; a clumpy dust distribution of the same mass as the diffuse dust disk, together with a UV emitting component, half of which in the form of a diffuse disk and half in sources embedded in clouds. In total, it is found that about 35\\% of the bolometric radiation is absorbed (and emitted) by dust; and that absorption of starlight from the old population contributes to about 60\\% of the dust emission. A significant component of the dust emission from clouds is due to absorption of diffuse radiation. Radial profiles of dust emission in a clumpy disk are almost independent of the wavelength, with the exception of the wavelength range on the Wien side of the thermal equilibrium peak. } {} ",
        "introduction": "The last two decades have seen a good effort in studying the role of dust in the appearance of a spiral galaxy\\footnote{A parallel and more considerable development has occurred in radiative transfer studies of dusty circumstellar environments \\citep[see, e.g., the various models described in ][]{PascucciA&A2004}.}. Radiative transfer models have shown that it is necessary to include scattering and consider appropriate geometries for a proper analysis of the internal extinction \\citep[][to cite only the most recent works]{PieriniApJ2004,TuffsA&A2004, RochaMNRAS2008}. Building on models of extinction in the optical, simulations of dust emission in the Far-Infrared (FIR) have been produced \\citep{BianchiA&A2000b,PopescuA&A2000,BaesProc2004}.  Puzzling has been the comparison between models and observations. Dust disks of moderate optical depth can reproduce observations of edge-on spirals \\citep{XilourisSub1998,BianchiA&A2007}, in accordance with the analysis of extinction of  background sources by less inclined disks \\citep[see, e.g., ][]{HolwerdaAJ2007}. Instead, the Spectral Energy Distribution (SED) in FIR/submm requires dust absorption and emission from a significantly larger mass \\citep{BianchiA&A2000b,PopescuA&A2000,MisiriotisA&A2001}.  The {\\em excess} mass derived from emission models could be associated with molecular clouds \\citep{PopescuA&A2000}. Thus, it becomes necessary to study the influence of inhomogeneities, or {\\em clumping}, on the radiative transfer. So far, the effects of clumping have been studied mainly in extinction for galactic disks \\citep{BianchiSub1999,MatthewsApJ2001,PieriniApJ2004} and in emission for geometries appropriate to starburst galaxies \\citep{MisseltApJ2001}. In all these cases, clumps have been modelled as higher density cells in a uniformly spaced dust grid. However, the lack of resolution inside a clumpy cell does not allow to study cases in which significant emission comes from within the cell, as when modelling young stellar sources embedded in a molecular cloud. To obviate this, \\citet{SilvaApJprep1998} use a complete radiative transfer solution for a single cloud, and add the cumulative output from clouds to the dust emission from the diffuse medium. A more self-consistent solution would require the usage of an adaptive mesh refinement grid, with a finer resolution for the denser medium.  I present in this paper the Monte Carlo (MC) radiative transfer code TRADING (Transfer of RAdiation through Dust IN Galaxies\\footnote{The acronym was suggested by a misprint in a notice at the {\\em Exploratory Workshop:Tracing Dust In Spiral Galaxies}. The workshop was funded by the European Science Foundation and took place in Ghent, Belgium on May 13-17, 2007.}). For selected viewing directions, the code produces images of dust-extinguished starlight in the Ultraviolet (UV), Optical and Near Infrared (NIR), and of dust emission from the Mid-Infrared (MIR) to the submm. Images can be integrated to obtain the total SED. TRADING builds on our previous models and has been developed from the regular grid version used in \\citet{BianchiA&A2007}.  The scheme of TRADING is similar to that of the code DIRTY \\citep{GordonApJ2001,MisseltApJ2001}. However, in TRADING the dust distribution is described with an adaptive grid, while DIRTY is based on a regular grid. An adaptive grid has been included in the dust radiative transfer codes SUNRISE~\\citep{JonssonMNRAS2006}, which has been used to study extinction in hydrodynamical simulations of isolated spirals \\citep{RochaMNRAS2008}, and ART$^2$ \\citep{LiApJ2008}. SUNRISE does not simulate dust emission, while ART$^2$ includes thermal equilibrium heating from a single grain. In TRADING, instead, emission is fully taken into account by using a distribution of grain sizes and materials; it includes both thermal and stochastic heating, and self-absorption. With DIRTY, SUNRISE and ART$^2$ (and many other codes in the radiative transfer literature) TRADING shares the use of the MC technique, as this is of easy implementation in any geometrical configuration and not necessarily heavier in computing time than analytical techniques \\citep{BaesMNRAS2001a}.  TRADING has been developed, and used in this paper, to study the influence of inhomogeneities (clouds) in the dust emission from a galactic disk. To my knowledge, this is the first attempt ever to model the influence of dust on a galactic SED by using a fully self consistent radiative transfer calculation over scales ranging from molecular clouds to the galactic disk. Despite the code is admittedly tuned to this problem, however, it is of general use and can be easily adapted to other astrophysical scenarios.  The scheme of the paper is as follows. The code is described in Sect.~\\ref{code}. In Sect.~\\ref{dust}, I present the implementation of the \\citet{DraineApJ2007b} dust model used in this paper. In Sect.~\\ref{results} TRADING is used to model the global SED of NGC891. The results of this application, together with the main features of TRADING, are summarised in Sect.~\\ref{summary}. ",
        "conclusions": "\\label{summary}  I have presented in this paper the Monte Carlo radiative transfer code TRADING. It is designed to study continuum extinction and emission in a dusty medium. The main features of TRADING are: \\renewcommand\\theenumi{\\roman{enumi}} \\begin{enumerate} \\item a description of the dust distribution based on a binary-tree adaptive grid;  \\item possibility of emitting stellar radiation from several diffuse distributions and collections of point sources/clouds, each with an independent spectrum;  \\item dust emission from a distribution of grains of different sizes and materials, under thermal equilibrium or stochastic heating conditions;  \\item self absorption through an iterative process.  \\end{enumerate}  The code shows an excellent agreement with benchmark cases and can be easily adapted to several environments of astrophysical interest. For the application shown in this paper, TRADING has been tailored to the case of a galactic disk. The disk includes: a diffuse dust component; a clumpy component made of a distribution of spherical clouds, with properties modelled on those of GMC in local galaxies. Building on the studies of \\citet{XilourisSub1998} and \\citet{PopescuA&A2000}, simulations have been produced for the well known case of the edge-on spiral NGC~891. The main results can be summarised as follows.  \\renewcommand\\theenumi{\\arabic{enumi}} \\begin{enumerate} \\item The global SED of NGC~891 from the UV to the mm can be reproduced by a model including: a diffuse dust disk, more widespread than the stellar and of moderate optical depth ($\\tau^\\mathrm{f. o.}_\\mathrm{V}$=1); a clumpy dust distribution, of mass similar to that of the diffuse disk; a population of old stars, smoothly distributed in an exponential disk and a bulge, emitting about 75\\% of the bolometric radiation; a population of young stars emitting the remaining fraction mainly in the UV, roughly half of which embedded in clouds. Despite the increased attenuation for UV radiation, the old stars are the main contributor to dust heating (about 60\\% of all radiation emitted by dust). This is in contrast with the result of \\citet{PopescuA&A2000}, which predict a dominant UV contribution to heating.  In the MIR, the current model produces less radiation than the observations. This can be only in part explained with a lack of spatial resolution for clouds. The difference can be possibly due to a more complex structure of the ISM and ISRF than modelled here. This could also be the reason for the larger extinction in the K-band image and the smoother appearence of the 850$\\mu$m than observed.  \\item Dust in clouds is found to be significantly heated by both the embedded UV radiation and the NIR radiation from the diffuse stellar distribution penetrating the outer shells of a cloud. This second source of heating results in cold dust emission which is not present in the template for localised emission used by \\citet{PopescuA&A2000} and \\citet{DasyraA&A2005}. In \\citet{PopescuA&A2000}, cold dust emission is provided by a second diffuse dust disk; however, contrary to this work, the radiative transfer calculation for the extra dust component in \\citet{PopescuA&A2000} is not done self-consistently with that for localised sources. This is the cause of the larger fraction of embedded UV radiation in the current work. Also, the difference between a clumpy and smooth disk results in a smaller attenuation of diffuse radiation at optical wavelengths in the current model. In \\citet{DasyraA&A2005}, the lack of cold dust emission associated with localised emission appears to be the main reason of their claim for a larger dust emissivity in the FIR. Other models which compute the contribution of clouds independently from that of the diffuse ISM, such as GRASIL \\citep{SilvaApJprep1998}, might need to include external heating of clouds by a diffuse ISRF. Because of the diffuse heating, cold emission in the submm comes from all the dust components.  \\item In a clumpy disk, and for the resolution of current and future instruments, color gradients are rather similar along all the dust emission spectrum. They follow the gradients of stars and of the main dust component. A steepening may be observed in the Wien side of the spectrum for thermal equilibrium heated grains (20$\\mu$m $\\la \\lambda \\la$ 100$\\mu$m), mostly because of the presence of a bulge.  \\end{enumerate}  The conclusions drawn on the energy budget for the single object NGC~891 might not be of general application for other galaxies, though NGC~891 is considered to be a prototypical spiral. Unfortunately, not many galaxies have detailed observations of the FIR/submm beyond the thermal peak. However, the advent of the SPIRE instrument aboard the Herschel satellite \\citep{GriffinAdSpR2007}, in conjunction with the new ground-based submm facilities, will soon permit a more detailed study of this spectral range, necessary for constraining the bulk of dust mass."
    },
    "0807/0807.2729.txt": {
        "abstract": "% context heading (optional) % {blabla} leave it empty if necessary {Hot cores in molecular clouds, such as Orion KL, Sgr B2, W51 e1/e2, are characterized by the presence of molecules at temperature high enough to significantly populate their low-frequency vibrationally excited states. For complex organic molecules, characterized by a dense spectrum both in the ground state and in the excited states, such as methyl formate, ethyl cyanide or dimethyl ether, lines from vibrationally excited states certainly participate to the spectral confusion.} % aims heading (mandatory) {Thanks to the recent laboratory study of the first torsional excited mode of methyl formate, we have searched for methyl formate, HCOOCH$_3$, in its first torsionally excited state ($\\upsilon_t$=1) in the molecular cloud W51 e2.} % methods heading (mandatory) {We have performed observations of the molecular cloud W51 e2 in different spectral regions at 1.3, 2 and 3 mm with the IRAM 30 m single dish antenna.  } % results heading (mandatory) { Methyl formate in its first torsionally excited state ($\\upsilon_t$=1 at 131 cm$^{-1}$) is detected for the first time toward W51 e2. 82 transitions have been detected among which 46 are unblended with other species. For a total of 16 A-E pairs in the observed spectrum, 9 are unblended; these 9 pairs are all detected. All  transitions from excited methyl formate within the observed spectral range are actually detected and no strong lines are missing. The column density of the excited state is comparable to that of the ground state. For a source size of 7\\arcsec we find that $\\mathrm{T_{rot}}$ = 104 $\\pm$ 14 K and N =  9.4$^{+4.0}_{-2.8}$ $\\times$10$^{16}$ cm$^{-2}$ for the excited state and  $\\mathrm{T_{rot}}$ = 176 $\\pm$ 24 K and N =  1.7$^{+.2}_{-.2}$ $\\times$ 10$^{17}$ cm$^{-2}$ for the ground state. Lines from ethyl cyanide in its two first excited states ($\\upsilon_t$=1, torsion mode at 212 cm$^{-1}$) and ($\\upsilon_b$=1, CCN in-plane bending mode at 206 cm$^{-1}$) are also present in the observed spectrum. However blending problems prevent a precise estimate of its abundance.  However as for methyl formate it should be comparable with the ground state for which we find $\\mathrm{T_{rot}}$ = 103 $\\pm$ 9 K and N =  3.7$^{+0.6}_{-0.4}$ $\\times$10$^{15}$ cm$^{-2}$ for a 7\\arcsec source size.} % conclusions heading (optional), leave it empty if necessary {With regard to the number of lines of excited methyl formate and ethyl cyanide detected in W51 e2, it appears that excited states of large molecules certainly account for a large number of unidentified lines in spectral survey of molecular clouds. } ",
        "introduction": "W51 e2 is a hot core part of the W51 H$\\mathrm{_{II}}$ region located in the Sagittarius spiral arm at a distance of about 7-8 kpc. It is a region of high-mass star formation. W51 e2 and W51 e1 are thought to be important star-forming cores. They exhibit a rich chemistry, comparable to that observed in Orion or in Sagittarius. Numerous large organic molecules have been observed towards them. CH$_3$CN and CS maps were studied by \\citet{zhang98}. Formic acid (HCOOH) was mapped by \\cite{liu2001}. Methyl formate (HCOOCH$_3$) and ethyl cyanide (CH$_3$CH$_2$CN) were observed in several studies \\citep{liu2001,ikeda2001,remijan2002}. \\citet{ikeda2001} studied ethylene oxide (c-C$_2$H$_4$O) and its isomer acetaldehyde (CH$_3$CHO). Acetic acid (CH$_3$COOH) was detected by \\citet{remijan2002} with a fractional abundance of (1-6) $\\times$ 10$^{-2}$ relative to HCOOCH$_3$. Glycine, whose presence may be suggested from the observations of acetic acid with which it shares common structural elements, was not detected in W51 e2 \\citep{snyder2005}. Recently, trans-ethyl methyl ether was detected in W51 e2 \\citep{fuchs05}.  The rotational temperature of most of these large molecular species is high. \\citet{liu2001} estimated the rotational temperature in W51 e2 to be in the 200-300 K range. From high resolution observations of CH$_3$CN analyzed with statistical equilibrium models, \\citet{remijan2004} derived the kinetic temperature in W51 e2 to be T$_{\\mathrm{kin}}$ =  153(21) K. At such temperature, low-energy vibrational excited states can be significantly populated. Transitions from vibrationally excited states have indeed been observed in other sources such as in Sgr B2(N-LMH) for C$_2$H$_3$CN, CH$_3$CH$_2$OH \\citep{numm98} and CH$_3$CH$_2$CN \\citep{mehringer2004}. Recently, lines from torsionally excited methyl formate have been identified in Orion KL \\citep{koba2007}.  In this study we present the first detection of excited methyl formate and ethyl cyanide in the molecular cloud W51 e2. The initial project was to look for methyl carbamate (NH$_2$COOCH$_3$), an isomer of glycine (NH$_2$CH$_2$COOH) which has a larger dipole moment, making its detection more favorable than glycine. It was not detected. However, few strong unidentified lines in the data attracted our attention and were possibly attributed to methyl formate in its first torsional excited state. Further observations confirmed this detection and also lead to the detection of excited ethyl cyanide in W51 e2. The first vibrationally excited state of methyl formate is the CH$_3$ torsion mode, $\\nu_{18}$, hereafter called $\\upsilon_t$=1, at 131 cm$^{-1}$ (188 K). The rotational spectrum in this excited state was measured and analyzed by \\citet{ogata2004}. Ethyl cyanide has two close vibrationally excited states, the CCN in-plane bending mode, $\\nu_{13}$, hereafter called $\\upsilon_b$=1, at 206 cm$^{-1}$  (296 K) and the CH$_3$ torsion mode, $\\nu_{21}$, hereafter called $\\upsilon_t$=1, at 212 cm$^{-1}$ (305 K). A preliminary analysis of the rotational spectrum in these two excited states is presented in the paper from \\citet{mehringer2004}.  The observations and the methods used for the data analysis are described in Sect.~\\ref{obs} and Sect.~\\ref{analysis}, respectively. The study of methyl formate and ethyl cyanide in the ground and excited state is presented in Sect.~\\ref{mtf} and Sect.~\\ref{propio}, respectively. The search for methyl carbamate and glycine is presented in Sect. ~\\ref{mtc}. The discussion is carried out in Sect.~\\ref{discussion}. ",
        "conclusions": "\\label{discussion}  We find that the rotational temperatures derived from the rotational diagram analysis are different from the kinetic temperature, T$_\\mathrm{kin}$ =  153(21) K derived by \\citet{remijan2004}. The rotational temperature of ethyl cyanide (103 - 114 K) is lower than T$_\\mathrm{kin}$ independently of the adopted source size. It is higher than T$_\\mathrm{kin}$ if we consider methyl formate in its ground state (176 - 199 K), lower if we consider only the excited state (104 - 131 K) and of the same order if we consider simultaneously the ground state and the first excited state of methyl formate (144 -154 K). All this tends to indicates that the LTE hypothesis is probably not fully valid. However, despite these possible departure from LTE, the rotational diagram method which implicitly assume that LTE is reached and that all temperature are equivalent (i.e. T$_\\mathrm{exc}$ = T$_\\mathrm{kin}$ = T$_\\mathrm{rot}$ = T$_\\mathrm{vib}$), is the only one left for us to get an estimation of the excitation temperature since a statistical analysis is ruled out by the absence of known collision rates for the studied molecules.  The relevance of the LTE assumption within each rotational level may also be evaluated by comparing the cloud density to the critical density of each vibrational state. Collisional rates are not known for molecules as complex as methyl formate or ethyl cyanide. However, it is possible to estimate roughly the critical density for these molecules by adopting the value of methanol interacting with H$_2$, of the order of a few 10$^{-11}$ cm$^3$.s$^{-1}$ \\citep{pottage2004} and assuming it is the same for the ground and vibrational excited states. The Einstein coefficients of the rotational transitions are in the range 10$^{-5}$-10$^{-6}$ s$^{-1}$. The critical density is thus of the order of 10$^{5}$-10$^{6}$ cm$^{-3}$, i.e., comparable to the hydrogen density in W51 e2 estimated to be n$_\\mathrm{H}$ = 5(2) $\\times$ 10$^5$ cm$^{-3}$ \\citep{remijan2004}. There is thus a competition between collisional and radiative excitation and it is clear that the levels are probably not all thermalized.  The excitation mechanism populating the observed excited state of methyl formate and ethyl cyanide may be infered in a similar way. The Einstein coefficients for ro-vibrational transitions from the excited states to the ground state are of the order of a few 10$^{-1}$ s$^{-1}$, much greater than for rotational transitions within the excited state (10$^{-5}$-10$^{-6}$ s$^{-1}$). Thus, adopting the same value for the collisional rate as before (a few 10$^{-11}$ cm$^3$.s$^{-1}$), the density needed to thermalize the levels in the excited state by collisions must be greater than the critical density which is of the order of 10$^{10}$ cm$^{-3}$. The dust in W51 e2 has a temperature of about 140 K \\citep{sollins2004} and thus emits efficiently at the wavelength of the excited states of both molecules. It is thus most probable that the excited states of methyl formate and ethyl cyanide are populated by radiative processes rather than by collisions.\\\\  We find that the rotational temperature of the gas decreases when the source size gets smaller. This is surprising since it is expected that the deepest regions of hot cores are also the warmest. The observed transitions are optically thin for the adopted source sizes of 12\\arcsec and 7 \\arcsec. BIMA observations of methyl formate in W51 e2 \\citep{remijan2002} show that the emission region is smaller for high energy transition than for low energy transition. The size of the emission region is around 12\\arcsec for the 7$_{2,5}$-6$_{2,4}$ transition at 90146 MHz corresponding to an energy of 10 cm$^{-1}$ (15 K) and around 7\\arcsec for the 18$_{5,13}$-17$_{5,12}$ transition at 228629 MHz corresponding to an energy of 75 cm$^{-1}$ (108 K). More interferometric observations are needed to locate precisely the emission region of methyl formate and ethyl cyanide in the ground and excited state and understand their temperature distribution.\\\\   The presence of methyl formate in torsional excited state in hot cores such as W51 e2 and Orion KL \\citep{koba2007} is not surprising. Methyl formate is very abundant and its torsional mode has a very low energy (188 K). By comparison, it is not surprising that ethyl cyanide is more difficult to detect: it is less abundant and its excited states lie at slightly higher energy, around 300 K. However it is clear that observations having better signal to noise, better spectral  and angular resolution, such as forthcoming observations from Herschel and ALMA, will reveal much more lines from excited states of these molecules but also of abundant large molecules possessing low-frequency vibrational states. For example dimethyl ether, CH$_3$OCH$_3$, has two torsional modes at about 203 and 242 cm$^{-1}$ and transitions from these modes should be present in the spectra of hot cores. More generally, a large number of unidentified lines reported in the spectra of hot molecular clouds should be due to transitions from abundant molecules in excited states.\\\\"
    },
    "0807/0807.3567_arXiv.txt": {
        "abstract": "Using the most recent results about white dwarfs in 10 open clusters, we revisit semi-empirical estimates of the initial-final mass relation in star clusters, with emphasis on the use of stellar evolution models. We discuss the influence of these models on each step of the derivation. One intention of our work is to use consistent sets of calculations both for the isochrones and the white dwarf cooling tracks. The second one is to derive the range of systematic errors arising from stellar evolution theory. This is achieved by using different sources for the stellar models and by varying physical assumptions and input data. We find that systematic errors, including the determination of the cluster age, are dominating the initial mass values, while observational uncertainties influence the final mass primarily. After having determined the systematic errors, the initial-final mass relation allows us finally to draw conclusions about the physics of the stellar models, in particular about convective overshooting. ",
        "introduction": "The initial-final mass relation (IFMR) for low- and intermediate-mass stars is an important input for many astrophysical problems. Given the initial Main Sequence (MS) mass of a formed star, the IFMR provides the expected mass during its final White Dwarf (WD) cooling stage, and is an estimate of the total mass lost by the star during its evolutionary history. A correct assessment of the IFMR is very important when predicting, for example, the chemical evolution history of stellar populations, or their mass-to-light ratio (defined as the ratio of the mass of evolving stars plus remnants -- WDs, neutron stars and black holes -- to the integrated luminosity of the population), and in general for any problem related to the origin and evolution of gas in stellar populations. It is also a crucial item when using the WD luminosity functions to determine the age of stellar populations \\citep[see, e.g.,][]{prad07}. It also provides an empirical test concerning the upper initial mass value $M_{\\rm up}$ of stars developing degenerate carbon-oxygen cores, and thus a lower limit of the mass range of core-collapse supernovae. For a thorough review of the history of the IFMR, we recommend \\citet{weide00}.  Theoretical estimates of the IFMR are still prone to large uncertainties. This is due to our poor knowledge of the efficiency of mass loss processes for low- and intermediate-mass stars but also to uncertainties in the predicted size of CO cores during the Asymptotic Giant Branch (AGB) evolution, resulting mainly from uncertainties in the treatment of the thermal pulse phase, with the associated third dredge-up, hot bottom burning \\citep[see, e.g.,][]{iberen83} and also the treatment of rotation \\citep{domi96}. In fact, the recognition that intermediate-mass stars may lose the largest part of their mass during the AGB phase stems from early IFMRs \\citep{koeweid80}.  Starting with the pioneering work by \\cite{weid77}, (semi-)empirical routes have been followed to establish the IFMR independently of theoretical modelling the AGB phase, leading to a series of global determinations of the IFMR by, e.g., \\citet{weidkoe83}, \\citet{weid87}, \\citet{weide00}, \\citet{fer05}, and many others.  Methods to estimate the IFMR that make use of the smallest number of assumptions are probably those based on the use of WDs harbored in star clusters \\citep[but see also][for a very recent study of the IFMR based on WDs in common proper motion pairs]{cata07}. Thanks to a large amount of observational effort by various authors \\cite[see, e.g.][]{kr93, kr96, clav01, dob04, kal05, bars05, dobbie06} the recent study of the IFMR  by \\citet{fer05} makes use of 40 DA (hydrogen-atmosphere) WDs belonging to 7 open clusters. Even more recently \\citet{kalir08} have added a few more data points below the low-mass end (initial masses below $\\sim$2.0 $M_{\\odot}$) of the \\citet{fer05} sample, by including WDs detected in the old open clusters NGC6819, NGC7789 and NGC6791, while \\citet{rubin08} have provided WD data for the $\\sim$200~Myr old cluster NGC1039. Therefore, the observational material is increasingly covering not only more objects, but also a larger range of cluster -- and thus WD -- ages and metallicities, such that even differential investigations may become feasible in the near future. At the same time, high-resolution, high signal-to-noise spectra and improved spectral analysis methods reduce the intrinsic errors in the determination of the WD masses. It is therefore time to investigate the systematic errors arising from stellar evolution theory and models that enter these semi-empirical methods. This is the main object of this work.  IFMR determinations based on cluster WDs work as follows: after detection, spectroscopic estimates of the WD surface gravity $g$ and \\teff\\ are needed. For a fixed $g$-\\teff\\ pair, interpolation within a grid of theoretical WD models covering a range of masses provides the mass $M_{\\rm f}$ and the age $t_{\\rm cool}$ of the WD. Independent theoretical isochrone fits to the Turn Off (TO) luminosity in the cluster Color-Magnitude-Diagram (CMD) provide an estimate of the cluster age $t_{\\rm clus}$.  The difference  $t_{\\rm clus}-t_{\\rm cool}$  is equal  to the lifetime $t_{\\rm prog}$ of the WD progenitor, from the MS until the start of the WD cooling. Making use of mass-lifetime relationships from theoretical stellar evolution models, the initial progenitor mass $M_{\\rm i}$ is immediately obtained from $t_{\\rm prog}$.  It is clear how and where theoretical models play an important role in this procedure. WD cooling models plus progenitor evolutionary tracks and isochrones are needed to determine the various $t_{\\rm clus}$, $t_{\\rm cool}$ and $t_{\\rm prog}$. Given the very short timescale of the AGB evolution (see also next section for a short discussion on this issue) AGB modelling and its related uncertainties do not play a role in these types of analyses, but uncertainites entering the previous evolutionary phases and the WD cooling stage can affect the derived IFMR. Another important point is the issue of consistency. First of all, the isochrones employed to determine $t_{\\rm clus}$ should be computed starting from the same evolutionary tracks adopted to determine $M_{\\rm i}$ from the estimated value of $t_{\\rm prog}$. Consistency is not usually achieved in existing estimates of the IFMR. Even in the most recent studies of the IFMR, in several cases cluster ages are obtained from earlier studies that have employed models different from the ones used to estimate $M_{\\rm i}$. Secondly, all cluster ages should be determined following the same method, and -- most importantly -- employing the same set of isochrones. Instead, results from different sources for the various clusters are used.  As an example, we summarize briefly the work by \\citet{fer05}.  The Hyades age has been taken from \\citet{perr98}, who had calculated their own isochrones with an updated version of the CESAM code \\citep{morel97}.  In case of M35 the age quoted in Table~1 of \\citet{hippel05} was employed; this age was derived summarizing the results of three different previous works by other authors. The same is true for the Pleiades, where the assumed cluster age comes from a synthesis of the results of several groups, obtained with different generations of stellar evolution models. The isochrones by \\citet{schall92} led \\citet{meynet93} to an age of 141~Myrs for NGC2516, which was quoted by \\citet{fer05} as 158~Myr. The remaining clusters in that sample also have ages from various isochrone sources, which span more than a decade of progress in stellar evolution modelling, most notably the significant improvements in opacities and the equation of state. The cooling tracks came from \\citet{font01}, and the $t_{\\rm prog}$-- $M_{\\rm i}$ relation from \\citet{gir02}, taken from the tracks with the metallicity closest to the adopted one, but without interpolation between compositions. It is obvious that there is a certain degree of arbitrariness in the use of stellar evolution theory. This inherent inconsistency is probably the result of the fact that most works on the IFMR are more focused on acquiring new and improved data with an emphasis on the internal errors of their derived \\teff\\ and $M_{\\rm f}$.  The aim of the present work is instead to concentrate on the stellar evolution models needed for deriving the IFMR. It is obvious from the previous discussion that a consistent use of stellar models and procedures about how to employ them is required, to achieve a higher accuracy. We therefore will make use of calculations of stellar models, that cover all phases from the MS to the WD cooling phase. They will be used for the appropriate metallicity of each cluster, in order to have at hand accurate and homogeneous isochrones and progenitor lifetimes. This will provide an IFMR with the highest internal consistency achieved so far.  To get a measure of the errors introduced by using different codes and input  physics, we  also repeat  the procedure  by using  alternative  sets of models for the WD cooling and/or progenitor and cluster ages. While this gives a first indication about the size of the systematic errors resulting form different stellar model sources, we finally quantify the influence of various physical effects important for both the isochrone and WD cooling calculations. Using Monte Carlo methods we determine the total extent of the systematic errors. This allows us also to verify which features  of the IFMR are robust and which ones depend on present model uncertainties.  Using the robust results, we can finally exclude some of the stellar models because they lead to internal contradictions. One of our conclusions is that models without overshooting from convective cores on the MS do not yield self-consistent results. On the other hand, the parametrization of overshooting during the AGB phase, which influences the growth of the carbon-oxygen core, cannot be the same in all convective layers as it is at the convective core during the MS phase, because it produces final masses inconsistent with observations. We thus can also reduce uncertainties in the physics of stellar interiors.  The structure of the paper is as follows: In Sect.~2 we outline the procedure to obtain an IFMR, summarize the observational input data we have used for the WDs and the parent clusters, and present the set of the stellar models used for determining the reference IFMR as well as for the evaluation of the systematics. Section~3 then concentrates on the estimates of the global error bars, while Sect.~4 contains the results and a comparison with theoretical IFMRs. A summary and conclusions follow in Sect.~5. ",
        "conclusions": "Since its earliest determination, the semi-empirical IFMR has indicated a monotonic rise of the final WD mass with the initial mass of the progenitor. It is also strong evidence that single stars of intermediate mass do not develop degenerate cores close to the Chandrasekhar limit, most likely due to the effect of strong mass loss on the AGB. These qualitative results have remained unchanged in spite of the significant improvement both in the amount and accuracy of observational data and in the theoretical models. Nevertheless there are still many open questions, such as a dependence on metallicity, the slope of the IFMR at the upper (initial) mass end, changes in slope, and spread around a mean relation. To answer these questions, not only increasingly accurate observations are needed, but also the systematic uncertainties coming from the theoretical input -- which is necessary to establish the semi-empirical IFMR --  have to be controlled. These are related to evolutionary tracks and isochrones that determine the total age of the observed WDs and cooling tracks used for deriving the WD age. Our paper was intended to investigate the influence of such systematic uncertainties and to derive robust features of the IFMR. We also investigated what uncertainties in the theory of stellar evolution would lead to inconsistent results, and therefore could be restricted or excluded.  From the available WD data we selected a total of 52 objects in 10 open clusters, plus Sirius B. There are, in fact, even more objects available, but we disregarded those for which our cluster distance determination (Sect.~2) would not be possible because of the lack of calibrating field stars at that metallicity (NGC6791 at ${\\rm [Fe/H]} = 0.45$). We also did not take into account the objects from \\citet{cata07}, because of the more indirect  way to obtain the age of the mostly unevolved binary companion. Inclusion of such objects would weaken our aim of highest possible self-consistency.  For the determination of mass and cooling age of the WDs we used different sets of cooling track, from \\citet{sala00} and the LPCODE \\citep{alt03} to evaluate the influence of different codes, input physics and chemical stratification (Sect.~3.1). We estimated the error ranges quite conservatively and to evaluate the global error budget we used Monte Carlo simulations (Sect.~3.3), for which we assumed either a normal or uniform distribution of the different error sources. We included the effect of observational errors, too.  Results concerning cooling ages and final masses are quite robust: the accuracy of the WD masses depend completely on the observational uncertainties, while their ages is influenced in most cases (in terms of number of clusters) by the uncertainties of the cooling tracks, mostly systematic differences between different sets of cooling tracks and changes in the neutrino emission rates.   This result implies that any spread in $M_{\\rm f}$, as seen in Fig.~\\ref{fig:ref_ifmr}, which is outside the error bar is real, in the sense that $M_{\\rm f}$ does vary at the same (within the errors in $M_{\\rm i}$) initial mass. Whether this spread is consistent with theoretical expectations, related to a steep increase of $M_{\\rm f}$ with $M_{\\rm i}$ in that mass range, or whether it is due to star-to-star variations in the total mass lost, cannot be decided yet, as the error bars in $M_{\\rm i}$ are too large even in the case of the Praesepe WDs, which we discussed in detail in Sect.~2.4.   Figure~\\ref{mimfs00aldo} demonstrates that uncertainties due to different WD codes with similar, but not identical physics and model details are essentially irrelevant  for the  estimate  of  $M_{\\rm f}$.  However,  differences in  the cooling ages have a significant influence on $M_{\\rm i}$, since they affect directly progenitor lifetimes, and the more massive progenitors evolve on very short timescales. Therefore the two different sets of cooling tracks lead  to quite  different values for  $M_{\\rm i}$, particularly for initial masses above 5~M$_\\odot$. In  this  regime, $M_{\\rm  i}$  differences are  between  $\\sim$  0.5 and  2.0 $M_{\\odot}$.  The determination of $M_{\\rm i}$ is more complex, as it involves several steps. The age determination of the cluster or binary system requires metallicity, reddening and distance. We have used literature values for the former two, and our own homogenous determination for the latter. Uncertainties from both steps have been taken into account. Using appropriate isochrones the age is then obtained, which, together with the WD cooling age, results in the pre-WD lifetime. As this is, to an accuracy of 1\\% or slightly more, identical to  the lifetime from  the MS to  the end of the  central He-burning phase, all uncertainties concerning the AGB evolution are irrelevant. We employed different sets of stellar models \\citep{pietr04,gir00}, which differ slightly in the detailed treatment of the input physics (but both include core overshooting during the MS) as well as additional models from the former source (BaSTI), which do not consider convective overshooting. The first effect -- different tracks -- has a very minor effect on $M_{\\rm i}$, as we showed in Fig.~\\ref{mibastileo}. Note that this includes both the cluster age determination and the relation between progenitor age and mass. The neglect of overshooting, however, leads to internal inconsistencies in the semi-empirical IFMR: at the high end, initial masses are predicted, for which the models themselves predict that no electron degenerate CO-WD would result from the evolution. By discriminating between the different error sources (but omitting the overshooting effect), Fig.~\\ref{fracsigmas} demonstrates that the cluster age uncertainties  (due  to uncertainties  in  the  cluster  [Fe/H] and  reddening estimates) dominate the error on $M_{\\rm i}$, although the above-mentioned WD age can be significant for the more massive progenitors, too.  The case of NGC2099 illustrates the importance of accurate metallicities and reddenings: Figure~\\ref{fig:m37} shows how drastically the semi-empirical IFMR changes, two different estimates of [Fe/H] and $E(B-V)$ are assumed. If one accepts this uncertainty as being real, NGC2099 cannot be used at all to learn about the IFMR.  After we had determined the full extent of the uncertainty associated to our semi-empirical IFMR, we compared it with the predictions from the theoretical calculations, both from BaSTI and LPCODE. Both differ in their treatment of convective overshooting during the AGB phase. One should recall that the BaSTI models treat that phase with synthetic models, while LPCODE follows the evolution in full detail. Figure~\\ref{mimfbastisoo} shows them in comparison with the semi-empirical IFMR and the relation between $M_{\\rm i}$ and the mass of the CO-core at the first thermal pulse, $M_{\\rm c1TP}$. As long as stars do not avoid the AGB phase, $M_{\\rm f} \\geq M_{\\rm c1TP}$ should be fulfilled. Therefore, this relation is a strong lower limit to the semi-empirical IFMR. The agreement between the $M_{\\rm i}-M_{\\rm c1TP}$ relations from both codes is excellent; there are no significant systematic uncertainties in its prediction. This again emphasizes the need for convective overshooting during the MS phase:  without it a significantly larger numbers of objects falls below this limiting line.  On the other hand, the theoretical IFMR from the LPCODE tracks is not consistent with the data: In these models overshooting during the AGB, mainly from the lower boundary of the pulse-driven convective zone, limits the growth of the CO-core such that the final core size is almost identical to that at the first pulse. However, most objects have final masses clearly above this line, such that we concluded that the extent of overshooting during the AGB phase must be significantly smaller than on the MS. This is in agreement with hydrodynamical simulations.  The   theoretical  IFMRs  displayed   in  Figure~\\ref{mimfbastisoo},   and  in particular those resulting from the BaSTI models display the largest gradient in the same mass range where a significant spread in $M_{\\rm f}$ is observed. We demonstrated that indeed that spread could arise solely from this local steep gradient, such that no other explanation would be needed. If it is not, one would have to invoke a spread in the total mass lost in a given, small mass range. This is not exceptional, as the morphology of horizontal branches in Galactic globular cluster shows, but the reasons for such a star-to-star variation are unclear. In other mass ranges the total error budget is too large to discuss the reality and reason for apparent spreads in $M_{\\rm f}$.  Figure~\\ref{fig:ref_ifmr} presents our reference semi-empirical IFMR, obtained using BaSTI isochrones and tracks with overshooting and S00 cooling models. The error bars include all errors we have investigated. Compared to \\citet{fer05} our error bars are necessarily larger (we consider more sources of error), but on average only by  less than a factor of two, and  mostly for $M_{\\rm i}$. We also present a simple linear fit through all the data as well as a piecewise fit, which reflects the properties of the theoretical IFMR based on the BaSTI models.  In summary, our most important results are \\begin{enumerate} \\item None of the WDs employed in current IFMR determinations is close to the Chandrasekhar mass, not even the progeny of the more massive intermediate mass stars; \\item stellar models without convective overshooting during core hydrogen burning lead to internal inconsistencies in the semi-empirical IFMR; \\item overshooting from convective boundaries during the AGB phase must be significantly reduced (compared to the case of convective cores along the MS) to reproduce the observed $M_{\\rm f}$ values; \\item the uncertainty in $M_{\\rm f}$ is dominated by observational errors; \\item  the  uncertainty in  $M_{\\rm  i}$  has  several reasons:  both  cluster parameters and isochrone details influence the cluster age and thus the progenitor mass; the uncertainty on the WD cooling age can sometimes also be the dominant factor; \\item the observed dispersion in $M_{\\rm f}$ at approximately constant $M_{\\rm i}$ (3.0--3.5 $M_{\\odot}$), in particular for the Praesepe objects, appears to be real. It may follow from the steep increase of $M_{\\rm f}$ with $M_{\\rm i}$ predicted by theoretical IFMRs in that $M_{\\rm i}$ range, rather than being caused by a spread in mass loss among cluster AGB stars; \\item from the general agreement  between the theoretical  IFMRs and the semiempirical data, we find no evidence that the mass loss prescriptions used in the stellar evolution calculations grossly disagree with the total mass actually lost by low- and intermediate- mass stars; \\item  the case  of NGC2099  illustrates  the necessity  for accurate  cluster parameters; reliable composition data and reddening are required; \\item as long as up-to-date input physics is used for the stellar models, systematic uncertainties do not change the overall appearance of the semi-empirical IFMR as determined recently by several groups. \\end{enumerate}  An extension of this analysis to a larger range of [Fe/H] is needed, to investigate a possible metallicity dependence of the IFMR, given the small [Fe/H] range spanned by the systems analyzed in this study (see Table~\\ref{Ages}).  In addition to the relevant observational data, this requires also a careful study of how to extend to a wider [Fe/H] range the empirical MS fitting technique adopted here to determine cluster distances, if we wish to maintain the same degree of homogeneity in the derivation of the cluster ages. An extension to lower metallicities seems to be preferable, as both calibrating objects and suitable clusters should exist, contrary to the case of an extension to higher, that is super-solar clusters.   As most of the clusters used for deriving the IFMR have nearly solar metallicities, the question arises, what the true solar chemical composition is? \\citet{ags:2005} have presented new determinations of the solar heavy element abundance, drastically lower than the standard composition assumed here and in the stellar evolution tracks employed. Both BaSTI and LPCODE progenitor models consider a standard solar mixture with $Z/X \\sim$0.0245, whereas \\citet{ags:2005} redetermination gives $Z/X \\sim$0.0165. Assuming that the differential abundance analyses for our open cluster sample provides the true abundances relative to the sun, also the clusters should be less metal-rich than thought. It would be interesting to see how this will affect the semi-empirical IFMR and its agreement with theoretical predictions. An extension of this investigation by using  in the computation of progenitor  models the new solar abundances is currently under way."
    },
    "0807/0807.4806_arXiv.txt": {
        "abstract": "In recent years there have been many reported detections of highly redshifted or blueshifted narrow spectral lines (both emission or absorption) in the X-ray spectra of active galaxies, but these are all modest detections in terms of their statistical significance. The aim of this paper is to review the issue of the significance of these detections and, in particular, take account of publication bias. A literature search revealed $38$ reported detections of narrow, strongly shifted ($v/c \\ge 0.05$) X-ray lines in the $1.5-20$~keV spectra of Seyfert galaxies and quasars. These published data show a close, linear relationship between the estimated line strength and its uncertainty, in the sense that better observations (with smaller uncertainties) only ever show the smallest lines. This result is consistent with many of the reported lines being false detections resulting from random fluctuations, drawn from a large body of data and filtered by publication bias such that only the most `significant' fluctuations are ever reported. The reality of many of these features, and certainly their prevalence in the population at large, therefore remains an open question that is best settled though uniform analysis (and reporting) of higher quality observations. ",
        "introduction": "\\label{sect:intro}  The arrival in recent years of large quantities of X-ray data from CCD and grating spectrometers has provided a vast increase in the amount of spectral information available to X-ray astronomers. These data have vastly improved our understanding of the X-ray properties of active galaxies such as Seyfert galaxies and quasars. Among the discoveries to have been made using these data were narrow line-like features in emission or absorption at unexpected energies in the canonical $2-10$  keV X-ray band \\citep[see e.g.][]{yaqoob99,turner02,pounds03a,porquet04,matt05,longi06,cappi06}. In many cases these have been identified with transitions in iron that have been strongly redshifted or blueshifted out of the usual iron line band ($6.4-6.9$ keV) by high bulk flow velocities ($v \\gsim 0.05c$) or gravitational effects. Yet they are narrow and line-like, indicating low velocity dispersion. The narrowness and highly shifted centroid energies of these features mark them as distinct from the relativistically broadened emission lines seen in Seyfert galaxies \\citep[see e.g.][for reviews]{fabian00,reynolds03}. The extreme nature of these features means they have profound implications for the structure and energetics of the nucleus \\citep{cappi06}.  Irrespective of their proposed physical importance the reported features are all rather modest detections, in the sense that the statistical significance is not outstanding in any single case. Indeed, the vast majority of reported cases appear to lie in the ``$2-3\\sigma$'' regime\\footnote{Usually this means the outcome of an hypothesis test was a $p$-value in the range from $\\sim 5\\times 10^{-2}$ down to $\\sim 3 \\times 10^{-3}$. Such results are usually reported as ``detected at $95 - 99.7$ per cent confidence'' or given in units of $\\sigma$ by comparison with the tail area under the Normal curve.}. The detection process usually involves modelling the featureless continuum emission of the target object and searching for large, localised, positive or negative residuals between the data and model (spectral lines may appear in emission or absorption); the larger the contribution of the residuals to the fit statistic (e.g. the change in $\\chi^2$ upon including a line in the fitted model) the more significant the feature. The wide bandpass and good spectral resolution of modern detectors, combined with the number of datasets that have been processed mean that data archives for recent spectroscopy missions will  contain many modest signal-to-noise ``features'' simply from random sampling fluctuations in the photon counting signal from otherwise featureless continua. For example, the Tartarus database\\footnote{{\\tt http://astro.ic.ac.uk/Research/Tartarus/}} contains some $661$ \\asca\\ observations of active galaxies, and the \\xmm\\ Science Archive (XSA\\footnote{{\\tt http://xmm.esac.esa.int/xsa/index.shtml}}) contains over 1,000 publically available ``Guest Observer'' observations listed under the proposal category ``AGNs, QSOs, BL Lacs and XRB''. The number of spectra is arguably even higher than this because each of the longer observations is often divided into multiple spectra as a function of time, source flux, and so on, and indeed many of the line detection papers report the features to be ``transient'' (i.e. detected in only a subset of the data).  The large number of moderately significant detections (see sect.~\\ref{sect:analysis}) might be considered as evidence to support the reality of the features. However, when considering results presented for individual datasets drawn from a much larger population of available data one should be aware of the distorting effects of ``publication bias,'' also known as the ``file-drawer effect'' -- the tendency for positive results to be published and negative results to go unreported (``filed away''). See \\citet{sterling59}, \\citet{rosenthal79} and \\citet{begg88} for general discussion of publication bias, and also \\citet{stern97} and \\citet{naylor97} for a more recent discussion of the importance of publication bias in the context of medical trials. One way to test for the presence of publication bias is through a ``funnel plot,'' originally proposed to aid meta-analyses of medical trials \\citep{egger97}, which compares the size of a trial\\footnote{The size of the sample used in the study; larger trials, i.e. those with larger sample sizes, tend to give higher signal-to-noise results.} to its estimate of the strength of the effect.  All estimates of the strength of an effect should be symmetrically scattered around the true value, with smaller trials providing less precise estimates and so larger scatter. In the absence of publication bias this will result in a symmetric funnel-shaped plot because the  estimates of the effect strength are independent of the sample sizes, but the scatter is larger for smaller samples. If publication bias is present, experiments or observations are less likely to be published if the estimate of the effect strength is low (or of low significance), i.e. the bias is against publishing non-detections, leading to an asymmetric funnel plot in which the strength of the effect is correlated with the sample size. The funnel plot produced from biased literature is the same as from the equivalent unbiased literature but with less ``interesting'' results (i.e. less significant or less strong results, which lie on one side of the funnel) systematically removed.  This paper describes a simple meta-analysis of the published detections of highly shifted, narrow X-ray lines in active galaxies using a funnel plot-like analysis. The conventional funnel plot would not be appropriate in the present context because there is expected to be considerable intrinsic heterogeneity in the strength and properties of the shifted lines, which means there is no single, true value for the strength of the ``effect.'' But the principle of the funnel plot should still hold: the estimated strengths of the lines should be independent of the quality of the data used to find them (in this context quality means essentially the signal-to-noise of the data, which is of course closely related to the size of the photon sample that constitutes the spectrum). The strength of a real line should be independent of the exposure time and detector sensitivity used to measure it. ",
        "conclusions": "\\label{sect:disco}  The tendency for stronger lines to be accompanied by proportionately larger uncertainties (as shown in Figure~\\ref{fig}), or equivalently the relatively constant $|EW|/error$ ratio over the large range in $|EW|$, requires explanation. The lines with larger (absolute) $EW$ should be easy to detect in more sensitive observations (and should give smaller uncertainties) and so should populate the upper-left region of Figure~\\ref{fig}, but this is not the case. One simple explanation is that many of the line detections are the most `significant' false detections from a large population of data covering a wide range of power to detect potential lines. But this begs the question of why are there so many false detections?  Examination of a spectrum in isolation may or may not provide a false detection, but an apparently stronger detection will be more strongly favoured for publication. The many hundreds of observations that have been examined and not provided detections (whether with strong or weak limits) would populate the shaded region of Fig.~\\ref{fig} if only they were published (very few observations are published with limits on the strength of undetected narrow, shifted lines). The reported detections may be in effect the ``tip of the iceberg'' -- the strongest or most significant of a population of random fluctuations, with the rest of the population unseen due to publication bias. If the lines are genuine, the challenge is to explain why all the detections are close to the detection limit despite the huge range in the quality of the data, or equivalently, why the largest (absolute) $EW$ lines appear only in the poorest data with the largest error bars. Of course it is plausible that Fig.~\\ref{fig} shows a mixture of false and true detections, with the true detections expected mostly among the weaker lines (e.g. $|EW| < 30$~eV) since they are necessarily limited to moderately weak detections even in the best data (which is not true of stronger lines).  \\subsection{Detection methods} \\label{sect:detect}  Virtually all these lines were justified on the basis of $p$-values from hypothesis tests being smaller than some threshold $\\alpha$, a process intended to limit the fraction of false detections (Type I errors) to $\\alpha$, and in almost all the cases listed above a reasonable detection significance level (e.g. $\\alpha = 0.05$ or $0.01$) was used. Perhaps the $p$-values were systematically underestimated, leading to an abundance of false detections?  Estimating the statistical significance of a possible line feature is indeed rather difficult \\citep[see][]{freeman99,protassov02}, and many authors have used and continue to use inappropriate tests (e.g. the $F$-test) which may inflate the number of false detections. Ideally, a thorough and uniform analysis of all the data would solve the problem that different data analysis techniques were used by different authors, but that is not the purpose of the present paper, which simply makes use of the published results as they are presented in the literature. However, even a uniform and comprehensive analysis of all the datasets listed in Table~\\ref{table}, using sophisticated statistical methods, would not change the basic result (the clustering of lines near the diagonal in Figure~\\ref{fig}) unless it was true that many of the uncertainties listed in the table are greatly exaggerated. If the re-calculated uncertainties were often much smaller the typical $|EW|/error$ would be much higher, and the tendency for strong lines to have proportionately larger uncertainties might be eroded -- the points would fill more of the top-left region of the figure rather than skirting the diagonal. But there is no clear reason to suspect this might be true. If a re-analysis found the claimed significances (e.g. the $p$-values used for detection) were slightly too high or too low this may affect the number of points in the figure, but not the trend it reveals. If the $p$-values were systematically far too small (i.e. too significant) there must be a large number of false detections present in the figure, whereas if the $p$-values are reliable (or even overestimated) the $|EW| - error$ relation still requires an explanation. Therefore, in the remainder of this discussion it is assumed that the statistical tests used in the papers listed in Table~\\ref{table} are sound, and that it is still necessary to seek an explanation for the relation shown in Figure~\\ref{fig}.  \\subsection{Confounding factors}  The observed effect could be produced from genuine line detections if the two variables, line strength and its uncertainty (i.e. data quality), were correlated with some other factor, such as source distance. Perhaps the more distant, luminous quasars, that often yield the poorest spectra, possess intrinsically stronger line features compared to the nearby, low-luminosity Seyferts that have high-quality spectra available. Indeed, there is a correlation between redshift and $EW$ (Spearman rank-correlation coefficient $\\rho = 0.62$, $p = 5.6 \\times 10^{-5}$), but it is much weaker than between $|EW|$ and its uncertainty ($\\rho = 0.96$, $p < 10^{-15}$), and is due  entirely to the four lowest redshift sources (NGC 4151, NGC 7314, NGC 3516 and NGC 3227 at $z \\le 0.01$), which all have long exposures and showed only very weak lines. If these sources are ignored (leaving $28$ lines) there is no significant correlation between redshift and $|EW|$ ($\\rho = 0.33$, $p=9.0\\times 10^{-2}$) or its uncertainty ($\\rho = 0.24$, $p = 0.22$). The tight, linear relation between $|EW|$ and its uncertainty involves all sources and so cannot be due to the $EW$ (or its uncertainty) being correlated with redshift. Also, as can be seen from Table~\\ref{table}, the strong lines from ESO 113-G010, ESO 198-G24 and Mrk 79 have large errors because the observations were very short ($<8$~ks), not because the targets are intrinsically faint. In principle longer observations of these strong lines should conclusively demonstrate the reality of these features. However, a subsequent, much longer ($100$~ks) observation of ESO 113-G010 by \\citet{porquet07b} did not detect the $5.4$ keV emission line previously found in a $\\sim 4$ ks exposure \\citep{porquet04}. If real, the $EW$ must have decreased tenfold between observations to avoid detection by \\citet{porquet07b}; either a coincidence or an indication that the original detection of a strong line was false.  Another possible confounding factor is line energy. The continuum spectrum of the active galaxies listed in Table~\\ref{table} is usually well described by a power law with a photon index typically in the range $\\Gamma \\sim 1.5 - 2.5$, meaning there are far fewer photons at higher energies than lower energies. If narrow lines appearing at higher energies (e.g. blueshifted Fe K$\\alpha$ at $E \\gsim 7.1$~keV) tended to have higher $|EW|$ than those at lower energies (e.g. redshifted Fe K$\\alpha$ at $E \\lsim 6.2$~keV), the strongest lines may have the  largest uncertainties simply because they occur preferentially at higher energies. There is very little correlation between line energy and $|EW|$ ($\\rho = 0.27$, $p=0.11$) or its uncertainty ($\\rho = 0.26$, $p = 0.13$). Indeed, Table~\\ref{table} shows that some of the smallest $EW$s (and uncertainties) were found in lines above $7.1$~keV, and the tight relation between strength and uncertainty is present in the subset of higher-energy lines. It is not clear what other source property could be the confounding factor needed to explain the absence of strong lines in the best observations.  There is one other point worth emphasizing about the line energies. Of the $38$ lines listed in Table~\\ref{table}, $14$ were found in the range $5-6$ keV. The excess of detections could indicate that modestly redshifted ($z \\sim 0.05 - 0.3$) emission or absorption are more significant or robust than the more strongly shifted features, i.e. the number of detections is enhanced because there are more `true' lines in this band. However, the findings of this paper remain unchanged if one considers only the $13$ lines in the $5-6$ keV range with errors: these lines follow the same $|EW| - error$ correlation as shown in Fig~\\ref{fig}, and the $|EW|/error$ is not higher for this subsample. Indeed, the mean $|EW|/error = 2.4$ for the $5-6$ keV subset, and $2.5$ for the remaining lines. (Excluding the $5.57$ keV line in NGC 3516, for reasons discussed below, the $5-6$ keV subsample gives a mean $|EW|/error = 2.1$.) Even for this subsample the problem remains to explain the tight, linear correlation between $|EW|$ and its uncertainty. There are at least two possible effects that might help explain an enhanced rate of false detections in this narrow band. The first is confusion with emission structure from a strong, broad (possibly asymmetric) line centred around 6.4 keV. Such emission may produce an excess of counts in the $5-6$ keV region, above the expected continuum, and so enhance the appearance of line-like residuals. The second effect is observer bias; it is plausible that individual observers preferentially attend to residuals in the $5-7$ keV region immediately around expected Fe structures. Very few of the papers listed in Table~\\ref{table} report that the detections were made on the basis of a systematic and uniform search over a wide spectral range (a notable exception is \\citet{nandra07}). Either or both these effects may enhance the number of false detections made in the $5-6$ keV range.  \\subsection{The strongest individual cases}  Of the $36$ lines listed in Table~\\ref{table} (ignoring the BALs) with strength and uncertainty estimates, only five have strengths that are greater than three times their uncertainty (half width of the 90 per cent confidence interval). These are: emission at $5.57$ keV in NGC 3516 \\citep[$EW \\approx 23$~eV;][]{turner02}, absorption at $6.2$ keV in E1821+643 \\citep[$EW \\approx -54$~eV;][]{yaqoob05}, absorption at $7.7$ keV in MCG-5-23-16 \\citep[$EW \\approx -33$~eV;][]{braito07}, $1.63$ keV and $2.94$ keV absorption in PG 1211+143 \\citep[$EW \\approx -14$~eV and $-36$~eV;][]{pounds03a,pounds05}. These are considered in turn below.  The emission line in NGC 3516 was found in \\chandra\\ HETGS data, but could not be detected in the partly simultaneous \\xmm\\ data, those data gave an upper limit on the line flux an order of magnitude smaller than the \\chandra\\ detection \\citep{turner02}. If real, the line flux must have decreased by at least one order of magnitude between the observations. Furthermore it should be noted that the line energy was held fixed during the evaluation of the confidence interval on the flux \\citep[see the footnote to Table~1 of][]{turner02}. As the line energy was not predicted but obtained from fitting the data, it should have remained a free parameter throughout the calculation of confidence intervals, otherwise the confidence interval may be artificially reduced. In any event, the procedure described by \\citet{turner02} differs from the standard procedure adopted in the other cases and so is perhaps best considered a lower limit on the size of the confidence region of that particular line (as indicated in Figure~\\ref{fig}).  The absorption line found in the \\chandra\\ observation of E1821+642 was seen in the HEG spectrum but could not be confirmed in the lower signal-to-noise MEG data. In the case of PG 1211+143 the simultaneous detection of multiple lines, and their identification at similar blueshifts, in both CCD (EPIC) and grating (RGS) data from \\xmm\\ \\citep{pounds03a} would seem to put this case on firmer ground. For completeness, it should be noted that the identification of the lines in terms of highly blueshifted features has been debated, e.g. \\citet{mckernan05,kaspi06,pounds07,reeves08}. \\citet{reeves05} presented \\chandra\\ grating data of the same object and again reported absorption lines, except redshifted not blueshifted. The $7.6$~keV absorption line found in the \\xmm\\ data was not detected in more recent \\suzaku\\ data \\citep[with an upper limit $\\sim 4$ times smaller than the original \\xmm\\ detection;][]{reeves08}.  \\citet{braito07} and \\citet{reeves07} studied MCG-5-23-16 using simultaneous \\xmm, \\chandra, \\suzaku\\ and \\xte\\ observations. The detection of $7.7$ keV absorption is based on the EPIC pn spectrum from \\xmm. The EPIC MOS spectrum is consistent with the pn spectrum but is unable to confirm the presence of the line due to the smaller photon sample size; the \\suzaku\\ data show a possible absorption feature, but it was poorly constrained compared to the EPIC pn data; and the \\chandra\\ data were unable to confirm the line detection.  The lines that gave the largest improvement in the $\\chi^2$ fit statistic were both absorption lines: at $5.9$~keV in NGC 3516 (with \\asca) and at $7.45$~keV in NGC 4151 (with \\xmm). The former has no published EW, while the latter has a surprisingly low $|EW|/error$ ratio given its apparent effect on the fit. However, as \\citet{nandra07} noted, there is some ambiguity over whether the NGC 4151 feature should be identified with a line or an edge (see their section 8.8.4).  The Seyfert galaxy NGC 3516 provides one more interesting example. \\citet{dovciak04} described a $6.08$~keV emission line in an \\xmm\\ EPIC pn spectrum of NGC 3516 taken in 2001 April \\citep[see also][]{bianchi04} and suggested the feature is varied in $EW$ and/or energy. The significance of the feature was assessed using an $F$-test (see section~\\ref{sect:detect}) but no uncertainties were given on the $EW$ and so the observation is not represented in Figure~\\ref{fig}. The same observation was analysed by \\cite{iwasawa04}, who claimed a periodic modulation in the spectral shape of a broad $5.6-6.5$ keV line based on $\\sim 3$ `cycles'. Although this claim is intriguing it is not an independent confirmation of the significance of the $\\sim 6.1$~keV line; the analysis is an attempt to better understand and model the feature reported by \\citet{dovciak04}, assuming its reality and using the same data, not an independent assessment of it.  \\subsection{The effect of selection and publication bias}  This leaves the possibility that many or most of the line detections are false detections caused by random sampling fluctuations. The number of false detections may at first sight appear large, but one must remember that each spectrum from \\xmm, \\chandra\\ etc. contains $\\gsim 50$ resolution elements, and there are many hundreds of spectra (especially considering that longer observations are routinely split into multiple spectra corresponding to different time intervals, flux levels, etc.). The non-detections from each resolution element of each spectrum contributes an (unpublished) point inside (or just above) the shaded region of Fig.~\\ref{fig}. The human analyst, or an automated search algorithm, has a tendency to focus on the largest fluctuations -- this is a selection bias. These may then be subjected to an hypothesis test, and those satisfying some conventional criterion ($p < \\alpha$ with e.g. $\\alpha = 0.05$) may be chosen for publication. The more `significant' the result (i.e. the smaller $p$ is), the more likely it is to be chosen for publication: publication bias. These two biases act in the same direction but at different stages in the process, and are examples of what Francis Bacon, in his \\emph{Novum Organum}, described as the tendency to ``notice the events where they are fulfilled, but where they fail, though this happens much more often, neglect and pass them by'' \\citep{bacon1620}. The result is that the vastly greater number of null results (whether with strong or weak limits) will go largely unpublished, making it difficult to estimate the global significance of any individual detection.  The question that needs to be addressed is whether any given excess or deficit in a spectrum is unlikely to be a sampling fluctuation given a large number of spectra each with many resolution elements. The only systematic attempts to address this specific problem are those of \\citet{nandra07} and \\citet{longi06}, both of which describe surveys of narrow, shifted lines in samples of observations.  One can make an order of magnitude estimate of the number of unpublished non-detections using the following simple argument. Let us assume that $\\gsim 500$ spectra have been examined in the last few years and each has $\\gsim 50$ resolution elements, the number of independent spectral resolution elements that have been examined must be $\\gsim 2.5 \\times 10^4$. If the residuals (after fitting a suitable continuum model) in each of these is approximately Normally distributed, the expected numbers of fluctuations at $|z| \\ge 2\\sigma$, $3\\sigma$ and $4 \\sigma$ are $\\gsim 1137.5$, $\\gsim 67.5$ and $\\gsim 1.6$. The data in Table~\\ref{table} record $23$ detections with $|EW|/error > 2$ and $5$ with $|EW|/error > 3$, given that the quoted errors are the half widths of the $90$ per cent confidence interval, these might better correspond to $\\sim 3\\sigma$ and $\\sim 4\\sigma$ detections\\footnote{This is a crude approximation. Interval estimation and hypothesis testing are different statistical procedures and the relative uncertainty of the strength of a line will not in general be simply related to its significance in an hypothesis test.}, suggesting the expected number of large fluctuations is consistent with the amount of available data. Furthermore, it seems reasonable that lower significance features are less likely to be reported, and therefore  the fraction of unreported $\\sim 3\\sigma$ detection would be larger than the fraction of unreported $\\sim 4\\sigma$ detections, hence the lower than expected number of (reported) detections with $|EW|/error$ in the range $2-3$. Another way to consider this is to compare the significance of individual lines, estimated accounting for the number of spectral channels, to the number of examined spectra. For example, one of the best detections is in E1821+643, where \\citet{yaqoob05} estimate a $2-3\\sigma$ detection over the entire spectrum, but given $\\gsim 500$ spectra one would expect $\\gsim 22.7$ detections at better than $2 \\sigma$ and $\\gsim 1.3$ at better than $3 \\sigma$, so perhaps this is not unexpected. See \\citet{scargle00} for a discussion of methods to estimate the number of unpublished observations.  \\subsection{\\emph{Post hoc} and \\emph{a priori} arguments}  One cannot argue \\emph{post hoc} that there are particular properties of individual datasets that allow them to be considered in isolation\\footnote{The possible exceptions are the lensed BAL quasars, which are exceptional sources for which there is prior knowledge of high velocity outflows (from their rest-frame ultraviolet spectra). One might expect stronger absorption systems from these faint sources, and so they might reasonably be considered separately (they were not included in the calculations performed above).}. For example, in the case of E1821+643 one might argue that the $6.2$~keV line was detected in high resolution \\chandra\\ HETGS data and there are far fewer of these observations, therefore the number of (unpublished) non-detections is much smaller. One could conceivably construe a similar argument in favour of the uniqueness of almost any observation. But if $N$ observations are analysed and subjected to hypothesis tests with a detection threshold of $\\alpha$ (it does not matter if the tests are the same), the expected number of false detections is $\\sim N \\alpha$, even if the data come from different target sources, missions, detectors etc.  Similarly, it is not valid to argue that certain time intervals of a particular observation should be treated as special unless they are selected on the basis of an explicit, prior criterion (i.e. derived and used independently of line detection). The expected number of false detections scales with the number of tests performed, therefore if a long, high quality observation is split into ten time intervals and each is examined (if only in a cursory fashion), this has increased the effective number of tests approximately tenfold. Additionally, one cannot engineer the data slicing to maximise the detection of a line, based on the detection of the line in the same data finely sliced, and consider this a fair test. This would be to test a hypothesis suggested by the data, as if the data were independent; one may equally well roll a die fifty times, find the most frequent number to be rolled and then claim the die is biased because there is only a $1/6$ chance of that number being the most frequent.  The plausibility of a line detection would be augmented if there was cogent prior information on the line properties that was confirmed in subsequent observations. For example, a small excess in a spectrum at $6.4$~keV might be considered a significant detection of an iron line, but the same excess at an arbitrary energy might not be significant. The coincidence of an excess appearing at the predicted energy adds to its plausibility as a real line (and this can be included in formal calculations of its `significance'). However, as discussed above, the lines appear at arbitrary energies and are often reported to be transient which, if true, makes prediction difficult.  If real, these lines represent quite extreme physical phenomena. But, at least in the case of absorption lines, low outflow-velocity absorption systems are routinely seen in the X-ray spectra of Seyfert galaxies and quasars, and high velocity outflows are observed in the rest-frame ultraviolet spectra of BAL quasars. Based on these one may argue that the existence of higher velocity X-ray outflows in other (i.e. non-BAL) sources is at least plausible. By contrast the existence of narrow, redshifted absorption, and highly red- or blue-shifted emission lines is unprecedented, and it is reasonable to demand high standards of evidence to support their existence. It is perhaps worth noting that the $15$ emission lines from Table~\\ref{table} have a lower average $|EW|/error$ ratio than the absorption lines, and none of them has $|EW|/error > 2.5$ (the $5.6$~keV emission line in NGC 3516 is excepted for the reasons given above).  \\subsection{Final remarks}  Of course, new and exciting discoveries are usually made at the limits of the available data, but these must be confirmed at higher significance by subsequent observations; when more detections are made but the significance does not improve despite more, longer observations and brighter targets one should not automatically consider the discovery confirmed. It is not the objective of this paper to argue that any specific detection is false; the argument based on Fig.~\\ref{fig} is purely statistical. Indeed, several of the existing detections may be genuine -- the case is arguably strongest for absorption lines with the smallest $EW$ (e.g $EW \\lsim 30$~eV) and largest $\\Delta \\chi^2$ -- but it is difficult to explain the tightness of the correlation between detected line strength and its uncertainty if all or most of the detections are genuine. The prevalence and importance of such features in the population at large therefore remains an open question."
    },
    "0807/0807.4037_arXiv.txt": {
        "abstract": "DUNE (Dark Universe Explorer) is a proposed mission to measure parameters of dark energy using weak gravitational lensing The particular challenges of both optical and infrared focal planes and the DUNE baseline solution is discussed. The DUNE visible Focal Plane Array (VFP) consists of 36 large format red-sensitive CCDs, arranged in a 9x4 array together with the associated mechanical support structure and electronics processing chains. Four additional CCDs dedicated to attitude control measurements are located at the edge of the array. All CCDs are 4096 pixel red-enhanced e2v CCD203-82 devices with square 12 $\\mu$m pixels, operating from 550-920nm. Combining four rows of CCDs provides a total exposure time of 1500s. The VFP will be used in a closed-loop system by the spacecraft, which operates in a drift scan mode, in order to synchronize the scan and readout rates. The Near Infrared (NIR) FPA consists of a 5 x 12 mosaic of 60 Hawaii 2RG detector arrays from Teledyne, NIR bandpass filters for the wavelength bands Y, J, and H, the mechanical support structure, and the detector readout and signal processing electronics. The FPA is operated at a maximum temperature of 140 K for low dark current of 0.02e$-$/s. Each sensor chip assembly has 2048 x 2048 square pixels of 18 $\\mu$m size (0.15 arcsec), sensitive in the 0.8 to 1.7 $\\mu$m wavelength range. As the spacecraft is scanning the sky, the image motion on the NIR FPA is stabilized by a de-scanning mirror during the integration time of 300 s per detector. The total integration time of 1500 seconds is split among the three NIR wavelengths bands. DUNE has been proposed to ESA's Cosmic Vision program and has been jointly selected with SPACE for an ESA Assessment Phase which has led to the joint Euclid mission concept. ",
        "introduction": "\\label{sec:intro}  %   The Dark Universe Explorer (DUNE) is a proposed wide-field space imager whose primary goal is the study of dark energy and dark matter with unprecedented precision using weak gravitational lensing. Immediate secondary goals focus on the evolution of galaxies, to be studied with unprecedented statistical power, the detailed structure of the Milky Way and nearby galaxies, and the demographics of Earth-mass planets. DUNE is a medium-class mission as defined by ESA\u2019s Cosmic Vision program, consisting of a 1.2m Korsch-like three-mirror telescope with a combined visible/near-infrared field-of-view of 1 square degree and 0.23 arcsec resolution (visible) in a geosynchronous orbit. DUNE plans to carry out an all-sky survey in one visible and three NIR bands which will form a unique legacy for astronomy. DUNE thus addresses multiple goals of the ESA Cosmic Vision program and would yield major advances in a broad range of fields in astrophysics and cosmology (http://www.dune-mission.net). DUNE is a realization of the wide-field imaging mission recommended by the ESO/ESA Working Group on Fundamental Cosmology (WGFC).  One of the most powerful tools to tackle dark energy is weak gravitational lensing of background galaxies by foreground dark matter: this forms the core of the DUNE mission\\cite{ref08,refrdous}. Gravitational deflection of light by intervening dark matter concentrations causes the images of background galaxies to acquire an additional ellipticity on the order of one percent. Utilization of this cosmological probe relies on the measurement of image shapes and redshifts, or distances, for several billion galaxies, both requiring space observations for point spread function (PSF) stablity and photometric measurements over a wide wavelength range in the visible and especially near-infrared (NIR). These are the driving requirements on the instrumentation of DUNE, and are solved by the separate, large visible and NIR focal plane arrays which will be discussed in more detail in this paper.  \\subsection*{Mission and orbital requirements}  The driving requirements for the mission design of DUNE are from the wide extragalactic survey and specifically the need for the stability of the PSF and the large coverage of the sky. PSF stability puts stringent requirements on pointing and thermal stability during the observation time. The 20,000 square degrees survey demands high operational efficiency, which can generally be achieved either in a step-and-stare imaging mode or by continuously scanning the sky during science acquisition. The pointing accuracy requirement is identical in both cases, namely 0.2 $\\mu$rad over 375 s integration time per CCD chip (to reach magnitude limits required by science goals), and is essentially driven by the elementary integration time and the spatial resolution. This can be achieved using a drift scanning mode (or Time Delay Integration, TDI, mode) for the CCDs in the visible focal plane, as for GAIA. A continuous scanning mode is by nature more efficient than the step-and-stare mode, since the science measurements can be continuously achieved over a relatively long fraction of a great circle, which reduces the number of maneuvers by one or two orders of magnitude, and is selected as the baseline for DUNE.  For the infrared focal plane, using available HgCdTe arrays, TDI mode not feasible and the proposed concept is to stop the image motion during the integration time in the NIR focal plane by using a small de-scan mirror located close to the telescope exit pupil. In practice, the de-scan design is such that this mirror is operated in quasi-static mode with low amplitude oscillations, about 1 degree and a period of 300 seconds. ",
        "conclusions": ""
    },
    "0807/0807.0422_arXiv.txt": {
        "abstract": "The low-mass end of the supermassive black hole mass function is unknown and difficult to determine. Here we discuss our successful program to find active nuclei of late type ``normal'' galaxies using X-ray detections and multiwavelength identifications. We conclude that most of the \\chandra detected nuclear X-ray sources are AGNs. We then outline methods of black hole mass determination when broad emission lines are unobservable. ",
        "introduction": "The local low-mass (below $10^6$ \\msun) supermassive black hole (SMBH) mass function is not well known. Current estimates of the SMBH mass function are based on host galaxy properties (luminosity of the bulge or bulge stellar velocity dispersion $\\sigma$) and known scaling relationships between the mass $\\mbh$ of the SMBH and these properties (most prominently $\\mbh - \\sigma$ and $\\mbh - \\mbul$). These estimates can be widely discrepant with each other at the low-mass end (see, for example, Fig.~8 in \\cite{gdea07}). Much of the uncertainty arises because it is unknown how the scaling relationships extrapolate to very late-type and very low mass galaxies. Yet SMBHs \\emph{do} exist in very late-type spirals, e.g.\\ NGC 4395, a spiral galaxy of type Sdm, with $\\mbh \\sim 3\\times 10^5$ \\msun\\ \\citep{pea05}, and in very low mass galaxies, e.g.\\ POX 52, a dwarf galaxy, with $\\mbh \\sim 3\\times 10^5$ \\msun\\ \\citep{bea04}.   Most of the measured SMBH masses are $\\sim 10^8 $ \\msun\\ or greater, however, as the sphere of influence of a less massive SMBH is extremely hard to resolve even at moderate distances, even with the Hubble Space Telescope (\\hst). For example, the sphere of influence of a $10^6 $ \\msun\\ SMBH at 15 Mpc is $\\sim 30$ milliarcseconds. Since we cannot detect low-mass SMBHs by their dynamical signature, looking for them by signs of their accretion activity may be the only viable way of detecting them. An accreting low-mass SMBH should be a low-luminosity active galactic nucleus (LLAGN).  How can an LLAGN be identified? Broad emission lines in the optical spectrum allow an unambiguous confirmation of the presence of an AGN. The sample of low-mass SMBHs in \\citet{gh07a} was identified this way. However, as the luminosity of an AGN decreases, the optical spectrum of the galaxy nucleus becomes more and more dominated by host galaxy light, and the signature of the AGN becomes difficult to detect. For the lowest-luminosity AGNs, therefore, it is possible that a system based on optical spectra would not classify the nuclei as AGNs at all. Even when the optical spectrum shows no clear evidence of an AGN, however, such evidence may still be present in other wavelengths, such as x-ray and radio \\citep{fea04}, and infrared \\citep{dea06,svea07,sea08}. ",
        "conclusions": "\\subsection{How to measure BH masses}  Reverberation mapping of broad emission lines provides a standard method for measuring BH masses in AGNs. Alternatively, scaling relations using the width of broad emission lines and AGN luminosity are used. How should we measure BH masses in our newly discovered AGNs which do not show broad emission lines? This can be done in a variety of ways:  \\noindent 1. As noted above, host galaxy contamination is a major contributor toward hiding the AGN signature of our low-luminosity sources. Broad emission lines of these AGNs might have been swamped by the host galaxy light in ground-based spectra. Therefore it is imperative to obtain spectra of these sources with the highest possible spatial resolution from space, e.g. with STIS on HST. This might uncover broad emission lines.   \\noindent 2. Construction of power-density spectra (PDS) using X-ray variability analysis provides a powerful tool for BH mass determination \\cite{mpea04}. The break frequency in PDS is found to correlate strongly with the BH mass; since our sources are X-ray detected, this is potentially a highly feasible method.  \\noindent 3. The X-ray power-law slope correlates strongly with the Eddington luminosity ration \\cite{wmp04}. Measuring the X-ray slope and luminosity would then estimate the BH mass.  \\noindent 4. SMBH mass is known to correlate with the bulge velocity dispersion \\cite{fm00,gea00} and the bulge luminosity \\cite{kr95,md01,mh03}. These relations can be used for BH mass determination, once we ascertain the existence of a SMBH. This method, however, will not work for bulge-less galaxies. Moreover, if we do not have independent measurements of BH masses, we cannot extend these known correlations to low BH masses.  \\noindent 5. SMBH mass is also known to correlate with the circular velocity of the host galaxy \\cite{f02,bea03}. This method will have the same caveat mentioned above, though it works for bulge-less galaxies as well.  Thus there are several ways to estimate masses of SMBHs which we uncover using X-ray detections. This is obviously a hard job; at this point, however, we are at the step of finding these low-mass SMBHs in the first place.  \\subsection{What's Next?}  In order to make progress in finding and identifying low-mass SMBHs and measuring their masses several steps need to be taken.  \\noindent 1. Through our \\chandra survey and the accompanying archival programs we have realized that traditional diagnostics for identifying AGNs fail for low-luminosity AGNs embedded in luminous galaxies. We need to develop new diagnostic tools using multi-wavelength data for this purpose.  \\noindent 2. Deep X-ray spectra of likely AGNs. This will help secure AGN identification. We have obtained XMM-{\\it Newton} data on one of our candidate.  \\noindent 3. HST STIS spectroscopy. As noted above, this will uncover broad emission lines if any and possibly change the narrow emission line flux ratios.  \\noindent 4. HST ACS imaging. This will help obtain true colors of the AGN, after removing a substantial contribution of the host galaxy. We can them identify an AGN with traditional ``UV excess'' and/or with broad-band spectral energy distribution.  \\noindent 5. IR diagnostics: {\\it Spitzer} IRAC colors. While AGNs are blue in the optical, they are red in mid-IR. AGNs occupy a distinct location in the IRAC color-color plot, allowing identification of X-ray sources \\cite{skea05}.  \\noindent 6. IR diagnostics: near IR spectra. High ionization IR emission lines with IRS on {\\it Spitzer} have been used to uncover AGNs by \\citet{sea08}. With the end of {\\it Spitzer} cryogenic mission, IRS is no longer available. Finding near-IR ``coronal'' lines might prove useful for the same purpose.  \\noindent 7. Radio observations. Finding a radio point source coincident with the X-ray point source will help secure the identification. VLA observations of our candidate AGNs is an obvious next step.  \\noindent 8. Bigger \\chandra/{\\it Spitzer} survey will provide better statistics on fraction of galaxies hosting AGNs. A deeper survey will allow us to determine AGN duty cycle to lower luminosities.  \\noindent 9. Hard X-ray observations. A survey of nearby galaxies in hard X-rays will detect and identify nuclear SMBHs (as well as off-nuclear accreting BHs). The Indian X-ray mission ASTROSAT, which has hard X-ray capabilities is due to launch in 2009. ASTROSAT observations will be invaluable in our BH quest.      \\begin{theacknowledgments} This work is supported in part by the National Aeronautics and Space Administration through Chandra Award Number GO7-8111X 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.  \\end{theacknowledgments}"
    },
    "0807/0807.0108_arXiv.txt": {
        "abstract": "We present Monte Carlo simulations on Eley-Rideal abstraction reactions of atomic hydrogen chemisorbed on graphite. The results are obtained via a hybrid approach using energy barriers derived from DFT calculations as input to Monte Carlo simulations. By comparing with experimental data we discriminate between contributions from different Eley-Rideal mechanisms. A combination of two different mechanisms yields, good quantitative and qualitative agreement between the experimentally derived and the simulated Eley-Rideal abstraction cross sections and surface configurations. These two mechanisms include a direct Eley-Rideal reaction with fast diffusing H atoms and a dimer mediated Eley-Rideal mechanism with increased cross section at low coverage. Such a dimer mediated Eley-Rideal mechanism has not previously been proposed and serves as an alternative explanation to the steering behavior often given as the cause of the coverage dependence observed in Eley-Rideal reaction cross sections. ",
        "introduction": "Molecular hydrogen is the most abundant molecule in the interstellar medium (ISM), where it serves as an important coolant and as a precursor for the formation of more complex molecules. In the cold and dilute interstellar medium no efficient gas phase routes exist for the formation of H$_2$. Hence, the most likely route is to form the molecule on the surfaces of small dust particles. This process has been intensely studied over the past years \\cite{Pirronello:1997,Pirronello:1997a, Pirronello:1999, Hornekaer:2003, Hornekaer:2005, Perry:2003, Cazaux:2004, Chang:2005, Cuppen:heating, Perets:2005, Katz:1999, Dulieu:2005}. The general consensus is that molecular hydrogen can be formed by a diffusive Langmuir-Hinshelwood mechanism, in the very cold regions below $\\sim$20 K \\cite{Hollenbach:1971}. At these low temperatures the hydrogen atoms are physisorbed on the dust particles, but can still easily move on the surface. Once two atoms meet, they will react to form molecular hydrogen. If the surface temperature becomes too high, the residence time of atoms on the surface becomes so short that, at the low fluxes in the ISM, the chance of two atoms meeting becomes negligible. Observations show that molecular hydrogen is also formed in warmer areas ($>$ 20 K) like Photon Dominated Regions (PDRs) and post-shock regions. A possible mechanism at these conditions would be through an Eley-Rideal reaction where an incoming hydrogen atom reacts with another H atom that is chemically bound to the surface. Since a considerable fraction of the interstellar grains is expected to consist of carbonaceous material, a popular model system for this process is hydrogen abstraction from graphite. This Eley-Rideal reaction has been studied by Density Functional Theory (DFT) and Quantum Wave packet calculations \\cite{Farebrother:2000, Meijer:2001, Sha:2002, Sha:2002a, Morisset:2004, Martinazzo:2006}. Based on these calculations several different mechanisms have been proposed to contribute to the Eley-Rideal abstraction process. These include: Direct Eley-Rideal \\cite{Farebrother:2000, Meijer:2001, Sha:2002, Sha:2002a, Morisset:2004, Martinazzo:2006}, barrier-less abstraction of one hydrogen atom forming part of a para-dimer configuration \\cite{Bachellerie:2007} and abstraction by rapidly diffusing H atoms in physisorbed states \\cite{Bonfanti:2007}.  However, due to computational limitations such detailed calculations are unable to take the full complexity of the H-graphite system into account. Hence, in the calculations the flat surface approximation is often used and generally either none or only a single C atom on the simulated graphite surface were allowed to relax during the Eley-Rideal reaction. These simplifications are potentially problematic. By not allowing for relaxation of the C atoms on the surface zero barrier reaction channels, such as sticking of hydrogen atoms into specific dimer configuration \\cite{Hornekaer:2006II,Rougeau:2006}, are in some cases artificially closed, making it impossible to evaluate the contributions from different proposed mechanisms to the Eley-Rideal abstraction process. Eg. calculations on the Eley-Rideal reaction by Martinazzo and Tantardini \\cite{Martinazzo:2006} show a higher probability for reaction if the incoming H atom is 0.5 - 1.5 \\AA~away from the target H atom instead of a direct hit on. This distance corresponds to the H-H distance in an ortho-dimer configuration (see Fig. 1) on the graphite surface (1.42 \\AA). Hence, surface corrugation and competing sticking reactions must be expected to influence the abstraction behavior. In these calculations, however, dimer formation was not a possible route since the flat surface approximation was used and only the carbon atom underneath the original hydrogen atom was allowed to relax. Impact parameters corresponding to the para-dimer (see Fig. 1) distance (2.84 \\AA) were not included in the study.  Both experimental observations and theoretical calculations show that hydrogen atoms preferentially form adsorbate clusters on the graphite surface. Scanning Tunneling Microscopy (STM) investigations show that at a hydrogen atom coverage of $\\sim$ 0.5 $ \\%$ more than 75 $\\%$ of all surface configurations are clusters \\cite{Hornekaer:2006II}. This indicates that $\\sim$ 85 $\\%$ of the hydrogen atoms are part of larger clusters. In particular, the existence of hydrogen dimer configurations has been studied theoretically \\cite{Ferro:2003,Miura:2003,Hornekaer:2006I,Hornekaer:2006II,Rougeau:2006} and experimentally \\cite{Hornekaer:2006I, Andree2006}. Hornek{\\ae}r et al. \\cite{Hornekaer:2006I} identified two stable hydrogen dimer configurations on the graphite surface, an ortho-dimer and a para-dimer (see Fig.~\\ref{dimers}). DFT calculations show that formation of the para-dimer is barrierless and that formation of the ortho-dimer has a reduced barrier \\cite{Hornekaer:2006II}, which makes dimer formation a competing channel to Eley-Rideal abstraction at non-zero impact parameter.  One set of experiments on the Eley-Rideal abstraction reaction forming HD has been reported \\cite{Zecho:2002, Zecho:2002a}. In these experiments a high cross section for the abstraction reaction was observed which varied from 17 \\AA$^2$ at low coverage to 4 \\AA $^2$ at high coverage. This variation in the cross section with coverage has been ascribed to a steering effect \\cite{Sha:2002, Sha:2002a}. However, other mechanisms could also contribute to the high cross section at low coverage.  Options include the barrier-less abstraction of one of the H atoms in a para-dimer configuration \\cite{Bachellerie:2007} and abstraction by fast diffusing H atoms \\cite{Bonfanti:2007}. We propose a further possibility, namely abstraction via a hydrogen dimer state, where the incoming atom is not immediately thermalized and some of its excess initial energy is used to overcome the barrier to H$_2$ formation and desorption. The finding that adsorption into the hydrogen para-dimer state is barrierless \\cite{Hornekaer:2006II} make this a strong competing reaction channel.  Hence, in total 5 different abstraction mechanisms have been proposed to contribute to Eley-Rideal abstraction of hydrogen on graphite:  \\begin{enumerate} \\item direct Eley-Rideal \\cite{Farebrother:2000,Morisset:2004}. \\item direct Eley-Rideal with steering \\cite{Zecho:2002,Sha:2002,Sha:2002a}. \\item preferred direct Eley-Rideal with hydrogen atoms part of a para-dimer \\cite{Bachellerie:2007}. \\item a dimer mediated reaction where the incoming atom is first adsorbed into a dimer configuration and, before thermalizing to the substrate temperature, reacts with the other atom in the dimer to form H$_2$ and desorb. \\item Eley-Rideal reactions by fast diffusing H atoms in the physisorption state \\cite{Bonfanti:2007}. \\end{enumerate}  However, due to the simplifications needed in the complex DFT and quantum wave packet calculations a quantitative comparison between experimental data and theory has not been possible. In this paper we employ a hybrid approach in which we include the findings of ab-initio DFT calculations in a Monte Carlo simulation program and then simulate Eley-Rideal abstraction experiments on a more realistic surface area with more complex hydrogen adsorbate configurations than what is possible in the DFT and wave packet calculations. Through this quantitative approach we aim to discriminate between the contributions of the different underlying mechanisms for hydrogen abstraction.  \\begin{figure} \\begin{center} \\includegraphics{802819jcp1.eps} \\end{center} \\caption{The two stable dimer configurations found by Hornek{\\ae}r et al. \\cite{Hornekaer:2006I}.} \\label{dimers} \\end{figure} ",
        "conclusions": "The Eley-Rideal abstraction of atomic hydrogen chemisorbed on the graphite surface has been studied via a hybrid approach using energy barriers derived from DFT calculations as input to Monte Carlo simulations. Through comparison with experimental data we discriminate between the contributions from different proposed Eley-Rideal mechanisms. Good quantitative and qualitative agreement between the experimentally derived and the simulated Eley-Rideal abstraction cross sections are found if two different Eley-Rideal abstraction mechanisms are included. One is a direct Eley-Rideal reaction with very fast diffusion of physisorbed H atoms leading to the formation of hydrogen dimer configurations, while the other is a dimer mediated Eley-Rideal mechanism with increased cross section at low coverage. Such a dimer mediated Eley-Rideal mechanism has not previously been proposed.  The effect on abstraction of fast diffusing physisorbed H atoms was first considered by Bonfanti et al. \\cite{Bonfanti:2007}, who suggested that diffusion of the physisorbed atoms could explain the high coverage dependence of the cross section. We tested this mechanism and it was indeed found that diffusion plays an important role, not only for the abstraction reaction but also to explain the high occurrence of dimers found by STM measurements. However, very little is known about the exact energetic landscape in the vicinity of a chemisorbed hydrogen atom, especially when the second atom is only weakly physisorbed. Since these weak interactions appear to be very important, further study would be desirable.  Furthermore, we investigated the effect of steering \\cite{Sha:2002,Sha:2002a} as a possibly alternative route to reproduce the experimental findings. The results presented here show that a simple steering mechanism which just results in increased cross sections for abstraction reactions is not sufficient to reproduce the experimental results. This finding does, however, not rule out steering as an important mechanism. It has for example been suggested that a coverage dependent steering effect could result from the creation of and interaction with electron-hole pairs on the surface \\cite{Hammer-Private-communications}. Further investigations of steering should, however, consider the presence of preferred binding sites in the vicinity of adsorbed hydrogen atoms. In the present case, sticking in a dimer position with low or no barrier is a competing reaction to the steered Eley-Rideal reaction. Hence, the described dimer mediated Eley-Rideal mechanism offers a competing mechanism to steering. Full quantum mechanical calculations where both channels, dimer formation and hydrogen abstraction, are accessible, are needed to settle this case."
    },
    "0807/0807.5085.txt": {
        "abstract": "Recently measured VHE gamma ray spectra from supernovae remnants (SNRs) are best fit by power laws with an exponential cutoff. But this feature does not occur at the `knee-equivalent' energy at which VHE gamma ray spectra are expected to reflect the `knee' in the otherwise featureless cosmic ray spectrum. In this article, the VHE gamma ray feature is explained as a consequence of a recently deduced quantum effect. The effect distinguishes `trajectory energy' from `particle state energy' and requires that the particle state energy depends strongly on gravitational potential at very high energies. Based on this effect and the observed CR spectrum, a tight two-parameter fit is obtained to the combined VHE gamma ray spectra of SNRs RX J1713.7-3946 and RX J0852.0-4622 and the Galactic Center ridge.  {\\textbf{Key words.}} cosmic rays, gamma rays: theory , elementary particles, gravitation,  supernovae: general , Poincar\\'{e} symmetry ",
        "introduction": "\\label{intro}  It is widely held that galactic cosmic rays are mostly atomic nuclei accelerated in supernovae remnants, that some CRs collide with local ambient matter producing gamma rays, and that CRs, being charged, are deflected as they travel to Earth. Being uncharged, the gamma rays that arrive at Earth retain directional information pointing back to their sources.  Since the galactic CR primaries have been deflected, the CR spectrum observed at Earth, see Fig. \\ref{CRearth}, represents an amalgam of SNR sources. The spectrum follows a near-constant power law from $10^{10}$ eV to the `knee' at $3\\times 10^{15}$ eV, where the spectral index of the power law steepens and is again constant to the `ankle' at $10^{18}$ eV. The `ultrahigh energy' CRs do not concern us here. The most energetic VHE gamma rays thus far observed have energies $10^{14}$ eV or so, implying that the relevant CR energies are below about $10^{15}$ eV.  Considering only proton primaries, CR protons at the SNR are expected to produce VHE gamma rays when the protons collide with ambient matter, most likely protons. The $p$-$p$ collisions create pions and etas that decay to produce the VHE gamma rays. The process is well understood and the only quantity that could provide the VHE gamma ray spectra with a distinct feature is the inelastic $p$-$p$ cross section $\\sigma_{inel}.$  However the log of $\\sigma_{inel}$ is a slowly rising function of incident proton energy and is remarkably featureless, see Fig. \\ref{sigmadEdE}a. Thus one expects the VHE gamma ray spectra from CR protons to be a power law up to some `knee equivalent' gamma ray energy, and at higher energies to again be a power law but with a probably steeper index. In short, the VHE gamma ray spectrum should be similar to the CR spectrum and a typical VHE spectrum should feature a knee equivalent somewhere around $0.17\\times 3 \\times 10^{15}$ =  $5.1\\times 10^{14}$ eV.  The H.E.S.S. collaboration reports a cutoff in the spectrum for the supernova remnant RX J1713.7-3946 (Aharonian et al. 2007a) at an energy of at most $1.8 \\times 10^{13}$ eV, a factor of about 30 below the expected knee equivalent. Therefore, it is thought that the feature reported is not the knee equivalent (Plaga 2008a). One must consider other explanations. It could be that the VHE gamma rays are emitted by electrons in the `lepton scenario' since electrons are also believed to be accelerated in supernovae remnants and electrons in certain conditions emit VHE gamma rays. For lepton scenario discussions, see Huang et al (2007), Katz \\& Waxman (2008), and Plaga (2008b).  It is shown in this paper that the hadron scenario fits the VHE gamma ray feature at $10^{13}$ eV if we apply the result of a recent calculation. The explanation distinguishes two different kinds of energy that are equal at low energies and diverge at very high energies. The `particle state energy' depends on gravitational potential, with less energy at the SNR and more energy at Earth (Shurtleff 2008). A summary of the calculation is included in Sec.~\\ref{PDPSE} for convenience.  Then a massive particle such as a proton has two energies, one that can be obtained by mass and a time-of-flight velocity, a `trajectory energy'. The other is the energy delivered in a collision due to the change in the particle's state. Call that the `particle state energy'.   The formula for the particle state energy $E$ is \\begin{equation} \\label{Eintro} E = E_{0} -4 \\phi \\frac{E_{0}^{3}}{m^2 c^4} \\, , \\end{equation} where $\\phi < 0$ is the gravitational potential and  $E_{0}$ is the energy at null potential $\\phi$ = 0. The gravitational field is assumed to be weak and the potential $\\phi$ is the Newtonian potential made unitless by dividing by the square of the speed of light $c.$ The mass of the particle is $m.$  The Solar System moves in the gravitational field of the Galaxy. Very roughly, the Galaxy should contribute a potential $\\phi_{Gal}$ due to $M_{Gal}$ = $10^{11}$ solar masses in effect collected some $R$ = $2.5\\times10^{20}$ m away from the Solar System at the Galactic Center. One has $\\phi_{Gal}$ = $-GM_{Gal}/Rc^{2} \\approx$ $-6\\times 10^{-7}.$  It follows that the effects of Eq. (\\ref{Eintro}) would be significant at about $E_{0} $ = $mc^{2}/2\\sqrt{-\\phi_{Gal}}$ = 0.6 TeV, where $mc^{2} \\approx$ 0.001 TeV is the rest energy of a proton. Since 0.6 TeV is well within the range of proton energies available at the Tevatron, and no such effects are reported from the Tevatron, the highest energy proton accelerator as of this writing, it follows that the potential due to the Galaxy and any larger potentials must be eliminated from consideration if Eq. (\\ref{Eintro}) is true for protons.  To eliminate the effects of the Galaxy's gravitational field, one may choose to work in a comoving coordinate system in free fall with the Solar System, see for example Weinberg (1972). In General Relativity, a particle's trajectory can be described in any coordinate system.  By making the particle state energy dependent on potential and requiring the potential to be calculated in a system in free fall with the Solar System, the description is valid in one coordinate system, the reference frame of the Solar System. This is allowed in General Relativity as long as the various quantities are properly transformed when the description is given in some other coordinate system.  While the coordinate system is in free fall with the Solar System, thereby canceling the Galactic and larger scale potentials, the local gravitational potential due to the matter in the Solar System does not cancel. Thus, at the Earth's surface, the $\\phi$ in (\\ref{Eintro}) is the potential due to the Sun and the Earth, a value of about $\\phi_{\\oplus} \\approx$ $-1\\times 10^{-8},$ which becomes significant at an energy $E_{0}\\approx$ 5 TeV. This is beyond the proton energies available at proton accelerators at the time of this writing, but is within the reach of the LHC which has not begun high energy runs at the time of this writing.  While the properties of electrons have no bearing on the explanation of the feature in the VHE gamma ray spectrum, one may note the energy at which the effects of (\\ref{Eintro}) would become significant for electrons is $E_{0} $ = $m_{e}c^{2}/2\\sqrt{-\\phi_{\\oplus}}$ = 2.5 GeV. Since the effects expected by applying (\\ref{Eintro}) have not been reported by accelerator experiments with electron energies many times the threshold 2.5 GeV, it follows that electrons do not obey (\\ref{Eintro}), at least for potentials $\\phi$ with magnitudes as large as the potential due to the Earth and Sun. This could mean the comoving coordinate system for electrons is in free fall with matter falling in the gravitational field due to the Earth and Sun (and due to the Galaxy and larger scales as well). Or it could mean that the assumptions that produce (\\ref{Eintro}) do not apply to electrons. This could be another way that electrons differ from protons.  One benefit of going to a system in free fall with the local matter is the resulting universality possible for the various SNRs. If the potential at each SNR includes the contribution from the Galaxy, then Eq. (\\ref{Eintro}) would have a different $\\phi_{S}$ for each supernova remnant. That would be a troubling problem. But if the potential $\\phi_{S}$ is the local potential in the system in free fall with the SNR, then the potentials at the various SNRs are likely to be nearly the same and Eq. (\\ref{Eintro}) can meaningfully be applied to SNRs anywhere in the Galaxy.  In this picture, to get a cosmic ray from the SNR to Earth, one must consider a number of overlapping patches of spacetime, each patch described by a coordinate system in free fall with the local matter at a scale comparable to objects like the Solar System or supernova remnants.   With Eq. (\\ref{Eintro}), a cosmic ray proton observed at Earth with energy $Ep_{\\oplus}$ in the potential $\\phi_{\\oplus}$ at the Earth has a different energy $Ep_{S}$ at the supernova in the potential $\\phi_{S}.$ In this notation, `$p$' stands for proton and `$\\oplus$' for Earth and `$S$' for supernova remnant.  Relating the CR spectrum at Earth,  $dN_{p}/dEp_{\\oplus},$ to the CR spectrum at a typical SNR, $dN_{p}/dEp_{S},$ involves the derivative $dEp_{\\oplus}/dEp_{S}$ relating the energy at the SNR and the corresponding energy at Earth. It is suggested here that this energy dependence is part of the feature found by the H.E.S.S. collaboration in the VHE spectrum of SNR RX J1713.7-3946.  Having obtained the expected CR spectrum at a typical SNR based on the CR spectrum at Earth, the expected VHE gamma ray spectrum at the SNR can be determined by the analysis by Kelner et al. (2006). Then the expected VHE gamma ray spectrum at Earth can be found if we know how it is related to the VHE gamma ray spectrum at the SNR.  The problem is how to relate the VHE gamma ray spectrum at the SNR to the VHE gamma ray spectrum at Earth. Since the basic equation (\\ref{Eintro}) applies to massive particles, one may entertain various arguments to get it to apply to massless gamma rays.  First of all, note that at 10 TeV, say, a proton's rest energy is just 0.01\\% of its total energy. The proton moves very much like a massless photon of the same energy. So a 10 TeV VHE gamma ray should behave much like a 10 TeV proton. Since the basis of the effect separates this `trajectory energy' from the `particle state energy', this argument may be spurious, even though it is intuitively appealing.  Secondly, consider the process of neutral pion decay that produces many of the VHE gamma rays in the hadronic scenario. A 10 TeV neutral pion decays to a 10 TeV VHE gamma ray plus a 4 keV X-ray. Within roundoff, the pion and the gamma ray each have an energy of 10 TeV. By time inversion, starting with a 10 TeV VHE gamma ray, an interaction with a 4 keV X-ray could produce a 10 TeV neutral pion. However unlikely the process, it could happen, so a 10 TeV VHE gamma ray is virtually a 10 TeV pion. Since the virtual transition from VHE gamma ray to neutral pion could occur anywhere from the SNR to Earth and a neutral pion traveling from the SNR to Earth would obey Eq. (\\ref{Eintro}), the energy of the VHE gamma should obey Eq. (\\ref{Eintro}) with $m$ as the neutral pion mass $m_{\\pi^{0}}.$  Imagine placing a shuttered X-ray source at each kilometer along the path of a VHE gamma ray from an SNR to Earth. Furthermore let there be a limitless supply of identical VHE gamma rays ready to be sent from the SNR to Earth one by one. Let the $n$ = 0 VHE gamma ray travel from the SNR to Earth without removing any of the shutters. This VHE gamma ray reaches Earth unchanged by the presence of the X-ray sources. For the next VHE gamma ray, i.e. $n$ = 1, release X-rays at the first kilometer at the SNR so that a neutral pion is created by inverse pion decay. The neutral pion has almost the same energy as the VHE gamma ray. For the $n$th VHE gamma ray release X-rays from just the one source at the $n$th kilometer making an $n$th neutral pion which has almost the same energy as the $n$th VHE gamma ray at the $n$th kilometer. Since the $n$ = 0 VHE gamma ray has the same energy as the $n$th VHE gamma ray just before the $n$th kilometer, and the $n$th VHE gamma ray has nearly the same energy as a neutral pion at the $n$th kilometer, and a neutral pion obeys (\\ref{Eintro}), it follows that the VHE gamma ray energy should also obey Eq. (\\ref{Eintro}) with $m$ equal to the mass of a neutral pion.  It should be clear that the result does not depend on there being any real X-rays along the path of the VHE gamma ray from the SNR to Earth. It should also be clear that none of this implies that a VHE gamma ray has mass, only that its energy could be transferred essentially undiminished to a neutral pion in some possibly imaginary circumstances.  But there are other ways a 10 TeV VHE gamma ray could interact with ambient matter of negligible energy, perhaps making the VHE gamma ray virtually an electron or a positron or an eta meson, etc. Again, there need not be any real ambient matter, imaginary ambient matter suffices. Not knowing how to calculate the likelihood of independent imaginary interactions means not knowing what mass to associate with the VHE gamma ray. Thus we let that mass be a free parameter.  By these arguments, we apply Eq. (\\ref{Eintro}) to VHE gamma rays by letting the mass $m$ = $m_{ray}$ stand for the average mass of the masses of massive particles that could be given the same energy as the VHE gamma ray by some possibly imaginary process. The mass $m_{ray}$ is a free parameter found by fitting the experimental VHE gamma ray spectra.  Allowing (\\ref{Eintro}) to apply to the VHE gamma rays allows us to use the expected VHE gamma ray flux at the SNR, $dN_{\\gamma}/dE\\gamma_{S},$  to determine the expected VHE gamma ray flux at Earth, $dN_{\\gamma}/dE\\gamma_{\\oplus}.$ The relationship involves the derivative $dE\\gamma_{S}/dE\\gamma_{\\oplus}$ which is energy dependent.  As discussed previously, one also has the derivative $dEp_{\\oplus}/dEp_{S}$ from the CR spectrum. The product of the two factors, $dEp_{\\oplus}/dEp_{S}$ times $dE\\gamma_{S}/dE\\gamma_{\\oplus},$ is an energy dependent quantity, see Fig. \\ref{sigmadEdE}b, that could describe the feature in the VHE gamma ray spectrum of the SNR RX J1713.7-3946 reported by the H.E.S.S. collaboration. ",
        "conclusions": "\\label{discus}  The explanation of the VHE gamma ray spectra of typical SNRs presented in this paper is based on the dependence of particle state energy on gravitational potential as displayed in (\\ref{Eintro}). A proton or a gamma ray in the near-null potential at an SNR has less energy than it has in the gravitational potential at the Earth. The quality of the fit to the data in Fig. \\ref{bestfit} supports this explanation.  The quality of the fit supports the application of formula (\\ref{Eintro}) to photons even though photons are massless and (\\ref{Eintro}) applies to massive particles. It is argued that the photon should have particle states with energies almost the same as the energies of various massive particle states. The coincidence of photon particle state energy with massive particle state energy is supported by the VHE gamma ray spectra.  For a given source of VHE gamma rays, the distance to the source and the number of CRs each contributes to the CR spectrum at Earth are not knowable to sufficient accuracy to be useful in setting the absolute normalization of the VHE gamma ray spectra. One might include as free parameters the $\\Delta_{i}$s needed to adjust the absolute normalizations of the three VHE gamma ray spectra.  One of the two parameters needed to obtain an expected VHE gamma ray spectrum from the observed CR spectrum at Earth is the mass $m_{ray}.$ Since the best fit mass $m_{ray}$ is less than the neutral pion mass, there must be some interaction that allows a VHE gamma ray to transfer almost all its energy to a particle with less mass than the neutral pion. Keep in mind that transferring energy to electrons and positrons does not contribute to $m_{ray}$ because they do not obey (\\ref{Eintro}).  The best fit gravitational potential at the SNR is about ten percent of the potential at the Earth's surface. Using this to imply something about the density of, the dimensions of, and the distance to the clouds of ambient matter that produce the VHE gamma rays is beyond the scope of this article.  Consider Fig. 7. Draw a straight line on the graph from a point near 0.3 TeV to a point near 30 TeV, on the graph this is from $\\log{E\\gamma_{\\oplus}}$ = 11.5 to 13.5. Note the `bump' in the VHE gamma ray spectra. The same bump is seen in all three observed spectra which implies that some common mechanism is in place to produce the bump. That the same bump can be replicated by the best fit values of $\\phi_{S}$ and $m_{ray}$ shows how well the explanation relates the CR spectrum at Earth with the VHE gamma ray spectrum at Earth.  The tightness of the fit in Fig. \\ref{bestfit} suggests that the feature found in the VHE gamma ray spectrum of SNR RX J1713.7-3946 and the other SNRs plotted in Fig. \\ref{bestfit} exhibits particle behavior consistent with the simple expression (\\ref{Eintro}) of a gravitational potential-dependent relationship between CR and VHE gamma ray energies at Earth and the energies at the SNR.     \\begin{table*}[!ht] \\caption{There are low goodness-of-fit values $g(\\phi_{S},m_{ray})$ for proposed VHE gamma ray spectra with $m_{ray}$ from 0.35 to $0.55m_{\\pi^{0}}$ and with $\\phi_{S}$ from 0.09 to $0.13 \\phi_{\\oplus},$ where $m_{ray}$ is the average mass of possible gamma ray produced particles and $\\phi_{S}$ is the gravitational potential at the SNR. The constant $\\phi_{\\oplus}$  is the potential at the Earth's surface due to the Earth and Sun referenced to zero potential at infinite distance, see (\\ref{phi}). The $\\Delta$s adjust the absolute normalizations of the three observed spectra to each proposed spectrum.  \\vspace{0.5cm} } \\label{table:Fits} \\centering \\begin{tabular*}{\\textwidth}{@{\\extracolsep{\\fill}}|c|c|c|c|c|c|c|} \\hline % \\# &  $g(\\phi_{S},m_{ray})$ & $m_{ray}/m_{\\pi^{0}}$  & $\\phi_{S}/\\phi_{\\oplus}$   & $\\Delta_{1}$  & $\\Delta_{2}$  & $\\Delta_{3}$ \\\\\\hline\\hline % 1    &  0.074  &  0.45  &  0.11  &    1.75  &  1.62  &  0.91  \\\\\\hline % 2    &  0.076  &  0.35  &  0.11  &    1.67  &  1.54  &  0.83  \\\\\\hline % 3    &  0.077  &  0.45  &  0.09  &    1.72  &  1.60  &  0.89  \\\\\\hline % 4    &  0.078  &  0.55  &  0.11  &    1.81  &  1.69  &  0.97  \\\\\\hline % 5    &  0.078  &  0.45  &  0.13  &    1.78  &  1.65  &  0.93  \\\\\\hline % 6    &  0.079  &  0.55  &  0.09  &   1.79  &  1.67  &  0.96  \\\\\\hline % 7    &  0.080  &  0.35  &  0.09  &    1.63  &  1.50  &  0.80  \\\\\\hline % 8    &  0.080  &  0.35  &  0.13  &   1.71  &  1.57  &  0.85  \\\\\\hline % 9    &  0.082  &  0.55  &  0.13  &   1.83  &  1.71  &  0.99  \\\\\\hline % \\end{tabular*} \\end{table*}   \\pagebreak \\clearpage \\begin{table*}[!h] \\caption{The particle state energies of 11 CR protons in the gravitational potential of the Earth ($\\oplus$), in a null potential ($0$), and in the potential at an SNR ($S$), providing examples of Eq. (\\ref{Eintro}) which show how particle state energy depends on gravitational potential. For the potential at the SNR, $\\phi_{S},$ the best fit value from Table \\ref{table:Fits} is used. \\vspace{0.5cm}} \\label{table:CRenergy} \\centering \\begin{tabular*}{0.5\\textwidth}{@{\\extracolsep{\\fill}}|c|c|c|c|} \\hline % \\# & $Ep_{\\oplus}$  (TeV) & $Ep_{0}$  (TeV) & $Ep_{S}$  (TeV)  \\\\\\hline\\hline % 1. &  1.00 &  0.96 &  0.96    \\\\\\hline % 2. &  2.00 &  1.74 &  1.77   \\\\\\hline % 3. &  4.00 &  2.87 &  2.99    \\\\\\hline % 4. &  8.00 &  4.27 &  4.68    \\\\\\hline % 5. &  16.0 &  5.94 &  7.05    \\\\\\hline % 6. &  32.0 &  7.94 &  10.6    \\\\\\hline % 7. &  64.0 &  10.4 &  16.3   \\\\\\hline % 8. &  128. &  13.4 &  26.0  \\\\\\hline % 9. &  256. &  17.1 &  43.4    \\\\\\hline % 10. &  512. &  21.7 &  75.6    \\\\\\hline % 11. &  1024. &  27.5 &  137.    \\\\\\hline % \\end{tabular*} \\end{table*}  \\begin{table*}[!ht] \\caption{The energies of 11 VHE gamma rays in the gravitational potential of the Earth ($\\oplus$), in a null potential ($0$), and in the potential at an SNR ($S$). While (\\ref{Eintro}) only applies to massive particles, it is applied here for massless gamma rays by considering virtual though possible processes that produce massive particles with essentially the same energies as the gamma rays. The energy dependence of a gamma ray should average the energy dependences of these virtual particles. There are many possible processes, i.e. inverse pion decay, the Compton effect, pair production, etc. An average mass called $m_{ray}$ is used in (\\ref{Eintro}) and is varied to fit the observed data. The best fit values of $m_{ray}$ and $\\phi_{S}$ in Table \\ref{table:Fits} are used with (\\ref{Eintro}) to obtain the values of $E\\gamma_{0}$ and $E\\gamma_{\\oplus}$ in this table.  \\vspace{0.5cm}} \\label{table:VHEenergy} \\centering \\begin{tabular*}{0.5\\textwidth}{@{\\extracolsep{\\fill}}|c|c|c|c|} \\hline % \\# & $E\\gamma_{S}$  (TeV) & $E\\gamma_{0}$  (TeV) & $E\\gamma_{\\oplus}$  (TeV) \\\\\\hline\\hline % 1. &    0.164 &  0.159 &  0.204   \\\\\\hline % 2. &    0.301 &  0.275 &  0.514   \\\\\\hline % 3. &   0.509 &  0.417 &  1.25   \\\\\\hline % 4. &   0.795 &  0.566 &  2.65    \\\\\\hline % 5. &    1.20 &  0.722 &  5.04   \\\\\\hline % 6. &    1.80 &  0.895 &  9.12   \\\\\\hline % 7. &    2.77 &  1.10 &  16.3   \\\\\\hline % 8. &   4.42 &  1.35 &  29.3   \\\\\\hline % 9. &    7.35 &  1.66 &  53.6   \\\\\\hline % 10. &    12.9 &  2.05 &  100.   \\\\\\hline % 11. &   23.3 &  2.54 &  191.   \\\\\\hline % \\end{tabular*} \\end{table*}  \\pagebreak \\clearpage \\begin{table*}[!h] \\caption{The CR and VHE gamma ray fluxes at Earth and at the typical SNR for the CRs and gamma rays in Tables \\ref{table:CRenergy} and \\ref{table:VHEenergy}.  The CR flux at Earth, $dNp/dEp_{\\oplus},$ represents experimental data. The CR flux at the SNR follows from the CR flux at Earth and particle state energy dependence, (\\ref{dNpdEpSNR}). The gamma ray flux at the SNR is the result of integrating over the CR proton energy at the SNR using the method of Kelner et al. (2006). The gamma ray flux at Earth is found by (\\ref{dNgdEgEarth}). The best fit values of $m_{ray}$ and $\\phi_{S}$ in Table \\ref{table:Fits} are used to obtain the calculated results in columns 3, 4 and 5.  \\vspace{0.5cm}} \\label{table:Flux} \\centering \\begin{tabular*}{\\textwidth}{@{\\extracolsep{\\fill}}|c|c|c||c|c|} \\hline % \\# &  $dNp/dEp_{\\oplus}$ & $dNp/dEp_{S}$  & $dN\\gamma/dE\\gamma_{S}$  & $dN\\gamma/dE\\gamma_{\\oplus}$   \\\\\\hline\\hline % % 1    &       $ 2.18 \\times  10^{-13} $  &\t $ 2.43 \\times  10^{-13} $   & \t $ 2.87 \\times  10^{-11} $   &\t $ 1.69 \\times  10^{-11} $     \\\\\\hline % 2    & \t $ 3.48 \\times  10^{-14} $   & \t $ 4.78 \\times  10^{-14} $   & \t $ 5.23 \\times  10^{-12} $   &\t $ 1.88 \\times  10^{-12} $     \\\\\\hline % 3   & \t $ 5.56 \\times  10^{-15} $   &\t $ 1.08 \\times  10^{-14} $   &\t $ 1.15 \\times  10^{-12} $   &\t $ 2.73 \\times  10^{-13} $     \\\\\\hline % 4   &  \t $ 8.87 \\times  10^{-16} $   &\t $ 2.49 \\times  10^{-15} $   & \t $ 2.92 \\times  10^{-13} $   &\t $ 5.37 \\times  10^{-14} $     \\\\\\hline % 5   & \t $ 1.41 \\times  10^{-16} $   &\t $ 5.52 \\times  10^{-16} $   & \t $ 7.70 \\times  10^{-14} $   &\t $ 1.21 \\times  10^{-14} $     \\\\\\hline % 6   & \t $ 2.25 \\times  10^{-17} $   & \t $ 1.14 \\times  10^{-16} $   & \t $ 1.90 \\times  10^{-14} $   &\t $ 2.69 \\times  10^{-15} $     \\\\\\hline % 7   &  \t $ 3.58 \\times  10^{-18} $   &\t $ 2.19 \\times  10^{-17} $   & \t $ 4.11 \\times  10^{-15} $   &\t $ 5.38 \\times  10^{-16} $     \\\\\\hline % 8   & \t $ 5.70 \\times  10^{-19} $   & \t $ 3.98 \\times  10^{-18} $   & \t $ 7.70 \\times  10^{-16} $   &\t $ 9.55 \\times  10^{-17} $     \\\\\\hline % 9   & \t $ 9.06 \\times  10^{-20} $   & \t $ 6.93 \\times  10^{-19} $   & \t $ 1.33 \\times  10^{-16} $   &\t $ 1.58 \\times  10^{-17} $     \\\\\\hline % 10   & \t $ 1.44 \\times  10^{-20} $   & \t $ 1.17 \\times  10^{-19} $   & \t $ 2.23 \\times  10^{-17} $   &\t $ 2.59 \\times  10^{-18} $     \\\\\\hline % 11   &   $ 2.27 \\times  10^{-21} $   & \t $ 1.92 \\times  10^{-20} $   & \t $ 3.70 \\times  10^{-18} $   &\t $ 4.22 \\times  10^{-19} $  \\\\\\hline % \\end{tabular*} \\end{table*}        \\pagebreak"
    },
    "0807/0807.1155_arXiv.txt": {
        "abstract": "In this work we test the assertion that the scatter in the mass of black holes which drive quasars should be luminosity dependent with less scatter in more luminous objects. To this end we measure the width of the \\mgii\\ $\\lambda2799$ line in quasar spectra from the SDSS, 2QZ and 2SLAQ surveys and, by invoking an unnormalised virial mass estimator, relate the scatter in line width to the scatter in mass in the underlying black hole population. We find conclusive evidence for a trend such that there is less scatter in line width, and hence black hole mass, in more luminous objects.  However, the most luminous objects in our sample show such a low degree of scatter in line width that, when combined with measures for the intrinsic scatter in the radius-luminosity relation for the broad-line region in active galaxies, an inconsistency arises in the virial technique for estimating black hole masses. This analysis implies that, at least for the most luminous quasars, either there is little-to-no intrinsic scatter in the radius-luminosity relation or the \\mgii\\ broad emission line region is not totally dominated by virial velocities.  Finally we exploit the measured scatter in line widths to constrain models for the velocity field of the broad-line region. We show that the lack of scatter in broad line-widths for luminous quasars is inconsistent with a pure planar/disk-like geometry for the broad-line region. In the case of a broad-line region with purely polar flows the opening angle to luminous quasars must be less than $\\sim55^{\\circ}$. We then explore the effects of adding a random or spherically symmetric component to the velocities of gas clouds in the broad-line region. Assuming an opening angle to quasars of $45^{\\circ}$ a planar field can be made consistent with our results if $\\sim40-50$\\,\\% of the velocities are randomly distributed. ",
        "introduction": "It is now apparent that the majority of massive galaxies harbour a super-massive black hole (SMBH) at their centre. Dynamical studies of the sphere of influence of these SMBHs have been successful in determining the mass for some tens of systems. In these cases the mass of the central SMBH has been observed to correlate strongly with properties of their host spheroid such as luminosity \\cite{mag98}, velocity dispersion ($\\sigma$; Ferrarese \\& Merritt 2000, Gebhardt et al. 2000) and concentration/Sersic index \\cite{gra01}.  More recently, Hopkins et al. (2007a) have demonstrated that these relationships can all be regarded as various projections of a ``black hole fundamental plane'' (BHFP), relating black hole mass to the potential well of the galaxy.  The existence of these correlations and, in particular, the similarity of the BHFP to the fundamental plane of elliptical galaxies points to an intimate link between the growth of SMBHs and galaxy evolution. Moreover, the BHFP is consistent with the notion that SMBHs grew in a self-regulated manner, from gas gravitationally confined in the galaxy centre, which was eventually expelled by feedback processes (Hopkins et al. 2007b). Recent hydrodynamical simulations that incorporate radiative cooling, star formation, black hole growth and feedback from both supernovae and nuclear activity have shown that mergers between gas-rich galaxies of comparable mass provide a mechanism for concentrating gas in galaxy centres through tidal effects (e.g. Barnes \\& Hernquist 1991, 1996), fueling the growth of SMBHs (e.g. Di Matteo et al. 2005), reproducing the correlations (e.g. Robertson et al. 2006), and explaining observed properties of quasars as a phase of evolution prior to the termination of black hole accretion (e.g. Hopkins et al. 2005a,b,c).  Further evidence for a merger-driven origin of quasars is provided by comparing the evolution of the abundance of luminous quasars with that of the cosmic star formation rate (SFR).  From an empirical determination of the bolometric quasar luminosity function, Hopkins, Richards \\& Hernquist~(2007) infer that the quasar luminosity density peaks at $z=2.15 \\pm 0.05$ and rapidly declines towards higher redshifts.  Some observational estimates suggest a similar behaviour for the evolution of the cosmic SFR (e.g. Fan et al.~2001; Hopkins \\& Beacom~2006).  However, at high redshifts, incompleteness, cosmic variance of the surveys, and uncertain corrections owing to dust extinction complicate this analysis.  From a measurement of the opacity of the Lyman-alpha forest (Faucher-Giguere et al. 2007a), Faucher-Giguere et al. (2007b) have shown that the optical depth at $z = 4$ is incompatible with a steep decline in the cosmic SFR at $z > 3$, as suggested by e.g. the results of Hopkins \\& Beacom~(2006), but in accord with theoretical modeling (e.g. Springel \\& Hernquist 2003; Hernquist \\& Springel 2003) which predicts that the SFR should peak at $z>4$.  The implied offset between the peak in the cosmic SFR and the quasar luminosity density indicates that the growth of SMBHs is not directly tied to star formation or gas density, but is related to a secondary process.  By employing estimates of the rate of halo mergers, Hopkins et al. (2007c) have shown that the quasar luminosity function can be reproduced if SMBH growth occurs primarily in major mergers of gas-rich galaxies.  Motivated by these various lines of argument, Hopkins et al. (2006a,b, 2007a,b) have developed a model for galaxy evolution in which mergers, starbursts, quasars, SMBH growth, and the formation of ellipticals are connected through an evolutionary sequence.  As part of this work, Hopkins et al. (2005a) used simulations of merging galaxies to quantify the phases of evolution associated with quasar activity and showed, in particular, that quasar lifetimes depend not only on the instantaneous luminosity of a quasar, but also its `peak' luminosity. Convolving these lifetimes with estimates of the merger rates of galaxies, Hopkins et al. (2005b,c; 2006a; 2007a) were able to reproduce the observed optical and X-ray luminosity functions of quasars.  In this interpretation of the quasar luminosity function, the brighter ($>L^*$) objects are all massive black holes, accreting near the Eddington rate towards the end of their growth phase. Less luminous quasars can be either low-mass systems accreting rapidly and undergoing significant growth, or larger black holes accreting at comparatively lower rates.  This model implies that the range of accretion rates in quasars should be luminosity dependent. That is, the range accretion rates (and by extension SMBH masses) at a given luminosity should be larger for lower luminosity objects, and decrease as the luminosity increases.  In this work we aim to derive the dispersion of the SMBH mass of quasars as a function of luminosity. We measure the \\mgii\\ broad emission-line width in spectra from three large spectroscopic surveys of Type I active galactic nuclei (AGN). Assuming there exists a virial relation for calculating SMBH masses from the \\mgii\\ line width we relate the measured distribution of emission line widths to the SMBH mass distribution. ",
        "conclusions": "We have measured the 50\\,\\% IPV width of the \\mgii\\ line in QSO spectra from the SDSS, 2QZ and 2SLAQ surveys and find a strong correlation between the dispersion in IPV widths and the optical luminosity of QSOs. If we assume there exists a virial relation of the form of equation~\\ref{equ_virial} this implies that there is an equivalent reduction in the dispersion in $\\mbh$. On face value this is in excellent agreement with models for the QSO population proposed by Hopkins et al.~(2005b).   However, the remarkably low scatter we find in our IPV width measurements, in particular for the more luminous objects, has implications as to the validity of virial mass estimators. We find less scatter in IPV widths of luminous QSOs than is intrinsic to the \\rl\\ relation. While the \\rl\\ relation is not defined for these very luminous objects these results are at odds with the practise of virial black hole estimation. Possible explanations for this observation include very low scatter in the \\rl\\ relation for bright QSOs or a BLR which is not fully virialised. In either case it is clear we are yet to gain a full understanding of the process of virial black hole mass estimation and one must be cautious when performing or interpreting these estimates.  Finally we show that the observed scatter in IPV widths can be used to constrain models for the velocity field of the BLR. We show how variations in observing angle to sources affects the scatter in measured IPV widths for a series of simplified geometries and derive constraints on these models from our data. Perhaps of most interest is the rejection of a pure planar BLR regardless of the assumed opening angle to the source.  A table containing the \\mgii\\ IPV widths for the quasars studied in this work is available through the 2SLAQ website ({\\sc www.2slaq.info})."
    },
    "0807/0807.4293_arXiv.txt": {
        "abstract": "{} {In this paper we discuss the physical conditions of clumpy nature in the IC\\,348 molecular cloud. We show that the millimeter and sub-millimeter line emission from the IC\\,348 molecular cloud can be modelled as originating from a photon dominated region (PDR).} {We combine new observations of fully sampled maps in [\\ion{C}{i}] at 492\\,GHz and $^{12}$CO 4--3, taken with the KOSMA 3~m telescope at about 1\\arcmin ~resolution, with FCRAO data of $^{12}$CO 1--0, $^{13}$CO 1--0 and far-infrared continuum data observed by HIRES/IRAS. To derive the physical parameters of the region we analyze the line ratios of [\\ion{C}{i}] $^{3}$P$_{1}$--$^{3}$P$_{0}/^{12}$CO 4--3, [\\ion{C}{i}]$^{3}$P$_{1}$--$^{3}$P$_{0}/^{13}$CO 1--0, and $^{12}$CO 4--3$/^{12}$CO 1--0. A first rough estimate of abundance is obtained from an LTE analysis. To understand the [\\ion{C}{i}] and CO emission from the PDRs in IC\\,348, we use a clumpy PDR model. With an ensemble of identical clumps, we constrain the total mass from the observed absolute intensities. Then we apply a more realistic clump distribution model with a power law index of 1.8 for clump-mass spectrum and a power law index of 2.3 for mass-size relation.} {We provide detailed fits to observations at seven representative positions in the cloud, revealing clump densities between 4\\,10$^{4}$\\,cm$^{-3}$ and 4\\,10$^{5}$\\,cm$^{-3}$ and C/CO column density ratios between 0.02 and 0.26. The derived FUV flux from the model fit is consistent with the field calculated from FIR continuum data, varying between 2 and 100 Draine units across the cloud. We find that both an ensemble of identical clumps and an ensemble with a power law clump mass distribution produce line intensities which are in good agreement (within a factor $\\sim 2$) with the observed intensities. The models confirm the anti-correlation between the C/CO abundance ratio and the hydrogen column density found in many regions.} {} ",
        "introduction": "Photon Dominated Regions (PDRs) are surfaces of molecular clouds where chemistry and heating are regulated by far-ultraviolet (FUV) photons (6.0 eV $<$ $h\\nu$ $<$ 13.6 eV ) \\citep{ht99}. Gas cooling happens mainly via the fine structure lines of [\\ion{O}{i}], [\\ion{C}{ii}], [\\ion{C}{i}] and the rotational lines of CO \\citep{kwhl99}. The form of carbon changes with increasing depth from the surface of the PDR from C$^{+}$ through C$^{0}$ to CO. Therefore emission from [\\ion{C}{ii}], [\\ion{C}{i}] and the rotational lines of CO can be used as probes of temperature, density and column density in the PDRs \\citep{kjmsbs04,mkrm06}. In particular, the two fine structure lines of neutral atomic carbon at 492 and 809\\,GHz give information on intermediate regions between atomic and molecular gas, since they trace both physical regimes. Initial homogeneous plane-parallel steady-state models of PDRs predicted that PDRs are confined to the FUV irradiated surfaces of molecular clouds. However, observations detected widespread [\\ion{C}{ii}] and [\\ion{C}{i}] emission, thus suggesting that PDR surfaces are formed deeply inside the molecular clouds. One possible explanation is in terms of the clumpiness of the molecular clouds which allows for deeper penetration of FUV radiation so that PDRs are formed at the surfaces of the clumps present throughout the molecular cloud \\citep{ssghjb88,hjgs91,mkrm06}. In addition, an explanation has also been proposed in terms of the dynamics of the ISM: a. clouds are suddenly shielded from FUV fields \\citep{sss97}; b. time-dependent chemistry of the PDRs \\citep{oim04}.  IC\\,348 is a young open cluster northeast of the Perseus cloud at a distance of 316\\,pc \\citep{her98}. The IC\\,348 cluster has about 300 stellar members, with spectral types from B5 to late M, as found from observations in X-rays \\citep{pzh96,pz01}, near-infrared \\citep{ll95,lrll98,ntc00,llme05}, and optical \\citep{her98,lsmrlbl03}. \\citet{lada06} constructed optical-infrared spectral energy distributions (SEDs) for all known members of the cluster by combining the Spitzer Space Telescope and ground-based observations. The IC\\,348 molecular cloud prominently shows up in the CO 3--2 survey of the Perseus molecular cloud complex \\citep{skobsm06}. The FUV radiation of its youngest member, HD\\,281159, a B\\,5\\,V star, heats the surrounding gas and dust. Surveys for molecular outflows revealed so far four bipolar outflows powered by deeply embedded young sources \\citep{efdm03,tkb06}. \\citet{jjkm07} estimated the current star formation efficiency of 10\\%--15\\% for the whole Perseus cloud using Spitzer mid-infrared data and SCUBA sub-millimeter maps.  \\citet{skobsm06} observed IC\\,348 in four low-$J$ CO transitions as part of their 7.1 square degree survey of the entire Perseus molecular cloud and quantified the spatial structure of the emission in terms of its power spectrum, applying the $\\Delta$-variance method \\citep{sbhoz98} to the integrated intensity maps. The $\\Delta$-variance spectra of the overall region was found to follow a power law with an index of $\\beta = 2.9-3.0$; while sub-regions of Perseus gave different power spectral indices. The active star forming regions showed higher indices typical for large condensations; lower indices were derived for dark clouds, representing more filamentary structure. IC\\,348 as an exception, showed a rather filamentary structure corresponding to the smallest value of $\\beta = 2.71$ in $^{13}$CO 2--1.  In this paper we present large-scale (20\\arcmin $\\times$ 20\\arcmin) fully-sampled maps of the [\\ion{C}{i}] $^{3}$P$_{1}$ -- $^{3}$P$_{0}$ (hereafter [\\ion{C}{i}]) and $^{12}$CO 4--3 emission from the IC\\,348 molecular cloud. At a common resolution of 70\\arcsec, we have combined our data with the $^{12}$CO 1--0, $^{13}$CO 1--0 data from the Five College Radio Astronomy Observatory (FCRAO) \\citep{rwmam03}, and the far-infrared (FIR) continuum data from HIRES/IRAS.  The goal of the analysis is to understand the extent to which the observed line emission from IC\\,348 can be understood in the framework of a photon dominated region.  Section~\\ref{sec_data} describes the KOSMA observations and the complementary datasets used for analysis. Section~\\ref{sec_obs} discusses the general observational results. A Local Thermodynamic Equilibrium (LTE) analysis of the observed line intensities is presented in Section~\\ref{sec_lte}. A comparison with clumpy PDR models is presented in Section~\\ref{sec_pdr}. Section~\\ref{sec_sum} summarizes the results. ",
        "conclusions": "\\label{sec_sum}  We have presented fully sampled [\\ion{C}{i}] and $^{12}$CO 4--3 maps of the IC\\,348 molecular cloud, covering a region of 20\\arcmin $\\times$ 20\\arcmin. The observed $^{12}$CO 4--3/$^{12}$CO 1--0 ratios vary between 0.2 and 1.5. High $^{12}$CO 4--3/$^{12}$CO 1--0 ratios occur near the B\\,5 star, at the cloud center and northern edge of the cloud.  We have estimated the FUV field from the FIR continuum intensities obtained from the 60 and 100\\,$\\mu$m HIRES/IRAS images. The FUV field in the whole observed region ranges between 1 and 100 Draine units. We also used HD\\,281159, the primary source of UV radiation in this region, to estimate the FUV field for the seven studied regions and the FUV field varies between 3 to 40 Draine units.  Using a clumpy PDR model based on the spherical KOSMA - $\\tau$ PDR model we fitted the observed line intensities and intensity ratios at seven selected positions. The observed intensity ratios provide a strong constraint on the clump densities. They fall between 4.4\\,10$^{4}$\\,cm$^{-3}$ and 4.3\\,10$^{5}$\\,cm$^{-3}$. The FUV field fitted in the model falls between 2 to 100 Draine units, consistent with independent estimates for the FUV field derived from the FIR continuum maps by IRAS and from the stellar radiation. We found that both an ensemble of identical clumps (Ensemble 1) and a model using a clump size and mass distribution following the typically observed scaling (Ensemble 2) produce line intensities which are in good agreement (within a factor $\\sim 2$) with the observed intensities. The models confirm the anti-correlation between C/CO abundance ratio and hydrogen column density found in many regions and explained by \\citet{mkrm06}.   The results of this study strengthen the picture of molecular clouds as highly fragmented, surface dominated objects. Although Ensemble 2 is the natural case in this picture, Ensemble 1 fits the observed intensities equally well. We predicted line intensities for [\\ion{C}{i}], CO, and $^{13}$CO transitions for both cases, which can be tested by future observations. The higher $J$ CO lines are needed to decide on the necessity of the clump mass and size distributions and to determine the clump mass limits \\citep{croks08}."
    },
    "0807/0807.0223_arXiv.txt": {
        "abstract": "Supermassive black holes ejected from galaxy nuclei by gravitational wave recoil will carry a retinue of bound stars, even in the absence of an accretion disk. We discuss the observable signatures related to these stars, with an emphasis on electromagnetic flares from stars which are tidally disrupted by the black hole. We calculate disruption rates for the bound, and the unbound, stars. The rates are smaller than, but comparable to, rates for non-recoiling black holes. A key observational consequence is the existence of powerful off-nuclear and intergalactic X-ray flares. We also discuss other observable signatures associated with the bound stars, including episodic X-ray emission from accretion due to stellar mass loss, intergalactic supernovae, and feedback trails. ",
        "introduction": "Gravitational waves emitted during gravitational collapse carry away linear momentum, causing the center of mass of the collapsing object to recoil (Peres 1962, Bekenstein 1973). In the case of binary black holes, coalescence can be accompanied by kicks as large as $\\sim 200$ \\kms~if the holes are nonspinning prior to coalescence, increasing to $\\sim 4000$ km\\,s$^{-1}$ in the case of maximally spinning, equal-mass holes with the optimum orientations (e.g., Campanelli et al. 2007; Gonz{\\'a}lez et al. 2007; Herrmann et al. 2007; Koppitz et al. 2007; Baker et al. 2008; Schnittman et al. 2008).  Recently, the first compelling candidate for a recoiling SMBH was discovered (Komossa et al. 2008). The quasar SDSSJ092712.65+294344.0 at $z=0.7$  exhibits two separate sets of unusual emission lines: a set of broad emission lines, presumably associated with gas bound to a SMBH, and a second set of atypically narrow emission lines. The broad lines are shifted by $\\sim 2650$ \\kms~relative to the narrow lines. This is one of the key predicted signatures of a recoiling SMBH, which would carry with it the gas from the broad line region while leaving the bulk of the narrow line gas behind (Merritt et al. 2006; Bonning et al. 2007).  Gas near to the SMBH is expected to respond to the kicks in a number of other, potentially observable ways. Accretion activity may be temporarily interrupted during the final phases of binary coalescence but would re-start thereafter with a certain time delay (Liu et al. 2003; Milosavljevi{\\'c} \\& Phinney 2005). The recoiling SMBH would then appear, temporarily, as a ``quasar'' that is spatially offset from the core of its host galaxy (Madau \\& Quataert 2004, Loeb 2007). UV (Lippai et al. 2008), soft X-ray (Shields \\& Bonning 2008) and IR flaring (Schnittman \\& Krolik 2008) could result from shocks in the accretion disk surrounding the SMBH just after recoil, or when the inner disk reforms.  All of these detection methods require the presence of an accretion disk around the coalescing SMBHs, and the duration of the predicted signal is limited to the time it would take the bound gas to be accreted, or much less. But such gas may be lacking in mergers of early type galaxies, and even in gas-rich mergers, inspiral of the binary SMBH may stall long enough that the nearby gas is mostly depleted. How can we observe such recoils ? Even in the absence of a gaseous accretion disk, a recoiling SMBH will always be accompanied by a retinue of tightly bound stars (Gualandris \\& Merritt 2008), and a fraction of these stars will eventually undergo tidal disruption. Stellar tidal disruptions (Luminet 1989) have been observed in the form of giant-amplitude, luminous X-ray flares (Komossa \\& Bade 1999; Halpern et al. 2004; Komossa et al. 2004). In the case of a recoiling SMBH, these flares will appear off-nuclear. In this {\\it Letter} we present tidal disruption of stars as one of the most universal observational signatures of recoiling black holes, and discuss other signatures related to bound stars. ",
        "conclusions": ""
    },
    "0807/0807.0709_arXiv.txt": {
        "abstract": "% We present results from our numerical simulations of collapsing massive molecular cloud cores. These numerical calculations show that massive stars assemble quickly with mass accretion rates exceeding $10^{-3} \\, \\Msol \\, \\yr^{-1}$ and confirm that the mass accretion during the collapsing phase is much more efficient than predicted by selfsimilar collapse solutions, $\\dot{M} \\sim c^3/G$. We find that during protostellar assembly out of a non-turbulent core, the mass accretion reaches $20 - 100 c^3/G$. Furthermore, we explore the self-consistent structure of bipolar outflows that are produced in our three dimensional magnetized collapse simulations. These outflows produce cavities out of which radiation pressure can be released, thereby reducing the limitations on the final mass of massive stars formed by gravitational collapse.  Additional enhancement of the mass accretion rate comes from accretion along filaments that are built up by supersonic turbulent motions. Our numerical calculations of collapsing turbulent cores result in mass accretion rates as high as $10^{-2} \\, \\Msol \\, \\yr^{-1}$. ",
        "introduction": "Our understanding of how massive stars form, while still far from complete has recently made significant strides~\\citep[see also Harold Yorke's contribution in this proceedings and the review by][]{ZinneckerYorke07}.  Low mass stars accrete the bulk of their mass through their circumstellar disks before nuclear burning turns on (e.g., \\citealt{Shu87}).  Massive stars on the other hand, have Kelvin-Helmholtz time scales that are smaller than the dynamical time so that young massive stars start to burn their nuclear fuels while still accreting the surrounding gas \\citep[e.g.,][and references herein]{Yorke02b, Yorke04}.  Infall, or flow of gas through a surrounding disk, therefore faces a major obstacle in the form of the radiative pressure that such massive stars will produce as they are still forming.  Early spherical accretion models suggest that the resulting radiation pressure could limit the final mass of the star to $\\sim 40 \\, \\Msol$ if the accretion rate is not high enough ($\\simge 10^{-3} \\, \\Msol \\, \\yr^{-1}$) \\citep{Kahn74, Wolfire87}. More recent two dimensional simulations by \\citet{Yorke02} showed that the limitations on the final mass of the massive star can be relaxed if accretion through the protostellar disk is included in models of massive star formation. Still, even this particular result yields mass limits of $\\sim 43 \\, \\Msol$. \\citet{Krumholz05} argued that the effects of radiation pressure are limited by the escape of radiation through an outflow cavity.  These results were based on simulations, using Monte-Carlo-diffusion radiative transfer models where the outflow cavities are parameterized by varying opening angles.  Alternatively, massive stars could form through coalescence of intermediate mass stars~\\citep[][]{Bonnell98}. This formation process would be starkly different from the formation of low mass stars which assemble quickly through accretion of the molecular gas. So far, observations of the intermediate state of massive star formation are rare and difficult to obtain. Nevertheless, a few massive objects which show evidence for an ongoing accretion process are known by now \\citep[][]{Chini04, Patel05, Chini06, Beltran06}.  Here we summarize our recent effort studying the collapse of massive cloud cores which final state will be one or more massive stars. These cores are either modeled as supercritical hydrostatic spheres or taken from simulations of supersonic turbulence in which the unstable cores are formed in shock compressions. The numerical setups and the technical details which also includes our description of various cooling and heating processes can be found in~\\cite{Banerjee06b} and~\\cite{Banerjee07a}.  We find that Bonnor-Ebert models for collapsing magnetized cores result in enormous accretion rates and outflows that are driven by toroidal magnetic pressure that is built up in the massive disk. When filamentary structure is added to this picture, infall becomes even stronger as gas is funneled from large scales into the filament, and then down onto the forming disk. ",
        "conclusions": ""
    },
    "0807/0807.2821_arXiv.txt": {
        "abstract": "Recently, \\citet{tiwarietal07} detected the bright component of the visual binary HD\\,151878 to exhibit rapid photometric oscillations through a Johnson $B$ filter with a period of 6\\,min (2.78\\,mHz) and a high, modulated amplitude up to 22\\,mmag peak-to-peak, making this star by far the highest amplitude roAp star known. As a new roAp star, HD\\,151878 is of additional particular interest as a scarce example of the class in the northern sky, and only the second known case of an evolved roAp star -- the other being HD\\,116114. We used the {\\it FIES} spectrograph at the Nordic Optical Telescope to obtain high time resolution spectra at high dispersion to attempt to verify the rapid oscillations. We show here that the star at this epoch is spectroscopically stable to rapid oscillations of no more than a few tens of \\ms. The high-resolution spectra furthermore show the star to be of type Am rather than Ap and we show the star lacks most of the known characteristics for rapidly oscillating Ap stars. We conclude that this is an Am star that does not pulsate with a 6-min period. The original discovery of pulsation is likely to be an instrumental artefact. ",
        "introduction": "\\subsection{Rapidly oscillating Ap stars}  After the discovery of low amplitude pulsation for the bright and well-studied cool, magnetic chemically peculiar A (Ap) star $\\beta$\\,CrB \\citep{hatzesetal04,kurtzetal07}, the question arose whether all cool Ap stars (with effective temperatures below $T_{\\rm eff}\\sim 8200$\\,K) are rapid oscillators. This still-unanswered question may be a pivotal point for theoretical modelling and understanding of the driving mechanism in rapidly oscillating (roAp) stars.  Ap stars are characterised by overabundances of Sr, Cr, Si, Mn and rare earth elements. Many searches for rapid oscillations in Ap stars have been made. \\citet{elkinetal08a} found that 25 known roAp stars that exhibit pulsations photometrically also show rapid radial velocity variations for the corresponding pulsation periods. However, several Ap stars with photometric indices typical for known roAp stars were tested photometrically for pulsations by \\citet{martinezetal94} and found to be stable to high precision; these stars are called non-oscillating Ap stars, or noAp stars.   As seen in, e.g., the astrometric HR-diagram by \\citeauthor{hubrigetal05}\\, (\\citeyear{hubrigetal05}, their figure 2) for roAp and noAp stars, the apparent noAp stars occupy essentially the same regions as the roAp stars. However, the noAp stars appear to be systematically more evolved than the roAp stars \\citep{northetal97,handleretal99,hubrigetal00}. Nevertheless, the theoretical roAp instability strip \\citep{cunha02} also predicts pulsations for the more evolved Ap stars near the terminal age main sequence, with periods in the range $16 - 25$\\,min (see Fig.\\,\\ref{fig:hrd} in Section~\\ref{sec:params}  below). HD\\,116114 \\citep{elkinetal05a} was indeed detected in this region of the HR-diagram, oscillating with the predicted frequency 0.79\\,mHz (the lowest frequency known for the roAp stars). Another somewhat less-evolved roAp star, HD\\,218994, has been recently discovered with a period of 14.2\\,min, which is typical of many roAp stars (\\citealt{gonzalez08}). Thus HD\\,116114 is still the only case of a luminous roAp star near the terminal age main sequence; it shows extremely low radial velocity pulsation amplitude only for a small number of chemical elements such as europium and lanthanum.  \\citet{freyhammeretal08a} searched for rapid pulsations among a group of 9 evolved Ap stars inside the roAp instability strip in the HR-diagram. Surprisingly, they did not detect any in the time-resolved radial velocity measurements for these stars, and showed that $3-5$ of the stars had magnetic field strengths considerably in excess of 2\\,kG. While more evolved stars are theoretically expected to require relatively stronger magnetic fields to suppress local surface convection and facilitate observable amplitudes of rapid pulsations (\\citealt{cunha02}), this study made a good case for the existence of noAp stars near the terminal age main sequence.  \\subsection{HD\\,151878}  HD\\,151878A ($\\alpha_{2000.0}= 16\\,48\\,41.5$, $\\delta_{2000.0}= +35\\,55\\,19$, $V=7\\fm22$) is the primary component of the visual binary HD\\,151878 and is separated by 5\\farcs8 from its 1.51\\,mag fainter companion. \\citet{bidelman85} used low resolution, objective prism plates and classified the star as Am:, where the colon means that the star is a marginal metallic line star with a difference between its spectral type determined from the Ca\\,\\textsc{ii} K\\,line and from the metals lines of less than 5 spectral subclasses. \\citet{grenier99} classified the star as A6mF0F5, where the A6 subtype refers to the Ca\\,\\textsc{ii} K\\,line, F0 refers to the H spectral type, and F5 refers to the metal line type. With 9 spectral subclasses difference between the Ca\\,\\textsc{ii} K\\,line type and metal line type, this classification indicates a strong classical Am star. Am stars are characterised by overabundance of iron-peak elements and underabundance of Ca or Sc. Importantly, the Ca deficiency of the Am stars is not a characteristic of the Ap stars (see, e.g., \\citealt{ryabchikovaetal00} and references therein for a study of 7 roAp stars), making it unlikely that this star is a magnetic Ap star. No medium or high dispersion spectroscopy has been published for the system prior to our study. The Str\\\"omgren indices of HD\\,151878 (see Section~\\ref{sec:params}) indicate spectral peculiarity and are consistent with either an Am or Ap spectral type.  \\citet{tiwarietal07} recently announced the discovery of rapid photometric oscillations in HD\\,151878 with a period of 6\\,min and an amplitude of as much as 22\\,mmag peak-to-peak, making this the highest amplitude roAp star known by a significant margin. The previous highest photometric amplitude observed for a roAp star was on the discovery night for HD\\,60435 when a peak-to-peak variation of 16\\,mmag was detected for this star  (\\citealt{kurtz84}). That star is the most multiperiodic of all roAp stars (\\citealt{matthewsetal87}) and it rarely shows an amplitude this high. For roAp stars with a single, or dominant pulsation mode, the highest photometric $B$ amplitudes are around 10\\,mmag peak-to-peak in, e.g., HD\\,99563 (\\citealt{handleretal06}), HD101065 (Przybylski's star; \\citealt{martinez90}) and HD83368 (HR\\,3831; \\citealt{kurtzetal97}).  Thus, the announcement by \\citet{tiwarietal07} of high amplitude photometric pulsations in HD\\,151878 was potentially very important for many reasons: (i) The photometric amplitude they found is the highest known for any rapidly oscillating A star; (ii) if the star is an Am star -- as two spectral classifications show -- then it would be the first known rapidly oscillating Am star; that would suggest that a strong magnetic field is not required for high overtone pulsation modes in chemically peculiar A stars, which would require a complete rethinking of theoretical models; (iii) if an error had been made in the spectral classification and HD\\,151878 is an Ap star, then it would be only the 6th known roAp star in the northern hemisphere; and (iv) the revised Hipparcos parallax ($5.81\\pm0.61$\\,mas, \\citealt{vanleeuwen07}), along with a spectroscopically estimated effective temperature of 7000\\,K, shows that HD\\,151878 has evolved past the terminal age main sequence (Section\\,\\ref{sec:params}). The most evolved roAp star, HD\\,116114, shows no photometric variability in Johnson $B$ above 1\\,mmag on 6 nights of observation (\\citealt{martinezetal94}), and has the lowest radial velocity amplitudes known for a roAp star. This supports theoretical expectations (\\citealt{cunha02}) of systematically lower frequencies and amplitudes for more evolved roAp stars (at fixed magnetic field strength). HD\\,151878 is inconsistent with this, unless it has a stronger field than that of HD\\,116114 (6\\,kG), which we show it does not.  Because of the importance of such high amplitude pulsation in HD\\,151878, we obtained the first high resolution spectra of the star to study the star's pulsation characteristics, and to study its physical properties -- particularly its spectral peculiarities. ",
        "conclusions": "The known roAp stars have radial velocity amplitudes in the range of few \\ms\\ to over 5\\,\\kms\\ for some rare earth element lines. Low amplitude oscillations in roAp stars are known from, e.g., $\\beta$\\,CrB \\citep{kurtzetal07}, HD\\,154708 \\citep{kurtzetal06} and HD\\,116114 \\citep{elkinetal05a}, the latter with amplitudes of $26-125$\\,\\changeb{\\ms}, measured in lines of the rare earth elements Eu, La and Zr and in the core of the H$\\alpha$ line. The highest radial velocity amplitudes yet detected are for lines of Eu\\,\\textsc{ii} in HD\\,99563 (\\citealt{elkinetal05b}) where individual components generated by surface spots can reach amplitudes of even 8\\,km\\,s$^{-1}$. The roAp stars also have strong magnetic fields that range from several hundred Gauss to 24.5\\,kG in HD\\,154708. Furthermore, they have strong abundances of rare earth elements and almost all of them show ionisation disequilibria for Pr and Nd.  HD\\,151878 fails to show these characteristics and is spectroscopically unlike the roAp stars. It shows no rapid oscillations with radial velocity amplitudes above a few tens of m\\,s$^{-1}$ at the time of our observations. The methods applied, and the length of the data set acquired, would be sufficient for detecting pulsations in most known roAp stars. The roAp stars are oblique pulsators with amplitude modulation of their pulsation with rotation. There is no clear correlation between photometric pulsation amplitude and radial velocity pulsation amplitude in the roAp stars, but the higher photometric amplitude roAp stars -- e.g. HD\\,83368, HD\\,60435, HD\\,99563, HD\\,101065 -- have H$\\alpha$ radial velocity amplitudes of hundreds of m\\,s$^{-1}$ to 2.5\\,km\\,s$^{-1}$; we would expect HD\\,151878 to be similar. Therefore, for the spectacularly high photometric amplitude reported by \\citet{tiwarietal07}, it is improbable that we observed at a rotation phase where the radial velocity amplitude was undetectable at the high precision of our measurements.  The evolutionary tracks in Fig.~\\ref{fig:hrd}, as well as those in figure 1 of \\citet{cunha02}, indicate that HD\\,151878 is past the terminal age main sequence. \\citet{cunha02} argued that more evolved roAp stars require stronger magnetic fields to suppress the upper envelope convection and enable the rapid oscillations to reach observable amplitudes. She further speculated that the magnetic field intensities needed for suppressing the convection in roAp stars more often are found in roAp stars with magnetically resolved lines. That typically requires a magnetic field modulus of $\\sim3$\\,kG for slowly rotating Ap stars ($v \\sin i=1- 3$\\,\\kms).  \\citet{freyhammeretal08a} searched for pulsations in 9 evolved cool Ap stars located inside the theoretical instability strip, the majority having estimated temperatures below $8100$\\,K. With only slightly better precision and upper limits on radial velocity amplitudes as in this study, these authors found 9 null results. Out of 7 stars, only 3 stars had magnetic fields significantly stronger than 2\\,kG, which possibly explains most of their null results for such evolved stars.  Based on comparison of synthetic spectra with HD\\,151878's Fe\\,\\textsc{ii}\\,6149\\,\\AA\\ line, we exclude magnetic fields stronger than $5.5$\\,kG in this star. The magnetic field seems too weak for the star's evolved state to suppress local convection and enable pulsational driving of rapid pulsations to reach observable amplitudes. In this respect, the small, or absent, magnetic field is another Am star characteristic.  Our abundance analysis found that HD\\,151878 has only mildly strong abundances of rare earth elements, no ionisation disequilibria for Nd and Pr and inhomogeneous surface distributions of these elements. In all these respects HD\\,151878 does not resemble known roAp stars. Its Ca\\,\\textsc{ii} and Sc deficiencies is a clear sign of an Am, rather than Ap, nature. Our abundance analysis agrees with the Am spectral classifications. We conclude that HD\\,151878 is a cool evolved Am star. As a such, rapid oscillations are unexpected.  Evolved Am stars are known to be potential pulsators (unlike their main sequence counterparts, which rarely are), exhibiting $\\delta$\\,Sct pulsations (such as $\\delta$\\,Delphini). In fact, the evolutionary state and dimensions of HD\\,151878 resemble that of the pulsating evolved cool Am star HD\\,98851 \\citep{joshietal03}. However, the time scales of pulsating Am stars are much longer ($\\sim$ hours) than those of roAp stars ($\\sim$ minutes). Our spectra only cover 2.55\\,h and are not suited for searching for $\\delta$\\,Sct pulsations. From our velocity measurements of \\halpha\\ and H$\\beta$ we see no indication of variability on timescales of $1- 3$\\,h and amplitudes above $60-250$\\,\\ms\\ (without removing linear trends).  Could HD\\,151878B be the source of the pulsations? This star was inside the diaphragm of the photometric observations of \\citet{tiwarietal07}, but to cause a 22\\,mmag change in the combined light, the much fainter secondary would have to pulsate with an amplitude of 100\\,mmag. While such an amplitude is characteristic of $\\delta$\\,Scuti stars, it is unheard of a roAp star (the highest known is 16\\,mmag for HD\\,60435). The secondary star is located at the red border of the roAp instability strip, and is even possibly outside of it. We conclude that the secondary star could not be the source of the observed 6-min, 22-mmag photometric variations.  This leads to the obvious question as to what is the cause of the observed variations reported by \\citet{tiwarietal07}. Their light curves are convincing; the peaks in their amplitude spectra are highly significant. Their observations were obtained with the 1-m Sampurnanand telescope of \\textsc{ARIES} at Naini Tal, India. The same telescope has been used for many years in the `Naini Tal -- Cape survey' for northern roAp stars (see, e.g., \\citealt{joshietal06}, and references therein). Joshi et al. discovered apparent roAp pulsations in another star classified as an Am star, HD\\,207561 (see their figure 3 and accompanying discussion). On two consecutive nights it showed significant high amplitude pulsation with a frequency of 2.75\\,mHz ($P = 6.1$\\,min) -- {\\it the same frequency as found by \\citet{tiwarietal07} for HD\\,151878} -- yet no further signal was seen again in observations obtained on 17 individual nights over three observing seasons. We therefore suggest that the signal published for HD\\,151878 is instrumental in origin."
    },
    "0807/0807.2156_arXiv.txt": {
        "abstract": "{The classical T Tauri star DG~Tau shows all typical signatures of X-ray activity and, in particular,  harbors a resolved X-ray jet. DG~Tau's jet is one of the most well studied jets of young stellar objects, having been observed for more than 25 years by a variety of instruments. We demonstrate that its soft and hard X-ray components are separated spatially by approximately 0.2~arcsec by deriving the spatial offset between both components from the event centroids of the soft and hard photons utilizing the intrinsic energy-resolution of the {\\it Chandra} ACIS-S detector. We also demonstrate that this offset is physical and cannot be attributed to an instrumental origin or to low counting statistics. Furthermore, the location of the derived soft X-ray emission peak coincides with emission peaks observed for optical emission lines, suggesting that both, soft X-rays and optical emission, have the same physical origin.} ",
        "introduction": "\\vspace{-2mm} The evolution of protostars to young stellar objects (YSOs) is accompanied by accretion from a circumstellar envelope and disk as well as the loss of angular momentum by a substantial, often jet-like mass-outflow perpendicular to the disk \\citep[e.g.][]{KP00}; however, neither the launching mechanism nor the collimation process leading to the observed jet-like outflows have been unambiguously identified. The jet of DG~Tau is among the most well studied jets of YSOs, and these observations have placed tight constraints on the nature of the relevant processes, e.g. the detection of rotation in the outer regions of the DG~Tau jet \\citep{Bacciotti2002} constrained the launch radius of the outflow.  DG~Tau is a classical T Tauri star (CTTS), whose basic properties were summarized by \\citet{2007A&A...468..515G}. Its mass-outflow (a few $10^{-7}M_\\odot\\,$yr$^{-1}$ extending out to $\\sim$~10\\arcsec$\\approx$2300~AU) was first resolved by \\citet{1983ApJ...274L..83M}. Most observations of DG~Tau's outflow were carried out in forbidden emission line regions (FELR), which trace material at temperatures of $\\sim10^{4}\\,$K and densities below $\\sim10^{7}\\,\\mbox{cm}^{-3}$. These studies indicated that at distances larger than $\\sim0.5$\\arcsec~ from the central source the forbidden line emission is concentrated in individual blobs moving at velocities of $\\sim300$~km/s (projected $\\sim 0.3$\\arcsec~yr$^{-1}$) approximately along the jet axis \\citep{Pyo2003}.  The structure of the innermost region of the DG~Tau system is subject to permanent variations. Several studies revealed evidence for material of different speeds and morphology in this region \\citep[e.g.][]{Kepner93,Baciotti2000,Takami2002}, indicating an evolution on time scales of years; as pointed out by \\citet{1993ApJ...410L..31S}, the material in the vicinity of the star is probably denser than in the more distant jet component.  In particular, the jet shows an onion-like structure, where the higher velocity material appears to be embedded in the more slowly moving material \\citep{Bacciotti2002}. The favored heating mechanism for jet emission is internal shocks, heating up the material to temperatures of $\\sim 10^{4}$~K \\citep{Lavalley2000}.  As many (if not all) CTTS, DG~Tau also is an X-ray source, first detected by \\citet{Feigelson81}. From the X-ray point of view the source is unusual in two aspects. First, DG~Tau is the only stellar X-ray source harboring a resolved X-ray jet \\citep{2008A&A...478..797G}, which can be traced out to a distance of $\\sim$~5\\arcsec~ from the central source with a luminosity of about 10\\% of the central soft X-ray component. Second, the X-ray properties of DG~Tau resemble that of the class of ``two-absorber-X-ray (TAX) sources'' \\citep{2007A&A...468..515G,2008A&A...478..797G}. X-ray spectra of TAX sources are basically the sum of two thermal components, differing not only in mean temperature but - in contrast to most other X-ray spectra - also in absorbing column density. In DG~Tau, the emission regions of the soft and the hard components appear to be disjoint spatially.  In an {\\it XMM-Newton} observation, \\citet{2007A&A...468..515G} found an increase in the hard component's count rate during a flare, while the soft component's count rate remained constant; they proposed therefore, supported by the spectral properties of the soft component, an interpretation of the soft component as internal shocks in a jet close to the star. Motivated by these indications that the soft X-ray component in DG~Tau might be spatially detached from the hard X-ray component, we performed a detailed position analysis of both components utilizing the superb angular resolution of the {\\it Chandra} telescope.  \\vspace*{-5mm} ",
        "conclusions": "\\vspace*{-2mm} By considering the measured offset of 0.2\\arcsec~ between the soft and the hard X-ray centroid position as physical, we can convert this offset into a physical distance of 48~AU from the central source, assuming a distance of 140~pc and a disk inclination of $38^\\circ\\,$\\citep{Eis98}.  This distance is an order of magnitude larger than reasonable launching regions of the jet \\citep[$\\sim$~1~AU,\\,][]{Anderson2003}, and our results therefore favor strongly the interpretation that the X-ray emission observed close to the star is originating from internal shocks of the jet, as proposed by \\citet{2007A&A...468..515G}. Internal shocks are incidentally also the preferred heating mechanism for the optically observed FELRs. With this interpretation, the total X-ray luminosity of the jet is an order of magnitude higher than that of the $Chandra$-resolved part of the jet \\citep{2008A&A...478..797G}, although even then its luminosity is far smaller than the optical jet luminosity; for example, \\citet{Lavalley2000} derived a luminosity, for the first emission peak of the [O~I]-line, of $1.1\\times 10^{-4} L_\\odot$, which is a factor of $\\sim40$ higher than the energy-loss by X-ray emission, suggesting that only a small amount of the outflowing material reaches X-ray temperatures. We now consider the following scenario:  we model the X-ray jet by a cylinder of radius $r$ and height $d$; the base of the cylinder is located at the shock region, where the material is heated to some temperature $T$, and the shocked material flows through this cylinder with some post-shock velocity $v$. We assume that the shocked material cools predominantly by radiation with a cooling time $\\tau_c$ given by $\\tau_c = 3k_B T / (n \\Lambda(T))$, where $n$ is the plasma density, $T$ the (post-shock) temperature, $\\Lambda(T)$ the cooling function, and $k_B$ Boltzmann's constant. The cooling distance $d$, i.e., the height of the cylinder, is then given by $d = \\tau_c \\cdot v = 3k_B T v/(n\\cdot\\Lambda(T)) $. We further know the total emission measure $EM$ of the X-ray emitting plasma and write \\begin{equation} EM =  f\\cdot n^{2} \\cdot A\\cdot d = 3 f\\cdot n\\cdot\\pi r^{2} \\cdot v k_B T/\\Lambda(T), \\end{equation} where $f$ denotes an unknown filling factor of the hot plasma. The mass outflow rate $\\dot{M}_{X-ray}$ of the X-ray emitting plasma can be computed from \\begin{equation} \\dot{M}_{X-ray} = m_H \\ f \\ n\\cdot \\pi r^{2} \\cdot v= m_H\\Lambda(T) EM/(3 k_B T), \\end{equation} i.e. the mass outflow rate of the X-ray emitting material is only determined by the observed quantities $T$ and $EM$. Our spectral fit provides a mean temperature of $T\\sim3.4$~MK for the shocked plasma and \\mbox{$EM \\sim 3.5\\times10^{52}\\,\\mbox{cm}^{-3}$} (APEC-models, metallicity at 0.3~solar) compatible with the values given by \\citet{2008A&A...478..797G}. With these numbers we find an outflow rate of $1.3\\times 10^{-11} M_\\odot/\\mbox{yr}$ using \\mbox{$\\Lambda(T)\\approx2\\times10^{-23}~\\mbox{erg cm}^{3}\\mbox{s}^{-1}$}. This value is indeed orders of magnitude smaller than the outflow rate of the high velocity material only \\citep[$4\\times 10^{-9} M_\\odot/\\mbox{yr}$, ][]{Lavalley2000}. Are such values physically reasonable? We note that the soft X-ray component is more or less point-like if we disregard for the time being the extended jet component described by \\citet{2008A&A...478..797G}.  Given the distance of 140~pc towards DG~Tau, this implies that the region $d_{max}$ has the approximate size of 112~AU (one ACIS pixel). If we estimate the outflow speed of the shocked material to be approximately 300~km/s, which is on the one hand the speed of the high-velocity material measured in FELRs and on the other hand approximately the speed required to produce the observed soft X-ray temperature by means of the strong shock formula \\citep[the optically observed shocks have speeds of only up to 100~km/s,][]{Lavalley2000} we can then derive $n > n_{min} =  3k_B T v/\\left(d_{max}\\cdot\\Lambda(T)\\right) = 1.3 \\times 10^{6} \\mbox{cm}^{-3}$.\\\\ Using these values of $n_{min}$ and $d_{max}$ to calculate the emission measure, we estimate the effective outflow cross sectional area \\mbox{$f \\pi r^{2} \\approx 1 \\times 10^{25}$ cm$^{2}$} or \\mbox{$f \\pi r^{2} \\approx 6 \\times 10^{-2}$ AU$^2$}. The launch and collimation distance of the jet in DG~Tau is believed usually to be approximately \\mbox{1 AU}.  Therefore, the filling factor of the X-ray emitting material must be small and we envisage a scenario of hot X-ray plasma with a small filling factor, immersed in cooler material with a far larger filling factor. A fraction of the shocked material is clearly observed radiating in the resolved jet at a distance of 5\\arcsec~ to DG Tau at essentially the same X-ray temperature as the ``inner jet''. The cooling time of this material may be sufficiently large to enable the material to move the distance required; this would descrease the density and the filling factor by one and two orders of magnitude, respectively, in comparison to the values above. We prefer, however, an interpretation in which the resolved jet is possibly ``re-shocked'' material. At any rate, only a minor fraction of the outflowing material in the ``inner jet'' experiences shocks at $\\sim 300$~km/s or more, while the densities and mean velocities of the hot (T $\\sim 3.4\\times 10^{6}$~K) and the cool (T $\\sim 10^{4}$~K) material are similar to within factors of a few."
    },
    "0807/0807.1069_arXiv.txt": {
        "abstract": "We use {\\it Spitzer} MIPS data from the FIDEL Legacy Project in the Extended Groth Strip to analyze the stellar mass assembly of massive ($\\mathcal{M}$$>$10$^{11}$~$\\mathcal{M}_\\sun$) galaxies at z$<$2 as a function of structural parameters. We find 24\\mic\\, emission for more than 85\\% of the massive galaxies morphologically classified as disks, and for more than 57\\% of the massive systems morphologically classified as spheroids at any redshift, with about 8\\% of sources harboring a bright X-ray and/or infrared emitting AGN. More noticeably, $\\sim$60\\% of all compact massive galaxies at z$=$1--2 are detected at 24\\mic, even when rest-frame optical colors reveal that they are dead and evolving passively. For spheroid-like galaxies at a given stellar mass, the sizes of MIPS non-detections are smaller by a factor of $\\sim$1.2 in comparison with IR-bright sources. We find that disk-like massive galaxies present specific SFRs ranging from 0.04 to 0.2~Gyr$^{-1}$ at z$<$1 (SFRs ranging from 1 to 10~$\\mathcal{M}_\\sun$~yr$^{-1}$), typically a factor of 3--6 higher than massive spheroid-like objects in the same redshift range. At z$>$1, and more pronouncedly at z$>$1.3, the median specific SFRs of the disks and spheroids detected by MIPS are very similar, ranging from 0.1 to 1~Gyr$^{-1}$ (SFR$=$10--200~$\\mathcal{M}_\\sun$~yr$^{-1}$). We estimate that massive spheroid-like galaxies may have doubled (at the most) their stellar mass from star-forming events at z$<$2: less than 20\\% mass increase at 1.7$<$z$<$2.0, up to 40\\% more at 1.1$<$z$<$1.7, and less than 20\\% additional increase at z$<$1. Disk-like galaxies may have tripled (at the most) their stellar mass at z$<$2 from star formation alone: up to $\\sim$40\\% mass increase at 1.7$<$z$<$2.0, and less than 180\\% additional increase below z$=$1.7 occurred at a steady rate. ",
        "introduction": "The formation and evolution of massive ($\\mathcal{M}$$>$10$^{11}$~$\\mathcal{M}_\\sun$) galaxies is one of the most studied topics in extragalactic astronomy during the last decade. Early expectations from hierarchical galaxy formation models considered that star formation began in low mass systems which built more massive galaxies through sequential merging \\citep{1996MNRAS.283.1361B,2000MNRAS.319..168C}, in a similar process to the growth of structures in Cold Dark Matter simulations \\citep{2005Natur.435..629S}. In apparent contradiction with hierarchical assembly, the finding of a substantial population of massive galaxies at z$>$1 \\citep{1988ApJ...331L..77E,1998Natur.394..241H, 2003ApJ...587L..79F, 2004Natur.430..181G}, some of them containing old stellar populations and evolving passively (according to their red optical colors --\\citealt{2004ApJ...617..746D}, \\citealt{2005ApJ...633..748R}--, and spectra --\\citealt{2006ApJ...649L..71K,2008ApJ...677..219K,2008A&A...482...21C}--), seems to favor a downsizing formation scenario \\citep{1996AJ....112..839C,2004Natur.428..625H,2005ApJ...619L.135J, 2005ApJ...630...82P,2006ApJ...651..120B}. This population of z$>$1 massive galaxies accounts for a significant fraction of the local stellar mass density ($\\sim$20\\% as early as z$\\sim$2, and $\\sim$10\\% at z$\\sim$4: \\citealt{2006A&A...459..745F,2007A&A...476..137A,2008ApJ...675..234P}). The discovery of such a population reinforced the idea that both stars and their host galaxies are coeval (resembling a monolithic--like collapse), and consequently, no expectations of structure evolution in these galaxies should be expected. For this reason, the recent observational evidence showing that the most massive galaxies were much more compact in the past \\citep{2005ApJ...626..680D,2006MNRAS.373L..36T, 2007MNRAS.374..614L,2008A&A...482...21C} has been surprising, and has again opened the question of how the stellar populations of these galaxies were assembled into their present shape.  The size evolution of the most massive objects since z$\\sim$2 has been characterized by \\citet{2007MNRAS.382..109T}. These authors found that, at a given stellar mass, disk--like objects at z$\\sim$1.5 were a factor of two smaller than their present-day counterparts. For spheroid--like objects, the evolution has been even stronger: they were a factor of four smaller at z$\\sim$1.5 than nearby similar mass ellipticals. In addition, the stellar mass densities of these high--z galaxies were almost two orders of magnitude higher than objects of the same mass today. These superdense galaxies have been found at even higher (z$\\sim$2.5) redshifts \\citep{2007ApJ...656...66Z,2007ApJ...671..285T,2008ApJ...677L...5V}, adding more controversy to the debate about the formation and evolution of massive galaxies.  Two processes have been proposed to allow the superdense high-z galaxies migrate to the local stellar mass--size relation. The first process is dissipationless (with absence of star formation) merging. Given the high metal abundances and old ages of the stellar population present in local massive elliptical galaxies (e.g., \\citealt{2006MNRAS.370.1106G,2006A&A...457..809S,2007ApJ...669..947J}), these mergers should preferentially be dry \\citep{2006MNRAS.366..499D}, and occur between z$\\sim$1.5 and z$=$0, the epoch when the red sequence appears \\citep{2007ApJ...665..944L}. In this context, a particular effective size evolutionary mechanism (r$_e$$\\sim$$\\mathcal{M}$$^{1.3}$) has been provided by \\citet{2006MNRAS.369.1081B} through head--on mergers of galaxies. The second possibility is the smooth envelope accretion scenario \\citep{2007ApJ...658..710N}, where accreted stars (mainly provided by minor mergers) form an envelope whose size increases smoothly at decreasing redshift.  The goal of this paper is to explore the evolutionary paths followed by the most massive galaxies and their dependence on the morphology. To do this, we quantify the growth in stellar mass via star formation events of massive ($\\mathcal{M}$$>$10$^{11}$~$\\mathcal{M}_\\sun$) galaxies as a function of size and brightness profile shape up to z$\\sim$2. We base our discussion on the characterization of the dust infrared (IR) emission of these systems, which is linked to the amount of recent star formation and/or the presence of obscured AGN. This IR-based study is complementary to the more classical approach to the characterization of the evolution of massive ellipticals based on rest-frame optical properties.  Throughout this paper, we use a cosmology with $\\mathrm H_{0}=70$~km\\,s$^{-1}$\\,Mpc$^{-1}$, $\\Omega_{\\mathrm M}=0.3$ and $\\Lambda=0.7$. All magnitudes refer to the AB system. The results for stellar masses and SFRs assume a \\citet{2003ApJ...586L.133C} initial mass function (IMF) with 0.1$<$$\\mathcal{M}$$<$100~$\\mathcal{M}_\\sun$. ",
        "conclusions": ""
    },
    "0807/0807.1296_arXiv.txt": {
        "abstract": "{ The Virtual Observatory (VO) is nearing maturity, and in Spain the Spanish VO (SVO) exists since June 2004. There have also been numerous attempts at providing more or less encompassing grid initiatives at the national level, and finally Spain has an official National Grid Initiative (NGI). In this talk we will show the VO and Grid development status of nationally funded initiatives in Spain, and we will hint at potential joint VO-Grid use-cases to be developed in Spain in the near future. ",
        "introduction": "% \\label{sec:introduction}  The Virtual Observatory, as proposed by \\citeauthinyear{2001Sci...293.2037S}, is today a tangible infrastructure used by thousands of astronomers (and non-astronomers) every day. It help scientists by letting them explore massive amounts of data, to derive statistical properties from thousands of objects, and even to find new kinds of objects; by providing a multi-wavelength view of particular objects or regions of the sky; and by providing time-lapsed views of any region of the sky from archived data.  As mentioned, the foundational paper of the Virtual Observatory by Szalay and Gray was written in 2001, and the first Astronomical Virtual Observatory (AVO) prototype, based on CDS's Aladin Sky Atlas\\urlnote{http://aladin.u-strasbg.fr/} is from 2002 (see \\citeauthyear{2005ASPC..347..183P}). By that time, our research team (the AMIGA group, Analysis of the interstellar Medium of Isolated GAlaxies, \\citeauthyear{2005A&A...436..443V}) was working on a multi-wavelength study for the 1,051 galaxies in the Catalogue of Isolated galaxies by \\citeauthinyear{1973SoSAO...8....3K} (see also her online catalogue, \\citeauthyear{1997yCat.7082....0K}). The group had already compiled quite a large amount of information (revised positions, IR fluxes in different wavelengths, images with different filters), but still needed more information to be retrieved, and wanted to make our different revised data products available to the community.  Using the existing International Virtual Observatory Alliance\\urlnote{http://ivoa.net/} (IVOA) Proposed Recommendations, such as the VOTABLE (\\citeauthyear{Ochsenbein:2004ty}) and the Simple ConeSearch (\\citeauthyear{Williams:2008fv}), the AMIGA group finally provided a public web site\\urlnote{http://amiga.iaa.csic.es:8080/DATABASE/} delivering data in VOTABLE format, providing at the same time a Simple ConeSearch interface.  \\subsection{Radio astronomy and the VO} % \\label{sub:radio_astronomy_and_the_vo}  As astronomy started when the first humans raised their heads and reckoned that the sky was very similar from night to night, with seasonal variations that repeated from year to year, the visible part of the spectrum has always been the most used, and more familiar, to all astronomers. However, the discovery by \\citeauthinyear{Jansky:1933db} of a extraterrestrial and extrasolar radiation with a wavelength of 14.6 meters, opened a new spectral window that even today is just 63 years old, a small fraction of the almost 400 years of instrumental optical astronomy.  The Virtual Observatory is also developing more rapidly for data from the visible parts of the electromagnetic spectrum, but radio astronomical data is of the utmost importance for understanding the most distant and obscured objects in the universe, and the integration of these data sets in the VO infrastructure has been a goal of our group since the beginning.   \\subsection{Grid and astronomy} % \\label{sub:grid_and_astronomy}  The computing grid was conceived by \\citeauthinyear{Foster:1999xe} as \\emph{a hardware and software infrastructure that provides dependable, consistent, pervasive and inexpensive access to high-end computational capabilities}. The similarity with the power grid arises from the fact that until the power grid became common-place, electricity based facilities depended on their own power generation. Providing computing-on-demand, as easy to access as the power grid, is the aim of grid computing initiatives.  The next generation of astronomical instruments, that we might call surveying instruments, such as the Large Synoptic Survey Telescope (LSST), the Square Kilometre Array (SKA), the LOw Frequency ARray (LOFAR), or even the Atacama Large Millimetre and submillimetre Array (ALMA), will be extremely sensitive, and at the same time their schedule will be completely automated, so that the raw data rates will overwhelm any conventional computing facility. For these telescopes, parallel processing of incoming data is needed, either by extensive pre-processing at the data generation site, or by running parallel pipelines for different levels of data products, or by splitting the different levels of processing between different \\emph{tiers}. An extreme, present case, is the LOFAR which, as can be seen from \\citeauthor{Valentijn:2008cr}'s proceeding, will need a super- computer class data reduction engine with a tiered architecture for data reanalysis.  At the same time, the storage needs for this instruments will be so demanding that for some applications no actual data access will be possible, but only streaming access to the data being provided. Again, the LOFAR is an example of this. In any case, data access has also to be distributed, and different grid techniques (grid-FTP, IBM's Grid Parallel File System...) are used.   In the following sections, we will learn about the SVO (the Spanish Virtual Observatory), grid developments in Spain, and joint grid and virtual observatory efforts in our country. We will end with a conclusions section. ",
        "conclusions": "% \\label{sec:conclusions}  We have shown that pure Virtual Observatory activities in Spain have a very good health, and there is ample experience both inside and (specially) outside the astrophysics community, but more and more services will be deployed on the grid as it becomes more pervasive. Spain will have in the near future a mature enough grid infrastructure within the European Grid Initiative, and different research groups are already porting, or planning to port, their code bases to the grid, taking as much as possible from existing VO infrastructures."
    },
    "0807/0807.4366_arXiv.txt": {
        "abstract": "We present a strong constraint on variation of the proton-to-electron mass ratio, $\\mu$, over cosmological time scales using molecular hydrogen transitions in optical quasar spectra. Using high quality spectra of quasars Q0405$-$443, Q0347$-$383 and Q0528$-$250, variation in $\\mu$ relative to the present day value is limited to $\\Delta\\mu/\\mu=(2.6 \\pm 3.0) \\times 10^{-6}$. We reduce systematic errors compared to previous works by substantially improving the spectral wavelength calibration method and by fitting absorption profiles to the forest of hydrogen Lyman $\\alpha$ transitions surrounding each H$_2$ transition. Our results are consistent with no variation, and inconsistent with a previous $\\approx4\\sigma$ detection of $\\mu$ variation involving Q0405$-$443 and Q0347$-$383. If the results of this work and those suggesting that $\\alpha$ may be varying are both correct, then this would tend to disfavour certain grand unification models. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0807/0807.0129_arXiv.txt": {
        "abstract": "Even though water is the main constituent in interstellar icy mantles, its chemical origin is not well understood. Three different formation routes have been proposed following hydrogenation of O, O$_2$, or O$_3$, but experimental evidence is largely lacking. We present a solid state astrochemical laboratory study in which one of these routes is tested. For this purpose O$_{2}$ ice is bombarded by H- or D-atoms under ultra high vacuum conditions at astronomically relevant temperatures ranging from 12 to 28 K. The use of reflection absorption infrared spectroscopy (RAIRS) permits derivation of reaction rates and shows efficient formation of H$_{2}$O (D$_{2}$O) with a rate that is surprisingly independent of temperature. This formation route  converts O$_2$ into H$_2$O via H$_2$O$_2$ and is found to be orders of magnitude more efficient than previously assumed. It should therefore be considered as an important channel for interstellar water ice formation as illustrated by astrochemical model calculations. ",
        "introduction": "Solid water ice has been observed on the surfaces of many different astronomical objects. In the Solar System it is found on planets and minor bodies such as comets, trans-Neptunian objects and Centaurs. In dense, cold interstellar clouds, infrared observations show that interstellar dust grains are covered with water-rich ices \\citep[e.g.]{Gillett:1973,Gibb:2004,Pontoppidan:2004}. The formation of these ice mantles is especially important in the process of star and planet formation, when a large fraction of heavy elements can be depleted onto grains. In the dense cloud phase water layers form on the bare grain surfaces. Then during the gravitational (pre-)collapse, virtually all gas phase species freeze-out on top of these water layers resulting in a CO dominated layer that likely also contains traces of O$_2$ but very little water.  The observed H$_2$O ice abundance cannot be explained by direct accretion from the gas phase only. The exact mechanism by which water ice is formed is not understood.  The Herschel Space Observatory, to be launched in the near future, will provide important new information on gaseous water in interstellar space and will measure quantitatively the water abundance as a function of temperature, UV field and other parameters. Furthermore, the Photodetector Array Camera and Spectrometer (PACS) on Herschel will cover the 62 $\\mu$m band of solid H$_2$O. In this way Herschel will provide a unique opportunity to observe the bulk of the water bands that are unobservable from the ground and relate them to Spitzer and groundbased mid-IR observations of ices in protostellar envelopes and protoplanetary disks. Understanding the processes by which water forms and why it is not formed under other circumstances will be essential for the interpretation of these data.  \\cite{Tielens:1982} proposed a reaction scheme in which water ice is formed on the surfaces of grains via three different routes: hydrogenation of O, O$_2$, and O$_3$. Models predict that water can indeed be formed through  such reactions in dense clouds \\citep[e.g.]{Tielens:1982,dHendecourt:1985,Hasegawa:1993,Cuppen:2007}. Using a Monte Carlo approach \\cite{Cuppen:2007} showed that the contributions of the different formation channels to water ice formation as well as its abundance strongly depend on the local environment. However, the initial reaction scheme with the corresponding rates as proposed by \\cite{Tielens:1982} is based on old, in some cases outdated, gas phase data of the equivalent reactions. Progress has been severely hampered by the lack of realistic experimental simulations of these low-temperature, solid state reactions. Preliminary laboratory studies of water synthesis  testing the first reaction channel have been reported by \\cite{Hiraoka:1998} and by \\cite{Dulieu:2007}. Both groups investigated the products of D- and O- reactions on an ice substrate (N$_{2}$O and H$_{2}$O, respectively) using temperature programmed desorption (TPD). In experiments exclusively using this technique, it is hard to rule out any H$_{2}$O formation during warm up. Furthermore, quantitative interpretation can be tricky because unstable species like H$_2$O$_2$ are destroyed in the mass spectrometer upon ionization leading to an artificially enhanced H$_2$O/H$_2$O$_2$ ratio.  The present work  focuses on the H + O$_2$ channel in which  O$_{2}$  is converted to H$_{2}$O via H$_2$O$_2$: \\begin{equation} \\rm{H} + \\rm{O}_{2} \\rightarrow \\rm{HO}_{2} \\label{H+O2->HO2} \\end{equation} \\begin{equation} \\rm{H} + \\rm{HO}_{2} \\rightarrow \\rm{H}_{2}\\rm{O}_{2} \\label{H+HO2->H2O2} \\end{equation} \\begin{equation} \\rm{H} + \\rm{H}_{2}\\rm{O}_{2} \\rightarrow \\rm{H}_{2}\\rm{O} + \\rm{OH} \\label{H+H2O2->H2O+OH} \\end{equation} \\begin{equation} \\rm{H} + \\rm{OH} \\rightarrow \\rm{H}_{2}\\rm{O} \\label{H+OH->H2O} \\end{equation} According to \\cite{Cuppen:2007}  this channel is, together with the O$_3$ channel, responsible for water formation in cold, dense clouds. The exposure of O$_{2}$ ice to hydrogen and deuterium atoms is investigated by means of reflection absorption infrared spectroscopy (RAIRS) and TPD. These techniques allow one to determine formation yields and  the corresponding reaction rates. The present work comprises a study of hydrogenation and deuteration reactions of O$_{2}$ ice for different temperatures between 12 and 28 K, i.~e.~, up to the desorption temperature of O$_{2}$ \\citep{Acharyya:2007}. The formation of H$_2$O and H$_2$O$_2$ is observed at all temperatures. An optimum yield is found at 28 K. ",
        "conclusions": "The hydrogenation and deuteration experiments presented in this letter show an efficient mechanism to convert O$_2$ ice to H$_2$O$_2$ and ultimately H$_2$O. In the model that describes these experiments, the rate limiting steps for formation are the trapping of the hydrogen into the O$_2$ ice and the penetration depth of the hydrogen into the H$_2$O$_2$. Astrochemical models take reaction, diffusion and desorption barriers as input. Our data shows that the formation of at least H$_2$O$_2$ proceeds via a reaction with a barrier that is much lower than the value of 1200 K previously assumed by \\cite{Tielens:1982}. It should therefore be  considered as a route for water formation on interstellar grain surfaces.  Care should however be taken when extending these findings from a laboratory environment directly to interstellar ices. The temperature independence of the reaction rate that is observed, for instance, is directly due to the unique property of the O$_2$ ice that allows hydrogen to penetrate its structure and thereby preventing desorption back into the gas phase. In interstellar clouds the mantle surfaces would not consist of pure O$_2$ ice but of a mixture of different species with water as its main constituent  \\citep{Whittet:1998}. The structure of these ``dirty'' ices and the binding energies to it would govern the desorption and diffusion behavior of adsorbates like O$_2$ and H. Experiments on the binding of atomic hydrogen on water ice surfaces showed that the surface abundance is generally dependent on temperature \\citep{Perets:2005, Hornekaer:2003, Dulieu:2005}. Also hydrogenation experiments of CO showed that the decreasing H coverage for increasing temperature results in lower effective rates \\citep{Watanabe:2003,Watanabe:2006A,Linnartz:2006,Cuppen:HK}. Here we discuss the implications of the present work to two different interstellar environments: hydrogenation of an apolar (water-poor) ice mantle after freeze-out and cold cloud conditions where ice is formed from direct deposition of H and O.  Observations show that interstellar ice mantles consist of polar (water-rich) and apolar (water-poor) layers \\cite{Tielens:1991}. Apolar ices are thought to form during freeze-out in the densest parts of the cloud. In dense gas, most of the atomic oxygen has converted to O$_2$ which can become subsequently a constituent of this polar phase. The lack of observed H$_2$O$_2$ and H$_2$O means that most of the O$_2$ in the ice does not react to produce H$_2$O$_2$ or H$_2$O. This laboratory work however shows that hydrogen atoms can penetrate deeply into O$_2$ ices and then react. In apolar interstellar ices the penetration depth is not determined by oxygen ice but by the main constituent of the ice mantle, CO. We therefore have performed additional laboratory experiments in which a mixture of CO and O$_2$ ice is exposed to H-atoms. These show indeed that only the top few layers are hydrogenated and that the main part of the ice stays intact in agreement with the observations. Details of these experiments are to be published in a future paper.  In cold and translucent clouds H- and O-atoms deposit onto the grain simultaneously. The species can then react  immediately and O$_2$ is converted all the way to water. Here the penetration depth observed in the laboratory becomes unimportant. What can we learn from the present experiments then? The fast reaction kinetics  justify treating H + O$_2$ and H + HO$_2$ in a similar way as H + H are treated in astrochemical models, with the difference that probably only hydrogen is mobile for reaction. The results further suggest that the continuing reactions leading to water proceed with high efficiency. The original grain surface network by \\citep{Tielens:1982} includes two more water formation routes: via OH and via ozone. Under dense cloud conditions the ozone route was proposed to be the most efficient as is shown in the left panel of Fig.~\\ref{model}. This figure plots the contributions for the three different H$_2$O formation channels using a reaction network limited to the three water formation routes and without any dissociation reactions. The model parameters are taken from model M1 by \\cite{Ruffle:2000} and the initial conditions from \\cite{Tielens:1982} are used ($n({\\rm H}_2) = 5 \\times 10^4$, $n({\\rm H}) = 1.41$ and $n({\\rm O}) = 2.36$ cm$^{-2}$). These calculations include a barrier for reactions 1 and 3. The present work shows that these barriers are negligible and the middle panel plots the same model calculation without these barriers. The figure clearly shows that the O$_2$ channel has a major contribution to the overall water formation rate.   Recent laboratory experiments by \\cite{Sivaraman:2007} on the temperature-dependent formation of ozone by irradiation of oxygen ices by high energy electrons suggest that mobile oxygen atoms prefer the O + O pathway over O + O$_2$ even though the O/O$_2$ ratio in the ice is very small. This implies the presence of a barrier for the formation of ozone. The right panel gives the model results when a small barrier of 500 K for this reaction is considered. The contribution of the ozone channel is now further reduced.  Figure \\ref{model} further shows the H$_2$O$_2$ abundance, $n$(H$_2$O$_2$), for all models as produced through the O$_2$ + H route. \\cite{Boudin:1998} set an observational constraint of the  H$_2$O$_2$ ice abundance toward \\object{NGC7538:IRS9} of 5.2\\% with respect to the H$_2$O ice abundance. All models are consistent with this number."
    },
    "0807/0807.0290_arXiv.txt": {
        "abstract": "Recently, it was suggested that a photospheric component that results from the internal dissipation occurring in the optically thick inner parts of relativistic outflows may be present in the prompt $\\gamma$/X-ray emission of gamma-ray bursts or X-ray flashes. We explore high-energy neutrino emission in this dissipative photosphere model, assuming that the composition of the outflow is baryon-dominated. We find that  neutrino emission from proton-proton collision process forms an interesting signature in the neutrino spectra. Under favorable conditions for the shock dissipation site, these low-energy neutrinos could be detected by ${\\rm km^3}$ detectors, such as Icecube. Higher energies ($\\ga10$ TeV) neutrino emission from proton-proton collision and photo-pion production processes could be significantly suppressed for dissipation at relatively small radii, due to efficient Bethe-Heitler cooling of protons and/or radiative cooling of the secondary mesons in the photosphere radiation. As the dissipation shocks continue further out, high energy neutrinos from photo-pion production process becomes dominant. ",
        "introduction": "The photosphere models have been widely discussed in relation to the prompt emission of GRBs (e.g. Thompson 1994; Ghisellini \\& Celotti 1999; Rees \\& \\Meszaros 2005; Thompson et al. 2007; Ioka et al. 2007). The potential advantage of photosphere models is that the peak energy can be stabilized, which is identified as the thermal or Comptonization thermal peak (see Ioka et al. 2007 and references therein).  The photosphere radiation may also produce a large number of electron-positron pairs, which may lead to a pair photosphere beyond the baryon-related photosphere (e.g. Rees \\& \\Meszaros 2005), and may also enhance the radiative efficiency (Ioka et al. 2007). On the other hand, it is also suggested that the number of pairs produced does not exceed the baryon related electrons by a factor larger than a few (Pe'er et al. 2006). For simplicity,  we here only consider the dissipation below the baryon-related photosphere, which is more favorable for $pp$ neutrino production.  Following Rees \\& \\Meszaros (2005) and Pe'er et al. (2006), we assume that during the early stage of the prompt emission, internal shocks of the outflow occur at radii below the baryonic photosphere. Initially, the internal energy is released at the base of the outflow, $r_0\\sim\\alpha r_g=2\\alpha GM/c^2$, where $\\alpha\\ga1$ and $r_g$ is the Schwarzschild radius of a central object of mass $M$. The internal energy is then converted to the kinetic energy of the flow, whose bulk Lorentz factor grows as $\\gamma\\sim r$ up to a saturation radius at $r_s\\sim r_0\\eta$, where $\\eta=L_0/(\\dot{M}c^2)$ is the initial dimensionless entropy, $L_0$ and $\\dot{M}$ are the total energy and mass outflow rates. Above the saturation radius, the observer-frame photospherical luminosity decreases as $L_\\gamma(r)=L_0(r/r_s)^{-2/3}$ and the greater part of energy is in kinetic form, $L_k\\sim L_0$. If the dissipation is maintained all the way to the photosphere, it will lead to an effective luminosity $L_\\gamma\\sim\\epsilon_d L_0$ and a temperature (Rees \\& \\Meszaros 2005) \\begin{equation} T_\\gamma=\\epsilon_d^{1/4}(r/r_s)^{-1/2}T_0=200 \\epsilon_d^{1/4} r_{11}^{-1/2} \\Gamma_2^{1/2} L_{0,52}^{1/4}{\\rm keV}, \\end{equation} where $\\epsilon_d$ is the dissipation efficiency, $T_0$ is the initial temperature of the fireball outflow and $\\Gamma$ is the bulk Lorentz factor of the outflow. The internal shock  occurs at $R\\simeq 2\\Gamma r_s =6\\times10^{10} \\alpha\\Gamma_2^2 (M/10 M_\\odot) {\\rm cm} $, for which the optical depth to Thomson scattering by the baryon-related electrons is $\\tau_{\\rm T}=\\sigma_{\\rm T} {L_k}/({4\\pi R \\Gamma^3 m_p c^3})=120 L_{k, 52} R_{11}^{-1} \\Gamma_2^{-3}$, where $\\sigma_{\\rm T}$ is the Thomson scattering cross section and $L_k$ is the kinetic energy luminosity.  The photosphere is further out, at radius $R_{ph}=1.2\\times10^{13} L_{k,52}\\Gamma_2^{-3} {\\rm cm}$. A detailed calculation taking into account of the electron/positron cooling and the Comptonization effect leads to a quasi-thermal emission which peaks at energy $\\sim300-500{\\rm keV}$ for dissipation at Thomson optical depth of $\\tau_{\\rm T}\\sim10-100$ (Pe'er et al. 2006). This temperature is consistent with the observed peak energies of prompt gamma-ray emission of a majority of GRBs.  Assuming that a fraction of $\\epsilon_B\\simeq 0.1$ of the shock internal energy is converted into magnetic fields, we have a magnetic field $B'=2.5\\times10^7 \\epsilon_{B,-1}^{1/2}L_{k,52}^{1/2} R_{11}^{-1} \\Gamma_2^{-1} {\\rm G}$. Protons  accelerated by internal shocks are assumed to have a spectrum $dn/d\\varepsilon_p\\sim \\varepsilon_p^{-p}$ with $p\\simeq2$, as often assumed for non-relativistic or mildly-relativistic shock acceleration. The maximum proton energy is set by comparing the acceleration time scale $t'_{acc}=\\alpha \\varepsilon'_p/(e B' c)=4.4\\times10^{-12} \\alpha (\\frac{\\varepsilon'_p}{1 {\\rm GeV}}) \\epsilon_{B,-1}^{-1/2}L_{k,52}^{-1/2} R_{11}\\Gamma_2 \\,{\\rm s}$ with the energy-loss time scales. The synchrotron loss time is $t'_{syn}=6\\pi m_p^4 c^3/(\\sigma_{\\rm T}m_e^2 {\\varepsilon'_p} B'^2)=10^{-4} \\epsilon_{B,-1}^{-1}L_{k,52}^{-1} R_{11}^2\\Gamma_2^2 (\\frac{\\varepsilon'_p}{10^{8} {\\rm GeV}})^{-1} \\,{\\rm s}$. Assuming that the sub-photosphere emission at the dissipation site peaks at $\\varepsilon_\\gamma=300{\\rm keV}$ with a thermal-like spectrum, the number density of photons in the comoving frame is $ n'_{\\gamma}={L_\\gamma}/({4\\pi R^2 \\Gamma^2 c \\varepsilon'_{\\gamma}})=5\\times10^{21} L_{\\gamma,51} R_{11}^{-2} \\Gamma_2^{-1} ({\\varepsilon_\\gamma}/{300 {\\rm keV}})^{-1} {\\rm cm^{-3}}$. The $p\\gamma$ cooling time is approximately $t'_{p\\gamma}={1}/({\\sigma_{p\\gamma} n'_{\\gamma} c K_{p\\gamma}})=10^{-4}L_{\\gamma,51}^{-1} R_{11}^2 \\Gamma_2(\\frac{\\varepsilon_\\gamma}{300 {\\rm keV}})\\, {\\rm s}$, where $K_{p\\gamma}\\simeq 0.2$ is the inelasticity and $\\sigma_{p\\gamma}=5\\times10^{-28} {\\rm cm^{2}}$ is the peak cross section at the $\\Delta$ resonance. By comparison with the synchrotron loss time, it is found that the most effective cooling mechanism for protons is the $p\\gamma$ process for protons with energies above the $p\\gamma$ threshold, but below $\\sim10^8 {\\rm GeV}$. Equating $t'_{acc}=t'_{p\\gamma}$, we obtain the maximum proton energy in the shock comoving frame \\begin{equation} \\varepsilon'_{p, max}=10^7 \\alpha_1^{-1} \\epsilon_{B,-1}^{1/2}L_{k,52}^{1/2}L_{\\gamma,51}^{-1}R_{11}(\\frac{\\varepsilon_\\gamma}{300 {\\rm keV}}) \\, {\\rm GeV} . \\end{equation} ",
        "conclusions": "Waxman \\& Bahcall (1997) as well as later works have studied the neutrino emission from the photomeson process during the prompt internal shocks of GRBs, assuming that the radiation in the shock region has a broken power-law nonthermal spectrum. It was found that the neutrino emission peaks at energies above $100$ TeV. Towards lower energies, the neutrino emission intensity decreases as $\\varepsilon_{\\nu}^2 \\Phi_{\\nu}\\sim \\varepsilon_\\nu$ (Waxman \\& Bahcall 1997). However, if internal shocks, especially at the early stage of the prompt emission, occur below the photosphere, a quasi-thermal spectrum will arise. In this {\\it Letter}, we have discussed the neutrino emission associated with the dissipative photosphere that produces such prompt thermal emission. We find that $pp$ interaction process becomes important for shock-accelerated protons and provides a new neutrino component, which dominates at energies below tens of TeV. The neutrino emission from photopion process of protons interacting with the sub-photosphere radiation could be significantly suppressed due to radiative cooling of secondary pions, when the dissipation radius is relatively small. Nevertheless, the total contribution by photopion process will not be suppressed since the shock dissipation could be continuous and occur at large radii as well. Although TeV neutrinos may also be produced during the early precursor stage of GRB, i.e. before the jet breaking out the progenitor star (e.g. Razzaque et al. 2004; Ando \\& Beacom 2005), we want to point out that the TeV neutrino component discussed here can be distinguished from them, because in our case the neutrino emission is associated in time with the prompt emission.   {"
    },
    "0807/0807.0773_arXiv.txt": {
        "abstract": "{ The WMAP satellite has made available high quality maps of the sky in five frequency bands ranging from 22 to 94 GHz, with the main scientific objective of studying the anisotropies of the Cosmic Microwave Background (CMB). These maps, however, contain a mixture of emissions from various astrophysical origins, superimposed on CMB emission.}{The objective of the present work is to make a high resolution CMB map in which contamination by such galactic and extra-galactic foregrounds, as well as by instrumental noise, is as low as possible.}{The method used is an implementation of a constrained linear combination of the channels with minimum error variance, and of Wiener filtering, on a frame of spherical wavelets called needlets, allowing localised filtering in both pixel space and harmonic space.}{We obtain a low contamination low noise CMB map at the resolution of the WMAP W channel, which can be used for a range of scientific studies. We obtain also a Wiener-filtered version with minimal integrated error.}{The resulting CMB maps offer significantly better rejection of galactic foregrounds than previous CMB maps from WMAP data. They can be considered as the most precise full--sky CMB temperature maps to-date.} ",
        "introduction": "The WMAP satellite is one of the most successful experiments dedicated to mapping the Cosmic Microwave Background (CMB). The all-sky maps obtained in the WMAP five frequency bands, in temperature and polarisation, offer the best data set to-date for making a sensitive all-sky map of the CMB anisotropies.  The CMB, however, is not the only source of emission at WMAP frequencies. Diffuse galactic emission from several processes contaminates the maps with an amplitude roughly proportional to the cosecant of the galactic latitude, compromising the observation of the CMB close to the galactic plane. In addition, a background of radio and infrared compact sources, galactic or extra-galactic,  contributes to the total emission even at high galactic latitudes.   Component separation consists in separating one or more of these sources of emissions  from the others in the data.  One of the main objectives of CMB experiments is the measurement of the CMB angular power spectrum $C_\\ell$ which, with the assumption of statistical isotropy, can be estimated on a fraction of the sky. For this reason, many ground-based and balloon-borne experiments have concentrated their observations in ``clean'' regions of the sky, where galactic emission is low enough to impact negligibly the observations. For power spectrum estimation from full-sky observations, a safe approach consists in masking regions at low galactic latitude, and estimating power spectra on the cleanest regions of the sky. The impact of extragalactic point sources (evenly spread on the sky) on power spectrum estimates can be evaluated and corrected for using ancillary data (catalogues of known point sources and priors on the statistical distribution of sources).  Besides the power spectrum, the CMB map itself is interesting for several additional purposes: \\begin{itemize} \\item As a CMB template, to be subtracted from millimetre--wave observations when the scientific focus is on other emissions, or to be used for the calibration of other instruments; \\item To assess the statistical isotropy of the CMB and check the homogeneity and isotropy of the Universe on the largest scales; \\item To search for signatures of non-trivial topology, as that of a multi-connected universe \\citep{2006MNRAS.369..240A,2007A&A...476..691C,2007PhRvL..99h1302N}; \\item To search for correlations of the CMB map with other emissions \\citep{2004ApJ...608...10N,2004MNRAS.350L..37F,2006MNRAS.372L..23C,2006MNRAS.369..645C,2006PhRvD..74d3524P,2007arXiv0704.0626M,2007MNRAS.377.1085R}; \\item To search for signatures of non Gaussianity in the CMB \\citep{2004ApJ...613...51M,2004ApJ...609...22V,2005MNRAS.362..826C,2006MNRAS.369..819C,2006MNRAS.369.1858M,2006MNRAS.365..265T,2006MNRAS.371L..50M,2007arXiv0706.2346W,2007arXiv0712.1148Y}. \\end{itemize}  Several CMB maps obtained with the WMAP data are available for such research projects. The WMAP team has released part--sky foreground--reduced maps in the Q, V and W bands, and maps obtained by an Internal Linear Combination (ILC) of all WMAP channels \\citep{2007ApJS..170..288H}. \\citet{2003PhRvD..68l3523T} have produced CMB maps with WMAP one year data, and subsequently with three year data, based on an harmonic space ILC method. \\citet{2004ApJ...612..633E} have an alternate version of the ILC CMB map at 1 degree resolution on one year data. \\citet{2007arXiv0709.1037E} use a Gibbs sampler to draw 100 realisations of the  CMB under the posterior of  a model of CMB and foregrounds; their estimated CMB map is the average of these realisations for three year data. On three year data again, \\citet{2007ApJ...660..959P} use an ILC technique on 400 different pixel ensembles, selected by the spectral index of the foreground emission as estimated by the WMAP team using a Maximum Entropy Method (MEM). More authors have addressed component separation on WMAP three year data and produced versions of a ``clean'' CMB map \\citep{2007MNRAS.374.1207M,2007arXiv0706.3567S,2007arXiv0707.0469B}. More recently, \\citet{2008arXiv0803.1394K} have obtained a CMB map from WMAP five year data.  All available maps suffer from limitations, some of which result from specific choices in the way the CMB map is produced. Several of these maps, for instance, do not fully exploit the resolution of the original observations. Some focus on cleaning the CMB from foregrounds at high galactic latitude, and are significantly contaminated by foregrounds in the galactic plane. Some are not full sky CMB maps. Finally, not all of the available maps have well characterised noise and effective beam. All these limitations impact their usefulness  for accurate CMB science.  In this paper, we address the problem of making a CMB temperature map which has the following properties : \\begin{itemize} \\item being full sky; \\item being as close as possible to the true CMB (minimum variance of the error) everywhere on the sky, and on all scales; \\item having the best resolution possible; \\item having well-characterised beam and noise. \\end{itemize}  In the following, we first review these requirements and their impact on a CMB cleaning strategy (Section~\\ref{sec:generalities}). We then review and compare available maps in Section~\\ref{sec:existing}. In Section~\\ref{sec:method}, we describe and explain our ILC needlet method. The approach is tested and validated on realistic simulated data sets (Section~\\ref{sec:simulation}) before applying it to WMAP data (Section~\\ref{sec:application}). We then compare our CMB maps to the other existing maps, discuss the results, and conclude.  This paper considers only \\emph{temperature} maps. ",
        "conclusions": "In this paper, we have described a new approach for implementing CMB extraction in WMAP data, using the ILC method on a needlet frame. Tests on simulations show excellent performance of the method, thanks to localisation both in pixel space and in harmonic space. Localisation in pixel space allows the ILC weights to adapt themselves to local conditions of foreground contamination and noise, whereas localisation in harmonic space allows to favour foreground rejection on large scales (where foregrounds dominate the total error) and noise rejection on small scales (where foregrounds are negligible but where, after beam deconvolution, the relative noise level between the various WMAP channels varies a lot as a function of scale).  Needlets permit to vary the weights smoothly on large scales, and rapidly on small scales, which is not possible by cutting the sky in zones prior to any processing.  As a further improvement to previous work on WMAP data, we include a dust template in the set of analysed observations. This is motivated by the fact that on the smallest scales, observed with reasonable signal to noise ratio by the W channel only, dust emission contributes a significant fraction of the total reconstruction of the map. Using the IRIS 100 micron map as an additional observation enables the ILC to reduce the final contamination by dust --thanks to correlations of dust emission between the W channel and the 100 micron map. Special care was also taken to subtract from the data, prior to making the ILC, a number of strong point sources.  As discussed at length in the main text and in the appendix, the implementation of a filter (the ILC) based on empirical estimates of covariance matrices leads to a bias. This is not particular to our map, but is the case for any ILC map. We have estimated the level of this systematic effect, both analytically and numerically, to be at the level of a few per cent on all scales. Our simulation tool permits to make accurate estimates of the amplitude of the effect, if needed for any scientific exploitation of the NILC5 map.  The application of the method to WMAP one year and three year data (in addition to five year data) permits to compare the needlet ILC solution to previous work. Our map is seen to be at least as good as others on large scales, while being significantly less contaminated by residual foregrounds and noise than others on small scales, in particular in the vicinity of the galactic plane.  The application of the method to WMAP five year data yields a CMB map which we believe to be the cleanest full sky map of the CMB to date. Contamination by noise, and the power loss due to the use of the ILC method, are characterised by means of Monte Carlo simulations. The map is available for download on the ADAMIS web site,\\footnote{http://www.apc.univ-paris7.fr/APC\\_CS/Recherche/Adamis/cmb\\_wmap-en.php} and can be used for a variety of science projects relying on accurate maps of the CMB."
    },
    "0807/0807.2776_arXiv.txt": {
        "abstract": "We have explored radial color and stellar surface mass density profiles for a sample of 85 late-type spiral galaxies with deep (down to $\\sim$27 mag arcsec$^{-2}$) SDSS $g'$ and $r'$ band surface brightness profiles. About $90\\%$ of the light profiles have been classified as broken exponentials, exhibiting either truncations (Type~II galaxies) or antitruncations (Type~III galaxies). The color profiles of Type~II galaxies show a ``U shape'' with a minimum of $(g' - r') = 0.47\\ \\pm\\ 0.02$ mag at the break radius. Around the break radius, Type~III galaxies have a plateau region with a color of $(g' - r') = 0.57\\ \\pm\\ 0.02$. Using the color to calculate the stellar surface mass density profiles reveals a surprising result. The breaks, well established in the light profiles of the truncated galaxies, are almost gone, and the mass profiles now resemble those of the pure exponential (Type~I) galaxies. This result suggests that the origin of the break in Type~II galaxies is more likely due to a radial change in stellar population than being associated with an actual drop in the distribution of mass. Type~III galaxies, however, seem to preserve their shape in the stellar mass density profiles. We find that the stellar surface mass density at the break for truncated galaxies is $13.6 \\pm 1.6 \\ {M}_{\\sun}$pc$^{-2}$ and for the antitruncated ones is $ 9.9 \\pm 1.3 \\ {M}_{\\sun}$pc$^{-2}$ for the antitruncated ones. We estimate that the fraction of stellar mass outside the break radius is $\\sim$15$\\%$ for truncated galaxies and $ \\sim$9$\\%$ for antitruncated galaxies. ",
        "introduction": "\\label{intro}  Our picture regarding the diversity of the radial surface brightness profiles of spiral galaxies has changed greatly since the early work of Patterson (1940) and de Vaucouleurs (1958), who showed that the disks of spiral galaxies generally follow an exponential decrease in their radial surface brightness profile. Nowadays, this view has become clearly insufficient, as not all the disk galaxies (indeed, only the minority for late-types) are well described with a single exponential fitting function as shown in several recent studies (Erwin et al.~2005; Pohlen \\& Trujillo 2006, hereafter PT06; Florido et al. 2006, 2007; Erwin et al. 2008), where they have identified three basic classes of surface brightness profiles depending on an apparent break feature or lack of one: (1) the pure exponential profiles (Type I) with no breaks, (2) Type II with a ``downbending break'' (revising and extending a previous classification introduced by Freeman [1970] to include the so-called truncations of the stellar populations at the edge of the disk discovered by van der Kruit [1979]) and (3) a completely new class (Type III), also described by a broken exponential but with an upbending profile. The latter, discovered by Erwin et al. (2005), is also termed antitruncated.  PT06 explored a sample of nearby late-type galaxies using the Sloan Digital Sky Survey to create a statistically representative set of radial surface brightness profiles. They found that about 60$\\%$ of the spirals are truncated (Type~II), 30$\\%$ are antitruncated (Type~III), and only 10$\\%$ have no detectable breaks (Type~I).  Still, little is known about the nature of the breaks or about the presence of stars beyond that feature. In such low-density ($\\le$ 10 $ {M}_{\\sun}$pc$^{-2}$) environments at the galaxy peripheries, current star formation theories forbid efficient star formation (Kennicutt 1989; Elmegreen \\& Parravano 1994; Schaye 2004). However, UV observations (Gil de Paz et al. 2005; Thilker et al. 2005) have shown ongoing star formation in these outer regions. In addition, it is clear that there are large number of stars in the outskirts of galaxies (see, e.g., PT06).  Several theories have investigated the formation of breaks in case of the Type~II morphology. Proposed models to explain the existence of truncations in stellar disks can be grouped into two branches depending on the relevant mechanism that causes the break: (a) models related to angular momentum conservation in the protogalactic cloud (van der Kruit 1987) or angular momentum cutoff in cooling gas (van den Bosch 2001), and (b) models that attribute the existence of breaks to star formation thresholds (Kennicutt 1989; Elmegreen \\& Parravano 1994; Schaye 2004). In agreement with this last scenario, Elmegreen \\& Hunter (2006) suggested a multicomponent star formation that would result in a double exponential surface brightness profile as observed for Type~II galaxies. Recent developments, some of them combining pieces of the aforementioned scenarios, conclude that secular evolution driven by bars or spiral arms or even clump disruptions can result in truncated exponential profiles (Debattista et al. 2006; Bournaud al. 2007; Ro\\v skar et al. 2007; Foyle et al. 2008). Magnetic fields have also been considered to explain the existence of truncations (Battaner et al. 2002). On the other hand, the Type~III morphology is proposed to be explained by a tidal stripping within a minor merger (Pe\\~narrubia et al. 2006; Younger et al. 2007), by a bombardment of the disk with dark matter subhalos (Kazantzidis et al. 2007) or by a high eccentricity flyby of a satellite galaxy (Younger et al. 2008).  In this Letter we show for the first time color and stellar surface mass profiles of a large sample of galaxies (the PT06 sample) to quantify the stellar mass density at the break position and the fraction of stellar mass beyond the break. These values become important while comparing observations to the results of numerical simulations for the outer-disk formation. In order to fully understand the galaxy formation and evolution process, it is necessary to perform a detailed study of the stellar population properties at the galaxies' outskirts. This kind of study gives insight into how the star formation progresses in the different parts of the disks, providing clues on the stellar mass buildup process. ",
        "conclusions": "How can the results found in this work be used to constrain the current models for the formation of breaks in the surface brightness profiles of disk galaxies?  In the case of Type~II galaxies, the traditional pictures of break formation (see \\S \\ \\ref{intro}) grouped into two families: angular momentum versus thresholds of star formation. Neither of these two ideas taken at face value can explain why we find so many stars beyond the break radius, so we will not go into more detail on these individual models. However, it is important to stress that the new rendition of models has been able to naturally explain the existence of stars beyond the break radius, and even more, the exponential nature of the shape of the surface brightness beyond this feature. In particular, in the case of the Ro\\v skar et al. (2008) model in which the breaks are the result of the interplay between a radial star formation cutoff and a redistribution of stellar mass by secular processes, a natural prediction is the existence of a minimum in the age of the stellar population at the break position, and a further aging (and consequently, a likely reddening) of these stars as we move farther and farther away from the break radius. This is in qualitative agreement with what we see in our color profiles for this kind of galaxy. However, what these models have not been able to reproduce is the absence of a clear break in the stellar mass density profile.  The near absence of a break in the stellar surface mass density profile for our galaxies gives a strong indication that the behavior of the surface brightness profile outside the break is basically due to a change in the ingredients of the stellar population. If the shape of the color we see is not caused by a change in metallicity, this behavior could be explained as a natural consequence of stochastic migration of young stars from the inner parts of the disk to the outskirts (Ro\\v skar et al. 2008). This will result in an age gradient where the oldest stars are the dominant component in the outskirts of the disks.  Unfortunately, other models (like Bournaud et al. 2007 and Foyle et al. 2008) that are capable of creating stellar mass density profiles that resemble the Type~II one we have found here (i.e., without a clear break) do not provide any prediction on the color (age) distribution the stars should have along the radial range.  Nevertheless, the fact that the stellar mass density of the break for this type of galaxy with $\\sim$13 ${M}_{\\sun}$pc$^{-2}$ is so close to the gas density threshold prediction of $\\sim$10 ${M}_{\\sun}$pc$^{-2}$ makes the case for a stellar population origin for the surface brightness break even stronger (with a 100\\% efficiency of transforming gas to stars).  Evidence for the same color phenomenology for Type~II galaxies also at high redshift is presented by Azzollini et al. (2008). They have shown that a similar minimum in the color profile can be found, at least, up to z $\\sim$ 1 and that the main source of the scatter of the color profiles is caused by the different stellar mass (in our case, absolute magnitude) of the galaxies in their sample. Following these findings, we can conclude that once the absolute magnitude of a galaxy is fixed, the color profiles within a given Type~(I, II or III) of galaxy are strikingly similar. A more massive galaxy has a redder global color but the same shape of the color gradient.  Combining the results found in Azzollini et al. (2008) with what we find here, one is tempted to claim that both the existence of the break in Type~II galaxies and the shape of their color profiles are long-lived features in the galaxy evolution, because it would be hard to imagine how the above features could be continuously destroyed and recreated maintaining the same properties over the last $\\sim$8 Gyr.  In the case of Type~III galaxies, the situation is less clear. On the one hand, our sample is smaller than in the case of Type~II galaxies and consequently our results less robust. And on the other hand, the theoretical models are less elaborated than for truncated galaxies, and no clear predictions have been made in particular for the color profiles. Taking into account that the shape of the stellar mass density profile does not differ too much from what we see in the surface brightness profiles, we are inclined to think that our data do not favor a sole origin in stellar population changes for this type of galaxy but an authentic change in the amount of stars from the exponential continuation of their inner region. It is interesting to note that in most of the proposed ideas summarized in \\S \\ \\ref{intro} to explain this kind of galaxy, the origin of the stars in the periphery are linked to a dynamical (in some case external) origin. So what we are seeing for these galaxies is maybe a combination of star formation combined with an infall of new stars from a external (satellite) source.  To corroborate our results, we plan to increase the number of galaxies as well as the number of observed filters in our next study."
    },
    "0807/0807.2892_arXiv.txt": {
        "abstract": "A general method for describing magnetic reconnection in arbitrary three-dimensional magnetic configurations is proposed. The method is based on the field-line mapping technique previously used only for the analysis of magnetic structure at a given time. This technique is extended here so as to analyze the evolution of magnetic structure. Such a generalization is made with the help of new dimensionless quantities called ``slip-squashing factors\". Their large values define the surfaces that border the reconnected or to-be-reconnected magnetic flux tubes for a given period of time during the magnetic evolution. The proposed method is universal, since it assumes only that the time sequence of evolving magnetic field and the tangential boundary flows are known. The application of the method is illustrated for simple examples, one of which was considered previously by Hesse and coworkers in the framework of the general magnetic reconnection theory. The examples help us to compare these two approaches; they reveal also that, just as for magnetic null points, hyperbolic and cusp minimum points of a magnetic field may serve as favorable sites for magnetic reconnection. The new method admits a straightforward numerical implementation and provides a powerful tool for the diagnostics of magnetic reconnection in numerical models of solar-flare-like phenomena in space and laboratory plasmas. ",
        "introduction": "\\label{s:intro}  In spite of the progress reached so far in understanding of the key role of magnetic reconnection in many phenomena of space and laboratory plasmas, its quantitative description in realistic three-dimensional (3D) configurations still remains a challenge \\citep{Priest2000, Schindler2006}. This concerns, in particular, the global aspects of reconnection such as an identification of reconnected magnetic fluxes and related structural changes in solar coronal configurations. There are basically two fundamental reasons why this important problem is so complicated in the 3D case.  First, by analogy with the two-dimensional (2D) case, one could assume that partitioning the configurations into distinct interacting magnetic fluxes by separatrix surfaces (SSs) and estimating their time variation might help to resolve this problem. And, indeed, this idea works very well whenever such SSs are due to the presence of magnetic null points (NPs) in the solar corona---the NP SSs are boundless surfaces, and so they provide the required partition. However, the situation is different if the SSs are due to the presence of so-called bald patches (BPs), which are the segments of the photospheric polarity inversion line where coronal field lines touch the photosphere \\citep{Seehafer1986a, Titov1993, Bungey1996, Delannee1999, Aulanier2002a}. The BP SSs are usually not boundless and, therefore, not able to provide the required partition of the respective configurations. It is also important that they are not exceptional structures, but are highly likely to be present in practice in the complex configurations that exist in the solar atmosphere \\citep{Titov1993, Bungey1996, Titov1999, Delannee1999, Aulanier2002a, Aulanier2002b, Pariat2004}. Moreover, the set of 3D structural features where reconnection may occur is even wider than mentioned above, because it includes also quasi-separatrix layers (QSLs) \\citep{Demoulin1996, Priest1995}. Similar to BP SSs, these QSLs also do not provide an exact partition of configurations into distinct magnetic fluxes.  The second fundamental difficulty in quantitative description of 3D reconnection is related to magnetic diffusion and the breakdown of frozen-flux behavior. Being an inherent part of the reconnection itself, magnetic diffusion in 3D can occur throughout the plasma volume in evolving configurations. This is of particular importance for numerical MHD simulations, which may presently be performed only with a relatively high numerical or physical diffusion compared to real solar conditions. Deviations from the frozen-in law condition are accumulated in this case throughout the volume to form a background against which it becomes difficult to distinguish between magnetic diffusion and reconnection.  Nevertheless, significant progress in understanding 3D reconnection has been achieved in the framework of general magnetic reconnection (GMR) theory \\citep{Schindler1988, Hesse1988, Hesse1991, Hesse1993}, which is focused on the analysis of magnetic field evolution constrained only by Maxwell's equations and Ohm's law. On the basis of this theory, \\citet{Hesse2005} have recently shown that, regardless of the presence of certain structural magnetic features, the time change of the reconnected magnetic flux is always directly related to the maximum of the field-line integral of parallel electric field. For a plasma-magnetic configuration contained in a closed volume, this means that the change of the magnetic connection between the boundary plasma elements is fully controlled by the voltage drop along these lines. It is clear, however, that such a change of the magnetic connection can always be found without calculating the electric field if the magnetic field in the volume and the plasma flows on the boundary are known at every moment. As will be shown, the theory based on this premise complements the GMR theory by providing new intriguing opportunities for the analysis of magnetic reconnection. The aim of this paper is to develop a general mathematical formalism that allows one to describe the change of magnetic connectivity and identify the reconnecting magnetic fluxes in any evolving configuration. The new theory has been outlined earlier by \\citet{Titov2007d}. This paper, however, provides its full exposition.  The new formalism can also be viewed as a generalization of the theory of coronal magnetic connectivity, whose spatial variation is measured with the help of the so-called squashing factor $Q$. This quantity was introduced first for closed magnetic configurations with a plane photospheric boundary \\citep{Titov1999a, Titov2002a, Titov2002} and then extended to other configurations with arbitrary shapes of the boundaries \\citep{Titov2007a}. The $Q$ factor allows one to identify in a given configuration at a fixed time all the structural magnetic features, such as SSs and QSLs as well as their hybrids. Here we generalize this quantity in order to characterize the evolution of magnetic connectivity by applying the $Q$ factor to so-called slip mapping---the mapping that describes the slippage between boundary plasma elements and footpoints of magnetic field lines.  It was previously noticed by \\citet{Forbes1984a} for 2D magnetic configurations that the speed with which the photosphere is swept by a magnetic separatrix line allows one to estimate the electric field at the reconnection site. \\citet{Forbes2000} have also proposed strong arguments that this principle can be extended to 3D as well, excluding the above-mentioned configurations with QSLs, if one focuses only on the amount of magnetic flux reconnected up to a particular instant of time. Our paper provides exactly this type of generalization of the principle for 3D configurations with QSLs included.  Section \\ref{s:evol} describes the general theory of evolution of magnetic connectivity, which is illustrated in section \\ref{s:exmpl1} for the example of magnetic evolution proposed earlier by \\citet{Hesse2005}. {{\\hl Section \\ref{s:exmpl2} presents another example of magnetic evolution in a current layer configuration to show the difference between reconnection and magnetic diffusion. }} In section \\ref{s:min}, we discuss the general properties of magnetic minimum points, which happen to be closely related to the considered theory. The obtained results are summarized in section \\ref{s:sum}. ",
        "conclusions": "}  We have developed a general theory for the evolution of magnetic connections between boundary plasma elements in an arbitrary 3D magnetic configuration with an arbitrarily shaped boundary. The theory assumes that at each instant the magnetic field is known everywhere throughout the volume domain including the boundary, where the tangential flow of plasma elements is also known.  The major quantity that is used for characterizing the magnetic connectivity and its evolution is the squashing factor $Q$ \\citep{Titov1999a, Titov2007a}. After being applied to the magnetic field-line mapping, the $Q$ factor allows one to identify in a given configuration all separatrix and quasi-separatrix surfaces, as well as their hybrids. We have demonstrated first how the $Q$ factor is determined for any ideal evolution of the magnetic configuration in terms of the initial magnetic field and tangential boundary flows without reference to the magnetic field at other moments in time. This result is itself of great importance for the further development of the theory of current-layer formations in magnetic configurations with a highly conducting plasma via its motion at the boundary.  In order to describe the nonideal evolution of magnetic connectivity, we have introduced the so-called ``slip mapping\", which characterizes the change of magnetic connections between the boundary plasma elements due to violation of the frozen-in law in the plasma volume. This mapping can be defined either forward or backward in time to give correspondingly slip-forth or slip-back mappings, which operate on the initial or final configurations, respectively, within a chosen time interval. By calculating the squashing factor for these new mappings, we have defined new quantities called slip-forth and slip-back squashing factors, $Q_{\\mathrm{sf}}$ and $Q_{\\mathrm{sb}}$, respectively. The large values of these quantities identify the reconnecting magnetic flux tubes, called reconnection fronts (RFs), at the initial and final time instances. The boundary areas that are swept by RFs within a chosen time interval correspond to the footprints of reconnecting magnetic flux tubes. In this way, the $Q_{\\mathrm{sf}}$ factor determines in the initial configuration the to-be-reconnected magnetic flux tubes, while the $Q_{\\mathrm{sb}}$ factor determines in the final configuration the reconnected magnetic flux tubes. The dependence of the properties of the RFs on tangential boundary flows is also analyzed.  {{\\hl The general results are illustrated with two specific examples based on evolving magnetic configurations. The first of these was considered earlier by \\citet{Hesse2005} in the framework of the general magnetic reconnection (GMR) theory \\citep{Schindler1988, Hesse1988, Hesse1991, Hesse1993}. Revisiting this example, we have demonstrated that our approach is consistent with GMR theory but complements it in several important aspects. The second example describes a magnetic evolution that is driven by a thin current layer with a field-line voltage drop at a small patch inside it. This example enables us to show that the localization of such a nonideal process in a small region may indeed cause the structural changes that are characteristic for reconnection rather than for magnetic diffusion. The latter example, in particular, demonstrates the formation of magnetic flux tubes with narrow, ribbon-like footprints, which is a typical feature of reconnection process.   In comparison with GMR theory, our approach handles regions with small or large amounts of reconnected flux on an equal footing. This is because the slip-squashing factors, in contrast to the voltage drop of GMR, are dimensionless geometrical quantities measuring only a relative spatial rate of magnetic slippage at the boundary. For the same reason, the slip-squashing factors discriminate between regions of reconnection and simple diffusion much better than the voltage drop method. Moreover, since the slip mappings are composed from the field-line mappings and both types of mappings are characterized by the same quantity, the analysis of reconnection based on our theory is intimately related to the analysis of magnetic structure. As we have demonstrated here, the latter provides intriguing opportunities for analyzing magnetic reconnection in 3D configurations with an unprecedented level of detail. The use of slip-squashing factors opens new perspectives for the comparing reconnecting flux tubes predicted by the theoretical models with the different morphological features observed in the solar flares and coronal mass ejections [e.g. brightening coronal loops, chromospheric kernels and ribbons, see the recent review by \\citet{Demoulin2007}].   Our analysis of these examples also reveals the special role of hyperbolic and cusp minimum points in 3D magnetic field configurations. The elemental magnetic flux tubes passing in the vicinity of such points may become extremely squashed, which implies the formation of the respective QSLs. We have extended these examples by studying the field-line structure in the neighborhood of generic magnetic minima and confirmed this particular result. }} Hence, similar to magnetic null points, hyperbolic and cusp minima provide an important localization of plausible reconnection sites. More detailed investigations in this direction promise to lead to even more interesting results.  It should be emphasized that, similar to GMR theory, our approach is applicable to the analysis of reconnection in magnetic configurations described not only by MHD, but also by any other plasma model. The algorithm of calculating slip-squashing factors admits a straightforward implementation and it is highly parallelizable. Our experience shows that, even without parallelization, the squashing factor may be computed for known magnetic field data with a modest demand on computing resources. Thus, the whole approach is practical enough to pursue the implementation of the proposed algorithms for calculating slip-squashing factors for the solar corona. The solar physics group at SAIC is presently working on this project."
    },
    "0807/0807.2540_arXiv.txt": {
        "abstract": "Numerical simulations of radiative processes in magnetized compact sources such as hot accretion disks around black holes, relativistic jets in active galaxies and gamma-ray bursts are complicated because the particle and photon distributions span many orders of magnitude in energy, they also strongly depend on each other, the  radiative processes behave significantly differently depending on the energy regime, and finally due to the enormous difference in the time-scales of the processes. We have developed a novel computer code for the time-dependent simulations that overcomes these problems.  The processes taken into account are Compton scattering, electron-positron pair production and annihilation, Coulomb scattering as well as synchrotron emission and absorption. No approximation has been made on the corresponding rates.   For the first time, we solve coupled integro-differential kinetic equations for photons and electrons/positrons without any limitations on the photon   and lepton energies. A numerical scheme is proposed to guarantee energy conservation when dealing with  synchrotron processes in electron and photon equations. We apply the code to model non-thermal pair cascades in the blackbody radiation field, to study the synchrotron self-absorption as particle thermalization mechanism, and to simulate time evolution of stochastically heated pairs and corresponding synchrotron self-Compton photon spectra which might be responsible for the prompt emission of gamma-ray bursts. Good agreement with previous works is found in the parameter regimes where comparison is feasible, with the differences attributable to our improved treatment of the microphysics. ",
        "introduction": "Spectral energy distributions of a number of compact, magnetized, high-energy sources such as relativistic jets from active galaxies, gamma-ray bursts, black hole accretion disk-coronae are strongly affected and shaped by Compton scattering, synchrotron radiation and electron-positron pair production \\citep[see e.g.][]{G99,gcc02,ZG04,SP04}.  Understanding the physical conditions in these sources requires detailed modeling of the interactions between the particles and photons, which is not an easy task. The basic problem is that we cannot compute radiative processes from a given a priori lepton distribution (e.g. Maxwellian or a power-law), because it  depends strongly on the radiation field, which in its turn is determined by the particle distribution. Another problem is that the time-scales for various processes differ by orders of magnitude. The energy range of particles and photons responsible for the emission also spans many orders of magnitude, with different processes making dominant contributions to the emergent spectrum in different bands. One of the main difficulties in calculating radiative processes over a wide range of energies is that a particular radiative process may behave significantly differently depending on the energy regime, the most well-known example of such processes being also the most important in relativistic plasma, namely Compton scattering. Depending on the energies of the interacting particles, an electron or a photon can lose (or gain) a significant or negligible fraction of its initial energy in one scattering. The former case has to be accounted for by the integral scattering terms in the kinetic equations, while the latter necessitates the Fokker-Planck treatment.   The treatment of radiative processes in relativistic plasmas has been the subject of several works. There are two basic approaches: Monte Carlo methods \\citep[e.g.][]{SBS95,PS97} and solving the relevant kinetic equations \\citep[e.g.][]{LZ87,coppi92,coppi99,NM98,PW05}. Both have their own advantages and disadvantages. Monte Carlo treatment makes it easy to take into account radiative transfer effects, on the other hand it usually suffers from poor photon statistics at high energies. Another problem can arise at very low energies, where the optical thickness to synchrotron absorption can be enormous. In the kinetic theory approach photon statistics is not an issue, the sole difficulty lies in solving the relevant integro-differential equations. In this work we have chosen to follow the second approach.  Due to the difficulties in solving the exact Boltzmann equations of the kinetic theory, different simplifying approximations have been made in earlier works. They fall in three basic categories: ad hoc assumptions about the particle energy distributions, approximate treatment of different physical processes, and  simplified treatment of radiative transport. Various approximations invoked to simplify the treatment of radiative processes at the same time limit the range of their applicability. One commonly employed approximation concerns Compton scattering, which is assumed to take place in the Thomson regime \\citep[e.g.][ hereafter GHS98]{GHS98} and is accounted for by a simple cooling term in the electron equation. This sets two restrictions that the photon energy in the electron rest frame is smaller than the electron rest mass and the average photon energy is much lower than the electron kinetic energy. Otherwise all photons would not contribute to electron cooling, the higher energy ones being downscattered via Compton scattering. This means that one is unable to treat cases when Comptonization approaches saturation, which may be relevant at high compactnesses, and to study electron heating by external radiation. Other works account for Klein-Nishina corrections to the electron cooling rate, but still neglect the diffusive nature of the process when electron and photon energies are comparable \\citep{coppi92,MSC05,PW05}, which works towards establishing an equilibrium Maxwellian distribution. Another useful approximation, when the integral terms describing Compton scattering are  accounted for, is to consider ultrarelativistic electrons and very low energy photons \\citep[e.g.][]{ZDZ88,MSC05}. This, however, becomes increasingly inaccurate when the electrons cool to sufficiently low energies.  The cyclo-synchrotron process also exhibits qualitatively different behavior depending on the energy of the radiating particles. If the emitting particles are relativistic, the emission spectrum is smooth and can span several orders of magnitude in energy, while in the nonrelativistic case the energy is radiated at discrete cyclotron harmonics and most of this radiation might be strongly self-absorbed. In the first case, the radiating particle (electron or positron) mostly loses its energy in a continuous fashion, while in the second case it can gain energy by absorbing the  cyclo-synchrotron photons emitted by other particles. This process  is a dominant particle thermalization mechanism in compact magnetized sources \\citep{GGS88}. Its proper account requires accurate emissivities in the transrelativistic regime, because electron thermalization usually takes place at mildly relativistic energies. Some codes for computing radiative processes in relativistic plasma \\citep[e.g. {\\sc eqpair} described in][]{coppi92,coppi99} neglect this process completely as the electrons are {\\it assumed} to be thermal at low energies or account for thermalization   by Coulomb  collisions only \\citep{NM98}. In other approaches synchrotron thermalization is computed (GHS98), but Compton scattering is then considered only approximately.  Owing to the fact that proper treatment of transport processes for all types of particles would make the task prohibitively difficult, and partly due to our ignorance of the exact geometry of the problem, it is rather common practice to neglect radiative transport altogether \\citep[e.g.][]{LZ87,coppi92} and assume spatially homogeneous and isotropic particle distributions. In this case particle and photon loss from the system is modeled in terms of simple escape probabilities. We also follow this approach here.  In this paper we introduce and describe in details a numerical code that can deal with Compton scattering, synchrotron emission and absorption, electron-positron pair production and annihilation without limitations on the energies of the photons and electrons/positrons. We solve coupled  integro-differential kinetic equations describing time evolution of the photons and lepton distributions simultaneously. When necessary, the Fokker-Planck  differential terms are substituted instead of the  integral terms with coefficients computed exactly from the moments of the integral equation.  Particle thermalization by synchrotron self-absorption, Coulomb (M{\\o}ller) scattering as well as Compton scattering is considered. Extreme caution is taken when dealing with  synchrotron self-absorption, because of cancellation of large, almost equal terms, which can result in inaccuracies and huge energy sinks. Numerical simulations show that our code conserves energy with about 1\\% accuracy. We present an extensive testing of the code using some problems described previously in the literature. Processes involving bremsstrahlung can be easily added to the code, while for the conditions considered in the paper, they are not important. ",
        "conclusions": "We have developed a new computer code for simulations of the radiative processes in magnetized rarefied plasmas encountered in the vicinities of accreting black holes and relativistic jets in active galaxies and gamma-ray bursts.  We take into account Compton scattering, pair production and annihilation,  synchrotron processes and Coulomb scattering without any limitations on the energies of the photons and electrons/positrons. We solve coupled  integro-differential kinetic equations describing time evolution of the photon and electron/positron  distributions simultaneously. The equations contain both integral and up to second order differential terms. The Fokker-Planck  differential terms are substituted when necessary instead of the  integral terms with coefficients computed exactly from the moments of the integral equation. This allows us to study thermalization of the lepton distribution by Compton and Coulomb scattering and synchrotron self-absorption. Processes involving bremsstrahlung can be easily added to the code, while for the conditions considered in the paper they are not important.  The presented technique guarantees energy (and particle, when relevant) conservation with high accuracy which is especially important when dealing with strongly self-absorbed synchrotron radiation. The implementation of the Chang and Cooper scheme gives the correct shape of the particle distribution at low energies. The area of application of the code is enormous as it can deal with photons and leptons covering many orders of magnitude in momentum space, with no potential difficulty of extending it to even lower/higher energies. We present a number of test runs, where we consider problems previously solved by other methods. We compute non-thermal pair cascades, and study lepton thermalization by synchrotron self-absorption, as well as model the emission from the stochastically heated pairs that might have a relation to the prompt emission of gamma-ray bursts. We find a good agreement in the parameter space where comparison is feasible while the differences can be explained by our improved treatment of microphysics."
    },
    "0807/0807.2281_arXiv.txt": {
        "abstract": "We  report on  a statistical  study of  the 51  radio galaxies  at the millijansky  flux level  from the  Faint Images  of the  Radio  Sky at Twenty centimeters, including their optical morphologies and structure obtained  with the  Hubble Space  Telescope.  Our  optical  imaging is significantly  deeper ($\\sim\\!2$~mag) than  previous studies  with the superior angular resolution of space-based imaging.  We that find 8/51 (16\\%) of the radio sources have no optically identifiable counterpart to ${\\rm  AB}\\!\\sim\\!24$~mag.  For the  remaining 43 sources,  only 25 are  sufficiently resolved  in the  HST  images to  reliably assign  a visual classification: 15 (60\\%)  are elliptical galaxies, 2 (8\\%) are late-type spiral galaxies,  1 (4\\%) is an S0,  3 (12\\%) are point-like objects (quasars), and 4 (16\\%) are merger systems.  We find a similar distribution of optical types with measurements of the S\\'ersic index. The  optical  magnitude  distribution   of  these  galaxies  peaks  at $I\\!\\simeq\\!\\alli$~AB  mag, which is  $\\sim\\!3$~mag brighter  than the depth  of our  typical HST  field and  is thus  not due  to  the WFPC2 detection  limit.   This  supports  the  luminosity-dependent  density evolutionary  model,  where  the  majority  of  faint  radio  galaxies typically  have $L^*$-optical  luminosities and  a median  redshift of $z\\!\\simeq\\!0.8$  with   a  relatively  abrupt   redshift  cut-off  at $z\\!\\gtrsim\\!2$.   We  discuss  our  results  in the  context  of  the evolution of elliptical galaxies and active galactic nuclei. ",
        "introduction": "\\label{intro}  Faint  radio galaxies  are good  probes of  typical galaxies  at large distances since they lack  a significant contribution from non-thermal light   or   strong   emission   lines  in   their   optical   spectra \\citep[eg.][]{kro85,ben93,ham95a}.   Furthermore,  their selection  is relatively  unaffected  by dust  absorption  when  these galaxies  are identified  by their  synchrotron  emission.  In  the local  Universe, powerful  radio  sources  are  primarily  found  in  giant  elliptical galaxies,  which suggests  that radio  selection may  be  an efficient method  of identifying high-redshift  ellipticals.  However,  at lower radio  powers,  typically below  the  break  in  the radio  luminosity function, there  is an increasing fraction of  star-forming and active galaxies      \\citep[such       as      Seyferts      and      LINERS; eg.][]{con89,ben93,low97,hop03,bes05,afo06,mai08}.    Often   of   the galaxy identification  is not  based on deep,  high-resolution optical imaging, but rather on  broadband optical colors.  Such classification schemes are only reliable to the extent to which these colors directly map onto the Hubble sequence  at a given redshift.  Therefore, in this paper  we focus  on the  optical morphologies  and structure  of radio galaxies      with      fluxes     of      1~mJy$\\lesssim\\!S_{1.4~{\\rm GHz}}\\!\\lesssim\\!100$~mJy  (we  will  denote $S_{1.4}$  as  integrated radio  flux at  1.4~GHz  throughout) obtained  from  the Hubble  Space Telescope (HST).   Early studies indicate that the majority of millijansky radio galaxies typically  have  extended  radio  emission  with  optical  colors  and luminosities     comparable    to     local     elliptical    galaxies \\citep[][]{kro85,win85,ben93,ham95a}.   The remainder  of  these radio sources are generally  quasars or very blue galaxies,  which are often interpreted as  late-type, starburst systems.  Like  the most powerful radio sources, millijansky galaxies have a broad redshift distribution of  $0\\!\\lesssim\\!z\\!\\lesssim\\!1.5$ \\citep{ove03},  which  reflect the increase of  active galaxies and  radio emission from  $z\\!\\sim\\!0$ to $z\\!\\sim\\!2$ \\citep[eg.][]{hop06}.   The space density  of millijansky radio galaxies evolves strongly with redshift, which can be reproduced by a simple evolutionary model.  \\citet{con89} argue that the majority of these radio galaxies reside in a ``shell'' of radius $z\\!\\sim\\!0.8$ with thickness  of $\\delta  z\\!\\sim\\!0.8$.  In addition  to supporting this   model,  \\citet{wad01}   find   a  deficit   of  high   redshift ($z\\!\\gtrsim\\!2$) millijansky radio galaxies  in the Hercules field of the Leiden-Berkley Deep Survey  (LBDS).  They interpret this result as an  effective redshift ``cut-off,''  which varies  from $z\\!\\simeq\\!2$ for the most  luminous radio sources, to $z\\!\\simeq\\!1$  for the lower luminosity galaxies.  Millijansky radio sources are intrinsically rather luminous at optical wavelengths,  with absolute  magnitudes in  a relatively  narrow range ($\\pm\\!0.8$~mag)    centered    at   $M_{r^*}\\!\\simeq\\!-22$~mag    and $M_{r^*}\\!\\simeq\\!-23$~mag  for blue  and  red galaxies,  respectively \\citep{mac00,ive02}.   From photographic plate  and deep  CCD imaging, \\citet{win84b} and \\citet{wad00} found the that the apparent magnitude distribution from the LBDS  peaks around $I\\!\\simeq\\!22$~mag, which at $z\\!=\\!0.8$ corresponds to $M_B\\!\\simeq\\!-21.5$~mag.  While there is a good deal of deep ground-based imaging on millijansky radio sources, there is little high-resolution HST imaging.  Early works with HST are generally based on at most a  few objects \\citep{pas96,wad99,wad02}.    In the microjansky flux range, \\citet{afo06} identified 64 faint radio galaxies of which  57 had optical counterparts in  the Advanced Camera for  Surveys (ACS)  imaging of  the Great  Observatories  Origins Deep Survey,  South   \\citep[GOODS-S;][]{gia04}.   The  median  photometric redshift of the optically identified microjansky radio sources is also $z\\!\\sim\\!0.8$.   Since the  redshift distribution  of  radio galaxies changes    only     very    slowly    for     radio    fluxes    below $S_{1.4}\\!\\simeq\\!1$~Jy   \\citep{win90},   this   median   photometric redshift is  consistent with the evolutionary  model of \\citet{con89}. Based on  their $X$-ray  properties and limited  optical spectroscopy, the  majority of these  microjansky radio  galaxies are  either active galactic   nuclei  (AGN)   or  star-forming   galaxies  \\citep{afo06}. Furthermore,  the   optical  counterparts  generally   have  disturbed morphologies, which suggests many are in the later stages of merging.    In this  paper, we investigate  the optical properties  of millijansky radio  galaxies selected from  the Faint  Images of  the Radio  Sky at Twenty centimeters  \\citep[FIRST;][]{bec95}.  The FIRST  survey covers $\\simeq\\!10^4$  square degrees  of the  north Galactic  cap to  a 95\\% completeness  limit  of   $\\simeq\\!2.8$~mJy.   In  order  to  reliably determine   the  morphology   of   these  small   and  faint   optical counterparts,  we  require   the  angular  resolution  of  space-based imaging.  Therefore, our  sample is derived from the  overlap of FIRST sources  with  the  existing   HST-WFPC2  Archive  as  of  2002.   All magnitudes  quoted herein  are in  the AB  system  \\citep{oke83}. This paper  is  organized  as  follows:  In \\S~\\ref{proc}  we  outline  the critical   steps  to   define   our  final   millijansky  sample,   in \\S~\\ref{results}   we   discuss   our   four   primary   results,   in \\S~\\ref{comment}  we provide general  comments on  individual objects, and in \\S~\\ref{sum} we conclude  with a general summary and additional discussion of our results. ",
        "conclusions": "\\label{sum}  We have determined the optical morphology and structure from HST-WFPC2 imaging of a randomly-selected  sample of 51~millijansky radio sources from the FIRST survey.   Earlier studies of millijansky radio galaxies have often  distinguished early-  versus late-types based  on observed optical colors  or spectra.  However,  we have extended such  works by providing optical  classifications based  on HST imaging,  and confirm that elliptical  galaxies constitute  the majority of  the millijansky source population.    We find that the optical  flux distribution of these galaxies peaks at $I\\!\\simeq\\!\\alli$~mag, which is consistent with evolutionary model of \\citet{con89}: faint radio  galaxies generally have $L^*$-type optical luminosities  with  a  median  redshift  of  $z\\!\\simeq\\!0.8$  and  an effective  redshift decline  for $z\\!\\gtrsim\\!1.5$.   At  low redshift ($z\\!\\lesssim\\!1$),     most    radio     sources     brighter    than $S_{1.4}\\!\\gtrsim\\!1$~mJy are generally  found in massive ellipticals, and  there   is  a  clear   deficit  of  these  massive   galaxies  at $z\\!\\gtrsim\\!3$   \\citep[eg.][]{dri98}.    \\citet{del06}  argue   that $\\sim\\!50$\\%  of the  stars which  will  likely end  up in  elliptical galaxies in the  local Universe are formed at  $z\\!\\simeq\\!3$, and are not finally  assembled into  the galaxy until  $z\\!\\simeq\\!0.8$.  This down-sizing  picture  requires an  increased  major  merger rate  from $z\\!\\sim\\!0$ to  $z\\!\\sim\\!3$, which has  been observed in  the Hubble Ultra  Deep  Field  \\citep[eg.][]{ryan08}.   Finally,  this  effective redshift  cut-off may  arise, since  $\\sim\\!1$~Gyr may  be  on average required to trigger an active  nucleus after the major merger based on hydrodynamical models \\citep{spi05}.  This  must be further studied in the  critical redshift  range of  $1\\!\\lesssim\\!z\\!\\lesssim\\!3$, which may  prove  useful  in  tracing  the high-redshift  evolution  of  the supermassive black hole--bulge mass relation \\citep[eg.][]{fer05}.  At these  redshifts, the  optical  counterparts will  be extremely  faint (${\\rm  AB}\\!\\gtrsim\\!25$~mag), and may  have a  high dust  content or their luminosity  function may  evolves strongly or  episodically with redshift.   Consequently, dedicated observations  with the  Wide Field Camera~3 (WFC3) or the {\\it James Webb Space Telescope} will be needed for  the  continued  study  of  the population  of  millijansky  radio sources."
    },
    "0807/0807.2476.txt": {
        "abstract": "{}{}{}{}{} % 5 {} token are mandatory \\abstract % Context {The central region of our Galaxy contains a large population of young massive stars. These stars are concentrated in three large star clusters, as well as being scattered in the field. Strong ionizing radiation and stellar winds of massive stars are the essential feedback agents that determine the physics of the ISM in the Galactic center.} % % aims heading (mandatory) {The aim is to study relatively isolated massive WN-type stars in the Galactic center in order to explore their properties and their influence on the ISM. } % % methods heading (mandatory) {The K-band spectra of two WN stars in the Galactic center, \\ka\\ and \\wrc, are exploited to infer the stellar parameters and to compute synthetic stellar spectra using the Potsdam Wolf-Rayet (PoWR) model atmosphere code. These models are combined with dust-shell models for analyzing the {{\\em Spitzer}} IRS spectra of these objects. Archival IR images complement the interpretation.  } % % results heading (mandatory) {We report that \\ka\\ and \\wrc\\ are among the most luminous stars in the Milky Way. They critically influence their immediate environment by strong mass loss and intense UV radiation, and thus set the physical conditions for their compact circumstellar nebula.  The mid-IR continua for both objects are dominated by dust emission.  For the first time we report the presence of dust in the close vicinity of WN stars. Also for the first time, we have detected lines of pure-rotational transitions of molecular hydrogen in a massive-star nebula.  A peony-shaped nebula around \\ka\\ is resolved at 24\\mim\\ by the {{\\em Spitzer}} MIPS camera. We attribute the formation of this IR-bright nebula to the recent evolutionary history of \\ka.  } % % conclusions heading (optional), leave it empty if necessary {} ",
        "introduction": "Visual light cannot penetrate the dust clouds obscuring the innermost part of our Galaxy, whereas infrared astronomy opens extraordinary views on this environment.   Three very massive star clusters have been discovered in the Galactic center (GC) during the last decade.  The central cluster is located within 1\\,pc from the central black hole (Krabbe \\etal\\ \\cite{krab95}). Two other massive star clusters, the Arches and the Quintuplet, are located within 30\\,pc projected distance from Sgr\\,A* (Serabyn \\etal\\ \\cite{arch98}, Figer \\etal\\ \\cite{quin99}).  While the Arches cluster is younger and contains many OB-type stars, the more evolved Quintuplet cluster (3-5\\,Myr old) harbors many Wolf-Rayet (WR) stars. Besides these compact stellar conglomerates, many high-mass stars whose association with stellar clusters is not obvious are scattered in the GC. Among these are the rather isolated WR type stars \\ka\\ and \\wrc\\ -- our program stars.  Massive stars severely influence their environment by strong ionizing radiation, mass and kinetic energy input (Freyer \\etal\\ \\cite{fr03}). The thermal Arched Filaments in the vicinity of the Arches and the Sickle nebula in the vicinity of the Quintuplet are thought to be powered by the combined action of hot massive stars located in these clusters (Simpson \\etal\\ \\cite{sim07} and references therein). However, radio (e.g. Yusef-Zadeh \\& Morris \\cite{yz87}) and infrared (Price \\etal\\ \\cite{msx01}, Rodr{\\'i}guez-Fern{\\'a}ndez \\etal\\ \\cite{rf01}, Simpson \\etal\\ \\cite{sim07}) surveys of the GC reveal a complex structure of the ionized gas with many small-scale compact sources of thermal emission. In this paper we investigate the emission from such nebulae around our program stars.  A star with initial mass over $\\approx$20\\,\\Msun\\ settles on the main sequence as an O-type. Stars more massive than $\\approx$30\\,\\Msun\\ evolve off the main sequence at more or less constant luminosity, but their mass-loss rate increases significantly. Stars with initial masses above $\\approx$40\\,\\Msun\\ pass through a short ($\\sim10^5$\\,yr) luminous blue variable (LBV) evolutionary stage that is characterized by high mass-loss rates and violent eruptions. The ejected material, up to several $M_\\odot$, is often observed in the form of an associated nebula -- one famous example is the Homunculus nebula around the LBV star $\\eta$\\,Car.  The most massive stars with initial masses $M_{\\rm i} \\msim$90\\,\\Msun\\ are thought to lose enough mass during their life on the main-sequence to evolve to the WR stage without ever becoming an LBV.  Stars that display CNO-processed matter in a strong stellar wind are classified as WR stars of the nitrogen sequence (WN type).  The cooler, late WN subtypes (WNL) usually contain some rest of hydrogen in their atmospheres, while the hotter, early subtypes (WNE) are hydrogen free (Hamann \\etal\\ \\cite{ham91}).  {\\changed The WNL evolutionary stage can, in fact, precede the LBV stage (Langer et al.\\ \\cite{lan94}, Smith \\& Conti \\cite{smith08}). Typically, WNL stars are significantly more luminous than WNE stars (Hamann \\etal\\, \\cite{wrh06}).} The WN phase may be followed by the WC stage, when the products of helium burning appear in the stellar atmosphere.  Wolf-Rayet stars end their lives with a super- or hypernova explosion.  During its evolution a massive star loses a considerable fraction of its initial mass. This material accumulates in the circumstellar environment (e.g.\\ Freyer \\etal\\ \\cite{fr03}).  Recently, concerns were raised that the empirically derived stellar mass-loss rates need to be drastically reduced (Fullerton \\etal\\ \\cite{ful06}).  However, when the inhomogeneous nature of stellar winds is accounted for, the analyses of optical, UV, and X-ray spectra of massive stars consistently yield mass-loss rates that are only factor of two lower than inferred under the ``standard`` assumption of smooth winds (Oskinova \\etal\\ \\cite{osk07}).  Infrared (IR) spectra of H\\,{\\sc ii} regions around main-sequence stars are composed of nebular emission lines and a dust-dominated continuum (Dopita \\etal\\ \\cite{dop05}). However, in the vicinity of WR stars dust is rarely found, except for WC-type stars in the inner, metal-rich parts of galaxies (e.g.\\ Crowther \\etal\\, \\cite{wc9}). These stars are surrounded by dust shells, and it appears that the metal-rich environment of the GC is favorable to the formation of circumstellar dust. For WN stars, however, it is generally believed that the chemical composition and strong radiation prohibit the formation of dust in their vicinity.  %%%%%%%%%%%%%%%%%%%%%%%%%  FIG 1  %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %------------------------------------------------------------- \\begin{figure} \\centering \\includegraphics[width=\\columnwidth]{9568fig1.ps} \\caption{{\\em Spitzer} IRAC 8\\,\\mim\\ archive image of the GC. Arrows point to the locations of our program stars, the Arches, the Quintuplet, and the Central clusters. North is to the top and east to the left. }  \\label{fig:gc_reg} \\end{figure} % %_____________________________________________________________  We secured the mid-IR spectra of WN-type stars \\ka\\ and \\wrc\\ with the {\\em Spitzer} Space Telescope. Our program stars are sufficiently isolated to allow high-resolution spectroscopy with the {\\it Spitzer} infrared spectrograph IRS. The basic strategy for the analysis is to model the stellar spectrum using a stellar atmosphere code, and then use the synthetic stellar spectrum as input for modeling the circumstellar nebula.  This paper is organized as follows. Stellar parameters and synthetic stellar spectra of our program stars are obtained in Section\\,2 from the analysis of K-band spectra. Infrared spectroscopic and imaging observations are presented in Sect.\\,3. The analysis and modeling of the mid-IR {\\em Spitzer} IRS spectra are conducted in Sect.\\,4. A discussion is presented in Sect.\\,5, followed by the summary in Sect.\\,6.  %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %--------------------  Section 1 PoWR ------------------------------ %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% ",
        "conclusions": "}  \\subsection{Stellar mass and evolution of \\ka\\ and \\wrc}  Figer \\etal\\ (\\cite{fig98}) presented evolutionary tracks for very massive stars based on the code by Langer \\etal\\ (\\cite{lan94}). According to these tracks, the initial stellar mass is $150\\,\\lsim\\,M_{\\rm i}/\\Msun\\,\\lsim\\,200$ for \\ka\\ and $100\\,\\lsim\\,M_{\\rm i}/\\Msun\\,\\lsim\\,150$ for \\wrc. Thus, both our program stars were initially among the most massive stars of the Galaxy.  %%%%%%%%%%%%%%%%%%%%%%%%%  FIG 16 %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %------------------------------------------------------------- \\begin{figure} \\centering \\includegraphics[width=\\columnwidth]{9568fig16.ps} \\caption{ Hertzsprung-Russel diagram with the location of some known LBVs, LBV candidates and our program stars. The thick line indicates the Humphreys-Davidson limit (according Eq.\\,(2) in Lamers \\& Fitzpatrick \\cite{lam88}). The diamonds marks LBVs with time averaged observations, squares shows the location of the LBV at the maximum, circles at the minimum. The small triangles represent LBV candidates, while the big triangles indicate the locations of \\ka\\ and \\wrc. Except of \\ka\\ and \\wrc, stellar parameters are from Figer \\etal\\ (\\cite{fig98}), van Genderen (\\cite{vg01}), Najarro  (\\cite{naj06}), Groh \\etal\\ (\\cite{gro06}).} \\label{fig:lbvhrd} \\end{figure} % %_____________________________________________________________  Our analysis of \\ka\\ yields an unconventionally high luminosity at a relatively low stellar temperature (cf. Table\\,\\ref{tab:par}).  In the HR-diagram (see Fig.\\,\\ref{fig:lbvhrd}) \\ka\\ is located above the Humphreys-Davidson limit, in the region populated by the LBV stars. \\ka\\ is of spectral type Ofpe/WN9. This class of objects is often considered as either LBV candidates or LBVs in quiescence (Crowther \\etal\\ \\cite{cr95}, Morris \\etal\\ \\cite{M1996}).  From the analysis of the K-band spectrum, the surface mass fraction of hydrogen in \\ka\\ is $0.2$. This is lower than the hydrogen mass fraction 0.3\\,...\\,0.4 found in known LBVs (Stothers \\& Chin, 2000).  The complimentary helium mass fraction in \\ka\\ is significantly higher than found for e.g. the Pistol star ($Y_{\\rm surf}\\lsim 0.4$). This indicates a more advanced evolutionary stage of the former, compared to ``normal'' LBV stars.  Interestingly, Maeder \\etal\\ (\\cite{maed08}) discussed possible filiations of Pop.\\,I massive stars. They suggest that stars with initial masses $M_{\\rm i} >90\\,M_\\odot$ do not pass through an LBV stage, but have high enough mass loss to get rid of their envelopes on the main sequence. According to this scenario, it is possible that \\ka\\ has not been a classical LBV, but evolved directly from the Of to the WNL stage.  In the HR-diagram \\wrc\\ is located in the same region as the outstandingly luminous WN6(h)\\,...\\,WN8(h) type stars WR\\,22, WR\\,24, and WR\\,25 (Hamann \\etal\\ \\cite{wrh06}). An orbital solution for the eclipsing binary WR\\,22 yields $M_*\\approx 70\\,M_\\odot$ (Rauw \\etal\\ \\cite{rauw96}) in agreement with an elaborate analysis based on the theory of optically thick winds (Gr{\\\"a}fener \\& Hamann \\cite{gr08}).  A crude estimate of the present stellar mass of \\wrc\\ can be obtained by using the scaling between $\\vinf$ and the escape velocity $\\varv_{\\rm{esc}}$ known for WN stars, $\\vinf\\,\\approx\\,1.5\\,...\\,4\\,\\varv_{\\rm{esc}}$ (Lamers \\& Cassinelli \\cite{lam99}).  The escape velocity is given by \\mbox{$\\varv_{\\rm esc}\\,=(2{\\rm G}M_{\\rm eff}/R_\\ast)^{0.5}$}, where \\mbox{$M_{\\rm eff}=M_\\ast(1-\\Gamma)$}, and $\\Gamma$ is the Eddington factor. Using stellar parameters and abundances of \\wrc\\ from Table\\,\\ref{tab:par} we estimate the present stellar mass to 45\\,...\\,55\\,\\Msun.  Because the scaling between $\\vinf$ and the $\\varv_{\\rm{esc}}$ is not established for stars located above the Humphreys-Davidson limit, the same method cannot be applied to constrain the present mass of \\ka.  A vast amount of chemically enriched material has been lost during recent stellar evolution of \\ka\\ and \\wrc, and has contributed to the nebulae around these stars.  Interestingly, the evolutionary tracks (Meynet \\etal\\ \\cite{mae03}) predict that the carbon to oxygen mass-ratio on the surface is larger than unity already in the WN phase for stars with $M_{\\rm i} \\msim 85\\,\\Msun$. Thus one may speculate that the PAH features observed in the spectra of our program stars and elsewhere in the GC reflect the carbon enrichment of the ISM by the stellar winds from initially very massive stars.  Neither \\ka\\ nor \\wrc\\ belong to the central parts of stellar clusters. They may have been either dynamically ejected from parental clusters, or formed in isolation.  Conspicuously, \\ka\\ is located at the Bubble Rim, where the ISM could be pressurized by the expanding hot bubble. \\wrc\\ is located in the Sickle nebula at the edge of a dense molecular cloud that is ionized by the Quintuplet cluster (Simpson \\etal\\ \\cite{sim97}).  The Quintuplet cluster is  3-5\\,Myr old and contains a large population of WC-type stars that are older than \\wrc. {\\changed Moreover, the radial velocity of the Quintuplet cluster, $130$\\,km\\,s$^{-1}$ (Figer \\etal\\ \\cite{Fig99}), is quite different from the average radial velocity of the Sickle nebula, $\\sim 35$\\,km\\,s$^{-1}$ (Lang \\etal\\ \\cite{lang97}). Hence, most likely formation of \\wrc\\ is not linked to the Quituplet cluster.}  %%%%%%%%%%%%%%%%%%%%%%%%%  FIG 17 %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %%------------------------------------------------------------- \\begin{figure} \\centering \\includegraphics[width=\\columnwidth]{9568fig17.ps} \\caption{{\\em Spitzer} MIPS 24\\,\\mim\\  image  of \\ka\\ and its peony-shaped nebula. The spatial resolution is $2\\farcs 5$. The nebula has a radius of about $50$\\,\\arcsec. The colors are chosen such that the background appears black. The image is shown on linear scale, as well as the contours. Two central parts of the image are saturated, and are shown in white for cosmetic reasons. North is up, east to the left.} \\label{fig:mipska} \\end{figure} % %%_____________________________________________________________ %  \\subsection{Nature of the nebula around \\ka}  Recently, the {\\em Spitzer} MIPS survey of the GC at 24\\,\\mim\\ became publically available. The 24\\,\\mim\\ image of \\ka\\ and its nebula is shown in Fig.\\,\\ref{fig:mipska}.  It confirms what was already obvious from the less sensitive MSX observation -- the presence of a compact dusty nebula centered on \\ka. The MIPS image resolves its roughly spherical shape, resembling a Peony flower.  The radius of the ``peony nebula'' in 24\\,\\mim\\ image is about 50\\,\\arcsec ($\\approx$1.5\\,pc), i.e.  nearly two times larger than the size of the unresolved image in the MSX E-band.  This is because the MIPS camera is $\\msim$$10^4$ times more sensitive.  Unfortunately, the MIPS image is saturated in central parts, and therefore fluxes cannot be extracted for modeling purposes. Intriguingly, the two brightest, saturated parts of the ``peony nebula'' coincide with the point-like central star and with the fuzzy, feather-like feature north of the star, visible in the 8\\,\\mim\\ IRAC image (right panel of Fig.\\,\\ref{fig:irac-ka}). This spatial coincidence is interesting, because the IRAC 8\\,\\mim\\ images trace the emission from PAH, while the MIPS 24\\,\\mim\\ image shows the emission from small grains. For our {\\em Spitzer} observation the slit was centered on \\ka, and therefore our spectrum samples the central part of the nebula.  Unfortunately, the spectrum covers only the short-wavelength range and does not allow a detailed study of the chemical composition and the dynamics of the \\ka\\ nebula. The available part of the nebula spectrum is similar to the spectra of known LBV nebulae (cf.\\ Voors \\etal\\ \\cite{voors}). As discussed above, \\ka\\ is either a post-LBV star or a star that suffered high mass loss already on the main sequence. It is natural to suggest that the photoionized peony-shaped nebula contains stellar material that was lost by \\ka\\ during LBV-type eruptions and/or its strong stellar wind."
    },
    "0807/0807.4528_arXiv.txt": {
        "abstract": "We revisit the issue of ripples in the primordial power spectra caused by trans-Planckian physics, and the potential for their detection by future cosmological probes. We find that for reasonably large values of the first slow-roll parameter $\\epsilon$ ($\\gwig 0.001$), a positive detection of trans-Planckian ripples can be made even if the amplitude is as low as $10^{-4}$. Data from the Large Synoptic Survey Telescope (LSST) and the proposed future 21 cm survey with the Fast Fourier Transform Telescope (FFTT) will be particularly useful in this regard.  If the scale of inflation is close to its present upper bound, a scale of new physics as high as $\\sim 0.2\\ M_{\\rm P}$ could lead to observable signatures. ",
        "introduction": "Observations of the temperature and polarisation anisotropies in the cosmic microwave background (CMB) radiation strongly support the idea that the initial conditions for structure formation were set by an earlier period of cosmic inflation~\\cite{Starobinsky:1980te,Guth:1980zm,Linde:1981mu,Albrecht:1982wi}. During this period, the rapid expansion of space caused quantum fluctuations of the inflaton field to be blown up to macroscopic scales.  The initial state of these quantum fluctuations is typically assumed to be given by the Bunch-Davies vacuum of de Sitter space at a time when $k \\gg aH$, i.e., the physical wavenumber $k/a$ is much larger than the Hubble expansion rate.  Along with the usual assumptions of single-field, slow-roll inflation, this leads to the almost scale-invariant spectrum of Gaussian adiabatic scalar perturbations that constitutes the ``primordial'' state of perturbations in today's cosmological concordance model~\\cite{Mukhanov:1990me}.  However, in an expanding space-time the notion of a vacuum is not unique. This can easily be understood by observing that the expanding space-time makes the Hamiltonian depend explicitly on time, so there is no time-independent lowest energy state that can serve as a vacuum at all times. We might then worry about how sensitive our predictions are to the choice of initial state.  The choice of vacua other than the Bunch-Davies one leads in principle to unacceptable ultraviolet divergences in the observables~\\cite{Brunetti:2005pr}, although this deficiency can be cured by introducing a high-energy cutoff to the theory.  In practice, such a cutoff will be given by the Planck scale, or possibly even lower, by a \\emph{scale of new physics}~$\\Lambda$ (if the cutoff is too close to the Planck scale, an unorthodox vacuum choice may in fact lead to backreaction problems \\cite{Tanaka:2000jw,Starobinsky:2001kn,Starobinsky:2002rp,KeskiVakkuri:2003vj,Porrati:2004gz,Greene:2004np}). An ultraviolet cutoff in the inflaton fluctuation modes violates energy conservation and has to be associated with a source term for the fluctuations, in order to account for the modes redshifting across the new physics hypersurface~\\cite{KeskiVakkuri:2003vj}.   Unless inflation started only just before the wavelengths corresponding to today's observable scales left the horizon, these fluctuations can be mapped to wavenumbers larger than $\\Lambda$ during inflation and have emerged from above the new physics hyper-surface. It is thus very well conceivable that the fluctuations bear an imprint of the unknown new physics~\\cite{Martin:2000xs,Brandenberger:2000wr}. In particular, the modes might be created in a non-standard vacuum state at the new physics hypersurface. Danielsson proposed a prescription of how to construct this initial state by minimising the Hamiltonian on this hypersurface~\\cite{Danielsson:2002kx}. However, it was found in references~\\cite{Bozza:2003pr,Giovannini:2003it} that the state constructed in this fashion is somewhat ambiguous, since it is not invariant under canonical transformations of the Hamiltonian. In fact, it has been argued from a purely effective field theory point of view that the effect on the primordial spectrum should be smaller than that proposed by Danielsson~\\cite{Kaloper:2002cs}.  Eventually, only a better understanding of the thinning of ultraviolet degrees of freedom in quantum gravity~\\cite{Atick:1988si} can help us  solve the puzzle of the correct choice of initial conditions and how the fluctuations emerge during inflation, as in the proposed {\\it trans-Planckian damping} mechanism~\\cite{Hassan:2002qk} (see also~\\cite{Starobinsky:2002rp}).  Nonetheless, the generic signature of choosing a non-Bunch-Davies vacuum state appears to be a superimposed oscillation of the primordial spectra (``ripples''), whose amplitude is suppressed by a power of $H/\\Lambda$. For definiteness we will consider Danielsson's model as a case study in the following.    The question of whether traces of these oscillations can possibly be detected in present or future CMB anisotropy data has been the subject of a number of recent studies \\cite{Bergstrom:2002yd,Elgaroy:2003gq,Martin:2003sg,Okamoto:2003wk,Martin:2004iv,Martin:2004yi,bib:ekp,Spergel:2006hy,Martin:2006rs,Groeneboom:2007rf}. While few will disagree that the CMB anisotropies are at present the most powerful probe of the primordial power spectrum, the advent of high-redshift surveys of the large scale structure (LSS) distribution in the not-so-distant future may help to enlarge the detection window. Planned galaxy redshift surveys such as the Large Synoptic Survey Telescope (LSST) \\cite{Ivezic:2008fe}, or surveys of the distribution of neutral hydrogen using the 21~cm spin-flip line (e.g., the Murchison Widefield Array (MWA) \\cite{bib:mwa}, the Square Kilometre Array \\cite{bib:ska}, or the Fast Fourier Transform Telescope (FFTT) \\cite{Tegmark:2008au}) have the advantage of mapping out the density perturbations in three dimensions over large volumes. They are less affected by sampling variance (or cosmic variance), which is ultimately the limiting factor for CMB data over a large range of scales.  In this work, we investigate how a combination of these probes with CMB data can enhance our potential to discover the footprints of new physics beyond the energy scale of inflation.   We shall begin by refreshing the reader's memory of the Danielsson model in section~\\ref{sec:model}. Some general considerations about the phenomenology of oscillating spectra will be presented in section~\\ref{sec:precon}, before we turn to the questions of under which conditions their traces can be detected in CMB data (section~\\ref{sec:detect?}), or, if no detection is possible, what constraints can be placed on the parameters of the model (section~\\ref{sec:constraints}). In sections~\\ref{sec:grs} and~\\ref{sec:21cm} we discuss how future galaxy redshift surveys and 21~cm surveys can contribute to improving the chances for a detection of ripples.  We conclude in section~\\ref{sec:conclusions}. ",
        "conclusions": "}  We have investigated the sensitivity of various future cosmological probes to ``ripples'' in the primordial power spectrum due to ``trans-Planckian'' physics.  The magnitude of this effect is determined by $\\xi=H_0/\\Lambda$, i.e., the energy scale of inflation relative to the scale of new physics. As also noted in some previous works on the topic (e.g., \\cite{Elgaroy:2003gq}), the sensitivity to $\\xi$ of a given probe is determined by two experimental parameters.  One is the precision with which the power spectrum can be measured over the relevant range, specified by the error bars of the data and/or the width of the window functions.  The other important parameter is the ``lever arm'', i.e., the span in $k$-space over which the power spectrum can be measured, which is crucial for disentangling  ripples in the power spectrum caused by trans-Planckian physics from a spectral tilt or an amplitude shift.  Qualitatively, the precision of the data is the limiting factor in the case of very small~$\\xi$, while the lever arm determines, for a given $\\epsilon$, the largest $\\xi$ that can be detected uniquely as ripples in the primordial power spectrum. The limit imposed by the lever arm also means that the detectability of trans-Planckian ripples depends strongly on the magnitude of the first slow-roll parameter $\\epsilon$. If $\\epsilon$ is too small, no detection window exists for $\\xi$, as can be seen from equation~(\\ref{eq:osc}), as well as more quantitatively from tables~\\ref{tab:table1} to~\\ref{tab:table3}.  In the same vein, if $\\epsilon=0$ is consistent with data, then no bounds can be set on $\\xi$.  We have estimated the detection threshold of these ripples for data from future CMB experiments (Planck and a hypothetical cosmic variance limited experiment CVL), a galaxy redshift survey exemplified by the LSST, and finally a large scale 21 cm survey. Due to their very low noise and large volumes, both galaxy and 21 cm surveys will eventually reach a higher precision than CMB observations alone. This will allow us to reach down to smaller values of $\\xi$ than ever before. However, large scale structure surveys will not be able to compete with CMB observations in terms of wavelength coverage. Therefore, in order to maximise the lever arm for the detection of trans-Planckian ripples, these surveys should be combined with CMB measurements.  In an optimistic scenario with $\\epsilon=0.01$, we find that on the intermediate timescale, data from Planck+LSST will be open a detection window spanning $10^{-3} \\lwig \\xi \\lwig 0.02$. In the more distant future the lower limit of the window can possibly be pushed down to $10^{-4}$ using data from a cosmic variance limited CMB experiment combined with the FFTT 21 cm survey. This same combination of data will also create a window of detection for $\\epsilon=0.001$ and possibly lower for the first time ($10^{-4} \\lwig \\xi \\lwig 0.0025$).   So what are the implications for the scale of new physics $\\Lambda$, which is presumably close to or at the Planck scale?  Given that present data already constrains the scale of inflation to $H_0 \\lwig 2 \\times 10^{-5}\\ M_{\\rm P}$~\\cite{Valkenburg:2008cz}, we conclude that in the Danielsson model, there are no realistic prospects of ever detecting the traces of new physics in the primordial spectra if $\\Lambda \\gwig 0.2\\ M_{\\rm P}$.  For models in which the amplitude of the ripples is quadratic instead of linear in $\\xi$, as predicted by purely effective field theory arguments in reference~\\cite{Kaloper:2002cs}, the new physics would need to set in at energies smaller than $ \\sim 2 \\times 10^{-3}\\ M_{\\rm Pl}$.  Nonetheless, we have shown that the fantastic precision of future CMB surveys like Planck, and large scale structure surveys such as LSST and FFTT will extend the detection window of trans-Planckian ripples by more than two orders of magnitude compared with present data. After all, the prospect of probing physics close to the Planck scale is a most tantalising one; a detection of trans-Planckian signatures would truly be one of the most important accomplishments of modern precision cosmology."
    },
    "0807/0807.3414_arXiv.txt": {
        "abstract": "{This paper is motivated by the recent detection of many extremely metal-deficient ($[\\rm{Fe/H}]<-3$) stars in the Milky Way.} {We explore numerically the chemical, thermal, and dynamical evolution of a shell formed by a high-energy supernova explosion ($10^{53}$~erg) in dwarf protogalaxies with total (dark matter plus baryonic) mass $10^7~M_\\odot$ at a redshift $z=12$. We consider two initial configurations for the baryonic matter, one without rotation and the other having the ratio of rotational to gravitational energy $\\beta=0.17$. The (non-rotating) dark matter halo is described by a quasi-isothermal sphere. The latter choice is motivated by recently proposed mechanisms for rapid flattening of a central cuspy region in dwarf protogalaxies. } {We use a finite-difference numerical hydrodynamics code to simulate supernova explosions in dwarf protogalaxies with axial symmetry. The advection is treated using a third-order piecewise parabolic scheme. The heating and cooling processes in the gas are taken into account by solving numerically the rate equations for main atomic, molecular and ionic species in the primordial gas.} {We find that the dynamics of the shell is different in protogalaxies with and without rotation. For instance, the Rayleigh-Taylor instability in the shell develops faster in protogalaxies without rotation. The fraction of a blown-away baryonic mass is approximately twice as large in models with rotation ($\\sim 20\\%$) than in models without rotation. We argue that these differences are caused by different $initial$ gas density profiles in non-rotating and rotating protogalaxies. On the other hand, the chemical evolution of gas in protogalaxies with and without rotation is found to be similar. The relative number densities of molecular hydrogen and HD molecules in the cold gas ($T \\le 10^3$~K) saturate at typical values of $10^{-3}$ and $10^{-7}$, respectively. The saturation times in models with rotation are somewhat longer than in models without rotation. The clumps formed in the fragmented shell move with velocities that are at least twice as large as the escape velocity. The mass of the clumps is $\\sim 0.1-10~\\msun$, which is lower than the Jeans mass. We conclude that the clumps are pressure supported.} {A supernova explosion with energy $10^{53}$~ergs destructs our model protogalaxy. The clumps formed in the fragmented shell are pressure supported. We conclude that protogalaxies with total mass $\\sim 10^{7}~M_\\odot$ are unlikely to form stars due to high-energy supernova explosions of the first stars.} ",
        "introduction": "The detection of extremely metal-poor stars in our Galaxy with the iron abundance equal to or less than $10^{-3}$ of the solar value \\citep{beers,christlieb02} motivated the scientists to put forward possible scenarios for the formation of such stars. According to \\citet{tsuji99} extremely metal-poor (EMP) stars form in a dense shell produced by Type II supernova explosions of the first stars and accrete metals from the surrounding medium during the subsequent evolution. The formation of EMP stars is made possible by fragmentation of the primordial gas in a supernova shell due to efficient cooling by molecular hydrogen and HD molecules. Indeed, the formation of molecular hydrogen increases significantly in the wake of a strong shock wave \\citep{suchkov83,sk87}. Moreover, at temperatures lower than $150~$K, HD molecules become a much more efficient coolant than molecular hydrogen \\citep[e.g.,][]{sh86}.  Subsequent studies by many authors \\citep[see review by][]{nishi99,salva04,machida,greif} have shown that the likelihood for the formation of EMP stars in supernova-driven shells is sensitive to both the mass of the dark matter halo ($M_{\\rm h}$) in a primordial galaxy and the assumed type of supernova explosions. For instance, \\citet{nishi99} have shown that the shell fragmentation due to Type~II supernova explosions ($10^{51}$~erg) is possible only in dark matter halos with total mass $M_{\\rm h} \\ga {\\rm a~few} \\times 10^7~\\msun$. Galaxies with $M_{\\rm h} \\la 10^7~\\msun$ may lose its baryonic matter before it has time to fragment and form stars \\citep{ferrara98}. The likelihood for this blow-away scenario increases if the first stars explode as pair-instability supernovae with the energy release of the order of $10^{53}$~erg \\citep{bromm03,greif}. However, these numerical results were confronted by \\citet{machida}, who presented a semi-analytic model for the evolution of a gas shell produced by supernova explosions with energy $10^{51}-10^{52}$~erg. They took into account the H$_2$ and HD chemistry and found that supernova explosions can induce fragmentation of the gas shell and formation of EMP stars in dark matter halos with $M_{\\rm h} = {\\rm a~few} \\times 10^6~\\msun$. The typical mass of the fragments in their model is about $1.0~\\msun$, which is consistent with the result by \\citet{uehara00} and similar to the masses of the observed low-metallicity stars \\citep{christlieb02}.  The likelihood for the formation of EMP stars depends also on the radial distribution of gas in a primordial galaxy {\\it prior} to supernova explosions. This distribution can (partly) be determined by HII regions around the first stars. For instance, \\citet{whalen04} considered the formation of the HII region in a low-mass dark matter halo with mass $10^6~\\msun$ and found that the gas is efficiently blown away by a supersonic shock wave accociated with the R-type ionization front. Similar results were obtained by \\citet{kitayama04} for the same mass of the dark matter halo. They have also considered more massive dark matter halos and concluded that the HII region is confined within the virial radius in halos with masses $\\simgt 10^7~\\msun$ (the ionization front is of the D-type). More recently, \\citet{kitayama05} studied numerically the destruction of dark matter halos due to the formation of HII regions due to radiation of the first stars and subsequent supernova explosions. Using a one-dimensional Lagrangean hydrodynamics code, they found that the SN shock wave remains well inside the virial radius for dark matter halos with masses larger than $10^7~\\msun$. On the other hand, in low-mass dark matter halos with masses of the order of $10^6~\\msun$ the shock wave propagates outside the virial radius. These results  indicate that the efficiency of halo destruction (and, by implication, the efficiency for the formation of EMP stars) are determined not only by the explosion energy but also by the gas density distribution and radiative feedback from the first stars prior to their explosion.  It now becomes evident that the formation of EMP stars due to fragmentation of supernova driven shells is a complicated phenomenon that depends on a variety of physical conditions in a host protogalaxy, which may vary from allowing star formation to shutting it off completely. In such circumstances, the construction of increasingly more sophisticated numerical models is justified. The majority of abovementioned studies are based on semianalytic models \\citep[e.g.,][]{salva04,machida} or on one-dimensional numerical hydrodynamic simulations \\citep[e.g.,][]{kitayama05}. In the framework of multi-dimensional numerical simulations, the priority was often given to a careful investigation of dynamical processes leading to the destruction of a dwarf protogalaxy by supernova explosions, neglecting the chemical processes that control the gas thermal properties. In the present paper, we perform axially symmetric numerical hydrodynamics simulations of high-energy supernova explosions ($10^{53}$~erg) in a model dwarf protogalaxy with total (dark matter plus baryonic) mass $1.7\\times 10^7~M_\\odot$. We give a careful treatment to the heating and cooling processes and chemical evolution of main molecular, ionic, and atomic species. We seek to determine the effect of galactic rotation on the dynamical and chemical evolution of a supernova-driven shell. The dark matter halo profile in our models is defined by a modified isothermal sphere \\citep[as observed in local dwarf galaxies,][]{burkert}, rather than by a cuspy profile \\citep{nfw}.  The paper is organized as follows. In Sections 2 and 3 we describe the model and numerical techniques, respectively; in Section 4 we discuss several neglected processes; in Section 5 we present results of our numerical simulations. The discussion and conclusions are given in Sections 6 and 7, respectively. Throughout the paper we assume a $\\Lambda$CDM cosmology with the parameters $(\\Omega_0, \\Omega_{\\Lambda}, \\Omega_m, \\Omega_b, h ) = (1.0,\\ 0.76,\\ 0.24,\\ 0.041,\\ 0.73 )$ as inferred from the Wilkinson Microwave Anisotropy Probe (WMAP), and deuterium abundance $2.78\\times 10^{-5}$, consistent with the most recent measurements \\citep{spergel06}. ",
        "conclusions": "In this paper we have considered numerically the effect of energetic supernovae explosions ($10^{53}~$erg) in non-rotating and rotating protogalaxies with the total mass $10^7~\\msun$ at a redshift of $z = 12$. The assumption of axial symmetry allowed us to evolve the supernova driven shell for several tens of million years. Specifically, we find the following.  \\begin{itemize}  \\item The process of the shell destruction is different for non-rotating and rotating protogalaxies. Fingers or spurs formed due to various instabilities (mainly, due to the Rayleigh-Taylor instability) grow faster and can be found on larger distances in non-rotating galaxies. That is explained by a steeper initial gas density profile in non-rotating protogalaxies and, as a consequence, by a faster characteristic time-scale for the development of the Rayleigh-Taylor instability.  \\item The supernova evacuates a significant portion of the initial baryonic mass to radial distances comparable with or larger than the virial radius. At $t=12$~Myr after the supernova explosion some of the fragments of the destructed shell have attained mean mass-weighted velocities of the order of $30$~km~s$^{-1}$ and $25$~km~s$^{-1}$ in non-rotating and  rotating protogalaxies, respectively. These values are at least twice as large as the escape velocity at the virial radius, which implies that the fragments eventually may escape the protogalaxy. We estimate the fraction of this blown-away gas in the end of our numerical simulations to be about $10\\%$ and $20\\%$ of the initial gas mass for models without and with rotation, respectively. The re-collapse phase of the shell sets in faster in models without rotation than in models with rotation. These results are mostly consistent with the intuitive expectations as far as the difference between models with and without rotation is concerned.  \\item The relative number densities of molecular hydrogen and HD molecules in the fingers and spurs are found to be quite large, approximately $10^{-3}$ and $10^{-7}$, respectively. The typical temperature in the spur cores is of the order of 500~K at $t\\simgt 8$~Myr after the supernova explosion.  \\item The total gas mass with $x[{\\rm H_2}] \\ge 10^{-3}$ saturates at $2 \\times 10^5~\\msun$, which corresponds to approximately $10\\%$ of the total baryonic mass in our model galaxy. The corresponding total mass of molecular hydrogen is about $300- 400~\\msun$. The saturation time-scales in the model with rotation are systematically longer than in the model without rotation. This can be attributed to longer characteristic cooling times in the model with rotation due to its shallower initial gas density profile.  \\end{itemize}  Finally, we would like to note that the typical masses of most fragments (assuming their spherical symmetry) are $\\sim 0.1-10~\\msun$. For the typical densities and temperatures in the fragments to be $\\sim 0.5-1$~cm$^{-3}$ and $(0.5-1)\\times 10^3$~K, respectively, these masses are strongly sub-Jeans and the fragments are expected to be pressure-supported. Their further evolution depends on both the efficiency of cooling and destruction due to the Kelvin-Helmholtz instability. We do not expect that the low-mass stars can be formed in such conditions as suggested by MacKey et al. 2003, and Salvaterra et al. 2004. In our opinion, a more feasible mechanism for low-mass, metal-poor star formation is related with the re-collapse of a supernova bubble in protogalaxies, whose total gravitational binding energy is much larger than supernova energy."
    },
    "0807/0807.1761_arXiv.txt": {
        "abstract": "We study the conditions for collisions between planetesimals to be accretional or disruptive in turbulent disks, through analytical arguments based on fluid dynamical simulations and orbital integrations.  In turbulent disks, the velocity dispersion of planetesimals is pumped up by random gravitational perturbations from density fluctuations of the disk gas.  When the velocity dispersion is larger than the planetesimals' surface escape velocity, collisions between planetesimals do not result in accretion, and may even lead to their destruction. In disks with a surface density equal to that of the ``minimum mass solar nebula'' and with nominal MRI turbulence, we find that accretion proceeds only for planetesimals with sizes above $\\sim 300$\\,km at 1AU and $\\sim 1000$\\,km at 5AU.  We find that accretion is facilitated in disks with smaller masses. However, at 5AU and for nominal turbulence strength, km-sized planetesimals are in a highly erosive regime even for a disk mass as small as a fraction of the mass of Jupiter. The existence of giant planets implies that either turbulence was weaker than calculated by standard MRI models or some mechanism was capable of producing Ceres-mass planetesimals in very short timescales. In any case, our results show that in the presence of turbulence planetesimal accretion is most difficult in massive disks and at large orbital distances. ",
        "introduction": "It is often considered that the evolution of protoplanetary disks and the consequent accretion of gas by the central protostar are driven by turbulent viscosity due to a Magneto-Rotational-Instability (MRI) \\citep[e.g.,][]{bal91}.  \\citet{lau04} and \\citet{nel04} carried out fluid dynamical simulations of MRI and found that the random torques due to the turbulent density fluctuations give rise to a random walk in semimajor axes of planetesimals.  \\citet{ric03} pointed out through model calculations that the random walk expands the effective feeding zone of protoplanets, and may lead to rapid formation of large cores for gas giants.  Through a Fokker-Planck treatment, \\citet{johnson06} also pointed out the importance of the random walk in planet accretion. Adopting the semi-analytical formula for the random torque derived by \\citet{lau04}, \\citet{ogi07} performed N-body simulations for the late stages of terrestrial planet accretion with a disk significantly depleted in gas, starting from Mars-mass protoplanets. They found that the MRI turbulence indeed helps to reduce the number of accreted terrestrial planets which is otherwise too large compared to our Solar System.  However, \\citet{nel05} found through direct integrations of the orbits of protoplanets in a MRI turbulent disk that orbital eccentricities are also excited. \\citet{britsch08} also found a similar feature in a self-gravitating disk. While the random walk itself is favorable to the growth of protoplanets by avoiding isolation, the excitation of their eccentricities, which had been neglected in \\citet{ric03} and \\citet{johnson06}, is a threat for planetesimal accretion processes because of increased collision velocity. Unfortunately, \\citet{nel05}'s orbital integrations were limited to 100--150 Keplerian times and neglected collision processes, so that it is not possible to conclude from that work whether planetesimals should grow or be eroded in the presence of turbulence.  In the present article, we explore by which paths planetesimals may have grown to planet-sized bodies in turbulent disks. Because the level of density fluctuations due to the MRI turbulence is not well determined, we choose to study the qualitative effects of the turbulence on the accretion of planetesimals and their dependence on the key parameters of the problem, in particular the progressive removal of the gas disk. In \\S 2, we summarize the conditions for the accretion and destruction of planetesimals in terms of their orbital eccentricities. In \\S 3, we analytically derive the equilibrium eccentricity for which the excitation due to turbulence is balanced by damping due to tidal interactions with the disk gas, aerodynamic gas drag, and collisions. Comparing the equilibrium eccentricities with critical eccentricities for accretion and destruction, we derive critical physical radii and masses of planetesimals for accretion or destruction. The results are applied to viscously evolving disks (\\S 4). We then discuss possible solutions for the problem of the formation of planetesimals and planets (\\S 5). ",
        "conclusions": "We have investigated the critical physical radii for collisions between planetesimals to be accretional ($R > R_{\\rm acc}$) or disruptive ($R < R_{\\rm dis}$) in turbulent disks, as functions of turbulent strength ($\\gamma \\sim O(\\Delta \\rho/\\rho)$), disk gas surface density, and orbital radius. The results presented here highlight the fact that MRI turbulence poses a great problem for the growth of planetesimals: generally, only those with sizes larger than a few hundred km are in a clearly accretive regime for a nominal value of $\\gamma \\sim 10^{-3}$. The others generally collide with velocities greater than their own surface escape velocities. For some of them, more severely in the kilometer-size regime, collisions are likely to be disruptive. The problem is greater when the disk is still massive and at large orbital distances. Also, if turbulence is stronger than $\\gamma \\sim 10^{-2}$, planetesimal accretion becomes extremely difficult.  However, the rate of occurrence of extrasolar giant planets around solar-type stars is inferred to be as large as $\\sim 20$\\% \\citep{Cumming08}, and depends steeply on the metallicity of the host star \\citep{Fischer05,santos04}. This strongly suggests that the majority of extrasolar giant planets were formed by core accretion followed by gas accretion onto the cores \\citep{IL05}. Thus, planetesimals should commonly grow to planetary masses before the disappearance of gas in protoplanetary disks.  The possibilities to overcome the barrier are in principle as follows (their likelihood is commented below): \\begin{enumerate} \\item Large $M$: Large planetesimals with sizes of 100 to 1000\\,km are formed directly in turbulent environment by a mechanism other than collisional coagulation, jumping over the erosive regime for physical radii.  \\item Small $\\Sigma_g$: Planetesimals start their accretion to planet-size only after the disk surface density of gas has declined to sufficiently small values.  \\item Small $\\gamma$: Planetesimals form in MRI-inactive regions (``dead zones'') of protoplanetary disks. \\end{enumerate}  Concerning point 1, the first-born planetesimals with sizes larger than $R_{\\rm acc}$ may be formed rapidly by an efficient capture of $\\sim$meter-size boulders in vortexes \\citep{Johansen07}.  Such large planetesimals may be consistent with the size distribution of asteroids (Morbidelli et al. 2008). Even if the first-born planetesimals are not as large, a small fraction of them could continue to grow larger than $R_{\\rm acc}$ by accreting smaller bodies, because accretion is not completely cut off as soon as $v_{\\rm disp} \\ga v_{\\rm esc}$ (there is always a small possibility for accretion) and the large planetesimals would not be disrupted by smaller ones.  This possibility, however, must be examined by a more detailed growth model taking into account the effect of fragmentation and the size distribution of planetesimals, which we neglected in this paper.  Concerning point 2, we have shown that planetesimals are most fragile at early times, in massive disks, and at large orbital distances. We therefore suggest that the growth towards planet sizes may be delayed due to MRI turbulence, and then proceed from inside out: planetesimals should start accretion first close to the star, then progressively at larger orbital distances, as the gas surface density declines.  The possibility to delay planet formation while keeping non-migrating km-size planetesimals is noteworthy because it would help planetary systems resisting to type-I migration: they would grow in a gas disk that is less dense, and for which migration timescales may be considerably increased.  \\citet{IL08}, \\citet{Alibert05}, and \\citet{Kominami06} showed that type-I migration must be lowered by one to two orders of magnitude from the linear calculation \\citep{Tanaka02} to provide an explanation for the existence of a population of giant planets in agreement with observations.  This ``late formation'' scenario is consistent with the noble gas enrichment in Jupiter \\citep{Guillot06}. However, in order to form gas giants, core accretion and gas accretion onto the cores must proceed fast enough to capture Jupiter-mass amount of gas from the decaying gas disk. Once the size of the largest planetesimals exceeds $\\sim 1000$km, their eccentricities are damped by tidal drag and dynamical friction from small bodies. Most of the other small bodies may be ground into sizes smaller than 1\\,km and their eccentricities could be kept very small by gas drag and collision damping.  This could facilitate the runaway accretion of cores to become large enough ($\\ga 10 M_{\\oplus}$) for the onset of runaway gas accretion.  This issue also has to be addressed by a detailed planetesimal growth model taking into account a size distribution. The likelihood of relatively rapid gas accretion without long ``phase 2'' is discussed by \\citet{Shiraishi}. If planetesimal sizes are relatively small, gas drag damping opens up a gap in the planetesimal disk around the orbit of a core and truncates planetesimal accretion onto the core.  The truncation of heating due to planetesimal bombardment enables the core to efficiently accrete disk gas.  Concerning point 3, the MRI inactive region (``dead zones'') may exist in inner disk regions in which the surface density is large enough to prevent cosmic and X rays from penetrating the disk \\citep{gammie96,sano00}.  The preservation of a dead zone can also contribute to stall type-I migration by converting it to type-II migration \\citep{Matsumura07} or by creating a local region with a positive radial gradient of disk pressure near the ice line \\citep{IL08b}.  However, dead zones can be eliminated by turbulent mixing/overshoot \\citep{Varniere06,Turner07,Ilgner08}, a self-sustaining mechanism \\citep{Inutsuka05}, and dust growth \\citep{sano00}.  The last effect comes form the fact that small dust grains are the most efficient agents for charge recombination. According to grain growth, the ionization of the disk and its coupling with the magnetic field become stronger to activate MRI turbulence. We remark that if MRI turbulence is activated, collisions are disruptive and they re-produces small grains to decrease the ionization degree. This self-regulation process might maintain a marginally dead state and keep producing small dust grains.  This might be related with relative chronological age difference ($\\sim 2$Myr) between chondrules and CAIs \\citep[e.g.,][]{Kita05}.  Whether dead zones exist or not is one of the biggest issues in evolution of protoplanetary disks and planet formation.  A more detailed analysis of planetesimal accretion in turbulent disks could impose a constraint on this issue.  At large orbital distances (10's of AU), the existence of a highly erosive regime that lasts until late in the evolution of the protoplanetary disk is an important feature of this scenario. It shows that the entire mass of solids is highly reprocessed by collisions, in qualitative agreement with the paucity of presolar grains (intact remnants from the molecular cloud core) found in meteorites. It also prevents the growth of large planetesimals and helps to maintain a large population of small grains in the disks.  This is in qualitative agreement with observations that do not indicate a significant depletion of micron-sized grains with time, contrary to what would be predicted in the absence of turbulence \\citep{Dullemond05,Tanaka05}.  In conclusion, the existence of MRI turbulence may be a threat to planetesimal accretion.  Given the uncertainties related to these explanations, we cannot provide a definitive scenario for the formation of protoplanetary cores. However, it offers several promising hints to explain important features of planet formation as constrained by today's observations of protoplanetary disks, exoplanets and meteoritic samples in the Solar System."
    },
    "0807/0807.4144_arXiv.txt": {
        "abstract": "{{The effects that gamma-ray bursts have on their environments is an important and outstanding issue. Dust destruction in particular has long been predicted while observational evidence is difficult to obtain. We examine the evidence for dust destruction by GRB\\,980703, in which various inconsistent measurements of the host galaxy extinction have been made using the GRB afterglow emission.}{ We construct a spectral energy distribution from nIR to X-ray to measure the extinction at early times and compare this with previous findings. We also construct nIR/optical SEDs at intermediate epochs to examine a previously reported decrease in extinction.}{ The extinction is very high for a GRB host galaxy. The earliest extinction measurement is likely to be lower than previously estimated, and consistent with most later measurements. In a series of SEDs we do not find any evidence of variable extinction. We therefore conclude that there is no clear evidence of dust destruction in this case.}{} ",
        "introduction": "The search for the predicted effects of gamma-ray bursts (GRBs) on their environments is key to understanding the properties of their host galaxies, found both nearby and in the early Universe. Ionisation by the powerful GRB and afterglow emission has now been seen in the form of optical absorption line variability (e.g. GRB\\,020813, Dessauges-Zavadsky et al., 2006; GRB\\,060418, Vreeswijk et al., 2007) and hinted at from several reports of low significance variations in X-ray column density for some GRBs, for example GRB\\,011121 \\citep{piro}, GRB\\,050730 \\citep{050730}, GRB 050904 \\citep{boer,campana,gendre2} and GRB 060729 \\citep{grupe} and from comparison of optical and X-ray hydrogen column densities \\citep{watson}.  Optical extinction should also decrease as we expect dust to be destroyed by the GRB jet out to 10--30 pc \\citep{drainesal,waxdraine,fruchter,perna}. This should occur very shortly after the onset of the GRB, but so far only one such GRB had potentially shown this effect, GRB\\,980703, and the decrease in extinction would have occurred from approximately 1 day after the GRB onwards \\citep[e.g.][]{vreeswijk1,holland,columnsI}. The evidence of this comes from two sources. Firstly the suggestion by \\cite{castrotirado} that the spectral slope estimated from their $R$- and $H$-band photometric data from around 1 day after the burst indicated a redder value than would normally be expected for GRBs, translating into $A_V$ $\\sim$ 2.2. This is large when compared with the average value of $A_V= 1.07$ found from afterglow studies at later times. Secondly, \\cite{vreeswijk1} compiled spectral energy distributions (SEDs) at four epochs covering 2.2--5.2 days after the burst and derive a potential decrease in the intrinsic optical extinction at the 2--3$\\sigma$ level.  There have been optical extinction measurements of GRB\\,980703 by five groups using a number of different methods and instrumentation, and measurements of $A_{V}$ at different times are far from consistent (Figure \\ref{literature}; Table \\ref{tab:literature}). There is, therefore, a need to return to the apparent variability of $A_V$ in GRB\\,980703, which if confirmed would be the first observational evidence of this kind. We do this via broadband spectral energy distributions, where possible combining an X-ray spectrum with early near-infrared (nIR) and optical data, and thereby discuss the observational manifestations of dust destruction by GRBs. This issue is particularly important to re-examine in the current era of rapid follow-up by the {\\it Swift} satellite \\citep{gehrels}, robotic telescopes (e.g. Liverpool Telescope and Faulkes Telescopes, 2m in diameter) and rapid response mode on larger ground-based telescopes such as the 4m William Herschel Telescope and the 8m Very Large Telescopes.  \\begin{figure} \\centering \\includegraphics[angle=90,width=8cm]{fig1.ps} \\vspace{0.6cm} \\caption{Optical extinction ($A_V$) values reported in the literature (squares). All errors are plotted at the 1$\\sigma$ confidence level.} \\label{literature} \\end{figure} \\begin{table*} \\begin{center} \\renewcommand{\\arraystretch}{2.0} \\begin{tabular}{|cc|ll|} \\hline Time since burst (days) & $A_V$ & Method & Reference \\\\ \\hline 0.94 & 2.2 & nIR/optical colour & Castro-Tirado et al. 1999\\\\ 1.22 & 1.5$\\pm$0.1 & nIR/optical/X-ray SED continuum & Vreeswijk et al. 1999 \\\\ 1.33 & 0.9$\\pm$0.1 & nIR/optical/X-ray SED continuum & Starling et al. 2007 \\\\ 3.22& 1.38$\\pm$0.35 & nIR/optical SED continuum & Vreeswijk et al. 1999 \\\\ 4.4 & 0.3$\\pm$0.3 & host spectrum Balmer decrement\\footnotemark & Djorgovski et al. 1998\\\\ 4.42& 0.84$\\pm$0.29 & nIR/optical SED continuum & Vreeswijk et al. 1999 \\\\ 5.22& 0.90$\\pm$0.25 & nIR/optical SED continuum & Vreeswijk et al. 1999 \\\\ 5.3 & 0.9$\\pm$0.2 & nIR/optical SED continuum & Bloom et al. 1998 \\\\ \\hline \\end{tabular} \\caption{Measurements of optical extinction gathered from the literature. Errors are either given at the 1$\\sigma$ level or assumed to be so where no confidence level is stated.} \\label{tab:literature} \\end{center} \\end{table*} ",
        "conclusions": "The optical extinction measurements in the literature for GRB\\,980703 are inconsistent, and suggested a decrease in extinction with time indicative of dust destruction which is predicted to occur in GRBs. We have re-examined the greatest outlier in these measurements, using broadband SED fits rather than optical colour, to estimate the extinction. We find a lower value than previously estimated, in line with most of the later measurements. We have also investigated the apparent decrease in extinction by recreating the nIR/optical SEDs at these epochs with extended wavelength coverage and a wider range of models for the extinction law. Both from our new SED fit results and given that the extinction measurements begin as late as 0.9 days after the burst, we conclude that there is no evidence of dust destruction by this GRB. In the current era of rapid accurate GRB localisations and follow-up we now have an opportunity to observe some of the effects of GRBs on their environments, and studies such as this using multiple early-time SEDs will be of great value."
    },
    "0807/0807.1885_arXiv.txt": {
        "abstract": "Atomic hydrogen escaping from the extrasolar giant planet HD\\,209458b provides the largest observational signature ever detected for an extrasolar planet atmosphere. In fact, the upper atmosphere of this planet is evaporating. Observational evidences and interpretations coming from various models are reviewed. Implications for exoplanetology are discussed. ",
        "introduction": "\\label{sec:intro} Extrasolar planets are now commonly found around main sequence stars, usually by measuring the variations of the stellar velocity projected on the line of sight. If the inclination of the planetary system is close to 90\\deg, the planet is seen from Earth passing in front of its host star: this transiting event is the current best tool used to characterize extrasolar planets. The transit of HD\\,209458b, producing a $\\sim 1\\%$ dip in the stellar light curve, has permitted to confirm the gaseous nature of extrasolar giant planets (Charbonneau \\etal\\ 2000; Henry \\etal\\ 2000). Using transmission spectroscopy, it was further possible to probe the composition and structure of the planet atmosphere.  Transmission spectroscopy consists in measuring the depth of the transit light curve -- which is related to the planetary radius -- as a function of wavelength (see the contribution of Marley in this volume). Atmospheric sodium is for instance detected as a supplementary absorption of $\\sim 10^{-4}$ in the \\ion{Na}{i}~D lines at 589~nm (Charbonneau \\etal\\ 2002). Since the absorption is located in the core of the lines, it appears stronger ($\\lesssim 10^{-3}$) at a better spectral resolution, thus probing higher altitudes in the planet atmosphere (Sing \\etal\\ 2008; Snellen \\etal\\ 2008), typically $\\sim 1\\,000$~km above the optically thick radius $r_p$ resulting from Rayleigh scattering by molecular hydrogen (Lecavelier des Etangs \\etal\\ 2008). Yet, it remains a tenuous signature in the visible.  On the other hand, intriguing strong absorptions were observed during the planetary transit in some atomic lines located in the ultraviolet part of the spectrum: the Lyman~$\\alpha$ (\\lya) line of neutral hydrogen (\\ion{H}{i}) at 121.6~nm yields a dimming of 5 to 15\\% depending on the spectral resolution (Vidal-Madjar \\etal\\ 2003, 2004, 2008; Ben-Jaffel 2007; Ehrenreich \\etal\\ 2008) and lines of atomic oxygen (\\ion{O}{i}) around 130.5~nm and singly ionized carbon (\\ion{C}{ii}) close to 133.5~nm show a dip between 7 and 13\\% (Vidal-Madjar \\etal\\ 2004).  In fact, these absorptions are larger than the signature of the planet itself during the transit. Hence, they must originate in the planet upper atmosphere at the level of the Roche lobe. Consequently, this gas must be escaping the planet gravity: HD\\,209458b is evaporating. After giving some keys to the understanding of the escape process (\\S\\ref{sec:defs}), I will review the observational results (\\S\\ref{sec:obs}) and their possible interpretations (\\S\\ref{sec:models}). Finally, I will discuss the implications and prospects of evaporation in the frame of comparative exoplanetology (\\S\\ref{sec:exoplaneto}). ",
        "conclusions": "A planetary transit in the ultraviolet produces remarkably large spectroscopic signatures compared to what is observed in the visible. So far, the observational results obtained on planet HD\\,209458b have been obtained from two different instruments and strengthened by independent data analysis. The evaporation of close-in extrasolar planets also relies on a solid theoretical ground. Progress will come from the observations of other cases of extended exospheres and evaporation remnants. This will be feasible after \\emph{HST} Servicing Mission 4, which aims are -- among other few -- to fix STIS and install the Cosmic Origins Spectrograph (COS; Green, Wilkinson \\& Morse 2003)."
    },
    "0807/0807.2697_arXiv.txt": {
        "abstract": "{A recent, detailed study of the well-known ``{\\it {solar wind disappearance event}}\" of 11 May 1999 traced its origin to a coronal hole (CH) lying adjacent to a large active region (AR), AR8525 in Carrington rotation 1949. The AR was located at central meridian on 05 May 1999 when the flows responsible for this event began.  We examine the evolution of the AR-CH complex during 5$-$6 May 1999 to study the changes that apparently played a key role in causing this disappearance event.} {To study the evolution of the solar source region of the disappearance event of 11 May 1999.} {Using images from the Soft X-ray Telescope (SXT), the Extreme-ultraviolet Imaging Telescope (EIT) and the Michelson Doppler Imager (MDI) to examine the evolution of the CH and AR complex at the source region of the disappearance event.} {We find a dynamic evolution taking place in the CH-AR boundary at the source region of the disappearance event of 11 May 1999.  This evolution, which is found to reduce the area of the CH, is accompanied by the formation of new loops in EUV images that are spatially-and-temporally correlated with emerging flux regions as seen in MDI data.}{In the period leading up to the disappearance event of 11 May 1999, our observations, during quiet solar conditions and in the absence of CMEs, provide the first clear evidence for Sun-Earth connection originating from an evolving AR-CH region located at central meridian. With the exception of corotating interacting regions (CIR), these observations provide the first link between the Sun and space weather effects at 1 AU, arising from non-explosive solar events.} ",
        "introduction": "During the period 11-12 May 1999 solar wind densities, as observed by the Advanced Composition Explorer spacecraft (ACE; \\citet{StF98}) located upstream of the Earth at the Lagrangian point L1, dropped to unusually low values ($<$ 1 cm${^{-3}}$) for extended periods of time ($>$ 24 hours). This unusual and extended density depletion was also accompanied by very low-velocity solar wind flows ($<$300 km s${^{-1}}$) and caused a dramatic expansion of the Earth's magnetosphere and bow shock \\citep{LeC00}. It has been estimated that the expanding bow shock moved outwards to a distance of $\\sim$ 60 Earth radii, the lunar orbit, from its normal location of $\\sim$10 Earth radii.  The extremely spectacular nature of this event has caused it to be referred to as the ``{\\it{day the solar wind nearly died}}\" \\citep{Loc01}.  In a recent study of this event, \\cite{JaF05} traced the solar wind outflows, observed at 1 AU, back to the Sun and showed that the flows responsible for the event began on 05 May 1999 from an active region-coronal hole (AR-CH) complex located at central meridian. They suggested that a continuously evolving CH boundary could cause a pinch-off, leading in turn to a separation of the CH outflow.  The expansion of this detached solar wind flow, by a factor of 6-7, could then give the desired low densities at 1 AU.  Given that the travel time between the Sun and the Earth, at the low velocities observed, was $\\sim$5$-$6 days, it was argued that a pinch-off taking place $\\sim$24$-$48 hours after the start of the coronal hole outflow would cause typical particle densities of 20$-$25 particles cm${^{-3}}$ at $\\sim$0.5 AU (the approximate distance the CH outflow would have moved outwards in 48 hours) to be reduced to 0.1 particle cm${^{-3}}$ at 1 AU.  Thus, the expansion of a large, detached, low-velocity flow region from a small CH (as it propagated out to 1 AU) could give rise to an extremely large, low-density cloud that engulfed the Earth on 11 May 1999, as seen from interplanetary scintillation (IPS) observations \\citep{BaJ03}.  Isotopic ratios of O${^{7+}}$/O${^{6+}}$ are known to be good proxies for associating solar wind outflows to either AR or CH \\citep{LNZ04}.  In a recent study it was shown that the solar source of this event could not be pinned down to either an AR or a CH \\citep{JaF08}, as the O${^{7+}}$/O${^{6+}}$ ratios were sometimes indicative of a CH origin and at other times indicative of an AR origin.  Such a fluctuating O${^{7+}}$/O${^{6+}}$ signature was attributed by \\cite{JaF08} to a dynamic and rapid evolution taking place at the AR-CH boundary region, which was at the solar source of the event.  Although many studies have attempted to understand the disappearance event of 11 May 1999 \\citep{CrS00, FaS00, RiB00, UGF00, BaJ03, JaF05, JaF08}, none have examined the source region of this event.  We study the \\begin{figure} \\centering \\includegraphics[width=0.85\\textwidth]{xarchive-f01.eps} \\caption{Left: Full-disk EIT 195~{\\AA} image on 05 May 1999.  The boxed region contains AR8525 and the small CH lying against its westernmost boundary.  Right: Three-dimensional structure of the coronal magnetic field on 05 May 1999, with the field lines projected on to a source surface at 2.5R${_{\\odot}}$.  The thick wavy line is the magnetic neutral line.  See text for details.} \\label{fig1} \\end{figure} source region of the 11 May 1999 disappearance event to understand the evolution and dynamics of this AR-CH complex and to try and pin-point its implications to solar-terrestrial relationships in the absence of explosive solar events. ",
        "conclusions": "Using both spacecraft observations and tomographic IPS observations \\cite{JaF05} located the solar source region of the disappearance event of 11 May 1999 and showed that the flows responsible for the event originated around 05 May 1999 from a small CH lying adjacent to AR8525.  We examined the AR-CH complex at the source region of this event using EIT, SXT, and MDI observations.  The observations have clearly shown the rapid evolution and changes taking place in the CH lying adjacent to AR8525.  The changes are seen to take place in a $\\sim$24 hour interval starting from 05 May 1999, the approximate launch time of the disappearance event.  \\begin{figure} \\centering \\includegraphics[width=0.85\\textwidth]{xarchive-f04.eps} \\caption{MDI magnetograms of the boxed region of the solar disk from Fig. \\ref{fig1}.  The panels differ in time by a little over 24 hours.  Arrow 1 shows a small region of newly emerging negative flux while arrows 2 and 3 show two evolving bipolar regions.} \\label{fig4} \\end{figure} Based on the combined observations, it appears that the rapid evolution seen in the CH is due to  a process of interchange reconnection taking place between the open CH fields and the closed fields from either the bipolar regions to their west or other small closed field regions as described above.   The exact magnetic topology of the AR-CH region is however, complicated and would require a much more detailed study to isolate and pinpoint the reconnection sites and locations of the opposing polarities involved. What is clear however, is that the interchange reconnection process causes the formation of new bright loops within the CH that can be perceived as a progressive constriction taking place across the CH.  Since there is a a high degree of correlation between solar wind speed and size of the CH from which it emanates \\citep{NoK76, Neu94, Wan94, NeF98, KoF99}, we believe that, in this event, the formation of the new EUV loops would cause a reduction in the CH area and lead to a suppression of CH outflow.  This would then give rise to slower velocity flows form regions that earlier produced faster flows, as has been observed in this event.  As stated above, the rapid changes taking place can be seen to be producing a progressive reduction in the area of the CH by causing a clear constriction or pinch-off across it.  The observations thus provide support for the mechanism suggested by \\cite{JaF05} for causing the long lasting low density anomaly or ``{\\it{disappearance event}}\" at 1 AU.  Since this disappearance event was known to have had significant space weather effects \\citep{Ros00, PaC00}, these observations clearly link the observed effects at 1 AU to a sequence of discernible changes taking place in an AR-CH complex on the Sun.  Not considering CIRs, these observations thus provide, to the best of our knowledge, the first evidence for the so-called Sun-Earth ``transmission-line\" arising from non-explosive solar events.  Whether or not AR open fields connect to the interplanetary medium to produce solar wind outflows has been debated for some years now \\citep{KoF99, LuL02, ArH03, SDe03}.  However, the first actual observations of solar wind outflows from AR open fields located at central meridian and lying at the boundary of an AR and a CH have recently been reported \\citep{SaK07}.  These authors have shown that the observed solar wind outflows came from regions that showed large flux expansion factors and low-velocity solar wind, as identified by tomographic IPS observations.  It must be noted here that the work by \\cite{JaF05} has reported both large flux expansion factors and low velocities from the source region of the 11 May 1999 disappearance event.  As opposed to the well-known drivers of space weather phenomena like CME's or large flares, disappearance events are not associated with explosive solar phenomena.  However, they do produce other observable effects that are not fully understood.  For example, \\cite{BaJ03} have reported very unusual IPS power spectra attributed to high-energy Strahl electrons.  The study of such events is therefore important in establishing and understanding solar terrestrial relationships in absence of explosive solar events.  With the exception of CIR's our observations, as stated earlier, provide the first evidence for solar terrestrial connection caused by a non-explosive solar event.  Solar wind disappearance events constitute extreme deviations from the average conditions expected in the solar wind at 1 AU.  It would therefore be important to continue such studies using both ground and space-based data to gain a better understanding of the dynamics and evolution of AR-CH boundary fields."
    },
    "0807/0807.2835_arXiv.txt": {
        "abstract": "We carried out extensive numerical orbit integrations to probe the long-term chaotic dynamics of the two strongest mean motion resonances of Neptune in the Kuiper belt, the 3:2 (Plutinos) and 2:1 (Twotinos).  Our primary results include a computation of the relative volumes of phase space characterized by large- and small-resonance libration amplitudes, and maps of resonance stability measured by mean chaotic diffusion rate.  We find that Neptune's 2:1 resonance has weaker overall long-term stability than the 3:2---only $\\sim 15\\%$ of Twotinos are projected to survive for 4~Gyr, compared to $\\sim 27\\%$ of Plutinos, based on an extrapolation from our 1-Gyr integrations.  We find that Pluto has only a modest effect, causing a $\\sim 4\\%$ decrease in the Plutino population that survives to 4~Gyr. Given current observational estimates, and assuming an initial distribution of particles proportional to the local phase space volume in the resonance, we conclude that the primordial populations of Plutinos and Twotinos formerly made up more than half the population of the classical and resonant Kuiper Belt.  We also conclude that Twotinos were originally nearly as numerous as Plutinos; this is consistent with predictions from early models of smooth giant planet migration and resonance sweeping of the Kuiper Belt, and provides a useful constraint for more detailed models. ",
        "introduction": "Over the past 15 years of observations, it has become clear that the trans-Neptune population of Kuiper Belt objects (KBOs) consists of several distinct dynamical classes \\citep{JL00,MDL00,GMV08}. The ``classical'' KBOs, found primarily in the semimajor axis range of 40~AU to 48~AU, consist of two sub-classes: a dynamically cold population having relatively circular and low-inclination orbits, and a dynamically hot population with higher orbital eccentricities and inclinations \\citep{Brown01,LS01,TB02}.  The scattered disk objects (SDOs) have highly eccentric and inclined orbits, with perihelia a few~AU beyond Neptune's orbit \\citep{DL97}; a small number of the SDOs with perihelia $>40$~AU are thought to be a distinct population, the ``extended scattered disk'' (also known as the `detached objects') \\citep{Gladman02}.  Finally, there are the resonant KBOs, which have orbits in mean motion resonance with Neptune; approximately 23\\% of all known KBOs are in this class \\citep{Chiang07}. Most of the known resonant KBOs are in the 3:2 resonance at semimajor axis $a=39.4$~AU, and have been dubbed `Plutinos' in recognition of the largest and longest-observed member of this population; the 2:1 resonance at $a=47.7$~AU has the next-highest observed resonant population, and have been dubbed `Twotinos'; several other resonances are also populated \\citep{Chiang03,Chiang07}.  The particular interest of this paper is the dynamics of bodies within the 3:2 and 2:1 mean-motion resonances with Neptune.  Because the structure of these resonances provides a readily understandable mechanism for preserving a population for 4~Gyr or more, as well as dynamical pathways to less stable orbits where they may encounter Neptune and be transported elsewhere in the solar system, resonant KBOs have been proposed as the source populations for the Centaurs and Jupiter-family comets \\citep{DLB95,DL97,Morby97}. Resonant KBOs are also of great interest because their relative populations hold clues to the orbital migration history of the giant planets and the dynamical history of the outer solar system as a whole \\citep{Malhotra95,MB04,MLG08}.  In the present work, we use numerical analysis to obtain a detailed map of stability as it varies with location in the phase space of the 3:2 and 2:1 mean-motion resonances with Neptune. One major goal of this study is to estimate (by extrapolating the rates seen in our 1-Gyr numerical integrations) the Plutino and Twotino populations $\\sim4$~Gyr ago when the dynamical structure of the Kuiper belt was presumably established. We also investigate the effects of Pluto on the Plutino population, and the behavior of particles after they have escaped from resonance.  Section~\\ref{plutinos.Resonances} of this paper discusses the properties of resonant particles and our treatment of chaotic diffusion; Section~\\ref{NumModel} describes the numerical experiments that we carried out to study the long-term stability of particles in the 3:2 and 2:1 resonances; Section~\\ref{plutinos.Results} describes the results of our numerical models; Section~\\ref{Discussion} provides discussion, analysis, and comparison to previous work; and Section~\\ref{plutinos.Conclusions} gives a summary and conclusion. ",
        "conclusions": "}  We carried out extensive numerical integrations to probe the characteristics of the 3:2 (Plutinos) and 2:1 (Twotinos) mean motion resonances with Neptune.  In each case we obtained a representative uniform sample of initial conditions in the resonance phase space, and integrated these for 1~Gyr.  Our results can be summarized as follows:  \\begin{itemize} \\item[1.] In both resonances, for orbits that are stable on resonance libration timescales ($\\sim 10^5$~yr), the resonance phase space volume is described by the following distributions of the libration amplitudes ($\\Delta \\phi$).  For Plutinos, $\\Delta \\phi$ has uniform distribution between $45^{\\circ}$ and $160^{\\circ}$; only a small fraction, $<5\\%$, of the phase space volume has $\\Delta\\phi<30^\\circ$.  For Twotinos, symmetric librators have a nearly uniform distribution of $\\Delta\\phi$ between $145^{\\circ}$ and $165^{\\circ}$, while asymmetric librators (both leading and trailing) have a nearly uniform $\\Delta\\phi$ distribution between $20^{\\circ}$ and $75^{\\circ}$; smaller and larger libration amplitudes represent $<10\\%$. \\item[2.] The presence of Pluto has only modest effects on the Plutino population.  It disrupts the stability of objects in the Kozai resonance and causes some particles to undergo libration in a 1:1 (``Trojan'') resonance.  The overall fraction of Plutinos that survive for 1~Gyr is 39\\% when Pluto's perturbations are neglected, and 37\\% when they are included.  Projected to 4~Gyr, these two surviving fractions are 28\\% and 27\\%, respectively (\\Fig{}~\\ref{FracLeftPlot}a).  If, however, the initial distribution of $\\Delta \\phi$ were heavily weighted towards low-amplitude librators, Pluto's effects would likely be more important. \\item[3.] The 2:1 resonance has weaker overall long-term stability than the 3:2; only $\\sim 24\\%$ of Twotinos survive in our 1~Gyr integration, and only $\\sim 15\\%$ are projected to survive for 4~Gyr (\\Fig{}~\\ref{FracLeftPlot}a). \\item[4.] The population decay rates obtained in our models indicate that, 4~Gyr ago, resonant KBOs made up more than half of the Kuiper Belt, and the Plutino/Twotino ratio was closer to unity than it is now.  These estimates depend on our assumption of an initial distribution of particles proportional to the local phase space volume in the resonance and to the extrapolation from our 1-Gyr integrations.  These population ratio estimates for the 4-Gyr old Kuiper belt are roughly consistent with simple models of Neptune's migration and resonance sweeping of the Kuiper Belt \\citep{Malhotra95}, and provide a useful constraint for more detailed models \\citep{MB04,MLG08}. \\end{itemize}"
    },
    "0807/0807.0001_arXiv.txt": {
        "abstract": "Current models of galaxy formation predict satellite galaxies in groups and clusters that are redder than observed. We investigate the effect on the colours of satellite galaxies produced by the ram pressure stripping of their hot gaseous atmospheres as the satellites orbit within their parent halo. We incorporate a model of the stripping process based on detailed hydrodynamic simulations within the Durham semi-analytic model of galaxy formation. The simulations show that the environment in groups and clusters is less aggressive than previously assumed. The main uncertainty in the model is the treatment of gas expelled by supernovae. With reasonable assumptions for the stripping of this material, we find that satellite galaxies are able to retain a significant fraction of their hot gas for several Gigayears, thereby replenishing their reservoirs of cold, star forming gas and remaining blue for a relatively long period of time. A bimodal distribution of galaxy colours, similar to that observed in SDSS data, is established and the colours of the satellite galaxies are in good agreement with the data. In addition, our model naturally accounts for the observed dependence of satellite colours on environment, from small groups to high mass clusters. ",
        "introduction": "Recent multiwavelength imaging with the Sloan Digital Sky survey (SDSS; \\citealt{york00}) has convincingly demonstrated that the colour-magnitude distribution of galaxies is bimodal \\citep{strateva01,hogg02,blanton03,baldry06}. Other galaxy properties, such as the star formation rates, disk-to-bulge ratios, concentrations, stellar surface mass densities and gas mass fractions also show bimodal distributions \\citep{kauffmann03,brinchmann04,balogh04,hogg04,kannappan04}. Elucidating the origin of these distributions requires understanding whether they arise in situ or whether they are environmentally driven (i.e., the ``nature'' or ``nurture'' dichotomy).  While they are not the only means by which theorists attempt to tackle the problem of galaxy formation, at present semi-analytic models offer the best hope for understanding the origin of the bimodality. The level of sophistication of these models has grown rapidly in recent years and so too has their ability to match a wide variety of observational data \\citep{kauffmann93,lacey93,cole94,kauffmann99,somerville99,cole00,benson02,benson03,baugh05,croton06,bower06}.  The recent addition of feedback from active galactic nuclei (AGN) in the semi-analytic models (e.g., \\citealt{granato04,bower06,cattaneo06,croton06,menci06,kang06}) has provided an explanation for three well-known, yet puzzling observational results: (1) the steep cut-off at the bright end of the  galaxy luminosity function; (2) the fact that most massive galaxies today tend to be dominated by old, red stellar populations; and (3) the absence of classical cooling flows in the centres of massive X-ray clusters.  In addition, some of these models (e.g., that of \\citealt{bower06}, hereafter B06, which we adopt as the baseline model in this paper) are also able to reproduce the evolution of the K-band luminosity and the galaxy stellar mass functions out to high redshift (see \\citealt{baugh06} and references therein for a more in-depth discussion of the successes and limitations of semi-analytic models).  In terms of galaxy colours, semi-analytic models with AGN feedback are able to produce a clear bimodal separation of colours at high luminosities, match the slope of the red and blue sequences and explain the absence of massive blue galaxies. However, it has recently become clear that these models are unable to match the relative numbers of faint satellite galaxies on the red and blue sequences. Specifically, \\citet{weinmann06b} and \\citet{baldry06} have shown that the faint satellite galaxies in groups and clusters modelled semi-analytically are on average too red in comparison to satellite galaxies of similar luminosity in the SDSS sample. Using DEEP2 data, \\citet{coil08} have shown that a similar problem may exist at higher redshifts  ($z \\sim 0.7-0.9$), where the semi-analytic models predict a stronger clustering of red galaxies than is observed. These results indicate that one or more physical processes (either internal or environmentally-driven) that affect the ability of the galaxies to form stars are not yet accurately accounted for in the models.  In the present study, we concentrate on the effect of environmental processes on the colours of galaxies, particularly on the role that the ram pressure stripping of the hot gas from galaxy haloes plays.  At present, the semi-analytic models outlined above treat the stripping of the hot gas from the haloes of galaxies in a crude fashion, by assuming that the gaseous halo is completely and instantaneously stripped as the system crosses the virial radius of the (more massive) parent system and becomes a satellite. A direct consequence of this stripping is that the only fuel that is available for star formation once the galaxy becomes a satellite is that which resided in a cold disk when the galaxy first fell into the halo.  As a result, the satellite galaxy experiences a sharp decline in its star formation rate and its stellar population becomes red over time.  This process of cutting off the supply of hot gas that would otherwise cool and replenish the reservoir of cold, star-forming gas, is sometimes referred to as ``strangulation'' or ``starvation'' (\\citealt{larson80,balogh00}).  It is clear, however, that the maximally-efficient stripping assumed by most current semi-analytic models is not realistic, especially in cases where the mass of the satellite is comparable to that of the parent system \\citep{wang07}.  In a very recent study, \\citet{kang08} investigate the effect of modelling the stripping of satellites by an exponential decay. They adopt a stripping factor that is independent of the properties and orbit of the parent and satellite haloes and vary this to achieve the best match to the observations. However, although this empirical approach is a step forward, one expects the efficiency of the stripping to depend on the orbit of the satellite and on the structural properties of both the satellite and parent systems. In contrast, the model we present here aims to implement a physically motivated description for the stripping effect, taking into account the relevant properties of the system.  Recently, \\citet{mccarthy08} carried out a large suite of high resolution hydrodynamic simulations of the stripping of the hot gaseous haloes of galaxies and presented a simple, physically-motivated model that describes the simulation results remarkably well. These authors concluded that, typically, galaxies are able to retain a significant fraction of their hot haloes for long periods of time following virial crossing.  The results of these simulations are in qualitative agreement with a recent {\\it Chandra} X-ray survey of massive satellite galaxies in hot clusters by Sun et al.\\ (2007), who found that most massive satellites do indeed have detectable hot haloes (see also \\citealt{jeltema08}).  The implication of these results is that the stripping of the hot gaseous haloes is much less efficient than previously assumed.  Consequently, some replenishment of the cold star-forming reservoir (via cooling of the hot halo) in the satellite galaxies is expected to take place, this prolonging star formation and resulting in bluer satellite galaxies.  In the present study, we incorporate the prescription for ram pressure stripping of hot haloes of \\citet{mccarthy08} in the Durham semi-analytic code for galaxy formation, {\\small GALFORM}, and show that this improvement in environmental physics brings the colours of satellite galaxies into much better agreement with the observational data.  The layout of the paper is as follows:  In Section \\ref{sec:model} we present the details of the implementation of the new ram pressure model in {\\small GALFORM}. In Section \\ref{sec:results} we present results for the fraction of blue galaxies in the new model (\\S \\ref{sec:bluefrac}), discuss the environmental signatures in the colours of galaxies (\\S \\ref{sec:environ}) and present predictions for the red and blue luminosity functions at redshift $z=0.1$ for galaxies of different type (\\S \\ref{sec:lumfunc}). In Section \\S \\ref{sec:discuss} we summarize and discuss our findings.  \\section[]{A Semi-analytic Model of Galaxy Formation with Ram Pressure Stripping of Hot Gaseous Haloes} \\label{sec:model}  \\subsection{The B06 version of {\\small GALFORM}} \\label{sec:B06}  Apart from the minor changes described below, we use the version of the {\\small GALFORM} semi-analytic code described in B06.  This version of {\\small GALFORM} makes use of halo merger histories extracted from the {\\it Millennium Simulation} with the techniques of \\citet{helly03} (see also \\citealt{harker06}). The {\\it Millennium Simulation} \\citep{springel_etal05} was carried out by the Virgo Consortium and it is one of the largest simulations of the growth of structure in the $\\Lambda$CDM cosmology to date, containing approximately 10 billion dark matter particles in a cubic volume of $(500 h^{-1}$ Mpc)$^3$ with a particle mass of $8.6 \\times 10^{8} \\, M_{\\odot}/h$. Throughout the paper we adopt the same cosmological model assumed in the Millennium simulation (and in the B06 model) and quote our results in terms of the Hubble variable $h = H_0/100$~km$^{_1}$ Mpc$^{-1}$.  In terms of baryonic physics, the version of {\\small GALFORM} developed by B06 has the same basic structure as in \\citet{cole00}, but with the addition of an improved treatment of gas cooling and a new scheme for AGN feedback. In the present study, most of the basic parameters of the code (including, e.g., the efficiencies of supernovae and AGN feedback and the timescales for star formation and dynamical friction) are the same as those adopted in B06. However, in order to achieve better agreement with the zero-point colours of the observed red and blue sequences, we increase the value of the yield to $p=0.04$, which is a factor of two higher than the ``standard'' solar value adopted by B06. With this change, the $^{0.1}(g-r)$ colours are redder by $0.1$~mag compared with those in the B06 model (for galaxies on the red sequence). It also improves the metallicity of the intracluster medium (ICM): whereas in the standard B06 model the metallicity of the ICM was too low, $Z_{ICM} \\simeq 0.15 Z_{\\odot}$, in the model with double the yields $Z_{ICM} \\simeq 0.3 Z_{\\odot}$ (Bower et al. 2008), a value which is in better agreement with the X-ray measurements (e.g., \\citealt{baum05}).  We note that large yields such as these have been used in semi-analytical models in the past in order to get better agreement between the model predictions for galaxy colours and the observations (see \\citealt{kauffmann98,delucia04}; but see also the discussion of \\citealt{cole00} about increasing yields and consistency with stellar evolution models).  One further minor change from the version of B06 concerns the division between haloes in the rapid cooling and hydrostatic regimes. B06 adopted a sharp transition between these two regimes.  In particular, the BO6 model assumes that radio-mode feedback is only effective in hydrostatic haloes with  \\begin{equation} t_{\\rm cool}(r_{\\rm cool}) > \\alpha_{\\rm cool}^{-1} t_{\\rm ff}(r_{\\rm cool}), \\label{eq:acool} \\end{equation}  \\noindent where $t_{\\rm cool}$ and $r_{\\rm cool}$ are the cooling time and radius, $t_{\\rm ff}$ is the free-fall time (as defined by \\citealt{cole00}) and $\\alpha_{\\rm cool}$ is an adjustable parameter (set in this model to $0.7$) that rescales the freefall time of the halo and has the effect of controlling the position of the break in the luminosity function\\footnote{Equation (2) of B06 is incorrect; $\\alpha_{\\rm cool}$ should be replaced by $\\alpha_{\\rm cool}^{-1}$.}.  If the above condition is satisfied, it is then determined whether or not the central AGN is able to inject sufficient power to offset the energy being radiated away in the cooling flow and, if the available AGN power is greater than the cooling luminosity, the cooling flow is assumed to be completely quenched.  This strict dichotomy in halo properties has the undesirable effect that small changes in, for example, the gas mass fraction of haloes near the transition can greatly affect the star formation rate in their central galaxy. This leads to a small population of rapidly forming galaxies at the bright end of the blue sequence.  In the present paper, we have refined the criterion to make the transition more gradual. To do this, we reduce the cooling rate in haloes for which the cooling radius and effective freefall radius [$r_{\\rm cool}(t)$ and $r_{\\rm ff,eff}(t) \\equiv r_{\\rm ff}(\\alpha_{\\rm cool} t)$] are close to each other. Specifically, if  \\begin{equation} |r_{\\rm cool}-r_{\\rm ff,eff}| \\le 0.5 \\epsilon_{\\rm cool}^{-1}r_{\\rm cool}, \\end{equation}  \\noindent the net cooling rate is reduced:  \\begin{equation} \\dot{m}_{\\rm cool,eff} = \\dot{m}_{\\rm cool} \\left( 0.5 + \\epsilon_{\\rm cool} [1-r_{\\rm ff}/r_{\\rm cool}] \\right), \\end{equation}  \\noindent where $\\dot{m}_{cool}$ is the cooling rate in the absence of AGN feedback. This function leaves the cooling rate unchanged if the inequality above is not satisfied. For large values of $\\epsilon_{\\rm cool}$ the inequality applies for only a narrow range of $(r_{\\rm cool}-r_{\\rm ff,eff})$, and the behaviour of the B06 model is maintained. Here we adopt $\\epsilon_{\\rm cool}=10$, so that the transition is still quite sharp, and the luminosity function is little affected.  The effect on the colours of galaxies at the bright tip of the blue sequence is, however, noticeable: by suppressing gas cooling in these objects, the tip of the blue sequence becomes redder, tending to curl up towards the red sequence.  This provides a slightly better match to observational data in this part of the colour-magnitude diagram (see \\S \\ref{sec:bluefrac}).  With these basic parameters, we run a model with a complete and instantaneous ram pressure stripping of the hot gaseous haloes of satellites (the ``default'' model) and a model where the ram pressure stripping of the hot haloes of satellites is calculated using the prescription of  \\citet{mccarthy08} (henceforth called the ``hot ram pressure model'').  \\subsection{Implementation of Ram Pressure Stripping} \\label{sec:ram}  Below we give a brief description of the \\citet{mccarthy08} ram pressure stripping model and how it is incorporated into {\\small GALFORM}. The McCarthy et al.\\ model is analogous to the original formulation of \\citet{gunn72} for the stripping of a face-on cold disk, except that it applies to a spherical distribution of hot gas.  Specifically, the hot gaseous halo of the satellite will be stripped if the ram pressure ($P_{\\rm ram}$) exceeds the satellite's gravitational restoring force per unit area ($P_{\\rm grav}$):  \\begin{equation} \\label{eq:ram} P_{\\rm ram} \\equiv \\rho_{\\rm gas,p} v_{\\rm sat}^2 > P_{\\rm grav} \\equiv \\alpha_{\\rm rp} \\frac{G M_{\\rm tot,sat}(r) \\rho_{\\rm gas,sat}(r)}{r} \\, , \\end{equation}  \\noindent where $\\rho_{\\rm gas,p}$ is the gas density of the parent halo, $v_{\\rm sat}$ is the velocity of the satellite with respect to this medium, $M_{\\rm tot,sat}(r)$ is the total mass of the satellite within radius $r$ and $\\rho_{\\rm gas,sat}(r)$ is the density of the satellite's hot halo at this radius.  The coefficient $\\alpha_{\\rm rp}$ is a geometric constant of order unity; McCarthy et al.\\ find that $\\alpha_{\\rm rp} \\approx 2$ gives the best fit to their hydrodynamic simulations. Note that equation (\\ref{eq:ram}) has not introduced any new free parameters into the {\\small GALFORM} semi-analytical model, as $\\alpha_{\\rm rp}$ has been tuned to match the results of hydrodynamic simulations.  The satellite-centric radius where $P_{\\rm ram} = P_{\\rm grav}$ is referred to as the stripping radius and McCarthy et al.\\ assume that within this radius the satellite's gaseous halo remains intact while all gas exterior to this radius is stripped on approximately a sound crossing time.  This physically simple model has been shown to match the results of their high resolution hydrodynamic simulations\\footnote{This includes both Lagrangian smooth particle hydrodynamic simulations with the GADGET-2 code \\citep{springel05} and Eulerian adaptive mesh refinement simulations with the FLASH code \\citep{fryxell00}.} remarkably well for a wide range of orbits, mass ratios, and structural properties.  In our implementation of the hot ram pressure model we fix the stripping radius by setting the ram pressure to its maximum value, which occurs at the pericentre of the satellite's orbit, and the gas is stripped at the instant it crosses the virial radius. We adopt this simplification as, by default, {\\small GALFORM} does not track the full orbital evolution of the satellite galaxies (but see \\citealt{benson03})\\footnote{The Millennium simulation provides full dynamical information for subhaloes until they are tidally disrupted. In view of the limited resolution of the simulation, however, we choose to calculate the merging timescale of subhaloes using the standard \\citet{chandrasekhar43} formula rather than following the subhalo orbits explicitly, as explained in BO6. Since we are interested mainly in the broad statistical distribution of galaxy colours, our results should not be affected by the loss of information about individual orbits.}. In reality, the physical size of the stripping radius is a function of time, as the magnitude of the ram pressure varies along a non-circular orbit.  By setting the ram pressure to its maximum value, we overestimate the amount of stripping that occurs between the time when the satellite crosses the virial radius of the parent halo and when it reaches pericentre for the first time. However, for typical orbits this is not unreasonable, since this timescale is generally a small fraction of the total amount of time that the satellite spends in orbit about the parent halo.  In addition, since the full orbital evolution of the satellites is not followed, we neglect the effects of tidal heating and tidal stripping of the satellite system.  In terms of the removal of hot halo gas, however, ram pressure stripping is a more efficient mechanism than tidal stripping for the vast majority of satellites (i.e., the stripping radius is typically smaller than the satellite's tidal radius; see \\citealt{mccarthy08}).  The pericentres and the velocities at pericentre (both of which are required to compute the maximum ram pressure along the orbit) of the satellites are calculated by assuming the satellites have the same 2-D joint radial ($v_r$) and tangential ($v_\\theta$) velocity distribution of infalling substructure as measured by \\citet{benson05} from a large suite of Virgo Consortium cosmological simulations. \\citet{benson05} finds that the following functional form describes the 2-D distribution well:  \\begin{equation} f(v_{r},v_{\\theta})=a_{1} v_{\\theta} \\exp{[-a_{2}(v_{\\theta}-a_{9})^{2}-b_{1}(v_{\\theta})(v_{r}-b_{2}(v_{\\theta}))^{2}]} \\, , \\end{equation}  \\noindent where $a_{1}$, $a_{2}$, $a_{9}$, $b_{1}(v_{\\theta})$ and $b_{2}(v_{\\theta})$ are coefficients and functions tabulated in \\citet{benson05}. We assume that this distribution is independent of halo mass and redshift. For each satellite we randomly sample this distribution, extracting a radial and tangential velocity pair which, in turn, allows us to calculate the energy and angular momentum (per unit mass) of the orbit.  We compute the pericentre radius and velocity of the satellite at pericentre by assuming that the orbital energy and angular momentum are conserved and by treating the satellite as a point mass orbiting within a Navarro-Frenk-White (NFW) \\citep{navarro96,navarro97} potential with the same total mass and concentration as the parent halo (see, e.g., \\citealt{binney87}).  By assuming that the orbital energy and angular momentum are conserved, we are ignoring the dynamical friction force acting on the satellites.  However, in cases where the dynamical friction force is strong, the satellite will quickly sink to the centre of the parent halo and, in any case, this process is important for the most massive satellites which are less affected by ram pressure stripping.  By default, the version of {\\small GALFORM} presented in B06 assumes that the gas density profiles of the hot haloes of all systems (including satellite systems) follow a $\\beta$-model \\citep{cavaliere76,cavaliere78}, with $\\beta = 2/3$ (which provides a reasonable match to the X-ray surface brightness profiles of massive groups and clusters; Jones \\& Forman 1984), a fixed core radius $r_c = 0.1 R_{\\rm vir}$, and a normalization $\\rho_0$ that is set to yield the correct mass of hot gas within the virial radius.  We adopt the same distribution in the present study.  We note, however, that we have also experimented with NFW density profiles for the hot gas (for both the satellite and parent systems), and find that the resulting fraction of satellite galaxies that are blue is quite similar to the case of the $\\beta$-model. This is likely to stem from the fact that, typically, the hot gas is only stripped down to intermediate radii, where the $\\beta$-model and the NFW profile are similar.  In addition to a hot tenuous atmosphere, each halo has a reservoir of cold gas that developed from the cooling of the hot atmosphere and which can potentially form stars. As stars are formed, some of this cold gas is re-heated by feedback from, for example, supernovae winds.  Prior to stripping, the re-heated gas is assumed to follow the same spatial and thermodynamic distribution as the hot gas in the halo (as in B06). We therefore treat the initial stripping of re-heated gas in the same fashion as the hot gas, and transfer the same fraction (as computed with the ram pressure stripping algorithm described above) from the re-heated gas of the satellite to the re-heated gas reservoir of the parent halo. We record the total mass stripped from the halo as $M_{\\rm strip}$ and the fraction of the hot halo that is stripped as $f_{\\rm strip}$. Although this procedure is simple, it is worth noting that in reality the spatial and thermodynamic properties of the re-heated gas could be quite complicated (e.g., multiphase) and that it need not be distributed in the same way as the hot gas phase.  It would be worth revisiting this issue once a better treatment of the re-heated gas is incorporated into the semi-analytic model (Benson et al., in prep.).  The cooling rate of the remaining, unstripped gas is calculated by cooling only the gas within the stripping radius and assuming that stripping does not alter the mean density of gas within this radius. We implement this by giving the satellite a nominal hot gas mass $M_{\\rm hot}' = M_{\\rm hot} + M_{strip}$ (where $M_{\\rm hot}$ is the true hot gas content of the halo) and applying the same cooling algorithm as that used for central galaxies (except limiting the maximum cooling radius to $r_{\\rm strip}$ rather than $R_{\\rm vir}$). This step ensures self-consistency in the treatment of the gas cooling between stripped and unstripped galaxies, and therefore that the colours of satellites are predicted correctly.  So far, we have only been considering the stripping of gas which was in the hot or re-heated phases at the time the satellite halo reached pericentre in the parent halo for the first time. However, as long as star formation continues in the satellite galaxy, supernova feedback will continue to re-heat gas and eject it from the satellite galaxy. The colours of satellite galaxies turn out to be sensitive to how much of this ``secondary'' re-heated gas is stripped. The numerical simulations of \\citep{mccarthy08} do not provide any direct information about the stripping of this re-heated gas, since they only treat the stripping of the initial hot halo. One can imagine several possible scenarios for the fate of the re-heated gas: we adopt a picture in which the the re-heated gas is ballistically ejected as relatively cold material that subsequently mixes or evaporates to become part of the hot halo. In isolated galaxies, the re-heated gas is put back with the same radial distribution as for the original hot gas halo. In satellite galaxies, we can consider two extreme cases for the treatment of the secondary re-heated gas. Continuing to apply the initial stripping criterion to the satellite galaxy as it orbits suggests that the re-heated gas should be stripped by the same factor $f_{\\rm strip}$ as it is ejected from the galaxy. However, applying this at every timestep is too extreme: most of the re-heating occurs during the outer part of the galaxy's orbit where the ram-pressure force is small. It might therefore be more appropriate to consider a second case where little of this re-heated gas is stripped by ram-pressure effects. Further numerical simulations are required to elucidate which of these scenarios is physically more realistic and to determine a suitable parameterization of the time averaged stripping rate. At this point, the second (minimal stripping) case seems more appropriate since the typical orbital timescale exceeds several Gyr, and is comparable to the timescale for the mass-growth of the parent halo (see below).  To allow for these uncertainties, we adopt a partially empirical approach, and model the time-dependence of the hot gas mass in the satellite halo after its first pericentre passage as:  \\begin{equation} \\dot{M}_{\\rm hot} = (1 - \\epsilon_{\\rm strip}f_{\\rm strip}) \\frac{M_{\\rm reheat}}{\\tau_{\\rm reheat}} - \\dot{M}_{\\rm cool} \\end{equation}  Here ${M}_{\\rm hot}$ is the mass of hot gas available for cooling and $\\dot{M}_{\\rm cool}$ is the rate at which it cools.  $M_{\\rm reheat}$ is the mass of gas which has been re-heated by supernovae but not yet returned to the hot phase. In the absence of ram-pressure stripping, this re-heated gas is assumed to return to the hot phase on a timescale $\\tau_{\\rm reheat} = t_{\\rm dyn}/\\alpha_{\\rm reheat}$, as in B06, where the value $\\alpha_{\\rm reheat}=1.26$ was chosen to match the observed galaxy luminosity function. The effect of ram-pressure stripping is described by the the term $\\epsilon_{\\rm strip}f_{\\rm strip}$, where $f_{\\rm strip}$ is the stripping factor already calculated for the initial pericentre using equation~(4), and $\\epsilon_{\\rm strip}$ is a new parameter (representing the time averaged stripping rate after the initial pericentre) which we adjust to fit the observations. The first (maximal) stripping case corresponds to $\\epsilon_{\\rm strip}=1$ if the orbital timescale is much shorter than the re-heating timescale, and the second case (minimal stripping) to $\\epsilon_{\\rm strip}=0$. The gas which is stripped from each satellite is added to the hot gas component of the parent halo.  We find that if we take $\\epsilon_{\\rm strip}=1$, then the colour distribution of satellite galaxies looks very similar to the default model in which there is complete stripping of the initial hot gas and of the re-heated gas, with all satellites lying on the red sequence. For $0.2 \\la \\epsilon_{\\rm strip}<1$, the satellite colour distribution looks similar to the case $\\epsilon_{\\rm strip}=1$. For $\\epsilon_{\\rm strip} \\la 0.2$, a blue sequence appears for the satellites, which results in better agreement with the observed colour-magnitude distribution. Values in the range $0 \\leq \\epsilon_{\\rm strip} \\la 0.2$ result in similar colour-magnitude distributions. For this paper, we adopt the value $\\epsilon_{\\rm strip} =0.1$.  The stripping of satellites is also affected by the growth of the halo in which the satellite is orbiting. If we did not allow for the effect of halo growth, a small satellite accreted at high redshift would not feel the increasing ram-pressure effect as the parent halo grows in mass, perhaps becoming a galaxy cluster by the present-day. In order to take this effect into account,  we assign the satellite galaxy new orbital parameters and derive a new stripping factor every time the halo doubles in mass compared to the initial stripping event. If this factor exceeds that applied previously, additional hot and re-heated material is removed from the galaxy. This process tends additionally to suppress on-going star formation in the satellites of massive haloes.  An additional physical effect that may be relevant but that is not accounted for in our model is that the (ram + thermal) pressure force exerted on the hot halo of the satellite could, in turn, raise the pressure force exerted on the cold gas and potentially stimulate additional star formation \\citep{bekki03}. This would tend to result in bluer colours for the satellite galaxies.  The magnitude of this effect is presently unclear, however, and needs to be quantified either observationally or with the aid of hydrodynamic simulations that accurately include the effects of cooling, star formation, and feedback. ",
        "conclusions": "\\label{sec:discuss}  Until now, most current semi-analytic models of galaxy formation have adopted a crude modelling of the ram pressure stripping of the hot gaseous halos of satellite galaxies.  In particular, they typically assume complete and instantaneous stripping of the hot gas halo when the galaxy first falls in, without regard for the galaxy's mass, orbit, or structural properties. This is at odds with results of hydrodynamic simulations, and also with recent X-ray observations of galaxies in massive clusters.  In the present study, we have improved the treatment of stripping by implementing the model of \\citet{mccarthy08} (which has been shown to match simulations of ram pressure stripping to high accuracy) into the {\\small GALFORM} semi-analytic model for galaxy formation. Although the initial stripping event does not require us to add additional parameters to the model, subsequent stripping of the gas re-heated from the disk requires additional parameterization. We parameterize this process by assuming that most of the ejecta are retained in the galaxy after the first pericentre passage.  We find that the newly implemented treatment of stripping leads to a significant improvement in the ability of the model to match the colours of satellite galaxies in the SDSS. The new model is also able to account for the environmental dependence of the colours of satellite galaxies.  Our results suggest that for the majority of satellite galaxies, this environmental process can be as important in modifying the galaxy colours as intrinsic processes, such as AGN or supernovae feedback, which operate within the satellites themselves. This finding is in broad agreement with previous studies that found that internal  processes that quench star formation do not seem capable of explaining the full range of colour and morphology data (e.g., \\citealt{weinmann06a,baldry06}).  Although we have only focused on the stripping of hot gaseous halos in the present study, other environmental processes may be relevant as well.  We now briefly review some of these and conclude that none of them appear to be as important as the ram pressure stripping of the hot haloes and the feedback already incorporated into our present model.  In cases where the ram pressure stripping of the hot halo is complete, some stripping of the cold gaseous disks may also occur \\citep{gunn72,abadi99,quilis00}. However, \\citet{okamoto03} and \\citet{lanzoni05} have explored ram pressure stripping of disks in semi-analytic models and have concluded that the effect on the colours and star formation rates of satellite galaxies are minimal.  This most likely stems from the fact that disk stripping is only expected to be relevant for a minority of satellite galaxies whose orbits take them into the very centres of massive systems (e.g., \\citealt{bruggen08}).  Thermal evaporation of the hot gaseous haloes (and/or disks) could also be relevant \\citep{cowie77}, but observational studies of bubbles and cold fronts in X-ray groups and clusters have placed strong constraints on the efficiency of conduction, concluding that it is strongly suppressed \\citep{markevitch07,mcnamara07}. Turbulent stripping via the generation of Kelvin-Helmholtz and Rayleigh-Taylor instabilities at the interface between the hot gaseous halo of the satellite and the parent system is possible, but the timescale for this type of stripping is generally quite long (see \\citealt{mccarthy08}).  Other possibly relevant processes include viscous stripping (unfortunately at present the viscosity of the hot gas in groups and clusters is poorly constrained; \\citealt{mcnamara07}), tidal effects on the gas as the result of the interaction with the gravitational potential of the parent halo \\citep{byrd90,merritt83} or with the other satellites (i.e., mergers \\citep{mihos95} or harassment \\citep{moore96}), and shock heating of the satellite's hot gas as it falls in at transonic velocities. However, \\citet{mccarthy08} have argued that ram pressure stripping of the hot gaseous haloes is always more efficient than tidal stripping by the parent halo's potential or shock heating in cases where the satellite mass is less than about 10\\% of the parent halo's mass.  In this paper we have concentrated only on comparisons with data at low redshifts ($z<0.1$). In a future study, we intend to test the model at other redshifts and compare the redshift evolution of galaxy colours with data from current deep surveys. Another important application of the model is the study of the clustering of galaxies as a function of different physical quantities, such as colour. In addition to the large scale dependence driven by the relative importance of AGN feedback and ram pressure stripping as a function of halo mass, our model also predicts a radial dependence within a halo driven by the variation in the strength of stripping with the galaxy orbit. The colour dependence of small-scale clustering will constrain the model and perhaps suggest a way to improve our treatment of the interaction between galaxies and their environment.  \\vskip 6pt {\\sf Galaxy catalogues for this model are available for download from this http URL: http://galaxy-catalogue.dur.ac.uk:8080/Millennium/}"
    },
    "0807/0807.3658_arXiv.txt": {
        "abstract": "Map-making presents a significant computational challenge to the next generation of kilopixel CMB polarisation experiments. Years worth of time ordered data (TOD) from thousands of detectors will need to be compressed into maps of the T, Q and U Stokes parameters. Fundamental to the science goal of these experiments,  the observation of B-modes, is the ability to control noise and systematics.  In this paper, we consider an alternative to the maximum-likelihood method, called \\emph{destriping}, where the noise is modelled as a set of discrete offset functions and then subtracted from the time-stream.  We compare our destriping code (Descart: the DEStriping CARTographer) to a full maximum-likelihood map-maker, applying them to 200 Monte-Carlo simulations of time-ordered data from a ground based, partial-sky polarisation modulation experiment.  In these simulations, the noise is dominated by either detector or atmospheric $1/f$ noise.  Using prior information of the power spectrum of this noise, we produce destriped maps of T, Q and U which are negligibly different from optimal. The method does not filter the signal or bias the E or B-mode power spectra.  Depending on the length of the destriping baseline, the method delivers between 5 and 22 times improvement in computation time over the maximum-likelihood algorithm.   We find that, for the specific case of single detector maps, it is essential to destripe the atmospheric $1/f$ in order to detect B-modes, even though the Q and U signals are modulated by a half-wave plate spinning at $5$-Hz. ",
        "introduction": "\\label{sec:intro} Measurements of the temperature and polarisation anisotropy of the cosmic microwave background (CMB) have been used to constrain a number of cosmological parameters to high precision (\\citealt{smoot:1992}, \\citealt{hanany:2000}, \\citealt{masi:2006}, \\citealt{hinshaw:2008}).   There are, however, still open questions that haven't been answered with the new results. In particular there is still no evidence for or against the existence of primordial gravity waves left over from a period of inflation at very early times. A bath of gravitational radiation should leave its imprint on the polarisation of the CMB in the form of a unique B-mode pattern. Such a signal will be faint compared to the polarisation that arises from density perturbations.  If we are to make a convincing measurement of B-modes (or rule out their existence), we need to construct an experiment with far greater sensitivity than the current state of the art. Given the physical limitations on the properties of individual detectors, the only realistic way of doing this is by observing the sky with multiple detectors for long periods of time. Most experiments that are being proposed have these characteristics.  A notable example is the C$_{\\ell}$OVER experiment \\citep{north:2008}. This experiment will have 576 detectors and will focus on small fractions of the sky for periods of up to two years. With current planned technology, it is hoped that the noise per map pixel in the experiment will be below 0.8-$\\mu$K.  To achieve this, it is essential to supplement the increased number of observations that multiple detector pixels bring by removing correlated noise in the map-making stage.  Solutions to the map-making problem aim to compress terabytes of time ordered data (TOD), acquired over months or years of scanning the sky, into a pixelised sky map, with noise uncorrelated between sky pixels.   In each time ordered datum, the signal is typically swamped by noise, necessitating the process of binning and averaging at the heart of every map-making algorithm.  Unfortunately, simple averaging is highly sub-optimal for this task because the noise in the TOD is correlated due to $1/f$ noise.  One could attempt to high-pass filter the resulting low frequency $1/f$ noise drifts, thereby leaving random uncorrelated or \\emph{white} noise in the TOD.  Such a method is not lossless, as it filters the signal as well.  Whilst the effects of filtering can be mitigated in multipole space, by simulating its effects using Monte-Carlo simulations \\citep{hivon:2002}, the variance of the recovered power spectra are larger.  A number of previous experiments, including COBE, MAXIMA and BOOMERANG, have opted to use the optimal maximum likelihood map-making algorithm (eg: \\citealt{smoot:1992}, \\citealt{hanany:2000}, \\citealt{masi:2006}).  This algorithm has been well discussed in the literature (eg: \\cite{tegmark:1997a}, \\cite{natoli:2001}, \\cite{stompor:2002}, \\cite{de-gasperis:2005}) and is renowned for being slow and memory intensive.  In addition, the algorithm requires accurate knowledge of the noise statistics of the TOD.  Use of this algorithm is impossible for the next generation of high resolution, kilo-pixel, polarisation experiments, where we will require Monte-Carlo noise maps to reconstruct the noise covariance required for estimating B-modes on the cut sky \\citep{ksmith:2007}.  Destriping is a promising alternative to the maximum likelihood approach (\\cite{keihanen:2004}, \\cite{keihanen:2005}).  By modelling the correlated noise in the TOD as a series of offset functions mapped onto the time-series, a nearly optimal map can be returned very quickly by solving for the offset function amplitudes. This system is much smaller than the full maximum likelihood system and can therefore be solved at a fraction of the computational cost.  Destriping has been well discussed in the literature for the Planck experiment and favourably compared, in simulations, to the optimal solution: destriping can produce near-optimal maps in orders of magnitude less time (\\citealt{poutanen:2006}, \\citealt{ashdown:2007a}, \\citealt{ashdown:2007b}).  A variant of destriping has been successfully used in the Archeops balloon experiment analysis pipeline \\citep{macias-perez:2007}.  However, these experiments are very different to upcoming ground-based B-mode experiments in terms of scanning strategy, presence of atmospheric noise and in the number of detectors.  Planck and Archeops use a \\emph{circular} scanning strategy: scanning the CMB in overlapping great circles before re-pointing to scan a new great circle.  Scanning strategies are very different for ground based experiments, as a matter of necessity, because the scanning must be in azimuth only (constant elevation) in order to avoid extremely high atmospheric noise contamination from changing elevation.   This causes a large difference in the level of cross-linking, which has a strong effect on the performance of map-making.  The biggest difference between ground based B-mode experiments and Planck or Archeops is the presence of atmosphere (Archeops flew so high it was essentially \\emph{above} the atmosphere) .  Atmospheric $1/f$ noise dominates the correlated noise in the TOD, even for the current generation of ground based experiments. It has a higher knee frequency and higher spectral index than the detector noise.  For example, raw (undifferenced) TOD from the QUaD experiment \\citep{hinderks:2008} appears to exhibit $1/f^{2}$ atmospheric noise rather than simple $1/f$.  Furthermore, the long term noise drifts will be very correlated between detector pixels \\citep{bussman:2005}.  In this paper, we concentrate on experimental designs that modulate the faint polarisation signal to frequencies higher than the detector $1/f$  knee frequency $f_{knee}$, using a rotating half-wave plate, so that Q and U will be sampled in the white noise regime.  Upcoming experiments with this design include C$_{l}$OVER \\citep{north:2008}, EBEX \\citep{oxley:2005} and SPIDER \\citep{mactavish:2007}.  These modulation experiments aim to mitigate possible detector $1/f$ using hardware.  However, sophisticated map-makers will still be required to produce optimal Q and U maps.  If we are to measure B-modes, we must treat the systematics in the data properly.  One of these systematics will be instrumental polarisation, which causes $T \\rightarrow P$ leakage in the data \\citep{johnson:2007}.  If the leakage is significant, it will be vital to remove the leaked temperature signal, a process that will require a high-resolution, near-optimal map of the temperature anisotropies as observed by the instrument.  Such leakage may also mix atmospheric $1/f$ noise into the Q and U data, for which a sophisticated map-maker like destriping will be needed to remove leaked $1/f$ in demodulated Q and U time-streams.  The atmospheric noise itself may also contain a polarised component that would likewise have to be removed in the map-making \\citep{hanany:2003}.    This paper is the first in a series evaluating destriping as applied to ground-based, partial-sky, experiments.  We examine the speed and accuracy of destriping for modulated experiments using simulations, dominated either by detector or atmospheric correlated noise, of a single detector pixel.  This is akin to assuming that the noise is uncorrelated between detectors. We note that this is a poor assumption for the atmospheric $1/f$ component: we will generalise from this assumption in the next paper in the series (see also Appendix \\ref{multi-detector appendix}).   In Section \\ref{algorithms}, we put our discussion on firm footing by deriving the optimal and destriping solutions to the map-making problems and discussing details of the implementation of the algorithms.  In Section \\ref{simulations}, we describe the simulation of the TOD and the scanning strategies.  In Section \\ref{results}, we first apply the algorithms to a preliminary signal-only simulation to test for bias and then proceed to apply the methods to the full Monte-Carlo simulations of TOD containing both signal and noise, comparing the performance of the algorithms directly with emphasis on speed and accuracy. ",
        "conclusions": "In this paper, we set out to compare the destriping approach to map-making to the optimal maximum likelihood estimator.  In particular, we want to know if destriping is fast and accurate enough to replace the maximum likelihood estimator for partial-sky, polarisation modulation experiments.  We have written a new covariant destriping code named Descart (DEStriping CARTographer) and a new maximum-likelihood estimation code and applied them to non-circular scanning strategies for the first time.  We introduce variations in cross-linking through our realistic, but non-ideal, sabre scan and the unrealistic, but ideally cross-linked, fence scan.  Neither of these strategies allow the co-addition of circular scans before destriping, a process relied upon by  the early ``traditional\" destriping efforts (\\citealt{burigana:1997}, \\citealt{delabrouille:1998}, \\citealt{maino:1999}, \\citealt{keihanen:2004}).  We must use a ``covariant\" destriping algorithm \\citep{keihanen:2005}, incorporating prior knowledge of the noise power spectrum through $\\mathbf{C_{a}}$.  This allows the use of very short baselines that produce nearly optimal residuals in the output T, Q and U maps.  We have shown that, without this prior, the noise model clearly breaks down pathologically at shorter baselines and the returned map becomes junk.  The simulations we produced applied destriping to different noise regimes than those to which it has previously been applied.  We have probed regimes of both detector noise and atmospheric noise, the latter of which represents a much greater challenge as it dominates noise even in the current generation of ground-based polarisation experiments.  Previous work on the efficacy of destriping concentrated on the Planck satellite mission, assuming low levels of correlation in the $1/f$ noise, as is possible in the absence of the atmosphere, using knee frequencies of $0.03$-Hz for the HFI \\citep{ashdown:2007a} and $0.05$-Hz for the LFI \\citep{ashdown:2007b}.  We have simulated detector noise with $f_{knee}=0.1$-Hz and atmospheric noise with $f_{knee} = 0.2$-Hz  and a spectral index $\\alpha=1.9$, similar to the level experienced by QuAD \\citep{hinderks:2008}.  Destriping performs better in the more realistic and challenging atmospheric noise simulations than in the detector only noise simulations.  For the shortest baseline of 1-s, it achieves $99.9\\%$ of the possible improvement in T, Q and U residuals achievable with the full maximum-likelihood algorithm.  This amounts to reduction in T residual magnitude of $90.02\\%$ and in Q and U of $55.04\\%$.  One of the motivators of this study is the requirement to produce high quality T maps, in addition to the Q and U maps, in order to remove $T\\rightarrow P$ leakage from instrumental polarisation (eg: \\citealt{johnson:2007}), which will also leak atmospheric $1/f$ into the Q and U signal despite the modulation of those signal by the rotating half-wave plate.  We have found that even without the simulation of this leakage, it is vital to deal with the correlated noise in order to observe B-modes: with naive mapping, the significance of the B-mode detection was destroyed, falling from $8.22\\sigma$ for the optimal and 1-s Descart algorithms to $0.02\\sigma$.  This is a conclusion specific to the single detector case, as the atmospheric noise is completely correlated between detector pairs that share the same horn. Exploiting this correlation, by combining data from both channels, can completely remove the atmospheric noise.  The detector noise simulations maps also saw significant improvement, with the 1-s baseline Descart maps showing $96\\%$ of the possible improvement in T and $94.71\\%$ of the possible improvement in Q residuals, though the latter amounts to less than $0.1\\%$ of the noise power in this case.  Whilst this is not reducible by detector time-stream differencing, the effect on the B-mode noise was very small.  The 1-s Descart maps were returned with improvements in computing time of a factor of 6 for both detector and atmospheric noise.  However, for the atmospheric noise simulations, the $98.73\\%$ of the improvement in T and $99.9\\%$ of Q and U was returned using a longer baseline of 10-s, which delivered an improvement in speed of $10 \\times$.  Most significantly, $89.16\\%$ of the improvement in Q residuals were achieved with a very long baseline of 1200-s, returning a speed improvement of $22\\times$ over the maximum likelihood algorithm.  This can be understood by noting that the atmospheric noise has a lot more noise power at low frequencies than the detector noise simulations do.  As the resolution of $C_{a}$ is increased, improvements in the noise are reached more rapidly as the extra noise is resolved.  The destriping algorithm was shown not to distort the CMB signal, which is the prohibitive drawback of fast filtering methods, such as the $C_{N}^{-1/2}$ filter, which introduces extra variance into the power spectra.  Further, destriping is unbiased for B-modes, producing correct null detections for simulations with artificially zero B-modes.  The ideally cross-linked fence scan showed considerably quicker reduction in residuals for T with decreasing baseline length, than did the sabre scan with detector noise, despite identical noise in the two sets of simulations.  Improvement in residuals reached a trough at 10-s baselines: further improvements were negligible for shorter baselines.  Q and U residuals improved similarly to the sabre scan.  Covariant destriping produces very close to optimal maps, but much more quickly than the full maximum-likelihood algorithm, at speeds that make it applicable to the large datasets from upcoming B-mode experiments.  Destriping can also reduce the vast majority of atmospheric noise in single detector time-streams very quickly, using more approximate long baselines.  This paper is the first in a planned series investigating destriping for ground based polarisation experiments.  The focus of future work will be for experiments in which the Q and U signals are subject to $1/f$ noise (eg: QUIET).  For these experiments, the reduction in $1/f$ will be critically important to the science return.  The simulations will be extended to higher resolution and will investigate the effects of sub-pixel signal gradients by using CMB input maps at much higher resolution than the recovered maps.  The simulations will also be extended to multiple detectors, with noise correlated between detectors sourced both from detector fluctuations and from the atmosphere.  Destriping will be applied alongside and directly compared to the only realistic alternative: the Monte-Carlo based MASTER method where the time-stream is high-pass filtered prior to map-making \\citep{hivon:2002}.  \\emph"
    },
    "0807/0807.4287_arXiv.txt": {
        "abstract": "The catalysis of nuclear reactions by negatively charged relics leads to increased outputs of primordial \\lisx\\ and \\ben. In combination with observational constraints on the primordial fractions of \\lisx\\ and \\ben, this imposes strong restrictions on the primordial abundance and the lifetime of charged relics. We analyze the constraints from the catalysis of \\ben\\ on supersymmetric models in which the gravitino is the lightest supersymmetric particle and a charged slepton---such as the lighter stau---the next-to-lightest supersymmetric particle (NLSP). Barring the special cases in which the primordial fraction of the slepton NLSP is significantly depleted, we find that the \\ben\\ data require a slepton NLSP lifetime of less than $6\\times 10^3\\,\\seconds$. We also address the issue of the catalytic destruction of \\lisx\\ and \\ben\\ by late forming bound states of protons with negatively charged relics finding that it does not lead to any significant modification of the limit on the slepton lifetime. ",
        "introduction": "Physics of the early Universe keeps proving to be an invaluable tool for probing particle physics models and, in particular, models of New Physics beyond the Standard Model.  One of the most important viability checks of these models results from the epoch of Big Bang Nucleosynthesis (BBN), i.e., from the early Universe at temperatures of $T\\simlt 1$~MeV. The combination of Standard Model physics, general relativity, and the experimental determination of the baryon-to-photon ratio from the anisotropies of the cosmic microwave background~\\cite{WMAPplusX} forms the framework for the standard BBN (SBBN) which makes firm predictions for the primordial abundances of light elements such as deuterium, helium, and lithium.  The comparison of SBBN predictions with observationally determined primordial fractions of these elements provides an important consistency check of standard cosmology and serves as a remarkably powerful discriminator among models of New Physics~\\cite{SBBN}. The three most generic ways in which New Physics can affect the outcome of BBN are the change in the timing of the reactions caused, {\\em e.g.}, by new significant contributions to the Hubble expansion rate~\\cite{SBBN}, the non-thermal processes from late annihilation and decays of heavy particles~\\cite{nonthermalone,nonthermaltwo,Sigl:1995kk,Jedamzik:1999di,Cyburt:2002uv,Jedamzik:2004er,Kawasaki:2004qu,Jedamzik:2006xz,nonthermalthree}, and the thermal catalysis of nuclear reactions caused by electromagnetically or strongly interacting relics~\\cite{Pospelov:2006sc}.  Catalyzed Big Bang Nucleosynthesis (CBBN) was already discussed almost twenty years ago in Ref.~\\cite{DS,champtwo}. However, only in the last two years after the appearance of Ref.~\\cite{Pospelov:2006sc}, a lot of work has been done~\\cite{Kohri:2006cn,Kaplinghat:2006qr,Cyburt:2006uv,Steffen:2006wx,Pradler:2006hh,Hamaguchi:2007mp,Bird:2007ge,Kawasaki:2007xb,Takayama:2007du,Jittoh:2007fr,Jedamzik:2007cp,Pradler:2007is,Kersten:2007ab,Pradler:2007ar,Jedamzik:2007qk,KusakabeplusX,Pospelov:2007js,Kawasaki:2008qe,Steffen:2008bt} in order to refine various aspects of catalysis and to understand the implications of CBBN in the framework of specific models. The significant interest in CBBN is also fuelled  by its direct connection with collider physics. Indeed, the relics causing catalysis could be directly produced at the Large Hadron Collider (LHC) because of their electromagnetic and/or strong interactions. One of the most interesting and perhaps one of the most natural frameworks is in this context the supersymmetric (SUSY) extension of the Standard Model. There a spectrum with the gravitino $\\gravitino$ as the lightest supersymmetric particle (LSP) and a charged scalar lepton $\\slepton$---such as the lighter stau $\\stau$---as the next-to-lightest (NLSP) is a commonplace occurence even if one adopts restrictive assumptions on the soft SUSY breaking sector~\\cite{Ellis:2003dn,Cerdeno:2005eu,Jedamzik:2005dh,Cyburt:2006uv,Pradler:2006hh,Pradler:2007is,Kersten:2007ab,Pradler:2007ar}. While the gravitino LSP is a promising candidate for dark matter~\\cite{Asaka:2000zh,Bolz:2000fu,Fujii:2003nr,Feng:2004mt,Steffen:2006hw,Pradler:2006qh,Rychkov:2007uq,Steffen:2007sp}, the charged slepton NLSP can be long-lived and thus lead to CBBN~\\cite{Pospelov:2006sc,Cyburt:2006uv,Steffen:2006wx,Pradler:2006hh,Hamaguchi:2007mp,Bird:2007ge,Kawasaki:2007xb,Takayama:2007du,Jedamzik:2007cp,Pradler:2007is,Kersten:2007ab,Pradler:2007ar,Kawasaki:2008qe}. If such a scenario is realized in nature, each Standard Model superpartner produced at the LHC will cascade down to the long-lived $\\slepton$ NLSP. As the lightest Standard Model superpartner, the $\\slepton$ NLSP will then appear as a quasi-stable muon-like particle that can escape the collider detector before decaying into the gravitino; cf.~\\cite{Ambrosanio:2000ik,Feng:2004mt,Steffen:2006hw,Ellis:2006vu,Steffen:2007sp} and references therein.  Thus, one would find signatures that are very different from the excess in missing energy expected in the alternative neutralino LSP scenarios.  The most dramatic catalytic enhancement is seen in the \\lisx\\ and \\ben\\ production triggered by the formation of bound states of \\hef\\ with a (generic) negatively charged relic that we call \\xm.  The catalytic path to \\lisx\\ and \\ben\\ is shown by the following sequence of transformations~\\cite{Pospelov:2006sc,Pospelov:2007js}: \\begin{eqnarray} \\label{eq:traf-Li6} X^- \\to (\\hefm X^-) \\to \\lisxm, \\\\ \\label{eq:traf-Be9} X^- \\to (\\hefm X^-) \\to (\\beetm X^-) \\to  \\benm. \\end{eqnarray} Although the (\\hef\\xm) system has a binding energy of about 350~keV, its formation is delayed until $T=8$~keV by an overwhelmingly large number of energetic photons that photo-dissociate (\\hef\\xm). Thus, the (\\hef\\xm) bound state serves essentially as a ``bottleneck'' for \\lisx\\ production, whereas the path to \\ben\\ goes through the ``double bottleneck'' of (\\hef\\xm) and (\\beet\\xm).  The key for the nuclear catalysis is an enormous enhancement of the reaction rates in the photonless recoil reactions mediated by \\xm~\\cite{Pospelov:2006sc,Pospelov:2007js}: \\begin{eqnarray} (\\hefm X^-) + {\\rm D} & \\to & \\lisxm +X^- \\label{Eq:CBBNLiSix} \\\\ (\\beetm X^-) + n & \\to & \\benm +X^- \\label{Eq:CBBNBeNine} \\ . \\end{eqnarray} Indeed, since the rates of these catalyzed reactions exceed the SBBN rates for the production of \\lisx\\ and \\ben\\ by many orders of magnitude, one finds a strong sensitivity of the efficiency of the primordial \\lisx\\ and \\ben\\ production to the abundance of \\xm\\ at relevant times. For the case of \\xm\\ being a thermal relic, this abundance is governed by the \\xm\\ annihilation rate and by the \\xm\\ lifetime $\\tau_{X^-}$. In turn, observationally inferred limits on the primordial abundances of both \\lisx\\ and \\ben\\ will impose significant constraints on the lifetime of \\xm, its mass, and its interactions.  The limits on the lifetime and the abundance are particularly interesting in view of the possible catalytic solution to the so-called \\lisv\\ problem \\cite{Pospelov:2006sc,Bird:2007ge}, which is a persistent discrepancy between the predicted primordial amount of \\lisv\\ and a factor of 2--3 lower observations of \\lisv\\ in the atmospheres of the metal-poor stars.  For the gravitino LSP scenarios in which \\xm\\ is identified with a negatively charged slepton NLSP, $X^-=\\slepton^-$, the constraints imposed by primordial \\lisx\\ catalysis have already been analyzed in detail in a number of publications~\\cite{Pospelov:2006sc,Cyburt:2006uv,Steffen:2006wx,Pradler:2006hh,Hamaguchi:2007mp,Bird:2007ge,Kawasaki:2007xb,Pradler:2007is,Kersten:2007ab,Pradler:2007ar,Kawasaki:2008qe,Steffen:2008bt}. Assuming a standard cosmological history that leads to a typical thermal $\\slepton$ relic abundance, the bound from \\lisx\\ catalysis translates into an upper limit on the $\\slepton$ lifetime of $\\tau_{\\slepton} = \\tau_{\\champ}\\simlt 5\\times 10^3\\,\\seconds$~\\cite{Pospelov:2006sc,Hamaguchi:2007mp,Kawasaki:2007xb,Pradler:2007is}. In collider-accessible regions of the parameter space, this limit is found to be considerably more restrictive than the BBN constraints associated with electromagnetic/hadronic energy release from $\\slepton$ decays~\\cite{Cyburt:2006uv,Steffen:2006wx,Pradler:2006hh,Kawasaki:2007xb,Pradler:2007ar,Kawasaki:2008qe}. The $\\tau_{\\slepton}$ limit implies a gravitino mass $\\mgr$ well below 10\\% of the slepton NLSP mass $m_{\\slepton}$ for $\\mslepton \\simlt\\mathcal{O}(1\\ \\TeV)$~\\cite{Steffen:2006wx}.  This seems to exclude a kinematical determination of $\\mgr$~\\cite{Buchmuller:2004rq} at the next generation of particle accelerators, which might have been feasible for $0.1\\,{\\mslepton}\\simlt\\mgr< \\mslepton$~\\cite{Martyn:2006as,Hamaguchi:2006vu}. Consequently, the $\\tau_{\\slepton}$ limit puts a big question mark over an idea of a collider test of supergravity via the microscopic determination of the Planck scale~\\cite{Buchmuller:2004rq}, as well as the $\\mgr$-dependent collider test of thermal leptogenesis \\cite{Pradler:2006qh}. Moreover, a post-inflationary reheating temperature above $\\TR \\simeq 10^7\\,\\GeV$ and thereby even the viability of thermal leptogenesis with hierarchical right-handed heavy Majorana neutrinos seems to be disfavored by the \\lisx\\ constraint within the Constrained Minimal Supersymmetric Standard Model (CMSSM) for a standard cosmological history~\\cite{Pradler:2006hh,Pradler:2007is,Pradler:2007ar}. In the CMSSM, the splitting between $\\mstau$ and $\\mgr$ required to evade the $\\taustau$ limit translates also into a lower limit on the gaugino mass parameter, which is assumed to take on a universal value $\\monetwo$ at the scale of grand unification. Indeed, for the natural gravitino LSP mass range in gravity-mediated SUSY breaking scenarios, the cosmologically favored region can be associated with a mass range of the colored superparticles ({\\em e.g.}, a gluino mass of $m_{\\tilde{g}}\\simgt 2.5~\\TeV$) for which it will be very difficult to probe SUSY at the LHC~\\cite{Pradler:2007is,Pradler:2007ar}.  In this paper we analyze the constraints imposed on SUSY models by the catalysis of a primordial \\ben\\ abundance. While---based on the results of Ref.~\\cite{Pospelov:2007js}---we do not expect the constraints from \\ben\\ to be considerably tighter than those from \\lisx, this analysis is warranted for a number of reasons:  \\begin{enumerate}  \\item Observations of \\lisx\\ are extremely difficult because its lines are not resolved spectroscopically with respect to the lines of \\lisv. The claim of a ``\\lisx\\ plateau'' with metallicity \\cite{Asplund} at $\\sim 10\\%$ of the \\lisv\\ abundance is being challenged in the recent paper~\\cite{Cayrel2007}, and a new re-analysis of \\lisx\\ data is warranted as some of the observations may turn out to provide only upper limits. In fact, even the value of an upper limit on the primoridal \\lisx\\ abundance is subject to discussions: Many papers adopt upper limits on primordial $\\lisx/\\Hyd\\equiv n_{^6{\\mathrm{Li}}}/n_{\\mathrm{H}}$ within a range from $10^{-11}$ to $10^{-10}$.  Unlike \\lisx, \\ben\\ is firmly detected in a significant number of stars at low metallicity, and its observational status is not in doubt. For the latest data on the \\ben\\ abundance in metal-poor stars, see, {\\em e.g.}, \\cite{Primas:2000gc,Be9,Boesgaard:2005pf}.  \\item \\lisx\\ is more fragile than \\lisv\\ and would burn more efficiently at lower temperatures. Therefore, if there is a (yet unconfirmed) stellar mechanism (see {\\em e.g.}~\\cite{Korn:2006tv}) that resolves the lithium problem, i.e., that depletes \\lisv\\ by a factor of two or three, \\lisx\\ would have been depleted by an even larger factor. Such a stellar mechanism, however, would affect \\ben\\ less than either \\lisv\\ or \\lisx\\ since both \\lisv\\ and \\lisx\\ are more fragile than \\ben, which thereby provides a more robust bound on New Physics.   \\item The nuclear physics rates that enter in the calculation of \\ben\\ catalysis are dominated by resonances. Given the wealth of experimental information on the \\ben\\ resonances \\cite{Sumiyoshi}, this may eventually allow for very reliable calculations of the catalytic rates.  \\end{enumerate}  While stating strong bounds on primordial abundances/lifetimes of \\xm, a lingering question remains: Is there an ``island'' of allowed lifetimes around $\\tau_{X^-} \\sim 10^6~\\seconds$? One point emphasized in Refs.~\\cite{DS,Kohri:2006cn,Jedamzik:2007cp} is that bound states of \\xm\\ with protons, \\px, may have a significant impact on the primordial abundances of elements such as \\lisx, \\lisv, \\bes\\ (and \\ben). Another point is related to \\xm\\ decays with very energetic decay products and the associated possibility of an environment in which \\lisx\\ may be destroyed efficiently. The second point has already been addressed~\\cite{Cyburt:2006uv,Kawasaki:2007xb,Kawasaki:2008qe}, and it has been found that late energy injection cannot suppress \\lisx\\ down to an acceptable level if it is significantly overproduced at 8~keV~\\cite{Cyburt:2006uv,Kawasaki:2007xb,Kawasaki:2008qe} Despite its importance, the issue of possible ($p$\\xm) catalysis remained largely unresolved.  While the importance of charge exchange reactions of \\px\\ with \\hef\\ that may reduce the abundance of ($p$\\xm) was already mentioned in the early paper~\\cite{DS}, subsequent publications either ignored this issue~\\cite{Kohri:2006cn} or generally underestimated the impact of charge exchange reactions~\\cite{Jedamzik:2007cp}. Recognizing its importance for the whole CBBN paradigm, we revisit the catalysis by ($p$\\xm) bound states. Indeed, in this paper, we put this issue to rest by reaching the conclusion that less than 10\\% of \\lisx\\ and \\ben\\ synthesized at 8 keV could possibly be affected by \\px\\ catalysis, whereas typically suppression factors in excess of 100 are needed in order to evade the corresponding limits. This clarifies that the islands in the parameter region with large abundances/large lifetimes, which were suggested to remain viable in Ref.~\\cite{Jedamzik:2007cp}, cannot exist.  This paper is organized as follows. In the next section we estimate the charge exchange reaction rates relevant for ($p$\\xm)--mediated catalysis and consider their consequences for lithium and beryllium. In Sect.~\\ref{Sec:9BeConstraints} we analyze the bounds on the lifetime/abundance of \\xm\\ imposed by observations of \\ben\\ in stellar atmospheres at low metallicities. The resulting constraints on SUSY models with the gravitino LSP and a charged slepton NLSP are given in Sect.~\\ref{Sec:SUSYImplications}. We reach our conclusions in Sect.~\\ref{Sec:Conclusions}. ",
        "conclusions": "\\label{Sec:Conclusions}  From observations of beryllium in Population II halo stars at very low metallicities, we have extracted a nominal upper limit on primordial beryllium of $\\ben/\\Hyd \\le 2.1\\times 10^{-13}$.  This limit allows one to set interesting constraints on models in which the primordial $A=8$ divide is bridged by catalytic effects.  Considering the primordial catalysis of $^9$Be via bound-state effects of a negatively charged massive relic $X^-$~\\cite{Pospelov:2007js}, we have derived \\tauX-dependent upper limits on the $X^-$ yield prior to decay, \\YXdec. For a typical relic abundance $\\YXdec\\simgt 3 \\times 10^{-4}$ ($10^{-4}$), we find that this $\\ben$ limit translates into an upper limit on the $X^-$ lifetime of $\\tauX\\simlt 6\\times 10^{3}\\,\\seconds$ ($10^{4}\\,\\seconds$), which is quite comparable with the $\\tauX$ limit inferred from the primordial catalysis of $^6$Li. Moreover, in the region where CBBN is efficient, we confirm that the ratio of the synthesized elements of $\\ben/\\lisx$ lies in the range $10^{-3}-10^{-2}$~\\cite{Pospelov:2007js}, which provides perhaps the most model-independent prediction in the whole CBBN paradigm.  We have clarified that the presence of ($p$\\xm) bound states cannot relax the $\\YXdec$ limits at long lifetimes $\\tauX$ in any substantial way. Indeed, we have shown explicitly that late-time effects of ($p$\\xm) bound states can affect the lithium and beryllium abundances synthesized at $T\\simeq8\\ \\keV$ by not more than 10\\%. Any substantial formation of ($p$\\xm) at $T\\simeq 0.7\\ \\keV$ is immediately intercepted by the very efficient charge exchange reaction of ($p$\\xm) with \\hef. This comes as no surprise given the large size of the ($p$\\xm) system $\\sim 30~\\fm$ and the fact that the proton deconfinement probability approaches unity already for a \\hef--\\xm\\ distance of $\\sim 95~\\fm$.  In particular, we find that the fractional density of protons in bound states does not exceed the level of $\\sim 10^{-6}$ for $\\YX\\lesssim Y_{\\hef}$.  Correspondingly, even with a ($p$\\xm)--induced \\lisx-destruction-cross section as large as the unitarity limit, at most a few percent of the synthesized \\lisx\\ could be destroyed. By the same argument, the \\ben\\ yield also remains unaffected by late-time catalysis. Thus, we find that the possibility of allowed islands in the parameter region with typical $\\YXdec$ and large $\\tauX$---which was advocated in Ref.~\\cite{Jedamzik:2007cp}---does not exist.  Applying the $\\tauX$-dependent upper limits on $\\YXdec$ derived from the primordial catalysis of \\ben, we have analyzed the new \\ben\\ constraint in SUSY models in which the gravitino is the LSP and a long-lived charged slepton the NLSP, $\\slepton=X^-$. For typical values of the slepton NLSP yield after decoupling, the $^9$Be constraint obtained in this paper is found in close vicinity to the constraint from the primordial catalysis of $^6$Li. Accordingly, the implications of the $^9$Be constraint for SUSY models do not differ much from the case of $^6$Li.  The important advantage of \\ben-derived constraints is due to the fact that \\ben\\ is firmly detected in a significant number of stars at low metallicities, while the status of \\lisx\\ observations remains somewhat questionable. Another great virtue of \\ben-derived constraints is due to the expectations that an \\ben\\ abundance would remain less affected compared to the one of \\lisx\\ by any stellar mechanism that might have caused the depletion in \\lisv.  Therefore, one would not expect any serious depletion factors between the (hypothetical) primordial fraction of \\ben\\ and the actual observationally determined abundances of \\ben.  One could question the calculational status of the CBBN chain (\\ref{eq:traf-Be9}) leading to \\ben. The rates for the first two reactions, resulting in (\\hef\\xm) and (\\beet\\xm), are determined by electromagnetic interactions and thus are not associated with large nuclear uncertainties.  The final step on the way to \\ben\\ is the neutron capture by the (\\beet\\xm) bound state. It is dominated by a resonant transition, and the associated rate is believed to have a factor of a few uncertainty \\cite{Pospelov:2007js}.  However, given the wealth of existing experimental data on the $^9$Be resonances, and rapid progress in nuclear calculations of few nucleon systems, one could hope that reasonably precise calculations of the catalytic rates may become available.  The main shortcoming of our analysis is that only the catalytic effects below 10 keV are taken into account. In a generic framework of a hypothetical \\xm\\ particle, this is fully justified.  However, in the specific SUSY model with gravitino LSP/slepton NLSP, one could go one step further and by combining the catalytic effects with the energy injection effects. Fortunately for us, the question of \\ben\\ synthesis is somewhat decoupled from the question of energy injection. The connection is mainly due to the modification of the deuterium abundance: For example, hadronic energy input at $T\\sim 30$ keV would lead to a larger deuterium abundance, which in turn would enhance the neutron abundance resulting in a larger abundance of \\ben. Thus, the energy injection would lead to abundances of \\ben\\ that are somewhat {\\em larger} than the ones that we determined by using the standard input for the deuterium abundance.  Since the main idea of our paper is to derive a conservative limit on SUSY models from \\ben, we can stay on the conservative side and disregard this modification, which would make the \\ben\\ limit only stronger.  In comparison to other cosmological constraints on gravitino LSP scenarios with a long-lived charged slepton NLSP, we find that the $^9$Be constraint (together with the one from $^6$Li) is the most relevant one in the collider-accessible region of slepton masses below 1~TeV. Indeed, if a SUSY scenario with a long-lived $\\slepton$ is realized in nature, one might be able to determine the combination ($\\tau_{\\slepton}$,\\,$\\mslepton$) at collider experiments.  Assuming standard cosmological history, one might also ``invert'' the collider data and infer $\\Ysldec$. Independently of the assumption of the gravitino LSP, these quantities can then be confronted with constraints on the ($\\tau_{\\slepton}$, $\\Ysldec$) parameter space imposed by the primordial catalysis of $^6$Li and $^9$Be.  Thereby, the CBBN constraints can be considered as predictions that could be tested in upcoming high-energy experiments. It will be most remarkable if collider measurements point to a ($\\tau_{\\slepton}$,\\,$\\mslepton$) combination in the vicinity of the CBBN constraints. Indeed, this could provide an experimental hint for the primordial catalysis being the origin of existing abundances of both $^9$Be and $^6$Li at lowest metallicities.  However tenuous the BBN--LHC connection may seem at the moment, we expect a lot more clarity brought to this issue in the coming years.  {\\bf Note added} -- After the submission of this paper, a dedicated nuclear physics study of some CBBN reactions has appeared, Ref.~\\cite{Kamimura:2008fx}. It supports the conclusion of this paper about the large rate for the charge exchange reactions that remove $(pX^-)$. At the same time, this work finds non-negligible shifts, $O(100~{\\rm keV})$, of the resonant energy levels employed in the \\ben\\ production chain.  This may affect the overall efficiency of \\ben\\ production, and further investigations of the nuclear rates are needed.  {\\bf Acknowledgments} -- Research at the Perimeter Institute is supported in part by the Government of Canada through NSERC and by the Province of Ontario through MEDT. M.P.\\ would like to thank C. Bird for an independent calculation of the cross sections for the charge exchange reactions. J.P.\\ is grateful to R.~Lang and Y.Y.Y.~Wong for helpful exchange on data fitting and to F.~Hahn-Woernle for discussions on the Boltzmann code. The authors acknowledge useful conversations with all participants of the ``BBN and Particle Physics'' workshop held at Perimeter in May of 2008.  The research of J.P.\\ and F.D.S.\\ is supported in part by the Cluster of Excellence `Origin and Structure of the Universe.' Both would like to acknowledge the hospitality of the Perimeter Institute where part of this work was completed.   \\begin{appendix}  \\begin{section}{CBBN reaction rates below 10 keV} \\label{sec:appendix}  In the following, we collect the key reaction rates, $ N_{\\mathrm{A}} \\langle \\sigma v \\rangle$, used in the numerical solutions of the Boltzmann equations. They are given in units of $\\mathrm{cm}^3\\mathrm{s}^{-1}\\mathrm{mol}^{-1}$ and $T_9 = T/10^9~\\mathrm{K}$. \\begin{itemize} \\item Recombination and photo-dissociation of \\xm: \\begin{align*} \\hef + \\xm \\to (\\hef\\xm) + \\gamma :\\hspace{0.5cm}  & 7900\\; T_9^{-1/2} \\\\[0.1cm] (\\hef\\xm) + \\gamma_{\\mathrm{bg}} \\to \\hef + \\xm :\\hspace{0.5cm}  & 1.85\\times 10^{10}\\; T_9^{-2}\\; \\exp{(-4.03/T_9)} \\\\[0.1cm] p + \\xm \\to (p\\xm) + \\gamma :\\hspace{0.5cm}  & 3980\\; T_9^{-1/2} \\\\[0.1cm] (p\\xm) + \\gamma_{\\mathrm{bg}} \\to p + \\xm :\\hspace{0.5cm}  & 1.18\\times 10^{9}\\; T_9^{-2}\\; \\exp{(-0.29/T_9)} \\end{align*} \\item Charge exchange reactions: \\begin{align*} (p\\xm) + \\hef \\to (\\hef\\xm) + p  :\\hspace{0.5cm}  & 3.9\\times 10^{10}\\; T_9^{1/2}  \\\\[0.1cm] (p\\xm) + \\lisx \\to (\\lisx\\xm) + p  :\\hspace{0.5cm}  & 6.45\\times 10^{10}\\; T_9^{1/2}  \\\\[0.1cm] (p\\xm) + (\\lisx\\xm) \\to (\\lisx\\xm_2) + p  :\\hspace{0.5cm}  & 3.37\\times 10^{9}\\; T_9^{1/2} \\; (1\\ \\TeV/m_{\\xm})^{1/2} \\\\[0.1cm] (p\\xm) + (\\lisx\\xm_2) \\to (\\lisx\\xm_3) + p  :\\hspace{0.5cm}  & 5.25\\times 10^{8}\\; T_9^{1/2} \\; (1\\ \\TeV/m_{\\xm})^{1/2} \\end{align*} \\item \\lisx\\ and \\ben\\ catalysis (from \\cite{Hamaguchi:2007mp} and \\cite{Pospelov:2007js}, respectively): \\begin{align*} (\\hef\\xm) + \\Deut \\to \\lisx + \\xm  :\\hspace{0.5cm}  & 2.37\\times 10^8\\; (1-0.34\\, T_9)\\, T_9^{-2/3} \\exp{(-5.33\\, T_9^{-1/3})}  \\\\[0.1cm] \\hef + (\\hef\\xm) \\to (\\beet\\xm) + \\gamma :\\hspace{0.5cm} & 10^{5}\\; T_9^{-3/2}\\; [\\ 0.95 \\exp{(-1.02/T_9)} \\\\ & \\hphantom{10^{5}\\; T_9^{-3/2}\\; \\qquad}  + 0.66 \\exp{(-1.32/T_9)}\\ ]  \\\\[0.1cm] (\\beet\\xm) + n \\to \\ben + \\xm  :\\hspace{0.5cm}  & 2\\times 10^{9} \\end{align*} \\end{itemize}  \\end{section}  \\end{appendix}"
    },
    "0807/0807.4880_arXiv.txt": {
        "abstract": "The  MAGIC collaboration has  recently reported  the detection  of the strong  \\gray\\ blazar 3C279  during a  1-2 day  flare. They  have used their spectral observations to draw conclusions regarding upper limits on  the  opacity of  the  Universe  to high  energy  \\gray  s and,  by implication, upper limits on the extragalactic mid-infrared background radiation. In this paper we examine the effect of \\gray ~absorption by the  extragalactic infrared  radiation on  intrinsic spectra  for this blazar  and  compare  our  results  with  the  observational  data  on 3C279. We  find agreement with  our previous results, contrary  to the recent  assertion  of  the  MAGIC  group that  the  Universe  is  more transparent to  \\gray s than our calculations  indicate.  Our analysis indicates that  in the energy range  between $\\sim$ 80  and $\\sim$ 500 GeV, 3C279  has a  best-fit intrinsic spectrum  with a  spectral index $\\sim$ 1.78 using  our fast evolution model and  $\\sim$ 2.19 using our baseline model.   However, we also  find that spectral indices  in the range of 1.0  to 3.0 are almost as equally acceptable  as the best fit spectral indices.  Assuming the same intrinsic spectral index for this flare as for the 1991 flare  from 3C279 observed by EGRET, viz., 2.02, which lies  between our best fit  indeces, we estimate  that the MAGIC flare  was $\\sim$3  times brighter  than the  EGRET flare  observed 15 years earlier. ",
        "introduction": "\\label{sec:intro}  It has  long been  recognized that by  studying the spectra  of strong extragalactic  \\gray\\ sources  one  can obtain  information about  the density  and  energy  spectra  of intergalactic  photon  fields.   The luminous blazar 3C279 was discovered  by the EGRET detector aboard the Compton Gamma  Ray Observatory  to be a  strong flaring  \\gray\\ source (Hartman  et al. 1992).   Shortly after  this discovery,  (Stecker, de Jager \\& Salamon 1992) (hereafter  SDS) proposed that the study of the TeV spectra of  such sources could be used  to probe the intergalactic infrared radiation.  SDS proposed to look for  the energy dependent features indicating the mutual annihilation of very high energy \\gray s and low energy photons of  galactic   origin  via  the  process   of  electron-positron  pair production $\\gamma + \\gamma \\rightarrow  e^+ + e^-$. The cross section for this  process is  exactly determined; it  can be  calculated using quantum electrodynamics (Breit \\&  Wheeler 1934).  Thus, in principle, if one  knows the  emission spectrum of  an extragalactic source  at a given  redshift,  one can  determine  the  column  density of  photons between the source and the Earth as a function of redshift.  Since  the  EGRET  discovery  of  3C279, the  infrared  background  at wavelengths not totally dominated by galactic or zodiacal emission has been measured by the  Cosmic Background Explorer (COBE).  In addition, there have  been extensive observations  of IR emission  from galaxies themselves, whose total  emission is thought to make  up the cosmic IR background (see review  by Hauser \\& Dwek 2001).  The latest extensive observations  have been  made by  the Spitzer  satellite.  It  is thus appropriate to use a synoptic  approach combining the very high energy \\gray ~observations  with the extragalactic IR  observations, in order to best explore both the  \\gray ~emission from blazars and the diffuse extragalactic  IR  radiation.   Using  this approach,  one  must  take account of {\\it both} the high energy \\gray\\ observations and the data from the many galaxy observations presently available.  The  MAGIC collaboration  has recently  observed the  spectrum  of the blazar 3C279  during a  flare which occurred  on 22--23  February 2006 (Albert  et al. 2008).   The highly  luminous blazar  3C279 lies  at a redshift  of 0.536  (Marzioni et  al.1996). To  date, it  is  the most distant \\gray\\ source observed in  the sub-TeV energy range.  Thus, as noted  by  SDS, this  source  is  potentially  highly significant  for probing the intergalactic background radiation.  Albert et al. (2008) used their observational data to draw conclusions regarding  the maximum  opacity  of the  Universe  to \\gray  s in  the sub-TeV energy  range.  Their conclusions  regarding the extragalactic background  radiation would  appear  to disfavor  the  results of  the extensive   semi-empirical  calculations   of  the   extragalactic  IR background spectrum  given by Stecker,  Malkan \\& Scully  (2006, 2007) (hereafter SMS).   In particular,  the limit shown  in their  Figure 2 appears to be  inconsistent with one of the SMS  models that was based primarily  on  galaxy  studies   by  the  Spitzer  infrared  satellite telescope.  In this paper, we will reexamine both the analysis assumptions and the conclusions presented in  the paper of Albert \\etal  Using a different analysis  technique that  we show  to be  superior to  that  of Albert \\etal, we  find that  the observations of  3C279 are  fully consistent with  both diffuse  IR background  models  obtained by  SMS.  We  then discuss the implications of these results regarding both the intrinsic energy spectrum and  luminosity of the 3C279 flare  and the opacity of the Universe to \\gray s. ",
        "conclusions": "We  conclude  that  the observational data of  3C279 from the MAGIC collaboration do not significantly constrain the intergalactic low energy photon spectra, nor do they indicate that the Universe is more transparent to high energy \\gray s than previously thought or obtained by the SMS models. Including  the effects  of  the systematic  errors,  which only significantly affect the two highest energy  data points, would  only strengthen  our conclusions.  We show here, as well as in our analysis of 1ES0229 (Stecker \\& Scully 2008),   that  the  SMS   models  produce   reasonable  fits   to  the observational data.   We further demonstrate that other  models in the literature that  give lower transparencies  do not fit the  3C279 data better  than the SMS  models do.   This is  because the  magnitude and shape of the intrinsic spectrum of the flare are not well determined by the MAGIC data, as we have shown.  The spectral  energy distribution  of the intergalactic IR radiation is well  constrained by astronomical data, as discussed in detail in SMS. However, the form of the  intrinsic  source  spectrum  of  the  \\gray  ~flare  is  not  well constrained.  If  one should wish to  speculate on producing  a fit to the  MAGIC data  to go  through  {\\it all}  of the  points, one  would require an intrinsic flare spectrum which flattens at higher energies. Such a spectrum may  arise either from relativistic shock acceleration (Stecker,  Baring   \\&  Summerlin  2007; Resmi \\& Bhattacharya 2008)  or from  intrinsic  source absorption (Liu \\& Bai 2006; Liu, Bai \\& Ma 2008; Sitarek \\& Bednarek 2008). However, given  both the  statistical and  the systematic errors of  the MAGIC data,  particularly those for the  highest energy bin, the invocation of such processes is unnecessary.  If we apply our analysis technique to the MAGIC data and assume an intrinsic power-law spectrum with index 2.02 as observed for the earlier EGRET flare at lower energies where absorption is negligible, we find that the luminosity of the MAGIC flare was similar to that of the earlier EGRET flare. In fact, the MAGIC flare was $\\sim$3 times brighter. This luminosity estimate could not be obtained using the analysis adopted by the MAGIC group, since {\\it all} of the data points that they started from were affected by intergalactic absorption.  One would  need  to have  additional information   concerning  the  intrinsic   source  spectrum   such  as additional observations at lower \\gray\\ energies that are unaffected by pair-production absorption in order to better constrain the intergalactic IR radiation through an  analysis of  the resulting  optical depth at the  higher energies where absorption  is more significant. We note that the ideal situation for exploring the exact amount of intergalactic absorption would be to have simultaneously and with overlapping observational energy ranges (1) observations of a strong flare with large photon counts at lower energies unaffected by intergalactic absorption, as can be obtained by FGST (Fermi Gamma-ray Space Telescope, nee GLAST), and (2) ground based sub-TeV observations of such a strong flare with larger photon counts."
    },
    "0807/0807.3975_arXiv.txt": {
        "abstract": "We have obtained deep Adaptive Optics (AO) images of Vega and $\\epsilon$ Eri to search for planetary-mass companions.  We observed at the MMT in the $L'$ (3.8 $\\mu$m) and $M$ (4.8 $\\mu$m) bands using Clio, a recently commissioned imager optimized for these wavelengths. Observing at these long wavelengths represents a departure from the $H$ band (1.65 $\\mu$m) more commonly used for AO imaging searches for extrasolar planets.  The long wavelengths offer better predicted planet/star flux ratios and cleaner (higher Strehl) AO images, at the cost of lower diffraction limited resolution and higher sky background. We have not detected any planets or planet candidates around Vega or $\\epsilon$ Eri.  We report the sensitivities obtained around both stars, which correspond to upper limits on any planetary companions which may exist.  The sensitivities of our $L'$ and $M$ band observations are comparable to those of the best $H$-regime observations of these stars.  For $\\epsilon$ Eri our $M$ band observations deliver considerably better sensitivity to close-in planets than any previously published results, and we show that the $M$ band is by far the best wavelength choice for attempts at ground-based AO imaging of the known planet $\\epsilon$ Eri b.  The Clio camera itself with MMTAO may be capable of detecting $\\epsilon$ Eri b at its 2010 apastron, given a multi-night observing campaign.  Clio appears to be the only currently existing AO imager that has a realistic possibility of detecting $\\epsilon$ Eri b. ",
        "introduction": "Early space based observations with the IRAS satellite identified four bright, nearby stars with strong IR excesses: $\\beta$ Pic, Vega, Fomalhaut, and $\\epsilon$ Eri \\citep{IRASv1,IRASa1,IRASa2,IRASa3}.  The only reasonable explanation for these excesses is that the systems contain substantial dust, which is warmed by starlight until it radiates brightly in the IR because of the large total surface area of its numerous small grains (see for example \\citet{backman, HD4796, deller}).  The dust in these systems cannot be primordial but must be continually generated by the grinding down of larger bodies such as asteroids \\citep{backman, HD4796, deller}. The stars are therefore said to have `debris disks'.  The clear implication is that each of these stars has at least an asteroid belt, and probably a more extensive planetary system, because it is unlikely that an asteroid belt would form without planets also forming, or that it would continue to grind down without ongoing gravitational stirring due to planets.  Theoretical models (e.g. \\citet{bur} and \\citet{bar}) predict that it should be possible to make direct images of giant planets orbiting nearby, young stars, using the current generation of large ground-based telescopes with adaptive optics (AO).  These observations are only possible at near infrared wavelengths from about 1-5 $\\mu$m, where giant planets are self-luminous due to the gravitational energy converted to internal heat in their formation and subsequent slow contraction.  Because giant planets radiate this energy away over time, they become cooler and fainter as they age. The youngest nearby stars are therefore the most promising targets for AO surveys attempting to image self-luminous giant planets.  Each of the four debris-disk stars discovered using IRAS is relatively young, so orbiting giant planets might be detectable if any exist.  We have imaged the two stars most easily observable from Northern Hemisphere sites: Vega and $\\epsilon$ Eridani.  Vega's age is about 0.3 Gyr \\citep{agevega}, while the age of $\\epsilon$ Eri is about 0.56 Gyr \\citep{fischer}. Besides the dust-dispersion timescale argument mentioned above for the existence of planetary systems around these stars, asymmetries in the dust distributions around each have led to hypotheses that the dust is being gravitationally sculpted by giant planets orbiting at large distances \\citep{ozernoy,quillen,wyatt,wilner,deller,marsh}. In the case of Vega there are suggestions that the dust may reveal the mass and approximate position of a giant planet \\citep{wilner,deller}.  For $\\epsilon$ Eridani, in addition to evidence for a planet in a distant orbit that may be sculpting the dust \\citep{deller,benedict}, there is the the radial velocity and astrometric detection of the closer-in planet $\\epsilon$ Eri b \\citep{benedict}.  The combination of radial velocity and astrometry observations permits a full orbital solution yielding ephemerides for the separation and position angle of $\\epsilon$ Eri b \\citep{benedict}, making this the most promising case yet where attempts to image a known extrasolar planet can target a specific location.  Most imaging searches for extrasolar planets to date have used either the $H$ band (1.5 - 1.8 $\\mu$m) or other filters in the same wavelength regime (see for example \\citet{hsttwa7a, masciadri, sdi1, epsindia, biller, GDPS}). The magnitude vs. mass tables of \\citet{bar} and the theoretical spectra of \\citet{bur} show clearly why the $H$ band is usually chosen: giant planets are predicted to be very bright at these wavelengths, much brighter than black bodies at their effective temperatures.  Detector formats are large, technology well developed, and sky backgrounds faint at the $H$ band relative to longer wavelengths.  However, theoretical models indicate that planet/star flux ratios are much more favorable at the longer wavelength $L'$ and $M$ bands (3.4-4.1 $\\mu$m and 4.5-5.0 $\\mu$m, respectively). For planets at sufficiently large separations, or planets orbiting faint stars, the planet/star flux ratio is not relevant.  Rather, it is the planet's brightness relative to the sky background and/or detector read noise that matters.  In this regime the very high sky background in the $L'$ and $M$ bands prevents them from being as sensitive as the $H$ band regime.  However, close to very bright stars the background becomes irrelevant and only the planet/star flux ratio matters.  Under these circumstances using the longer wavelengths makes sense.  Vega is a magnitude 0.0 standard star and is among the brightest stars in the sky at almost any wavelength.  $\\epsilon$ Eri, while not impressive at visible wavelengths, is a very bright magnitude 1.9 at $H$ band.  The stars are therefore excellent targets for Clio, an $L'$ and $M$ band optimized AO camera that had its first light on the MMT in June 2005 \\citep{oldvega}.  We have made deep $\\sim$1 hour integrations in both the $L'$ and $M$ bands on both stars.  Our $M$-band observations are the deepest ground-based images yet published in this band.  In Section \\ref{obs} we present our observations and data analysis strategy. In Section \\ref{sensanal}, we describe our methods of analyzing our sensitivity, and present our sensitivity results.  Blind sensitivity tests in which simulated planet images were inserted directly into the raw data show that we have obtained 100\\% completeness for sources at 10$\\sigma$ signficance, 77\\% completeness for 7$\\sigma$ sources, and 41\\% completeness for 5$\\sigma$ sources, where $\\sigma$ is an estimate of the RMS noise amplitude in the image at the spatial scale of the PSF core (the relevant scale for detection of faint point sources).  We note that no other planet-imaging papers to date present such careful blind tests in their sensitivity analyses, and that the fact that our tests result in somewhat lower completeness values at each significance level than might have been expected suggests such tests should always be attempted and may result in a need to revise some sensitivity estimates to more conservative values.  In Section \\ref{vega} we compare the sensitivity we have obtained around Vega to that of other deep observations of Vega, and to the expected brightness of planets that have been hypothesized to explain the dust distribution.  In Section \\ref{eri}, we present the same comparisons for $\\epsilon$ Eri, and in Section \\ref{conclusion} we present the conclusions of our study. ",
        "conclusions": "\\label{conclusion} We have taken very deep $L'$ and $M$ band images of the interesting debris disk stars Vega and $\\epsilon$ Eri to search each system for orbiting planets and brown dwarfs.  For both stars we obtained better sensitivity than shorter-wavelength observations at small separations from the star.  The sensitivity of our observations compared more favorably to the sensitivity of $H$-regime observations in the case of $\\epsilon$ Eri than in the case of Vega.  For $\\epsilon$ Eri, our $M$ band observation appears to set the best upper limits yet for planets out to a separation of about 7 arcseconds, beyond which the sensitivity of the \\citet{GDPS} $H$-regime observations becomes very slightly superior.   The reason our $\\epsilon$ Eri observations have a greater sensitivity advantage over $H$ regime observations than do our images of Vega is the smaller distance to the $\\epsilon$ Eri system. This is another instance of the same physical reality we discussed in Section \\ref{finsens} above, when explaining why our $M$ band sensitivity is much better than our $L'$ results on $\\epsilon$ Eri but not on Vega.  As we stated above, for $\\epsilon$ Eri, the sensitivity of a given observation at any wavelength extends down to less luminous, lower $\\mathrm{T_{eff}}$ planets than for Vega.  The $H-L'$ and $H-M$ colors, as well as the $L' - M$ color, of low $\\mathrm{T_{eff}}$ giant planets are more red than those of hotter ones.  Therefore the faintest detectable objects in the $\\epsilon$ Eri system would be more red than those in the more distant Vega system, and longer wavelength observations are most useful for the nearer system.  This is a general and important principle for planning optimal planet search strategies: the faintest detectable planets will be more red, and therefore the relative advantage of long wavelengths over short ones will be higher, for the nearest stars.  For distant stars where only hot objects with blue infrared colors can be detected, long wavelengths observations are not as useful.  For very nearby stars such as $\\epsilon$ Eri, where very interesting, extremely low-mass, low $\\mathrm{T_{eff}}$ planets can be detected, the long wavelengths are very useful because the planets being sought have such red colors.  Planet-search observations at the $L'$ and $M$ bands have a considerable advantage over those in the more commonly used $H$ band regime for $\\epsilon$ Eri and a handful of other bright, very nearby stars.  For more distant bright stars such as Vega, $L'$ and $M$ band observations give markedly better results only at separations inside about 3 arcsec, and in this regime no currently employed method gives sensitivity to any but the highest mass planets.  Observations in the bands we have employed are still useful on Vega, but their use tends toward a diversification of planet-search effort in case theoretical models are overpredicting planets' $H$ band brightnesses.  For nearer systems such as $\\epsilon$ Eri, by contrast, $L'$ and $M$ band observations clearly provide the best sensitivity at the most interesting separations, and it is the $H$ regime images that naturally take the role of diversifying effort under the supposition that the models may overpredict planet brightness at longer wavelengths.  We have set a limit on the Vega planet hypothesized by \\citet{wilner} that is close to the 3 MJ mass he suggested for it.  It appears that \\citet{marois} could set a similar limit.  The evidence seems fairly strong that no 3 MJ planet exists at this location.  This favors alternative models involving smaller planets, such as those of \\citet{marsh} and \\citet{wyatt}, or a 3 MJ planet in a larger orbit, such as that of \\citet{deller}.  Since a 3 MJ planet around Vega could be imaged in multiple wavelength regimes with current technology, more submillimeter observations and further modeling to determine if such a planet is required to explain the observed dust distribution is very desirable.  If this does turn out to be the case, deep AO observations to detect the planet could be strongly prioritized, and a sucessful detection in one or more wavelength bands would be very likely.  Our $\\epsilon$ Eri observation was not timed to catch the known planet $\\epsilon$ Eri b at a large separation, and therefore our current data do not allow us to set an interesting limit on its mass.  \\citet{janson} observed $\\epsilon$ Eri at several epochs of more promising separation using SDI, and set limits in the 9-10 MJ range.  We have set a limit of 4-5 MJ for additional planets in more distant orbits around $\\epsilon$ Eri.  The existence of a planet in such an orbit may be indicated by a long term RV trend \\citep{benedict} and by a clearing of dust from the inner disk \\citep{deller}. \\citet{benedict} and \\citet{deller} suggest a mass of around 1 MJ for this hypothetical outer planet, so our non-detection is not surprising.  We have explored the question of whether SDI imaging \\citep{janson, biller} or $L'$ and $M$ band imaging (this work) is the method most likely ultimately to detect $\\epsilon$ Eri b. Our $M$ band images were much more sensitive at small separations than our $L'$ results, so we have not considered the latter.  We find that the sensitivity of the \\citet{janson} observations at the best epoch, where the planet was near the optimal separation for SDI imaging, was still insufficiently sensitive to detect the planet by a factor of at least a thousand.  By contrast our observations, if carried out at apastron, would have missed the planet by a factor of only about 12.  This striking difference suggests that it is at $M$ band that the planet $\\epsilon$ Eri b will first be imaged.  A several-night observing campaign using Clio at the MMT might detect it during the 2010 apastron passage, since we have observed the sensitivity in the speckle-dominated regions of $M$ band images does go up approximately as the square root of the exposure time.  More advanced PSF subtraction, or coronagraphic capability in Clio \\citep{phaseplate}, might reduce the required exposure time to detect the planet to as little as 2 nights. At present, we believe Clio with MMTAO is the only system capable of deep planet imaging integrations in the $M$ band.  Spitzer, despite its enormously lower background and correspondingly excellent sensitivity, does not have sufficient resolution to detect objects at the separations expected for orbiting planets.  $\\epsilon$ Eri b could be studied in more detail using new giant telescopes such as the LBT and GMT with planned adaptive secondary AO systems.  The $M$ band will remain the best wavelength choice for observations using these larger telescopes, so adaptive secondaries will remain essential: conventional AO systems even on giant telescopes will likely still have too high a thermal background for efficient, deep $M$ band images.  An $L'$ and $M$ band imager called LMIRCam is planned for the LBT \\citep{elmer}. When JWST is launched, it should also deliver interesting scientific results on $\\epsilon$ Eri b.  However, it is possible that the first image of this planet --- the first direct image of any mature extrasolar planet --- will be obtained using Clio at the MMT in 2010."
    },
    "0807/0807.2899_arXiv.txt": {
        "abstract": "{Water fountain stars are very young post-AGB stars with high velocity water maser jets. They are the best objects to study the onset of bipolar jets from evolved stars due to their young dynamical ages. However, none of them have been observed in any thermal lines.} {Search for CO lines in the water fountain star IRAS\\,16342-3814 and investigate the properties of its thermal gas.} {The proximity, peculiar stellar velocity and high Galactic latitude of IRAS\\,16342-3814 make a single dish observation possible. We use the Arizona Radio Observatory 10m telescope to observe the CO \\,J=2$-$1 line and compare the line parameters with that of masers.} {We report the detection of $^{12}$CO and $^{13}$CO\\,J=2$-$1 lines from IRAS\\,16342-3814. The inferred $^{12}$CO mass loss rate is an order of magnitude lower than the infrared and OH mass loss rates, indicating a very cold and thick O-rich circumstellar envelope around the star. We also find a $^{12}$CO expansion velocity of $V_{\\rm exp}=46\\pm 1$\\,km\\,s$^{-1}$ that is too high for an AGB wind and confirm the systemic velocity of $44\\pm1$ km/s. In addition we measure a very low $^{12}$CO/$^{13}$CO line ratio of 1.7.} {The first detection of CO lines has provided a new way to investigate the water fountain stars. Given the high expansion velocity of the CO gas and its relation to maser velocities, we infer that the CO emission region is co-located with the OH mainline masers in the warm base of the optical bipolar lobes, while the high velocity OH\\,1612\\,MHz and H$_2$O masers are located in the side walls and at the farthest ends of the bipolar lobes, respectively. Further observations are highly desired to understand the very low $^{12}$CO/$^{13}$CO line ratio.} ",
        "introduction": "Water fountain stars are a group of rare objects that display a thick circumstellar envelope (CSE) that is typical for asymptotic giant branch (AGB) stars and also high velocity bipolar jets ($\\gtrsim 100$\\,km\\,s$^{-1}$) detected via their H$_2$O masers. They could be stars near the end of AGB or post-AGB evolution stage when non-spherical circumstellar structures begin to develop inside the relic of the spherical AGB circumstellar envelope. Currently only about 10 such objects have been discovered through single dish and interferometry observations of their OH, H$_2$O and SiO masers~\\citep[see the review by][]{Imai07a}. However, none of them have ever been detected in a thermal molecular line. In this letter, we report the first detection of the $J=2-1$ transition of $^{12}$CO and $^{13}$CO in the water fountain star \\object{IRAS\\,16342-3814}.  \\object{IRAS\\,16342-3814} was discovered by \\citet{Likk88} to show H$_2$O masers with a line-of-sight expansion velocity up to $117.4$\\,km\\,s$^{-1}$, much higher than the typical speed of $\\sim 15$\\,km\\,s$^{-1}$ for a spherical AGB wind \\citep[see, e.g.,][]{Chen01}. They also found a progressive increase of expansion velocity in OH masers, from 36\\,km\\,s$^{-1}$ (1665\\,MHz) through 42\\,km\\,s$^{-1}$ (1667\\,MHz) to 58\\,km\\,s$^{-1}$ (1612\\,MHz). The symmetrical H$_2$O maser velocities found by \\citet{Likk92} allowed them to determine a star velocity of $V_{\\rm sys}=43.2\\pm 0.9$\\,km\\,s$^{-1}$, which further constrained a higher line-of-sight H$_2$O maser expansion velocity of $130$\\,km\\,s$^{-1}$.  Optical and infrared imaging to this star has revealed striking structures in its CSE. The Hubble Space Telescope (HST) observations by \\citet{Saha99} revealed a bipolar nebula. Corkscrew patterns were found in the bipolar lobes in follow-up infrared $L_{\\rm p}$ band imaging with the Keck telescope \\citep{Saha05}, indicating that the bipolar lobes and the H$_2$O masers might be created by precessing jets. Infrared imaging in the $3.8\\sim 20\\mu$ range by \\citet{Dijk03} showed the variation of the CSE morphologies from bipolar at near infrared to spherical at middle infrared, demonstrating that the outer CSE ejected earlier was spherical and the star is currently undergoing CSE morphology transformation. \\citet{Dijk03} also found from the $2\\sim 200\\mu$ Infrared Space Observatory (ISO) spectra that the star is the only object known to show absorption in both amorphous and crystalline silicate features. The very red spectral energy distribution possibly indicates a very cold and thick dusty CSE with an AGB mass loss rate on the order of $10^{-3}\\,M_{\\odot}{\\rm yr}^{-1}$ \\citep{Dijk03}.  Interferometric measurements of masers have supplied kinematic information. The mapping of OH masers in \\object{IRAS\\,16342-3814} by \\citet{Saha99} and \\citet{Zijl01} showed linear increase of line-of-sight outflow velocity with projected distance from the central star. High-velocity OH maser spots are located around the optical bipolar lobes while low-velocity maser spots appear near a dark waist close to the central star. These facts are in accord with a wind interaction model of the bipolar structure \\citep{Zijl01}. However, the corkscrew pattern in the infrared images also suggests that precessing jets may play some role. Very Long Baseline Array (VLBA) mapping of H$_2$O masers by \\citet{Morr03} and \\citet{Clau04} showed bipolar distribution of maser spot clusters with more than $2\\arcsec$ separation along the optical bipolar lobes and quasi-linear alignment of maser spots that delineates shock front of the jets. The interplay of the jets with magnetic field is still unclear, although extremely high degree of polarisation has been found in the OH masers \\citep{TeLi88,Szym04}.  The post-AGB nature of \\object{IRAS\\,16342-3814} has been critically investigated by \\citet{Saha99}. However the status of thermal gas in the CSE of this star is still unclear due to the lack of thermal molecular line observations. Unlike other known water fountain stars, \\object{IRAS\\,16342-3814} is the easiest object in which to search for its CO lines with a single dish telescope, because it is relatively near with a distance of $0.7\\sim2$\\,kpc \\citep{Saha99,Zijl01}, it has a high galactic latitude of $5.^\\circ8$ and a large line-of-sight velocity difference of at least 40\\,km\\,s$^{-1}$ from that of nearby interstellar molecular clouds. ",
        "conclusions": "\\subsection{The low $^{12}$CO mass loss rate} \\label{mlr}  The detection of the $^{12}$CO\\,2$-$1 line allows us to directly estimate the mass loss rate of \\object{IRAS\\,16342-3814}. Using the $^{12}$CO\\,2$-$1 formula scaled by \\citet{Hesk90} from a $^{12}$CO\\,1-0 formula postulated by \\citet{Knap85}, we derive $\\dot{M}(^{12}{\\rm CO}) = 0.7\\sim 5.6\\times 10^{-6}M_{\\odot}{\\rm yr}^{-1}$ for a source distance of $0.7\\sim 2$\\,kpc. As discussed in \\S~\\ref{vel}, the detected CO gas may not represent the spherical AGB wind. The true AGB wind CO mass loss rate, if also estimated with Heske's formula, should be even lower than this value, because the CO gas emission from the slow AGB wind is not definitely detected yet, while the mass loss rate is $\\propto V_{\\rm exp}^2 T_{\\rm mb}(CO)$. However, the formula of \\citet{Baud83} applied to an OH\\,1612\\,MHz maser peak flux of 9\\,Jy \\citep{Likk88} yields a much higher mass loss rate of $\\dot{M} {\\rm (OH)} = 1.7\\sim 4.9\\times 10^{-5}M_{\\odot}{\\rm yr}^{-1}$ (with $V_{\\rm exp}=46$\\,km\\,s$^{-1}$ from our $^{12}$CO line). \\citet{Do07} fitted the Spitzer MIPS $70\\mu$ images with a two shell model, and obtained $\\dot{M} (70\\mu) = 0.4\\sim 3\\times 10^{-5}M_{\\odot}{\\rm yr}^{-1}$ for an inner smooth CSE and $\\dot{M} ({\\rm shell}) = 0.4\\sim 3\\times 10^{-4}M_{\\odot}{\\rm yr}^{-1}$ for an outer extended dust shell (assuming a gas-to-dust mass ratio of 200 and $V_{\\rm exp}=15$\\,km\\,s$^{-1}$. Here we note that our measured $^{12}$CO expansion velocity of $46$\\,km\\,s$^{-1}$ may not be that of the dusty CSE, see the discussion in \\S~\\ref{vel}). The inner shell mass loss rate agrees with the OH mass loss rate. Due to its huge radius of $1.4\\sim4$\\,pc, the outer shell could be explained by sweep-up of interstellar material, as suggested by the recent hydrodynamic simulation of \\citet{Ware07}. To summarise, we suggest that \\object{IRAS\\,16342-3814}, while located in a very extended and cold detached outer dust shell, has an inner CSE whose mass loss rate (estimated from IR and OH emission) is roughly an order of magnitude higher than its $^{12}$CO mass loss rate. The weakness of $^{12}$CO\\,2$-$1 line could be interpreted by the coldness of the CO gas, as suggested for some other very red OH/IR stars by \\citet{Hesk90} .   \\subsection{The high $^{12}$CO expansion velocity} \\label{vel}  The $^{12}$CO expansion velocity $V_{\\rm exp}(^{12}{\\rm CO})=46$\\,km\\,s$^{-1}$ is too high for a typical OH/IR star \\citep[typically $\\sim 15\\pm 5$\\,km\\,s$^{-1}$ for long period OH/IR stars, see e.g.,][]{Chen01}. Another possible source of the broad $^{12}$CO emission is the swept-up bipolar bubble walls. According to a momentum driven wind-interaction model suggested by \\citet{Zijl01}, the radial motion velocity of different parts of the bubble wall will be proportional to distance to the central star. This has been confirmed by the linear OH maser velocity offset-radius (V-r) relationship found by \\citet{Saha99} and \\citet{Zijl01}. Assuming that $V_{\\rm exp}(^{12}{\\rm CO})$ follows the same V-r relationship in Fig.~3 of \\citet{Saha99}, we find a rough angular radius of $0.5\\arcsec$ for the CO emission region, similar to that of OH mainline masers. Further assuming a Gaussian source shape, a brightness temperature of the $^{12}$CO\\,2$-$1 line can be estimated to be $\\sim 30$\\,K. Therefore, the $^{12}$CO\\,2$-$1 line could be optically thin if the gas temperature is much higher than this.  Further insight can be gained by comparing the velocities of $^{12}$CO\\,2$-$1 line, OH and H$_2$O masers. We find several interesting facts by comparing the CO velocities with the V-r relationship of OH masers in the Fig.~3 of \\citet{Saha99}: 1) The 1612\\,MHz maser velocity offsets are all $\\gtrsim V_{\\rm exp}(^{12}{\\rm CO})$, while most of the OH\\,1665/1667\\,MHz (mainline) maser velocity offsets $\\lesssim V_{\\rm exp}(^{12}{\\rm CO})$; 2) The OH\\,1612\\,MHz masers show two components -- a pair of bipolar high velocity clumps that tightly follow the linear V-r relation and a peculiar group of points with $V_{\\rm LSR}\\sim 0$\\,km\\,s$^{-1}$ but scattered in a large radius range $(-1.1\\arcsec\\sim 0\\arcsec)$; 3) The position of the bipolar 1612\\,MHz maser component is not symmetrical with respect to the star position, perhaps due to 3D projection effect. The peculiar group of 1612\\,MHz masers are not only scattered in radius, but also scattered in position angle with respect to the central star in the VLA maps \\citep[see Fig.~1b and Fig.~9 of][respectively]{Saha99,Zijl01}. This is difficult to explain with a spherical expanding shell model in which masers should show an elliptical relation between line-of-sight velocity and projected radial distance. Taking away this peculiar 1612\\,MHz maser component, one can see a better V-r relationship among the rest 1612, 1665 and 1667\\,MHz masers. If we further take a velocity span of $\\sim 248$\\,km\\,s$^{-1}$ and a spatial separation of $\\sim 3\\arcsec$ of H$_2$O masers from \\citet{Morr03} and overlap their symmetry center with that of OH masers, we can see that the  H$_2$O masers follow the same linear V-r relationship quite well. Upon these arguments, we can sketch a picture for the CSE as follows:  Both the OH mainline masers and the CO emission arise from the same region near the base of the bipolar outflow, the two clumps of high velocity OH 1612 maser spots are located on the farther and faster part of the bipolar lobe side walls, the H$_2$O masers are stimulated by the shock fronts at the two farthest and fastest ends of the bipolar lobes, while the peculiar group of spatially scattered 1612\\,MHz maser spots are located in an independent structure that is blue-shifted by $\\sim -46$\\,km\\,s$^{-1}$ with respect to the star but not necessarily associated with the bipolar outflow. The dark waist and the slow AGB wind might be simply too cold to be detected in our CO\\,2$-$1 line spectra. Actually, the pair of faint spikes on the top of the broad $^{12}$CO\\,2$-$1 line profile near the line centre in Fig.~\\ref{f1} could be the emission from the slow AGB wind. If this is true, we may estimate a crude AGB wind speed of $\\sim 13$\\,km\\,s$^{-1}$, typical for an AGB star.   \\subsection{On the origin of the extremely low $^{12}$CO/$^{13}$CO ratio}  The detection of the $^{13}$CO~J=2$-$1 line yields an extremely small $^{12}$CO/$^{13}$CO peak flux ratio of 1.7. A range of $^{12}$CO/$^{13}$CO ratios have been found in PPNe and PNe by other researchers. \\citet{Hase05} concluded from a survey of published $^{13}$CO observations of PNe \\citep{Pall00,Bals02,Bach97,Joss03} that the $^{12}$CO/$^{13}$CO ratios are about $10\\sim 30$ in most PNe, but are $2\\sim 3$ in a few and $>60$ in a few others. Some examples of low ratios are: $^{12}{\\rm CO}/^{13}{\\rm CO}\\sim 4.6$ in \\object{AFGL\\,618}, $\\sim 2.2$ in \\object{M\\,1-16} \\citep{Bals02}, $\\sim 3$ in the PN \\object{NGC\\,6302} \\citep{Hase03}, and $\\sim 3.2$ in \\object{HD\\,17982} \\citep{Joss01}. \\citet{Bach97} pointed out that the $^{12}$CO/$^{13}$CO ratios in PPNe and PNe are affected by two competing effects: selective photodissociation of $^{13}$CO and isotopic fractionation of $^{13}$C in CO. Another issue is the opacity in the $^{12}$CO lines, e.g., the very low ratio in \\object{NGC\\,6302} was later found by interferometry work of \\citet{Pere07} to be affected by opacity effects in a massive torus and the true $^{12}$C/$^{13}$C should be $>15$. For an embedded dusty post-AGB star like \\object{IRAS\\,16342-3814}, we do not need to worry about the selective photodissociation in the central region. If the detected CO emission does come from the accelerated bipolar bubble walls (see the discussion in \\S~\\ref{vel}), neither isotopic fractionation nor $^{12}$CO line opacity should be an issue, since the dynamical age of the bipolar jets is short ($\\sim 100$\\,yr) and the high expansion velocity and high radial velocity gradient (according to wind interaction model) might render the $^{12}$CO emission not too optically thick. A rough but conservative calculation of the optical depth under an LVG approximation gives a value smaller than unity. However, better CO line data is required to confirm this. We also find that our $^{12}$CO/$^{13}$CO ratio is about 1.1 and 3.2 on the red and blue line wings ($46>|V_{\\rm exp}|>15$\\,km\\,s$^{-1}$) respectively, but the uncertainties are larger there. One more issue is that the $^{12}$CO and $^{13}$CO lines might come from different clumps of gas. Again, better CO line profiles and, better still, CO mapping data are needed to clarify this point and thus to constrain a reliable $^{12}$C/$^{13}$C ratio. Therefore, later in this section, we will discuss different opacity cases separately, using only the assumption that the $^{12}$CO and $^{13}$CO lines are co-located in the CSE.  A precise measurement of the $^{12}$C/$^{13}$C ratio has the potential to cast light on the evolutionary status of the star. Stellar models of low mass ($\\textrm{M}\\lesssim 2.0 \\textrm{M}_\\odot$) and intermediate mass ($2.0\\textrm{M}_\\odot\\lesssim\\textrm{M}\\lesssim 8.0 \\textrm{M}_\\odot$) stars \\citep[e.g.,][]{Elei94,Stra97,Kara07} show that there are quite a few phases of evolution where this ratio is altered from the initial value (assumed to be solar at 89). The first occurs during the first dredge up (FDU) at the beginning of the Red Giant Branch (RGB).  Here the ratio declines from the initial value to $\\sim 20$ as the convective envelope moves in and mixes up partially burnt material (we note that the second dredge-up has only a minor effect on the ratio). The second evolutionary phase that alters the $^{12}$C/$^{13}$C ratio is not predicted by stellar models. Observations have shown that the isotopic ratio reduces along the RGB in low-mass stars ($\\textrm{M} \\lesssim 2.0 \\textrm{M}_\\odot$), via a process known as Deep Mixing \\citep[also known as $\\delta_\\mu$ mixing - see the recent theoretical models by][]{Eggl08}. This phenomenon reduces the ratio down to a value of $\\sim 15$ by the end of the RGB in Population I stars. The third phase that significantly alters the ratio actually increases it -- the Third Dredge-Up (TDU). Here $^{12}$C is periodically mixed into the envelope during the thermally-pulsing AGB stage (TPAGB), raising the $^{12}$C/$^{13}$C ratio from the FDU value to values of $\\sim60$ or more \\citep[see e.g.,][for recent observations of this effect]{Lebz08}. The final phase to alter the ratio is Hot Bottom Burning (HBB). However this only occurs in stars more massive than $\\sim 4$ or $5 M_\\odot$ (at solar metallicity). HBB is active during the later phases of the TPAGB and acts to burn the carbon and nitrogen in the convective envelope via the CNO cycles. This results in the $^{12}$C/$^{13}$C ratio approaching the equilibrium value of $\\sim 3$ \\citep[see e.g.,][for some observational evidence of this phenomenon]{Plez93,Mcsa07}. The very low $^{12}$C/$^{13}$C ratio in some Galactic PNe, PPNs and extremely young post-AGB stars like \\object{IRAS\\,16342-3814} could be also indicative of the HBB process.  If \\object{IRAS\\,16342-3814} has a high $^{12}$C/$^{13}$C ratio close to the solar value of 89 (i.e., $\\sim50$ times larger than the $^{12}$CO/$^{13}$CO line ratio, e.g., as a 3\\,M$_\\odot$ star experiencing TDU), the CO line optical depths can be derived from the observed $^{12}$CO/$^{13}$CO ratio (=1.7) to be $\\tau(^{12}{\\rm CO})\\sim82$ and $\\tau(^{13}{\\rm CO})\\sim0.93$ (according to $T\\propto(1-e^{-\\tau})$ under the assumption of co-location of $^{12}$CO and $^{13}$CO gas and the same excitation temperature). However, the $^{12}$CO mass loss rate should also be raised by a factor of $\\tau/(1-e^{-\\tau})=82$ to $5.7\\sim45.9\\times10^{-5}{\\rm M}_\\odot$/yr, which is significantly higher than the OH and IR mass loss rates.  If the true $^{12}$C/$^{13}$C ratio is an intermediate value of 20, i.e., a factor of $\\sim10$ higher than the $^{12}$CO/$^{13}$CO line ratio, the similarly determined opacities are $\\tau(^{12}{\\rm CO})\\sim19$ and $\\tau(^{13}{\\rm CO})\\sim0.93$. In this case, the $^{12}$CO mass loss rate is comparable to the OH and IR mass loss rates. Therefore, a low mass AGB star ($\\lesssim1.5$\\,M$_\\odot$) that didn't experience TDU or HBB and thus retains its FDU plus Deep Mixing surface abundances would be a possible progenitor of \\object{IRAS\\,16342-3814}.  If the $^{12}$CO\\,2-1 line is only slightly optically thick so that the $^{12}$C/$^{13}$C ratio is $\\sim3$ (the CNO equilibrium value), the CO line opacities are $\\tau(^{12}{\\rm CO})\\sim2.3$ and $\\tau(^{13}{\\rm CO})\\sim0.77$. In this case, we can see that the only process (to the best of our knowledge) that could have lowered the $^{12}$C/$^{13}$C ratio to such a value is HBB. Thus the inferred mass of the star that produced the material that we observe today would have been $> 4$ or 5 M$_\\odot$. Interestingly the low $^{12}$CO/$^{13}$CO ratio, bipolar outflow, and the properties of the OH masers of \\object{IRAS\\,16342-3814} are reminiscent of the massive PPN \\object{IRAS\\,22036+5306} investigated by \\citet{Saha03,Saha06}, although \\object{IRAS\\,16342-3814} is probably in an earlier phase of evolution due to the presence of the young `water fountain'. Alternatively, a binary system in which a massive AGB star had transfered low $^{12}$C/$^{13}$C matter to the current AGB star is also possible. We note that a low isotopic ratio is supported by our rough estimation of the small $^{12}$CO\\,2-1 opacity at the beginning of this section.  However, both the measured CO and H$_2$O maser systemic velocities \\citep[$\\sim$44\\,km\\,\\,s$^{-1}$,][]{Likk92} confirm that the spatial motion of \\object{IRAS\\,16342-3814} is in the reverse direction of the Galactic disk rotation \\citep[interstellar CO cloud velocity near the same line of sight is $\\sim -17\\,{\\rm km\\,s}^{-1}$, ][]{Dame01}. The peculiar stellar motion lends weight to \\object{IRAS\\,16342-3814} being a low mass stars, although in rare cases kick-out of stars from the crowded galactic plane can give rise to peculiar velocities.  Given the preceeding discussions it is obvious that further observations are highly desirable to investigate the different opacity possibilities. Molecular line interferometry and maser proper motion measurements \\citep[such as the work of][]{Imai07b} are also needed to check the interstellar contamination, confirm the co-location of the $^{12}$CO and $^{13}$CO gas, and to locate the birthplace of the possibly escaped AGB progenitor on the Galactic plane."
    },
    "0807/0807.1548_arXiv.txt": {
        "abstract": "\\spider\\ is a balloon-borne experiment that will measure the polarization of the Cosmic Microwave Background over a large fraction of a sky at $\\sim1^\\circ$ resolution. Six monochromatic refracting millimeter-wave telescopes with large arrays of antenna-coupled transition-edge superconducting bolometers will provide system sensitivities of 4.2 and 3.1 $\\mukcmbrts$ at 100 and 150 GHz, respectively.  A rotating half-wave plate will modulate the polarization sensitivity of each telescope, controlling systematics.  Bolometer arrays operating at 225 GHz and 275 GHz will allow removal of polarized galactic foregrounds. In a 2-6 day first flight from Alice Springs, Australia in 2010, \\spider\\ will map 50\\% of the sky to a depth necessary to improve our knowledge of the reionization optical depth by a large factor. ",
        "introduction": "\\label{sec:intro}  Rapid progress in millimeter-wave receiver technology enabled major advances in cosmology over the past decade.   Extremely sensitive receivers measured the very faint anisotropies and polarization of the Cosmic Microwave Background (CMB), which are a mere part in $10^5$ and $10^6$ of the background, respectively.  CMB observations from the ground, stratospheric balloons, and space have confirmed and powerfully constrained the standard model of cosmology.  Two periods of the history of the universe still elude complete description: the period of inflation in the very early universe, and the period of first star formation that reionized the intergalactic medium in the more recent universe.  The theory of inflation postulates a period of very rapid expansion of the Universe ($\\sim 10^{-36}$ seconds after the Big Bang) that seeded the formation of structure.  The theory remains incomplete and current observations only provide general constraints\\cite{2008arXiv0803.0547K}.  Inflation also  predicts that the universe is filled with gravitational waves generated in the early universe.  The gravitational waves have extremely long wavelength, making them difficult to detect directly, but they leave a unique imprint in the polarization pattern in the CMB: the B-mode polarization.  Measuring the B-mode signal (a part in $10^7$ of the background) requires another leap forward in technology.  The CMB also provides a way to investigate the period of reionization of the universe.  Reionization enhances polarization and anisotropies on scales larger than the horizon size at reionization and damps them on smaller scales.  Measuring the scale of damping and the amount of tells us the optical depth of reionized universe through which the CMB photons have travelled.  Finding the signature of reionization is easiest in  polarization, as other cosmological parameters can create a similar signature in the temperature anisotropy.  The 5 year results from WMAP included a measurement of the total integrated optical depth, but tighter constraints could be found with a more sensitive instrument.  Here we describe \\spider, a balloon-borne instrument designed to measure the polarization of the CMB over a large fraction of the sky at large angular scales.  The optical design of \\spider, including a rotating  half-wave plate at the telescope aperture and on-axis monochromatic refracting telescopes, will provide the required control of systematics.  Large arrays of antenna-coupled transition-edge superconducting bolometers with multiplexed SQUID readouts provide the raw sensitivity for probing inflation and reionization.  The choice of frequency bands in \\spider\\ will minimize confusion from galactic foregrounds.  \\spider\\ complements \\planck\\  with its greater control of polarization systematics and closely spaced passbands for polarized foreground  characterization.  The instrument will make a 2-6 day first flight from Alice Springs, Australia in the austral autumn of 2010 that will probe the reionization of the universe with an unprecedented measurement of the large scale polarization of the CMB.  Subsequent flights will provide the longer integration time necessary for investigating inflation.  The balloon-borne platform is an ideal environment for these observations, allowing observation of nearly half the sky.  The \\spider\\ instrument was introduced in Montroy et al. (2006) \\cite{2006SPIE.6267E..24M}.    MacTavish et al. (2007) \\cite{2007arXiv0710.0375M} made a detailed investigation of systematics in \\spider\\ observations.  Here we discuss modifications in the instrument design and discuss the prospects for scientific results from our first flight. ",
        "conclusions": "\\label{sec:conclusions}  \\spider\\ is unique as a sub-orbital millimeter-wave polarimeter that can observe a large fraction of the sky. Large-format multiplexed TES bolometer arrays provide excellent raw sensitivity,  and the refracting optical design and a rotating half-wave plate will minimize systematics.  In its first flight, \\spider\\ will make a precise measurement of the optical depth of reionization.  Later flights will measure the B-mode signal from primordial gravitational waves.  \\spider\\ acts a pathfinder to the proposed CMBpol space mission more than any other planned B-mode experiment.   \\spider\\ will test a receiver with a large format array and multiplexed readout electronics in a high-altitude environment as well as observing strategies and systematics of deep polarimetery on a large fraction of the sky.  The design of \\spider emphases lightweighting, power savings, and autonomous operation all of which are directly applicable to a space mission."
    },
    "0807/0807.1062_arXiv.txt": {
        "abstract": "The VLTI Spectro Imager (VSI) was proposed as a second-generation instrument of the \\emph{Very Large Telescope Interferometer} providing the ESO community with spectrally-resolved, near-infrared images at angular resolutions down to 1.1 milliarcsecond and spectral resolutions up to $R=12000$. Targets as faint as $K=13$ will be imaged without requiring a brighter nearby reference object; fainter targets can be accessed if a suitable reference is available. The unique combination of high-dynamic-range imaging at high angular resolution and high spectral resolution enables a scientific program which serves a broad user community and at the same time provides the opportunity for breakthroughs in many areas of astrophysics.  The high level specifications of the instrument are derived from a detailed science case based on the capability to obtain, for the first time, milliarcsecond-resolution images of a wide range of targets including: probing the initial conditions for planet formation in the AU-scale environments of young stars; imaging convective cells and other phenomena on the surfaces of stars; mapping the chemical and physical environments of evolved stars, stellar remnants, and stellar winds; and disentangling the central regions of active galactic nuclei and supermassive black holes. VSI will provide these new capabilities using technologies which have been extensively tested in the past and VSI requires little in terms of new infrastructure on the VLTI. At the same time, VSI will be able to make maximum use of new infrastructure as it becomes available; for example, by combining 4, 6 and eventually 8 telescopes, enabling rapid imaging through the measurement of up to 28 visibilities in every wavelength channel within a few minutes. The current studies are focused on a 4-telescope version with an upgrade to a 6-telescope one. The instrument contains its own fringe tracker and tip-tilt control in order to reduce the constraints on the VLTI infrastructure and maximize the scientific return. ",
        "introduction": "\\label{sec:introduction}  At the beginning of the 21st century, infrared observations performed at the milli-arcsecond scale are essential for many astrophysical investigations either to compare the same physical phenomena at different wavelengths (like sources already observed with the VLBI or soon to be observed by ALMA) or to get finer details on observations carried out with the \\emph{Hubble Space Telescope} (HST) or 10-m class telescopes equipped with adaptive optics.  The astrophysical science cases at milli-arcsecond scales which range from planetary physics to extragalactic studies can only be studied using interferometric aperture synthesis imaging with several optical telescopes.  In this respect, the \\emph{Very Large Telescope} (VLT) observatory of the \\emph{European Southern Observatory} (ESO) is a unique site world-wide with $4\\times8$-m unit telescopes (UTs), $4\\times1.8$-m auxiliary telescopes (ATs) and all the required infrastructure, in particular delay lines (DLs), to combine up to 6 telescopes. The \\emph{VLT Interferometer} (VLTI) infrastructure can be directly compared to the \\emph{Plateau de Bure Interferometer} (PdBI) which combines $6\\times15$-m antenna over 500-m in the millimeter-wave domain. The quality of the foreseen images can be directly compared to the images provided by the PdBI. However, the angular resolution of the VLTI is a few hundred times higher due to the observation at shorter wavelengths.  The large apertures of the VLTI telescopes and the availability of fringe tracking allow sensitivity and spectral resolution to be added to the imaging capability of the VLTI.  In April 2005, at the ESO workshop on \\emph{``The power of optical/infrared interferometry: recent scientific results and second generation VLTI instrumentation''}, two independent teams have proposed two different concepts for an imaging near-infrared instrument for the VLTI: BOBCAT \\cite{2008poii.conf..407B} and VITRUV \\cite{2008poii.conf..357M}. In October 2005, the science cases of these instruments were approved by the ESO \\emph{Science and Technical Committee}. In January 2006, the two projects merged in order to propose the \\emph{VLTI spectro-imager} (VSI) as a response to the ESO call for phase A proposals for second generation VLTI instruments. The phase A study ended in September 2007 after an ESO board review. The VSI instrument has been recommended by the ESO \\emph{Science and Technical Committee} in October 2007 and ESO is currently proposing to proceed with the construction in 2009-2010 for an installation of the 4/6 telescope version at the VLT in 2015. In this volume, we present the result of the VSI phase A study. ",
        "conclusions": "The VLTI Spectro-Imager is an instrument whose science objectives is to image astrophysical sources at the milli-arcsecond scale with both high angular and high spectral resolutions in the near-infrared. The image fidelity requirements drive a 6-beam design which makes use of the full potential of the VLTI site (including the use of 4 or 6 telescopes, use of PRIMA,...). VSI will deliver spectro-images as final data product like in the ALMA scheme.  The design of the instrument is relatively simple and self-consistent with its own fringe tracker in addition to other service modules like tip-tilt, ADC,... The main science instrument is based on integrated optics technology which has been shown to be as performant as a solution bulk optics and has been sky validated. The fringe tracking is an integral part of the imaging strategy.  In conclusion, the design of VSI is the simplest one that maximizes science return for a large number of astrophysical domains interested to phenomena at the milli-arcsecond scale."
    },
    "0807/0807.0932.txt": {
        "abstract": " ",
        "introduction": "$\\Lambda$ cold dark matter ($\\Lambda$CDM) models have achieved wide acceptance as concordance models due to the results of observations over the last decade. In particular, cosmic microwave background (CMB)\\cite{Spergel:2003cb} and supernovae (SNe)\\cite{Riess:1998cb,Perlmutter:1998np,Knop:2003iy,Riess:2004nr} observations have played critical roles in this acceptance. The isotropy of our universe is strongly supported by the CMB observations. Thus, if we assume that our universe is homogeneous, results from SNe observation suggest that the volume expansion of our universe is accelerating. The accelerating expansion of a homogeneous and isotropic universe means the existence of exotic energy components, the so-called dark energy, within the framework of general relativity (GR). If we consider dark energy as a perfect fluid, it has negative pressure. One of the candidates for dark energy is the cosmological constant. However, there are several crucial problems regarding the existence of the cosmological constant or other dark energy candidates (see, for example, Ref.\\citen{Kolb:2007gb}). No one knows the origin of dark energy and there has not yet been a conclusive illustration of dark energy.  There are other possibilities that explain the CMB and SNe observations. observations. In the above logic, we have imposed two assumptions. One is validity of GR at all cosmological scales. The other is the homogeneity of our universe. Hence, in order to describe the observational results without dark energy, we need to discard GR or homogeneous models of the universe. In this paper, we attempt to describe the observational results without dark energy or any modification of gravitational theory, but with inhomogeneous universe models.  The basic idea is that we are in a large underdense region, i.e., a large void; we reject the Copernican principle, which states that we live at a typical position in the universe. Pioneering works include those Zehavi et al.\\cite{Zehavi:1998gz} in 1998 and Tomita\\cite{2000ApJ...529...38T,Tomita:2000jj,Tomita:2001gh} in 2000 and 2001. Zehavi et al. analyze early SNe data and suggested that such a large void might exist around us without dark energy. Tomita proposed the void universe model and discussed the possibility of explaining the observed magnitude-redshift relations of SNe \\cite{2000ApJ...529...38T,Tomita:2000jj,Tomita:2001gh}. There are several works in the same direction \\cite{Celerier:1999hp,Iguchi:2001sq,Bolejko:2005fp, Chung:2006xh,Enqvist:2006cg,Garfinkle:2006sb, Kasai:2007fn,Biswas:2007gi} (See also the reviews.\\cite{Celerier:2007jc,Enqvist:2007vb,Mattsson:2007tj}). In these works,  Lema\\^itre-Tolman-Bondi (LTB) solutions\\cite{Lemaitre:1933gd,Tolman:1934za,Bondi:1947av} are often employed. LTB solutions are exact solutions of the Einstein equations, which describe the dynamics of a spherically symmetric dust fluid and are useful for constructing universe models with a spherically symmetric void.  Recently, several authors have discussed the possibility of explaining the CMB observations\\cite{Alnes:2005rw,Alexander:2007xx,GarciaBellido:2008nz}. They used asymptotically homogeneous LTB models and reported that these models may be consistent with the observed CMB anisotropy. In order to obtain more precise predictions on whether we are located near the center of a large void, we need to study the evolution of nonspherical density perturbations in the LTB universes\\cite{Zibin:2008vj}. Other methods of observationally investigating the inhomogeneity of our universe have been proposed in several papers \\cite{Clarkson:2007pz,Caldwell:2007yu,Uzan:2008qp, Bolejko:2008cm,Clifton:2008hv}.  The aforementioned works are very important for strengthening the observational foundation of physical cosmology. At present, the Copernican principle is not based on sufficient observational facts. However, by virtue of the recently improved observational technologies, we might reach the stage at which we are able to investigate through observation whether our location in the universe is unusual or not . In order to test inhomogeneous universe models by observations, it is important to know the types of inhomogeneous universe models that can be used to explain current observations, and to reveal what predictions are given by these universe models. This is the subject of this paper.  The inverse problem using LTB universe models is a useful method for investigating the possibility of explaining observational results with inhomogeneities in the universe \\cite{Mustapha:1998jb,Iguchi:2001sq,Chung:2006xh,Vanderveld:2006rb}. LTB dust solutions contain three  arbitrary functions of the radial coordinate: the mass function $M$, the big-bang time $t_{\\rm B}$ and the curvature function $k$, and the inverse problem means determining $M$, $t_{\\rm B}$ and $k$ so that a given distance-redshift relation is realized on the past light-cone of an observer. % In order to specify the three arbitrary functions, we need three conditions. One of these conditions corresponds to the choice of the radial coordinate and thus has no physical meaning.  Hence, one more condition in addition to the distance-redshift relation is necessary.  Iguchi et al.\\cite{Iguchi:2001sq} showed that the distance-redshift relation in the standard $\\Lambda$CDM model can be reproduced in LTB dust universe models. When they solved the inverse problem, they imposed two kinds of the additional condition: one is a uniform big-bang time $t_{\\rm B}=0$ and the other is a vanishing curvature function $k=0$. It should be noted that their models failed to reproduce the distance-redshift relation for $z\\gtrsim 1.7$. The reason why they could not continue their computation beyond $z\\sim1.7$ is the existence of the ``critical point'' discussed by Vanderveld et.al\\cite{Vanderveld:2006rb}. They reported the difficulty in solving the inverse problem and concluded that it is unlikely that a solution to the inverse problem can be found in all the redshift domain. However, Tanimoto and Nambu recently solved the inverse problem in all the redshift domain with a nonuniform $t_{\\rm B}$\\cite{TNpriv}. In this paper, we show that it is possible to obtain a solution to the inverse problem with the condition of the uniform big-bang time $t_{\\rm B}$. % \\footnote{ Mustapha et al.\\cite{Mustapha:1998jb} proposed the following theorem :{\\it Subject to the conditions of Appendix B in Ref.\\citen{Mustapha:1998jb} for any given isotropic observation of apparent luminosity $l(z)$ and number count $n(z)$ with any given source evolution function $\\hat L(z)$ and total density over source number density $\\hat m(z)$, a set of LTB functions can be found to make the LTB observational relations fit the observations. } It is nontrivial whether the same statement holds for any given set of some quantities different from $l(z)$, $n(z)$, $\\hat{L}(z)$ and $\\hat{m}(z)$. In this paper, the big-bang time $t_{\\rm B}(r)$ is given as one of the quantities. Thus, the situation is different from the case in Ref.\\citen{Mustapha:1998jb}. } In terms of perturbation theory, the inhomogeneity of the big-bang time corresponds to the decaying mode. The condition of uniform big-bang time might guarantee the consistency of the present model with the inflationary scenario, since a universe that experiences inflation is almost homogeneous immediately after the inflationary period is over.  In this paper, we also discuss the effects of local clumpiness. In addition to the effect of the large-scale void structure, small-scale clumpiness may affect the observed distance-redshift relation. If there are clumpy objects such as galaxies in the foreground of light sources, data for the apparent luminosities of these sources are often contaminated with gravitational lensing or absorption. For this reason, supernova teams specifically search for objects with a minimum amount of intervening material in the foreground\\cite{Pain:2002wj,Mattsson:2007tj}. The observed sources might be biased by this selection, since the light of the observed sources might propagate in a lower-density region of our universe. The simplest way to take account of this effect is to introduce the so-called smoothness parameter $\\alpha$, which is the ratio of the smoothly distributed matter density except for clumps to the mean energy density $\\rho$ for all matter, assuming that the energy density on the paths of observed light rays is given by $\\alpha\\rho$ \\cite{1969ApJ...155...89K,1973ApJ...180L..31D,Kantowski1998}. Tomita has studied the effects of small-scale clumpiness in a local void model by introducing the smoothness parameter $\\alpha$. In this paper, we also investigate the corresponding effects by introducing $\\alpha$.   The time derivative of the cosmological redshift is a very important observable quantity in distinguishing LTB models from concordance $\\Lambda$CDM universe models. We study it and show that it gives a criterion for observationally deciding whether the universe is described by an LTB model with uniform big-bang time.  This paper is organized as follows. In \\S\\ref{sec2}, we briefly review LTB dust universes and list basic equations that we have to solve. The temporal variation of the cosmological redshift of comoving sources is discussed in \\S\\ref{sec3}. Then, we show details of our numerical method, particularly, focusing on singular points of the equations at the center and the critical point in \\S\\ref{sec4}. Numerical results are given in \\S\\ref{sec5}. Section\\ref{sec6} is devoted to a summary and discussion.  Throughout this paper, we use the unit of $c=G=1$, where $c$ and $G$ are the speed of light and the gravitational constant, respectively. %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% ",
        "conclusions": "\\label{sec6} %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%   In this paper, we have attempted to solve the inverse problem to construct an LTB universe model that has the same distance-redshift relation as that of the concordance $\\Lambda$CDM model with $(\\Omega_{\\rm M0},\\Omega_{\\Lambda 0})=(0.3, 0.7)$, and we obtained solutions by numerical integration. In the present study, assuming an inflationary period in the early universe, we have restricted ourselves to models with uniform big-bang time, and the resultant universe model has a very large void whose symmetry center is at the observer's position. We have also studied the effects of local clumpiness by introducing the smoothness parameter $\\alpha$, and have shown that the local clumpiness may reduce the amplitude of the large-scale inhomogeneity of the void. Our results imply that it is possible to construct an inhomogeneous but isotropic universe model with a distance-redshift relation that agrees quite well with the observational data of the distance-redshift relation.  Our LTB universe model is regarded as an unnatural model from the viewpoint of the Copernican principle, because the observer stands exactly at the center of the isotropic universe. The extent to which we can separate ourselves from the center and remain consistent with current observations is discussed in Refs.\\citen{2000ApJ...529...38T} and \\citen{2000ApJ...529...26T,Tomita:2000rf,Alnes:2006pf,Alnes:2006uk}. The strictest limit is given by Alnes and Amarzguioui in Ref.\\citen{Alnes:2006pf} as 15 Mpc from the observation of CMB. This value is much smaller than the cosmological scale, and thus we should remain at a special position if the void universe is real. However, no observational data has yet been reported that entirely excludes inhomogeneous universe models. Therefore, it is important to know the type of inhomogeneous universes that can explain current observations and to propose observational methods for testing inhomogeneous universe models.  In this paper, we have also studied the temporal variation of the distance-redshift relation in our LTB universe model whose distance-redshift relation is the same as the concordance $\\Lambda$CDM model with $(\\Omega_{\\rm M0},\\Omega_{\\Lambda 0})=(0.3, 0.7)$. The result implies that if we can observe the time derivative of the cosmological redshift with sufficient accuracy, we can distinguish our LTB model from the concordance $\\Lambda$CDM universe model. Innovations in observational technology might provide us with data on the time derivative of the cosmological redshift in the near future \\cite{Uzan:2008qp,Li:2008cja,2008MNRAS.386.1192L}.  Finally, we should note that if we ignore the inflationary paradigm, the functional freedom of the big-bang time $t_{\\rm B}(r)$ is returned, and we might be able to construct an LTB universe model with the same redshift-distance relation and, furthermore, the same $dz/dt_0$ as those of the concordance $\\Lambda$CDM model by choosing an appropriate big-bang time $t_{\\rm B}(r)$\\cite{Uzan:2008qp}. However, in this case, we may need an other mechanism to explain the results in CMB observations and other cosmological problems than the standard cosmology starting from the inflation. This will be the subject of a future work and will be discussed elsewhere."
    },
    "0807/0807.1124_arXiv.txt": {
        "abstract": "In this paper we derive a possible mass profile for the low surface brightness galaxy, Malin 1, based upon previously published space-based and ground-based photometric properties and kinematics. We use properties of the bulge, normal disk, outer extended disk and \\ion{H}{1} mass as inputs into mass profile models. We find that the dark matter halo model of Malin 1 is best described by a halo profile that has undergone adiabatic contraction, inconsistent with the findings for most disk galaxies to date, yet consistent with rotation curve studies of M31. More importantly, we find that Malin 1 is baryon dominated in its central regions out to a radius of $\\sim10$ kpc (in the bulge region). Low-surface brightness galaxies are often referred to as being dark matter dominated at all radii. If this is the case, then Malin 1 would seem to have characteristics similar to those of normal barred disk galaxies, as suggested by other recent work. We also find that Malin 1 also falls on the rotation curve shear versus spiral arm pitch angle relation for normal galaxies, although more LSB galaxies need to be studied to determine if this is typical. ",
        "introduction": "\\begin{deluxetable}{ll} \\tabletypesize{\\scriptsize} \\tablecaption{Properties of the bulge, inner and outer disk and dark matter halo of Malin 1} \\tablewidth{0pt} \\label{tab1} \\tablehead{ \\colhead{Property} & \\colhead{Measurement}\\\\ } \\startdata Bulge effective radius$^1$, $R_e$                      & 0.6 kpc \\\\ Bulge effective surface brightness$^1$, $\\mu_e$        & 16.8 mag arcsec$^{-2}$\\\\ Inner disk scalelength$^1$, $h_{\\rm in}$                   & 4.8 kpc \\\\ Inner disk central surface brightness$^1$, $I_{\\rm in_0}$  & 20.1 mag arcsec$^{-2}$\\\\ Outer disk scalelength$^2$, $h_{\\rm out}$                  & 53 kpc\\\\ Outer disk central surface brightness$^2$, $I_{\\rm out_0}$\\hspace*{1.5cm} & 24.8 mag arcsec$^{-2}$\\\\ \\ion{H}{1} mass$^3$, $M_{\\rm HI}$                          & $(6.8\\pm0.7)\\times10^{10} M_{\\odot}$\\\\ Spiral arm pitch angle, $P$                            & $25\\fdg0\\pm1\\fdg0$\\\\ Halo concentration, $c_{\\rm vir}$                      & 8\\\\ Halo virial mass, $M_{\\rm vir}$                        & $2.6\\times10^{12} M_{\\odot}$\\\\ \\enddata \\tablenotetext{1}{From Barth (2007) who used an HST/WFPC F814W ($I$-band) image of Malin 1 to perform a two-dimensional structural decomposition into bulge and disk components.} \\tablenotetext{2}{From Moore \\& Parker (2007) who determine properties of the outer disk from a deep ground-based $R$-band image of Malin 1.} \\tablenotetext{3}{From Pickering et al.\\ (1997).} \\end{deluxetable}              Malin 1 is a highly unusual disk galaxy characterized by an enormous \\ion{H}{1} rich and extremely low surface brightness disk (Bothun et al.\\ 1987; Pickering et al.\\ 1997). It has the largest radial extent of any known spiral galaxy, with low surface brightness emission extending out to $\\sim100$ kpc, and its disk was found to have an extrapolated central surface brighness of only $\\mu_0\\simeq25.5$ mag arcsec$^{-2}$ in the $V$-band (Bothun et al.\\ 1987; Impey \\& Bothun 1989), with an exponential disk scalelength of $\\sim50-70$ kpc (e.g., Moore \\& Parker 2007). Although it has a very low surface brightness, its optical luminosity is $M_v\\simeq-22.9$ mag (Pickering et al.\\ 1997), due to its large extent. It also has an extremely high gas mass, with an estimated \\ion{H}{1} mass of $\\sim7\\times10^{10}M_{\\odot}$ (Pickering et al.\\ 1997; Matthews et al.\\ 2001). As a result Malin 1 is often considered a Low Surface Brightness (LSB) galaxy. However, recent studies of Malin 1 have started to highlight features more typical of normal disk galaxies (e.g., Barth 2007). The analysis of a {\\em Hubble Space Telescope} (HST) WFPC2 F814W ($I$-band) image presented by Barth (2007) shows that Malin 1 possesses an inner normal stellar disk, with characteristics similar to those in regular disk galaxies. They calculate an exponential disk scalelength of $\\sim 5$ kpc and a disk central surface brightness of $\\sim20$ mag arcsec$^{-2}$. These data suggest that Malin 1 has characteristics similar to those of normal disk galaxies, in particular barred lenticular galaxies (SBOs), which typically show an outer disk with a larger disk scalelength (Aguirre et al.\\ 2005). Moore \\& Parker (2007) have also recently presented a deep ground-based image of Malin 1, which shows spiral structure in its inner disk, another hint that Malin 1 may be more closely related to normal disk galaxies than originally thought. Indeed, Malin 1 may have much in common with the recently discovered class of objects that host extended ultraviolet (XUV) disks, such as M83 (Thilker et al.\\ 2005) and NGC 4625 (Gil de Paz et al.\\ 2005). All of these objects have apparently normal disks, but are surrounded by very extended low surface brightness emission (sometimes missed entirely in the optical) that shows up in the UV as a result of recent star formation.  In this paper we make use of a recently published \\ion{H}{1} rotation curve (Sancisi \\& Fraternali 2007) to determine a possible mass profile for Malin 1. We use the bulge/disk decomposition from Barth (2007) and the properties of the outer disk from Moore \\& Parker (2007) in our model. We also take into account the \\ion{H}{1} mass from Pickering et al.\\ (1997). The mass models produced show that Malin 1 is baryon dominated out to a radius of $\\sim$15 kpc. As LSB galaxies are typically dark matter dominated down to small radii (e.g., de Blok \\& McGaugh 1997; Kuzio de Naray et al.\\ 2008) on the surface it would appear that LSB may not be typical of LSB galaxies. However, it should be noted that the studies presented by de Blok \\& McGaugh (1997) and Kuzio de Naray et al.\\ (2006, 2008) consisted of dwarf LSB galaxies, which seem to be dark matter dominated beyond the inner $\\sim$1 kpc and may be baryon dominated within this radius. By extrapolation to much large giant LSB galaxies, such as Malin 1, it is not implausible that these objects would also be baryon dominated out to 10$-$15 kpc. We also find that the rotation curve shear and spiral arm structure of Malin 1 show that it sits nicely on the spiral pitch angle versus shear relation for normal disk galaxies reported by Seigar et al.\\ (2005, 2006). However, we also note that more LSB galaxies need to be studied to determine if they typically fall on the same relation. ",
        "conclusions": "We conclude that Malin 1 is not as atypical as originally thought. We highlight the fact that its $B$ band disk central surface brightness of $\\mu_B(0)=22.3$ mag arcsec$^{-2}$ as determined by Barth (2007) seems to place it in the category of intermediate brightness galaxies (McGaugh 1997). Taken together with our result here, that Malin 1 appears to be baryon dominated to $\\sim$10 kpc, this may suggest that Malin 1 has characteristics typical of normal galaxies. However, it still remains a very unusual galaxy, as it is also embedded in a very extended, gas-rich, outer disk. While Barth (2007) compared Malin 1 to SBO galaxies, the discovery of spiral structure in its disk (Moore \\& Parker 2007) would suggest that Malin 1 may very well be of later-type than this. The break in the outer disk of Malin 1 to that of a disk with a larger scalelength is not unusual for disk galaxies (e.g., Pohlen et al.\\ 2002; Erwin et al.\\ 2005, 2007). Malin 1 may just exhibit an extreme case of this phenomenon.  The spiral structure and rotation curve shear of Malin 1 are both consistent with those of normal disk galaxies, and they both fall nicely on the rotation curve shear versus spiral arm pitch angle relation reported by Seigar et al.\\ (2005, 2006). It is possible that a comparison of shear values and pitch angles for LSB galaxies reveal that they follow the same relation as normal galaxies. For this reason, in the future, we intend to make these measurement for a large sample of LSB galaxies."
    },
    "0807/0807.0534_arXiv.txt": {
        "abstract": "We have obtained the most extensive and most accurate photometric data of a Blazhko variable MW Lyr during the 2006-2007 observing seasons. The data within each 0.05 phase bin of the modulation period ($P_m=f_m^{-1}$) cover the entire light cycle of the primary pulsation period ($P_0=f_0^{-1}$), making possible a very rigorous and complete analysis. The modulation period is found to be 16.5462 d, which is about half of that was reported earlier from visual observations. Previously unknown features of the modulation have been detected. Besides the main modulation frequency $f_\\mathrm{m}$,  sidelobe modulation frequencies around the pulsation frequency and its harmonics appear at $\\pm2f_\\mathrm{m}$, $\\pm4f_\\mathrm{m}$, and $\\pm12.5f_\\mathrm{m}$ separations as well. Residual signals in the prewhitened light curve larger than the observational noise appear at the minimum-rising branch-maximum phase of the pulsation, which most probably arise from some stochastic/chaotic behaviour of the pulsation/modulation. The Fourier parameters of the mean light curve differ significantly from the averages of the Fourier parameters of the observed light curves in the different phases of the Blazhko cycle. Consequently, the mean light curve of MW Lyrae never matches its actual light variation. The $\\Phi_{21}$, $\\Phi_{31}$ phase differences in different phases of the modulation show unexpected stability during the Blazhko cycle. A new phenomenological description of the light curve variation is defined that separates the amplitude and phase (period) modulations utilising the phase coherency of the lower order Fourier phases. ",
        "introduction": "The Blazhko modulation of RR Lyrae stars, a phenomenon known for about a century, is still one of the open questions in astrophysics. The problems and the inconsistencies of the existing models with observational facts were recently discussed in details by \\cite{stothers}, therefore we only briefly mention here that none of the suggested models (magnetic oblique rotator, \\citet{shiba}; resonant excitation of nonradial modes, \\citet{dz}) can explain the complexity of all the observed properties of the modulation.  Alternatively,  \\cite{stothers} suggests that the modulation may be explained by continuous amplitude and period changes of the pulsation due to the action of a turbulent convective dynamo in the lower envelope of the star, with dynamo cycle identical with the Blazhko period. This explanation does not involve any nonradial mode component, the phenomenon is interpreted in the framework that Blazhko RRab stars are purely fundamental mode radial pulsators. However, \\cite{stothers} gives only a qualitative, rough picture of his model, that have to be checked both observationally and theoretically, thus the investigation of the phenomenon still remains  an important and valid task.  A lot of efforts have been already made in studying the light curve changes of RR Lyrae stars but most of the available photometries of Blazhko variables have some defects: inaccuracy (visual, photographic data), biased data sampling (the photoelectric observations focused mostly on the rising branch, maximum phase of the light curve), data inhomogeneity due to rare data sampling on a too long timebase (the time scale of period changes and changes in the modulation properties can be as short as a few years), etc. Most of the recent CCD and/or photoelectric observations of individual Blazhko stars are not extended enough to study the modulation properties in full detail. The first multicolour photometric data that covered each phase of both the modulation and the pulsation within a short time-base (one season) were published for RR Gem \\citep{rrgI}. During the period of the CCD observations RR Gem showed, however, small amplitude modulation, that limits the available information due to the low S/N ratios of the modulation signals.  We started a systematic search for previously unknown Blazhko variables at Konkoly Observatory in 2004 using the advantage of  full access to an automated 60 cm telescope \\citep{sodor}. The aim has also been to clear up questionable cases. In \\citet{ibvs} we have revised the list of known Blazhko variables \\citep{smith}. The modulation of MW~Lyrae was found to be ambiguous as the photographic observations \\citep{gessner} did not confirm the modulation detected in visual data by \\cite{mandel}. Our CCD observations of MW Lyrae affirm the light curve modulation of the star, but with a modulation period about half of that \\cite{mandel} announced.  Spectroscopic observation of MW Lyr has never been obtained, however, based on its short period ($P_{\\mathrm puls}=0.397$~days) it should be a relatively hot and metal rich RR Lyrae star.  In the 2006 and 2007 seasons we obtained CCD photometric observations of MW~Lyrae \\hbox{($\\alpha=18^h 19^m 53\\fs8, \\delta=+31\\degr 58\\arcmin 54\\arcsec$, J2000)}. This is the first multicolour photometric data set of a large modulation amplitude Blazhko variable that is  condensed, extended, and accurate enough to detect previously unknown properties of the modulation.  In this  paper  the photometric data of MW Lyr are published and analysed. The light curve solution (frequency analysis) is discussed using mostly the $V$ band data. A second paper is going to be devoted to the study of the colour behaviour during the course of the modulation cycle.  \\begin{table*} \\caption{Data of comparison and check stars} \\label{comps} \\begin{tabular}{llclcccc} \\hline Name &GSC2.2 ID &USNO-B1.0 ID   & \\multicolumn{1}{c}{  coord } & $V$[mag]$^*$&$B-V$[mag]$^*$& $V-R_{\\rm C}$[mag]$^*$&$V-I_{\\rm C}$[mag]$^*$\\\\ &&&      RA \\, \\, \\, \\, \\, \\, \\, \\, \\, Dec     & &&&\\\\ \\hline C0&N0223233663 &1219-0328347 &18 19 56.80 \\,\\, +31 58 13.3&$11\\fm841$&$0\\fm582$&$0\\fm315$&$0\\fm666$\\\\ C1&N0223233659 &1219-0328429 &18 20 06.56 \\,\\, +31 58 56.4&$12\\fm986$&$0\\fm682$&$0\\fm383$&$0\\fm788$\\\\ C2&N0223233671 &1219-0328258 &18 19 48.61 \\,\\, +31 57 14.6&$13\\fm921$&$0\\fm858$&$0\\fm480$&$0\\fm950$\\\\ C3&N02232332613&1219-0328267 &18 19 48.91 \\,\\, +31 56 54.4&$14\\fm363$&$0\\fm603$&$0\\fm338$&$0\\fm671$\\\\ \\hline \\multicolumn{8}{l}{\\scriptsize $^*$ Standard magnitudes measured by A. Henden}\\\\  \\end{tabular} \\end{table*}  \\begin{figure} \\includegraphics[width=7.5cm]{Figure01.eps} \\center \\caption{Map of MW~Lyr and the comparison (C0) and check stars (C1,C2, and C3). } \\label{chart} \\end{figure} ",
        "conclusions": "The photometric observations of MW~Lyr analyzed in this paper comprise the most extended and accurate data set of a Blazhko variable ever obtained. Utilizing this unique opportunity a detailed and circumspect phenomenological description of the modulation is given, which may provide  crucial information  to find the correct explanation of the phenomenon.  The main results and conclusions of the analysis are the followings:  \\begin{itemize}  \\item{In the Fourier spectrum of the light curve besides the $kf_0\\pm f_\\mathrm{m}$ triplet frequencies $kf_0\\pm 2f_\\mathrm{m}$ quintuplet, and  $kf_0\\pm 4f_\\mathrm{m}$ septuplet components also appear.}  \\item{Both $f_\\mathrm{m}$ and $2f_\\mathrm{m}$ frequencies can be detected in the spectrum.}  \\item{Frequency components at $kf_0\\pm 12.5f_\\mathrm{m}$ are  detected. If these frequencies are not `just by chance' at $\\pm 12.5f_\\mathrm{m}$ separations but somehow they are indeed connected to the main modulation frequency, then it is a great challenge to find an answer to their origin. }  \\item{Modulation with $12.5f_\\mathrm{m}$ frequency can also be detected in the maximum brightness data but not in the maximum phase observations. Consequently, the modulation connected to the $12.5f_\\mathrm{m}$ frequency is dominantly amplitude modulation. }  \\item{\\cite{stothers} mentioned the lack of large amplitude modulation around minimum phase of the pulsation as a failure of his model. The amplitude of the modulation around minimum phase of the pulsation in MW Lyr is, however, commensurable with the amplitude of the modulation in maximum brightness. On the contrary, the modulation of MW Lyrae is more strictly periodic and regular than it would be expected to be if the triggering mechanism behind the modulation were the cyclic weakening and strengthening of the turbulent convection in the ionization zones as \\cite{stothers} proposes. }  \\item{Though the modulation shows high degree of regularity both in the phase (period) and in the amplitude changes, the light curve cannot be fitted with the required accuracy even with 66 identified and further 30 frequencies appearing in the residual spectrum. Significant deviations in the residual light curve are concentrated at the minimum-rising branch-maximum phase of the pulsation. These residuals, however, do not show any periodicity, most probably they can be explained with some stochastic and/or chaotic behaviour of the modulation itself.}  \\item{The mean pulsation light curve defined by the Fourier parameters of the $kf_0$ pulsation frequency components of the full light curve solution differs significantly from the light curve in any phase of the modulation. Especially the $\\Phi_{k1}$ phase differences of the mean light curve are discrepant, they are out of the range of the phase difference values measured in any phase of the Blazhko cycle.}  \\item{The light curves in the small and large amplitude phases of the modulation can be fitted with 6-10 and 11-15 order harmonic fits, respectively. If nonlinear effects (e.g., shock waves) account for the occurrence of the higher order harmonic components of the pulsation, then their diminishing amplitudes in the low amplitude phase of the modulation may indicate that these nonlinear effects are not so important in this phase of the modulation.}  \\item{\\cite{stothers} proposed that the enhanced convection lowers the pulsation amplitude in the small amplitude phase of the modulation, while the phase relation betwen the period changes and the amplitude variations depends on the physical parameters of the individual variable in a complex way. The pulsation period of MW Lyr is the longest about 1-2 days (0.1 phases of the modulation cycle) later than  the minimal amplitude phase, while it is the shortest about 3-4 days (0.2 phases of the modulation cycle) later than  the maximal amplitude phase.}  \\item{The small variations in the phase differences of the lower order harmonic components of the light curves during the Blazhko cycle indicate that the phase variations of these frequency components are very coherent. If the modulation were caused by the interaction of close radial mode and nonradial mode frequencies \\citep{dz}, then no such phase coherency would be expected.}  \\item{The modulation of the light curve of MW~Lyr can be separated into amplitude and phase modulation components using only one parameter, the phase of the $f_0$ pulsation frequency in each phase bins of the modulation. This is in agreement with the explanation of the Blazhko effect with period and amplitude changes as recently proposed by \\citet{stothers}.}  \\end{itemize}  It would also be important to know how common the detected properties of the light curve modulation of MW~Lyrae are. Only further similar observations of other Blazhko stars can give an answer to this question."
    },
    "0807/0807.0228_arXiv.txt": {
        "abstract": "We present high resolution spectroscopic VLT observations of the outflow seen in QSO~2359--1241. These data contain absorption troughs from five resonance \\feii\\ lines with a resolution of $\\sim$7 \\kms\\ and signal-to-noise ratio per resolution element of order 100.  We use this unprecedented high quality data set to investigate the physical distribution of the material in front of the source, and by that determine the column densities of the absorbed troughs.  We find that the apparent optical depth model gives a very poor fit to the data and greatly underestimates the column density measurements.  Power-law distributions and partial covering models give much better fits with some advantage to power-law models, while both models yield similar column density estimates.  The better fit of the power-law model solves a long standing problem plaguing the partial covering model when applied to large distance scale outflow: How to obtain a velocity dependent covering factor for an outflow situated at distances thousands of time greater than the size of the AGN emission source.  This problem does not affect power-law models.  Therefore, based on the better fit and plausibility of the physical model, we conclude that in QSO~2359--1241, the outflow covers the full extent of the emission source but in a non-homogeneous way. ",
        "introduction": "In recent years, the potential impact of quasar outflows on their environment has become widely recognized.  (e.g., Blandford \\& Begelman 2004; Scannapieco \\& Oh 2004; Vernaleo \\& Reynolds 2006). Observationally, these outflows are detected as absorption troughs in quasar spectra that are blueshifted with respect to the systemic redshift of their emission line counterparts.  The absorption troughs are mainly associated with UV resonance lines of various ionic species (e.g., \\mgii~$\\lambda\\lambda$2796.35,2803.53 \\civ~$\\lambda\\lambda$1548.20,1550.77, \\siiv~$\\lambda\\lambda$1393.75,1402.77 \\nv~$\\lambda\\lambda$1238.82,1242.80).  A large spread in maximum  velocity is detected for different quasar outflows: from several hundred \\kms\\ to more than 30,000 \\kms.  Outflow troughs are known as intrinsic absorbers to distinguish them from intervening and associated absorber, which are also detected in quasar spectra (Hamann et al.\\ 1997; and Barlow 1997).  The lines we study here are intrinsic narrow lines (see Arav et~al.\\ 2002 for full discussion) and are similar to the ones analyzed by de~Kool et~al.\\ (2001).  Reliable measurements of the absorption column densities in the troughs are crucial for determining almost every physical aspect of the outflows: the ionization equilibrium and abundances; number density; distance; and mass flux and kinetic luminosity.  Column density measurements depend on the absorption model one adopts for the observed troughs in the spectrum.  In the case of Interstellar Medium (ISM) and Intergalactic Medium (IGM) absorbers (the latter also known as intervening absorbers), the apparent optical depth method is an excellent absorption model.  This method postulates that the absorber covers the full size of the emission source homogeneously and that the relationship between the the residual intensity in the trough ($I$) and the optical depth of the line ($\\tau$) is give by $\\tau_{ap}\\equiv-\\ln(I)$. The derived $\\tau$ is then converted to column density using a standard formula (e.g., equation 1 in Arav et~al.\\ 2001a).  However, our group (Arav 1997; Arav et~al.\\ 1999a,  1999b; de~Kool et~al.\\ 2001; Arav et~al.\\ 2001b, 2002, 2003; Scott et al.\\ (2004) Gabel et al.\\ (2005a)) and others (Barlow 1997, Telfer et~al.\\ 1998, Churchill et~al.\\ 1999, Ganguly et~al.\\ 1999) have shown that the apparent optical depth method, is not a good approximation for outflow troughs (see also \\S~3).  Most of this evidence came from data sets that showed fully resolved troughs from unblended doublets (e.g., the \\nv, \\siiv, \\civ\\ and \\mgii\\ lines mentioned above).  For all these doublets, the oscillator strength of the blue (the shorter wavelength) transition is  twice that of the red transition.  Therefore, the expected optical depth ratio is 2:1 in favor of the blue transition. (It is not exactly 2:1 since the wavelengths differ by a small amount.)  For the apparent optical depth method to hold, we must have $\\tau_{ap}$(blue)=$2\\tau_{ap}$(red) which requires: $I_B=I_R^2$, where $I_B$ and $I_R$ are the residual intensities of the blue and red absorption troughs, respectively. Using high signal-to-noise (S/N) observations of fully resolved and unblended outflow doublets, it became clear that in most cases $I_B$ is significantly smaller than $I_R^2$, demonstrating that the $\\tau_{ap}$ model is not applicable for AGN outflows.  To avoid this contradiction for cases of unblended doublets, most works use a pure partial covering model to determine the real optical depth and hence the actual column density.  This model assumes that only a fraction $C$ of the emission source is covered by the absorber and then solves for a combination of $C$ and $\\tau$ that will fit the data of both doublet troughs while maintaining the intrinsic 1:2 optical depth ratio (see \\S~3 of Arav et~al. 2005 for the full formalism, including velocity dependence). The weakness in this method is that we solve for two unknowns ($C$ and $\\tau$) given two residual intensity equations.  As long as the ratio of the residual intensities is in the permitted physical range (see \\S~3 of Arav et~al. 2005), such a procedure will always yield a solution. However, this is also the case for other two-parameter fitting algorithms, which yield different estimates of column densities.  A main alternative to the pure partial covering model is the so called inhomogeneous absorber model (de Kool, Korista \\& Arav 2002, hereafter dKKA; see \\S~3.2.2 here).  This model successfully matched the data of two quasar outflows (dKKA), while proving less applicable for the outflows troughs of Mrk~279 (Arav et~al.\\ 2005; but see the Discussion here).  Being able to distinguish between these two (and other) absorption models is important not only for column density measurements, but also for obtaining a more accurate geometrical picture of the outflow, including its internal structure.  Since doublet observations are inherently inconclusive in distinguishing between these methods (or any other two parameter models), to resolve this issue we need data sets that show outflow troughs from more than two lines of the same ion. These troughs must be unblended, fully resolved and have a high signal-to-noise ratio (S/N).  In this paper we present unprecedented high quality observations of a quasar outflow, which cover troughs from five \\feii\\ resonance lines associated with \\feii\\ UV multiplets 1, 2 and 3 (see \\S~2.3). We use the strong constraints available from these troughs to test models of absorption material distribution, and to determine the ionic column densities in all the troughs seen in this outflow.  The object we targeted is the extensively studied broad absorption line (BAL) quasar: QSO 2359--1241 (Arav et~al.\\ 2001a; Brotherton et~al.\\ 2001). Having detailed knowledge of the existing Keck/HIRES data, allowed us to tailor the new observations towards a definitive test for models of absorption material distribution. In future papers we will use these measured ionic column densities to determine the ionization equilibrium of the outflow, it's number density, distance and ultimately the mass flux and kinetic luminosity associated with the outflow.  The plan of the paper is as follows: In \\S~2 we describe the observations, data reduction and the phenomenology of the outflow; in \\S~3 we test the goodness of the fit for the five \\feii\\ resonance lines obtained from different models of absorption material distribution; in \\S~4 we discuss our results, and in \\S~5 we summarize them. ",
        "conclusions": "Understanding the nature of quasar outflows and their influence on the nuclear and galactic environments, begins with deciphering of their geometry on both small and large scales. Since all of the observational information is embedded within the absorption troughs, the first step is to determine the distribution of absorbing material that creates the troughs. Advances in our understanding of this issue will shed light on the microphysics of the outflows, and more importantly will yield reliable measurements of the ionic column densities. These measurements in turn are crucial to determining almost every physical aspect of the outflows: the ionization equilibrium and abundances, gas density, distance, mass flux and kinetic luminosity.  Until the late 1990s, all the work published on quasar outflows (e.g., Korista et al.\\ 1992; Turnshek et al.\\ 1996; Hamann 1998) used the apparent optical depth ($\\tau_{ap}$) model  to extract column densities.  This method works very well for ISM and IGM absorbers and corresponds to a simple gemetrical interpretation for the absorbing medium: a homogeneous screen in front of the light source.  The well demonstrated failure of the $\\tau_{ap}$ model when applied to AGN outflows troughs (see sections 3 and 1) motivate us to look more carefully at the geometry involved.  Relatively secure is the following picture:  For quasars, the size of the continuum source is  $\\sim10^{16}$ cm and the size of the  broad emission line region is $\\sim10^{18}$ cm (Kaspi et al.\\ 2005). Along our line of sight, there is material flowing towards us with radial velocities ranging from a few hundred to more than 30,000 \\kms.  This material is, cosmologically speaking, in the vicinity of the quasar, within 0.1-10,000 pc.\\ from the central source: 0.1 pc.\\ in BALQSO 1603+3002 (Arav et al 1999b); 25 pc.\\ for NGC 3783, (Gabel et al.\\ 2005b); 1000 pc.\\ for BALQSO 1044+3656 (de Kool et al.\\ 2001); 28 kpc in quasar 3C 191 (Hamann et al.\\ 2001). The question we try to address here is: what is the distribution of the absorbing  material across the quasar emission region  Most of the work that tried to advance beyond the $\\tau_{ap}$ model has used the pure covering factor method (see \\S~1 and \\S~3). This method has three advantages: 1) It lends itself to a simple geometrical picture where the absorbing material covers only part of the source. 2) It yields perfect fits to doublet troughs, which comprise the vast majority of high-resolution data that is needed to advance beyond the $\\tau_{ap}$ model.  However, as mentioned in \\S~3.2, the perfect fit arises mainly from fitting two equations with two unknowns. 3) Changes in the covering factor as a function of velocity were interpreted as accelerated motion with a small non-radial component (the dynamical-geometrical model of Arav et al.\\ 1999a).  However, the pure covering factor model has some significant weaknesses as well.  1) It has been known for a long time that troughs of different ions show different covering factors at the same velocity (e.g., BALQSO 1603+3002, Arav et al 1999b). Explaining such behavior involves more complicated variants of the model, chief among them is abandoning the notion of homogeneous material distribution over the covered surface.  We see a strong case of the above problem in QSO~2359--1241 where the best fit covering factor solution for the excited levels of \\feii\\ is different (smaller) than the the covering factor for the E=0 level. 2) While the issue above can be somewhat addressed using more complicated variants of the model, the following issue poses a more serious problem. We now have several cases where the distance of the outflow from the source was obtained by determining the number density of the outflow and combing it with the ionization equilibrium solution (knowledge of the ionization parameter).  Some outflows, including the one studied here, are found to be at distances of $\\sim1000$ pc from the source, and at the same time show clear evidence that the $\\tau_{ap}$ model is inapplicable to their troughs (e.g., BALQSO 1044+3656, de Kool et al.\\ 2001).  At distance up to $\\sim$10 times the size of the emission source, the dynamical-geometrical model for the changes in the covering factor as a function of velocity can work well. However, it is very difficult to envision a scenario where this model can work when the distance of the outflow from the emission source is $10^3-10^4$ times the size of the emission source.  We note that the covering factor model does not suffer from these difficulties for outflows at much closer distances to the nucleus (e.g., BALQSO 1603+3002 and NGC~3783, mentioned above).  Two main efforts were undertaken to address this serious flaw.  One effort concentrated on finding a way to interpret the number density using a model that reduces the inferred distance by 2-3 orders of magnitude (Everett et al 2002).  This was done by invoking a shielded, multiphase gas that is made of a continuous low-density wind with embedded high-density clouds.  The model also requires that the clouds are dusty and have strong differential dust depletion of iron compared to magnesium.  While the Everett et al (2002) model solves the covering factor problem by reducing the inferred distance substantially, this achievement comes at the heavy price of invoking several special conditions (an ionization shield, multiphase gas, and dusty clouds with strong differential dust depletion).  The second effort studied inhomogeneous distribution of absorbing material across the emission source as an alternative to the pure partial covering model.  The strong feature of these models is that they do not require a finely tuned covering factor at large distances and at the same time do not invoke special physical conditions to do so. Before we discuss the physical picture of these models we need to address the history of these models, which at first glance looks contradictory.  de Kool et al (2002, dKKA), showed that such models (similar to the ones used here) can yield adequate fits for two quasar outflows.  In contrast, Arav et al (2005), found that the same models could not fit the outflow troughs of the Seyfert galaxy Mrk~279. However, in the Mrk~279 data analyzed by Arav et al (2005), more than half the flux at the spectral location of the outflow troughs is contributed by the broad emission line (BEL) of the same transition. Gabel et al (2005a) determined that for Mrk~279 the outflow covers the entire continuum source but only a portion of the BEL. This picture looks physically plausible if the distance of the outflow is co-spatial or just outside the BEL region. Supporting evidence to this distance scale comes from the fact that the troughs in Mrk~279 vary on timescale of $\\sim1$ year.  Arav et al (2005) tried to find a simple inhomogeneous power-law model that will fit the data without differentiating between the BEL and continuum flux distributions, and concluded that this is not feasible.  It is not known if more complicated inhomogeneous models that will differentiate between the BEL and continuum flux sources could give a good fit to the Mrk~279 data.  In contrast, for QSO~2359--1241 the contribution of the BELs at the spectral location of the \\feii\\ troughs is less than 10\\%, and the distance scale is $\\sim1000$ pc. (Based on the inferred number density from \\feii* troughs combined with a determination of the ionization parameter, Arav et al 2008). The combination of these conditions permits us to ignore the contribution of the separate BELs and yields a better fit for the power-law model than the pure covering factor model does.   Finally, we need to address the fundamental question: What is the physical distribution of absorbing material in quasar outflows? In this work we have showed that a power-law distribution of material is a better model both statistically and physically than the pure covering factor model. A power-law model is of course a highly simplified version of what the real distribution might be. A hint about the actual situation can be inferred from resolved images of other astrophysical outflows. Planetary nebula, super nova remnants and other resolved outflows, tend to show a fragmented or web-like structure of the emitting material. We speculate that quasar outflows show similar structures on scales smaller than the size of the emission source. This is a straight forward way to explain the break down of the homogeneous absorption screen hypothesis. It is our hope that future dynamical models are able to yield physically-motivated predictions with regards to the distribution of outflowing absorbing material. These predictions could then be tested against the QSO QSO~2359--1241 and similar data sets."
    },
    "0807/0807.5061_arXiv.txt": {
        "abstract": "Mass loss is a determinant factor which strongly affects the evolution and the fate of massive stars. At low metallicity, stars are supposed to rotate faster than at the solar one. This favors the existence of stars near the critical velocity. In this rotation regime, the deformation of the stellar surface becomes important, and wind anisotropy develops. Polar winds are expected to be dominant for fast rotating hot stars.  These polar winds allow the star to lose large quantities of mass and still retain a high angular momentum, and  they modifie the evolution of the surface velocity and the final angular momentum kept in the star's core. We show here how these winds affect the final stages of massive stars, according to our knowledge about Gamma Ray Bursts. Computation of theoretical Gamma Ray Bursts rate indicates that our models have too fast rotating cores, and that we need to include an additional effect to spin them down. Magnetic fields in stars act in this direction, and we show how they modify the evolution of massive star up to the final stages. ",
        "introduction": "Rotation has a strong influence on the stellar surface. Indeed, it adds a centrifugal component to the gravity, which modifies the shape of the surface, and various quantities such as the effective temperature $T_\\mathrm{eff}$ and the mass loss flux. These effects can be derived from the von Zeipel theorem \\citep{vonZeipel1924a}, which is originally valid for conservative cases of angular momentum distribution, and was treated in the more general case of the so-called \"shellular rotation\"  by \\citet{Maeder1999a}. This theorem gives the relation between the local flux $\\mathbf{F}$ of the star and the local effective gravity: \\begin{equation} \\mathbf{F} = -\\frac{L(P)}{4\\pi G M_\\star (P)}\\mathbf{g_\\mathrm{eff}}\\left( 1 + \\zeta(\\theta)\\right) \\end{equation} \\noindent where $L(P)$ is the luminosity on the isobar and $\\mathbf{g_\\mathrm{eff}}$ the local effective gravity. The two remaining terms are given by \\begin{eqnarray} \\zeta(\\theta) &=& \\left[ \\left( 1-\\frac{\\chi_T}{\\delta}\\right)\\Theta + \\frac{H_T}{\\delta}\\frac{\\mathrm{d}\\Theta}{\\mathrm{d}r}\\right]P_2 (\\cos(\\theta))\\\\ M_\\star &=& M\\left( 1-\\frac{\\Omega^2}{2\\pi G\\rho_\\mathrm{m}}\\right) \\end{eqnarray} \\noindent Here, $\\rho_\\mathrm{m}$ is the internal average density, $M_\\star$ represents the effective mass, modified by the rotation velocity $\\Omega$, $\\chi = 4acT^3 / (3\\kappa\\rho)$ is the thermal conductivity coefficient and $\\chi_T$ is its partial derivative with respect to $T$, $\\Theta = \\frac{\\tilde{\\rho}}{\\overline{\\rho}}$ is the ratio of the horizontal density fluctuation to the average density on the isobar \\citep{Zahn1992a}. $\\delta$ is the thermodynamic coefficient $\\delta = -\\left( \\partial\\ln \\rho / \\partial\\ln T\\right)_{P,\\,\\mu}$ and $H_T$ is the temperature scale height. Generally the term $\\zeta(\\theta)$ is very small, and we can neglect it.  The total gravity at the surface of the star is given by $\\mathbf{g_\\mathrm{tot}} = \\mathbf{g_\\mathrm{eff}} + \\mathbf{g_\\mathrm{rad}} $ with $\\mathbf{g_\\mathrm{rad}} = \\frac{\\kappa\\mathbf{F}}{c}$ is the term due to radiative forces and $\\kappa$ the total Rossland mean opacity. The Eddington limit in a rotating star is defined by the vanishing of $\\mathbf{g_\\mathrm{tot}}$ and we find the limiting flux \\begin{equation} \\mathbf{F_\\mathrm{lim}} = -\\frac{c}{\\kappa}\\mathbf{g_\\mathrm{eff}}\\,. \\end{equation}  \\noindent The Eddington factor at a given colatitude, which is given by the ratio of the local flux to the limiting flux, becomes \\citep{Maeder1999a} \\begin{equation} \\Gamma_\\Omega(\\theta) = \\frac{L(P)}{L_\\mathrm{max}}\\,\\,\\,\\mathrm{with}\\,\\,\\, L_\\mathrm{max} = \\frac{4\\pi cGM}{\\kappa(\\theta)\\left( 1 + \\zeta(\\theta)\\right)}\\left( 1-\\frac{\\Omega^2}{2\\pi G\\rho_\\mathrm{m}}\\right)\\,\\, . \\end{equation} \\noindent A first interesting consequence is that the maximum luminosity of a star, given by $\\Gamma_\\Omega(\\theta) = 1$,  is lowered by rotation (compare $L_\\mathrm{max}$ given above with the non-rotating one $L_\\mathrm{max,\\,no\\,rot} = \\frac{4\\pi cGM}{\\kappa}$).  Rotation has also a strong impact on the mass loss rate of the star. \\citet{Maeder2000a} found the following relation between the mass loss rate at a given rotational velocity $\\dot{M}(\\Omega)$ and the non-rotating one:  \\begin{equation} \\frac{\\dot{M}(\\Omega)}{\\dot{M}(0)} =  \\frac{\\left( 1-\\Gamma\\right)^{\\frac{1}{\\alpha} - 1}}{\\left[ 1-\\frac{\\Omega^2}{2\\pi G\\rho_\\mathrm{m}}-\\Gamma \\right]^{\\frac{1}{\\alpha} - 1}} \\end{equation} \\noindent where $\\Gamma = \\frac{\\kappa L}{4\\pi cGM}$ is the \"non-rotating\" Eddington factor. We see that rotation will increase the total mass loss rate of the star.  Rotation has another effect on the mass loss rate: it is no long isotropic at the surface of a rotating star, but becomes colatitude-dependent. \\citet{Maeder2000a} give the following local mass loss rate $\\Delta\\dot{M}$ per unit surface $\\Delta\\sigma$: \\begin{equation} \\frac{\\Delta\\dot{M}}{\\Delta\\sigma} \\simeq \\left( k\\alpha\\right)^\\frac{1}{\\alpha}\\left( \\frac{1-\\alpha}{\\alpha}\\right)^\\frac{1-\\alpha}{\\alpha}\\left[\\frac{L(P)}{4\\pi GM_\\star(P)}\\right]^\\frac{1}{\\alpha}\\frac{g_\\mathrm{eff}}{\\left(1-\\Gamma_\\Omega(\\theta)\\right)^{\\frac{1}{\\alpha}-1}} \\end{equation} \\noindent with $\\kappa$ and $\\alpha$ the force multiplier parameters. We neglect here the small effect of $\\zeta(\\theta)$. Rotation favors mass loss through the terms $M_\\star$ and $\\Gamma_\\Omega$. Mass loss occurs preferentially where $g_\\mathrm{eff}$ is small, \\textit{i.e} at the poles. Then mass loss rate is thus varying as a function of the colatitude, producing the anisotropic wind phenomenon.  \\begin{figure} \\begin{center} \\includegraphics[angle = 0, width=9 cm]{P020z-5S7Flux.eps} \\caption{Effects of rotation on a $20\\,\\mathrm{M}_\\odot$ star at $Z=10^{-5}$ and with $\\frac{\\Omega}{\\Omega_\\mathrm{crit}}=0.95$. The star is seen equator-on, and the axis are in $\\mathrm{R}_\\odot$. The latitudinal variation of $T_\\mathrm{eff}$ is shown (gray scale), and the arrows represents the masse flux at a given colatitude.} \\end{center} \\label{shape} \\end{figure} Figure \\ref{shape} shows various effects of the rotation at the surface of a $20\\, \\mathrm{M}_\\odot$ star at a metallicity of $10^{-5}$ and at 95\\% of the critical rotation velocity. We can first remark that the star becomes oblate, with an equatorial-to-polar radius ratio $\\frac{R_\\mathrm{eq}}{R_\\mathrm{pol}}\\simeq 1.3$. Then we note the variation of the effective temperature with respect to colatitude: $T_\\mathrm{eff}$ at the pole is around $48000\\,\\mathrm{K}$, while it is only around $34000\\,\\mathrm{K}$ at the equator. Finally, we see that the mass loss flux is larger at the pole by a factor of $\\sim 3.8$. This allows the star to lose  $10\\%$ less angular momentum than if the same amount of mass was lost isotropically. ",
        "conclusions": ""
    },
    "0807/0807.3471_arXiv.txt": {
        "abstract": "We consider perturbations in a cosmological model with a small coupling between dark energy and dark matter. We prove that the stability of the curvature perturbation depends on the type of coupling between dark sectors. When the dark energy is of quintessence type, if the coupling is proportional to the dark matter energy density, it will drive the instability in the curvature perturbations; however if the coupling is proportional to the energy density of dark energy, there is room for the stability in the curvature perturbations. When the dark energy is of phantom type, the perturbations are always stable, no matter whether the coupling is proportional to the one or the other energy density. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0807/0807.3192_arXiv.txt": {
        "abstract": "Variations of the sunspot number are important indicators of the solar activity cycles. The sunspot formation is a result of a dynamo process inside the Sun, which is far from being understood. We use simple dynamical models of the dynamo process to simulate the magnetic field evolution and investigate general properties of the sunspot number variations during the solar cycles. We have found that the classical Parker's model with a standard kinetic helicity quenching cannot represent the typical profiles of the solar-cycle variations of the sunspot number,  and also does not give chaotic solutions. For modeling of the solar cycle properties we use a nonlinear dynamo model of Kleeorin and Ruzmaikin (1982), which takes into account dynamics of the turbulent magnetic helicity. We have obtained a series of periodic and chaotic solutions for different layers of the convective zone. The solutions qualitatively reproduce some basic observational features of the solar cycle properties, in particular, the relationship between the growth time and the cycle amplitude. Also, on the longer time scale the dynamo model with the magnetic helicity has intermittent solutions, which may be important for modeling long-term variations of the solar cycles. ",
        "introduction": "Investigation of solar activity has a long history. The 11-year sunspot cycles were discovered by Schwabe (1844). Detailed observations of sunspots locations on the solar disk and their sizes provided the latitudinal distribution during the cycle, known as the \"butterfly diagram\". Discovered by Hale (1919) the flip-flop phenomenon of the global solar magnetic field with a period of 22 years showed the connection between variations of the sunspots number and the magnetic field evolution. Taking into account the Hale's law Bracewell (1953) proposed to include the information about the periodical reversals of the magnetic field in the sunspot number series by assigning alternating positive and negative signs to the sunspot cycles, thus representing the sunspot cycles  as a 22-year oscillation. Bracewell suggested that the nonlinear temporal distribution of the sunspots number can be written in the form of a three-halfs law  \\citep{brace88}: $W(t)=100\\left[\\mid W_{lin}(t)\\mid / 83\\right]^{3/2}$, where $W_{lin}(t)$ is a new \"linearized\" sunspot number, which is supposedly proportional to a typical strength of the Sun's toroidal magnetic field.  For an explanation of the magnetic field generation Parker (1955) proposed a simple dynamo model, which describes the phenomenon as an action of two factors: the differential rotation and cyclonic convective vortices. The mean-field theory and discovery of the $\\alpha$-effect give us a general description of the process of magnetic field generation  \\citep{krause80}. For recent reviews we refer to Charbonneau (2005) and Brandenburg and Subramanian (2005).  It is known that the dynamo process is characterized by algebraic and dynamic nonlinearities \\citep[e.g.][]{klee07, sokol07}. The algebraic nonlinearity can be determined as influence of the magnetic field on fluid motions and on the kinetic helicity. This results in quenching of the electromotive force and limits the growth of the magnetic field. The evolution of the small-scale magnetic helicity in the turbulent plasma causes a dynamical nonlinearity in the dynamo process.  The turbulent helicity conditionally can be divided into two parts: hydrodynamic and magnetic. The kinetic helicity describes helical turbulent fluid motions; the magnetic helicity determines the order of twisted magnetic field lines.  Due to the fact that the kinetic helicity makes the magnetic field small-scaled, the back influence on the turbulent fluid motions can restrict the unlimited growth of the magnetic field. In the mean-field approach the magnetic helicity is separated into large- and small-scale components. Because of the conservation of the total helicity a growth of the large-scale magnetic helicity due to the dynamo action is compensated by the growth of the small-scale helicity of opposite sign \\citep{sokol07}. Thus, the small- and large-scale magnetic fields grow together and are mirror-asymmetrical. This means that the condition of magnetic helicity conservation is, perhaps, more severe for a restriction of the dynamo action than the condition of the energy conservation, which leads to quenching of the kinetic helicity. Since the concept of magnetic helicity was introduce into the dynamo theory by Pouquet {\\it et al.} (1976) many magnetic helicity models were suggested. For modeling the solar cycle we choose the formulation of Kleeorin and Ruzmaikin (1982), explicitly based on the idea of magnetic helicity conservation. Observational data of solar magnetic fields are in a reasonable agreement with the idea of the magnetic helicity conservation \\citep{kleeorin2003, zhang2006, sokol07, sokoloff2007}. Similar dynamo models with magnetic helicity have been also considered in the context of galactic dynamos \\citep{klee00} and have included helicity transport by mean flows \\citep{sur07}, which prevents the effect of `catastrophic' helicity quenching (e.g. Brandenburg and Subramanian, 2005). In our model, we don't include the helicity transport and consider the helicity quenching with a free parameter.  In this paper, we consider nonlinear behavior of the Parker's dynamo model (without magnetic helicity) and the Kleeorin-Ruzmaikin model (with magnetic helicity), and show that the latter can reproduce the qualitative behavior of the sunspot number variations during the solar cycles, including periodic and chaotic solutions for conditions of the solar convective zone. For simplicity, we use a \"low-order model\" approach \\citep[e.g.][]{weiss84, soknef07}, reducing the dynamo equations to a simple nonlinear dynamical system. In section 2, for consistency, we reproduce the formulations of the Parker's and Kleeorin-Ruzmaikin's models, and discuss conditions of linear stability and solar parameters. In section~3, we present the numerical solutions and compare these with observed properties of the sunspot number variations in solar cycles. ",
        "conclusions": "We have presented a numerical analysis of simple dynamical models describing the nonlinear behavior of two dynamo models, the classical Parker's dynamo model with the standard $\\alpha$-quenching and the model of Kleeorin and Ruzmaikin (1982), which takes into account effects of turbulent magnetic helicity. The models are analyzed in the low-mode approximation with the goal of representing variations of a sunspot number parameter during the solar cycles. We found that the Parker's model does not  reproduce the typical behavior of the sunspot number with a fast growth and slow decay or obtain a chaotic solution in the low-mode approximation, even in strong nonlinear regimes. The analysis of the Kleeorin-Ruzmaikin model showed the existence of nonlinear periodic and chaotic solutions for conditions of the solar convective zone. For this model we obtained the profiles of the sunspot number variations, which qualitatively reproduce the typical profile of the solar cycles.  It is interesting that the properties of the simulated cycles demonstrate good qualitative agreements with the properties of the observed solar cycles. In particular, the relationship between the cycle amplitude and the growth time looks similar. Of course, such simplified models cannot pretend to provide a realistic description of the solar cycles. More complicated models such as the models with flux transport \\citep[e.g.][]{dikpati06} should be considered for more detailed description of the solar magnetic field evolution.  Nevertheless, these results encourage further theoretical and observational investigations of the dynamo model with the turbulent magnetic helicity. We also hope that through the use of a simple dynamical model reproducing the basic properties of the solar cycle it will become possible to apply the data assimilation methods for predicting the solar cycles \\citep[e.g.][]{kiti08}."
    },
    "0807/0807.4317_arXiv.txt": {
        "abstract": "Possible deviations from a  radiation-dominated evolution, occurring prior to the synthesis of light nuclei, impacted on the spectral energy density of high-frequency gravitons. For a systematic scrutiny of this situation, the  $\\Lambda$CDM paradigm must be complemented by (at least two) physical parameters describing, respectively, a threshold frequency and a slope. The supplementary frequency scale sets the lower border of a  high-frequency domain where the spectral energy grows with a slope which depends, predominantly, upon  the total sound speed of the plasma right after inflation. While the infra-red region of the graviton energy spectrum is nearly scale-invariant, the expected signals for typical frequencies larger than $0.01$ nHz are hereby analyzed in a  model-independent framework by requiring that the total sound speed of the post-inflationary plasma be smaller than the speed of light.  Current (e.g. low-frequency)  upper limits on the tensor power spectra (determined from the combined analysis of the three large-scale  data sets) are shown to be compatible with a detectable signal in the frequency range of wide-band interferometers. In the present context, the scrutiny of the early evolution of the sound speed of the plasma can then be mapped onto a reliable strategy of parameter extraction including not only the well established cosmological observables but also the forthcoming data from wide band interferometers. ",
        "introduction": "\\label{sec1} Cosmological observations rely on three pivotal data sets, i.e. the Cosmic Microwave Background (CMB) data, the determinations of the matter power spectrum from galaxy surveys and the supernova light curve observations.  The large-scale measurements are inextricably bound to the model used to interpret the data.  Consequently the three aforementioned data sets are jointly analyzed in terms of a standard scenario which is often dubbed $\\Lambda$CDM paradigm, where $\\Lambda$ qualifies the dark energy component and CDM denotes the cold dark matter component. While all the current observations are based, directly or indirectly, on the electromagnetic spectrum, there is the hope, in the future, that the electromagnetic observations could be complemented by the analysis of the spectrum of the relic gravitons which have been produced both in the context of the $\\Lambda$CDM paradigm as well as in other related contexts.  The problem is, therefore, twofold: on the one hand reliable estimates of the spectrum of the relic gravitons arising in the $\\Lambda$CDM paradigm are needed. On the other hand it will be important to analyze other complementary scenarios. The purpose of the present paper is to address both issues in quantitative terms. The relic graviton background produced in the context of the $\\Lambda$CDM paradigm is expected to be rather minute and undetectable by wide-band interferometers \\cite{LIGO,VIRGO,TAMA,GEO} in one of their future realizations. There are, however, extensions of the $\\Lambda$CDM scenario where the signal potentially detectable by wide-band interferometers is much larger than in the current paradigm.  The recent WMAP 5-yr data \\cite{WMAP51,WMAP52,WMAP53,WMAP54,WMAP55}  set quite stringent bounds on the amplitude of the relic graviton spectral energy density for typical frequency scales \\footnote{Natural units $\\hbar = c = k_{\\mathrm{B}} =1$ will be consistently adopted all along the present investigation.}: \\begin{equation} \\nu_{\\mathrm{p}} = \\frac{k_{\\mathrm{p}}}{2\\pi}  = 3.092\\times 10^{-18} \\,\\, \\mathrm{Hz} \\equiv 3.092 \\,\\,\\mathrm{aHz}, \\label{EQ1} \\end{equation} where, according to the prefixes of the international system of units, $1\\,\\mathrm{aHz} = 10^{-18}$ Hz. The wavenumber $k_{\\mathrm{p}}$ is also called sometimes pivot scale\\footnote{The pivot wavenumber $k_{\\mathrm{p}}$ corresponds to an effective multipole $\\ell_{\\mathrm{eff}} \\simeq 30$.} since it is customary, in the experimental analyses of CMB data \\cite{WMAP54,WMAP55}, to assign the amplitude of the scalar and tensor modes exactly at $k_{\\mathrm{p}}$.  In the same units of Eq. (\\ref{EQ1}) the typical frequency interval potentially accessible to the observations of the wide-band interferometers ranges between few Hz and $10$ kHz with a peak of sensitivity around $100$Hz.  The logic followed in the present investigation will be to compute as accurately as possible the relic graviton spectra in terms of the parameters of the putative $\\Lambda$CDM paradigm. The CMB data will then be used to enforce the normalization of the spectra at $\\nu_{\\mathrm{p}}$. This will allow for the estimate of the spectral energy density at the frequency explored by wide-band interferometers (i.e., approximately, $100$ Hz) not only in the case of the $\\Lambda$CDM paradigm but also in the context of its extensions.  The latter extensions will be examined on the basis of their plausibility, e.g. the models which are already incompatible (or barely compatible) with CMB observations will not be analyzed and the attention will be focussed on those scenarios which are not ruled out by (current) large-scale observations and which may lead to a potentially large signal at the wide-band interferometer scale. In the present introductory section, after a general discussion of the typical frequencies of the graviton spectrum, the present status of CMB observations and wide-band interferometers observations will be swiftly discussed.  Specific attention will be paid to those quantitative aspects which are germane to our theme, i.e. the stochastic backgrounds of relic gravitons. Some of the concepts introduced here will also be more specifically addressed in the forthcoming sections.  At the end of this introduction, the purposes of the present investigation will be more specifically outlined.  \\subsection{Typical frequencies of the problem} To compare frequencies it is mandatory to specify the background and the appropriate conventions on the normalization of the scale factor. Consistently with the $\\Lambda$CDM paradigm, the background geometry will be taken to be conformally flat, i.e. \\begin{equation} ds^2 = \\overline{g}_{\\mu\\nu} dx^{\\mu} dx^{\\nu} \\equiv a^2(\\tau) [ d\\tau^2 - d\\vec{x}^2], \\qquad \\overline{g}_{\\mu\\nu} = a^2(\\tau) \\eta_{\\mu\\nu}, \\label{EQ2} \\end{equation} where $\\eta_{\\mu\\nu}$ is the Minkowski metric with signature mostly minus, i.e. $(+, -, -, -)$. The scale factor at the present time will be normalized to unity, i.e. $a_{0} = 1$. Within the latter convention, the comoving frequencies (or wavelengths) coincide, at the present time, with the physical frequencies (or wavelengths). The (conformal) time derivative of the logarithm of the scale factor will be used throughout the script and it is defined as \\begin{equation} {\\mathcal H} = \\frac{a'}{a}  = \\frac{d \\ln{a}}{d\\tau}, \\label{T7a} \\end{equation} Note that, in Eq. (\\ref{T7a}), the prime denotes a derivation with respect to the conformal time coordinate $\\tau$: this notation will be consistently enforced in the whole investigation.  The evolution of the background can be expressed in terms of ${\\mathcal H}$ and ${\\mathcal H}'$ and it is given by: \\begin{eqnarray} && 3 {\\mathcal H}^2 = a^2 \\ell_{\\mathrm{P}}^2 \\rho_{\\mathrm{t}}, \\label{FL1}\\\\ && 2({\\mathcal H}^2 - {\\mathcal H}') =  a^2 \\ell_{\\mathrm{P}}^2 (\\rho_{\\mathrm{t}} + p_{\\mathrm{t}}), \\label{FL2}\\\\ && \\rho_{\\mathrm{t}}' + 3 {\\mathcal H} (\\rho_{\\mathrm{t}} + p_{\\mathrm{t}}) =0, \\label{FL3} \\end{eqnarray} where $\\rho_{\\mathrm{t}}$ and $p_{\\mathrm{t}}$ denote, respectively, the total energy density and the total pressure of the plasma.  The frequency of Eq. (\\ref{EQ1}) can be usefully compared with two other important frequencies, i.e. the frequency of matter-radiation equality (be it $\\nu_{\\mathrm{eq}}$) and the frequency of neutrino decoupling (which also coincides, in loose terms, with the Hubble radius at the onset of big bang nucleosynthesis). These two frequencies can then be written, respectively, as: \\begin{eqnarray} \\nu_{\\mathrm{eq}} &=& \\frac{k_{\\mathrm{eq}}}{2 \\pi} = 1.281 \\times 10^{-17} \\biggl(\\frac{h_{0}^2 \\Omega_{\\mathrm{M}0}}{0.1326}\\biggr) \\biggl(\\frac{h_{0}^2 \\Omega_{\\mathrm{R}0}}{4.15 \\times 10^{-5}}\\biggr)^{-1/2}\\,\\, \\mathrm{Hz}, \\label{EQ3}\\\\ \\nu_{\\mathrm{bbn}}&=& 2.252\\times 10^{-11} \\biggl(\\frac{g_{\\rho}}{10.75}\\biggr)^{1/4} \\biggl(\\frac{T_{\\mathrm{bbn}}}{\\,\\,\\mathrm{MeV}}\\biggr) \\biggl(\\frac{h_{0}^2 \\Omega_{\\mathrm{R}0}}{4.15 \\times 10^{-5}}\\biggr)^{1/4}\\,\\,\\mathrm{Hz}\\simeq 0.01 \\,\\,\\mathrm{nHz}. \\label{EQ4} \\end{eqnarray} In Eqs. (\\ref{EQ3}) and (\\ref{EQ4}) $\\Omega_{\\mathrm{M}0}$ and $\\Omega_{\\mathrm{R}0}$ denote, respectively, the present critical fraction of matter and radiation with typical values drawn from the best fit to the WMAP 5-yr data alone and within the $\\Lambda$CDM paradigm. In Eq. (\\ref{EQ4}) $g_{\\rho}$ denotes the effective number of relativistic degrees of freedom entering the total energy density of the plasma.  While $\\nu_{\\mathrm{eq}}$  is still close to the aHz, $\\nu_{\\mathrm{bbn}}$ is rather in the nHz range.  The success of the CMB and BBN calculations implicitly demands that, after neutrino decoupling, the Universe was already dominated by radiation. If we assume that the radiation dominates right at the end of inflation, then the maximal frequency of the graviton spectrum can be computed and it is given by \\begin{equation} \\nu_{\\mathrm{max}}  = 0.346 \\,\\biggl(\\frac{\\epsilon}{0.01}\\biggr)^{1/4} \\biggl(\\frac{{\\mathcal A}_{\\mathcal R}}{2.41 \\times 10^{-9}}\\biggr)^{1/4} \\biggl(\\frac{h_{0}^2 \\Omega_{\\mathrm{R}0}}{4.15 \\times 10^{-5}}\\biggr)^{1/4} \\,\\mathrm{GHz}, \\label{EQ5} \\end{equation} where ${\\mathcal A}_{\\mathcal R}$ denotes the amplitude of the power spectrum of curvature perturbations evaluated at the pivot wavenumber $k_{\\mathrm{p}}$. Between $\\nu_{\\mathrm{bbn}}$ and $\\nu_{\\mathrm{max}}$ there are roughly 20 orders of magnitude in frequency. In the $\\Lambda$CDM scenario the relic graviton spectrum has, in this range, always the same slope.   \\subsection{CMB data and relic gravitons} As already mentioned, CMB experiments are sensitive to long wavelength gravitons with typical frequencies of the order of $\\nu_{\\mathrm{p}} \\sim \\mathrm{aHz}$ (see also Eq. (\\ref{EQ1})). The number of CMB parameters depends upon the specific model used to interpret (and fit) the data. The $\\Lambda$CDM scenario probably contains the fewest number of parameters required to have a consistent fit of CMB data.  The $\\Lambda$CDM parameters can be inferred from various experiments and, among them, a central role is played by  WMAP   \\cite{WMAP51,WMAP52,WMAP53,WMAP54,WMAP55}  (see also \\cite{WMAPfirst1,WMAPfirst2,WMAPfirst3} for first year data release and \\cite{WMAPthird1,WMAPthird2} for the third year data release) as well as other experiments (see, for instance, \\cite{AC1} in connection with the 5-yr WMAP data release). The TT, TE  and, partially EE angular power spectra\\footnote{Following the custom the TT correlations will  simply denote the angular power spectra of the temperature autocorrelations. The TE and the EE power spectra  denote, respectively, the cross power spectrum between temperature and polarization and the polarization autocorrelations.} have been measured by the WMAP experiment. Other (i.e. non space-borne) experiments are now measuring polarization observables, in particular there are the 3-yr Dasi release \\cite{dasi3yr},  the CAPMAP experiment \\cite{CAPMAP}, the recent QUAD data \\cite{Quad1,Quad2}, as well as various other experiments at different stages of development.  The TT, TE and EE power spectra are customarily analyzed in the light of the minimal $\\Lambda$CDM scenario but also other models are possible and they include, for instance, the addition of spatial curvature (i.e. the open-$\\Lambda$CDM), more general parametrizations for the equation of state of dark-energy and so on and so forth.  The combined analysis of the CMB data, of the large-scale structure data \\cite{LSS1,LSS2} and of the supernova data \\cite{SN1,SN2} can lead to quantitative upper limits on the possible contribution of the tensor modes to the initial conditions of the CMB temperature and polarization anisotropies. These upper limits can be phrased in terms of $r_{\\mathrm{T}}$, i.e.  the ratio between the power spectrum of tensor fluctuations and the power spectrum of the scalar fluctuations evaluated at the pivot wavenumber  $k_{\\mathrm{p}} = 0.002\\,\\,\\mathrm{Mpc}^{-1}$.  In the minimal paradigm (i.e. the $\\Lambda$CDM scenario) the tensor are not included in the fit.  If the inflationary phase is driven by a single scalar degree of freedom and if the radiation dominance kicks in almost suddenly after inflation, $r_{\\mathrm{T}}$ not only determines the tensor amplitude but also, thanks to the algebra obeyed by the slow-roll parameters, the  slope of the tensor power spectrum, customarily denoted by $n_{\\mathrm{T}}$. To lowest order in the slow-roll expansion, therefore, the tensor spectral index is slightly red and it is related to $r_{\\mathrm{T}}$ (and to the slow-roll parameter) as $n_{\\mathrm{T}} \\simeq - r_{\\mathrm{T}}/8 \\simeq  - 2 \\epsilon$, where $\\epsilon = - \\dot{H}/H^2$ measures the rate of decrease of the Hubble parameter during the inflationary epoch \\footnote{The overdot will denote throughout the paper a derivation with respect to the cosmic time coordinate $t$ while the prime will denote a derivation with respect to the conformal time coordinate $\\tau$.}. Within the established set of conventions the scalar spectral index $n_{\\mathrm{s}}$ is given by $n_{\\mathrm{s}} = (1 - 6 \\epsilon + 2 \\overline{\\eta})$ and it depends not only upon $\\epsilon$ but also upon the second slow-roll parameter $\\overline{\\eta} = \\overline{M}_{\\mathrm{P}}^2 V_{,\\varphi\\varphi}/V$ (where $V$ is the inflaton potential, $V_{,\\varphi\\varphi}$ denotes the second derivative  of the potential with respect to the inflaton field and $\\overline{M}_{\\mathrm{P}} = 1/\\sqrt{8\\pi G}$).  Depending upon the specific data used in the analysis, the upper limits on $r_{\\mathrm{T}}$ as well as the determination of the other cosmological parameters may change slightly. \\begin{table}[!ht] \\begin{center} \\begin{tabular}{||l|c|c|c|c|c||} \\hline Data & $r_{\\mathrm{T}}$ &$n_{\\mathrm{s}}$ &$\\Omega_{\\Lambda}$&$\\Omega_{\\mathrm{M}0}$&$k_{\\mathrm{eq}} \\mathrm{Mpc}$\\\\ \\hline WMAP5 alone& $<0.43$ &$0.986 \\pm 0.22$& $0.770_{-0.032}^{+0.033}$&$0.230_{-0.033}^{0.032}$ & $0.00936$ \\\\ WMAP5 + Acbar& $< 0.40  $& $0.985_{-0.020}^{0.019}$&$0.767 \\pm 0.032$&$0.233\\pm 0.032$ &$0.00944$\\\\ WMAP5+ LSS + SN &$<0.20$ &$0.968 \\pm 0.015$&$0.725 \\pm 0.015$&$0.275 \\pm 0.015$&$0.00999$\\\\ WMAP5+ other CMB data & $<0.36$ & $0.979\\pm 0.020$&$0.775\\pm 0.032$&$0.225\\pm 0.032$&$0.00922$\\\\ \\hline \\end{tabular} \\caption{The values of $r_{\\mathrm{T}}$ are reported as they have been estimated in the absence of any running of the (scalar) spectral index.} \\end{center} \\label{MOD1} \\end{table} In Tab. \\ref{MOD1} the upper limits on $r_{\\mathrm{T}}$ are illustrated as they are determined from the combination of different data sets. For illustration the determined values of the scalar spectral index (i.e. $n_{\\mathrm{s}}$), of the dark energy and dark matter fractions (i.e., respectively, $\\Omega_{\\Lambda}$ and $\\Omega_{\\mathrm{M}0}$), and of the typical wavenumber of equality $k_{\\mathrm{eq}}$ are also reported in the remaining  columns.  While different analyses can be performed, it is clear, by looking at Tab. \\ref{MOD1} that the typical upper bounds on $r_{\\mathrm{T}}$ range between, say, $0.2$ and $0.4$. Slightly more stringent limits can also be obtained by adding supplementary assumptions. Within a conservative perspective, the tensor power spectra are, at least, ten times  smaller than the power spectra of curvature perturbations. In the near future the Planck explorer satellite \\cite{planck} might be able to set  more direct limits on $r_{\\mathrm{T}}$ by measuring (hopefully) the BB angular power spectra \\footnote{Forthcoming projects like Clover \\cite{CLOVER}, Brain \\cite{BRAIN}, Quiet \\cite{QUIET} and Spider \\cite{SPIDER} have polarization as specific target.}. The E-mode power spectra and the B-mode power spectra arise as two orthogonal combinations of the Stokes parameters which are frame-dependent (i.e. $Q$ and $U$). While the adiabatic mode leads naturally to the E-mode polarization, the only way of obtaining the B-mode (in the standard $\\Lambda$CDM paradigm) is through the contribution of the tensor modes. Consequently, a detection o the BB angular power spectra would be equivalent, in the $\\Lambda$CDM framework, to a first determination of $r_{\\mathrm{T}}$. Having reviewed the essentials of CMB data and their connection with the relic graviton spectra, we can now move to higher frequencies and describe the status of the other devices which could shed light on the relic gravitons, i.e. the wide-band interferometers.  \\subsection{Wide-band interferometers}  The wide-band interferometers operate in a window ranging from few Hz up to $10$ kHz. The available interferometers are Ligo \\cite{LIGO}, Virgo \\cite{VIRGO}, Tama \\cite{TAMA} and Geo \\cite{GEO}. The sensitivity of a given pair of wide-band detectors to a stochastic background of relic gravitons depends upon the relative orientation of the instruments. The wideness of the band (important for the correlation among different instruments) is not as large as $10$ kHz  but typically narrower and, in an optimistic perspective,  it could range up to $100$ Hz. The putative frequency of wide-band detectors will therefore be indicated as $\\nu_{\\mathrm{LV}}$, i.e. in loose terms, the Ligo/Virgo frequency. There are daring projects of wide-band detectors in space like the Lisa \\cite{LISA}, the BBO \\cite{BBO} and the Decigo \\cite{DECIGO} projects. The common denominator of these three projects is that they are all space-borne missions and that they are all sensitive to frequencies smaller than the mHz. While wide-band interferometers are now operating and might even reach their advanced sensitivities along the incoming decade, the achievable sensitivities of space-borne interferometers are still on the edge of the achievable technologies.  Since $\\nu_{\\mathrm{bbn}} < \\nu_{\\mathrm{LV}} < \\nu_{\\mathrm{max}}$ the wide-band interferometers are an ideal instrument to investigate the relic graviton spectrum in the unknown territory where there are neither direct nor indirect tests on the thermal history of the plasma. The problem is that, as it will be carefully shown, the  spectral energy density of the relic gravitons produced within the $\\Lambda$CDM model is quite minute and it is undetectable by interferometers even in their advanced version where the sensitivity is expected to improve by 5 or even 6 orders of magnitude in comparison with the present performances \\cite{LIGOS1,LIGOS2,LIGOS3} (see also \\cite{virgoligo} and \\cite{auriga}). This impasse, as previously stressed, stems from the assumption that, right after inflation, the radiation-dominated evolution kicks in almost suddenly. At the moment, there are no evidences neither in favor of such a statement nor against it.  The main theme of the present investigation will be to reverse this problem. It will be argued that wide-band detectors, in their advanced version, will be certainly able to test definite deviations from a simplistic thermal history of the plasma, i.e. the one stipulating that, after inflation, the radiation was suddenly dominating the evolution.  \\subsection{Layout of the investigation}  In the present investigation the spectral energy density of the relic gravitons will be calculated first at small frequencies (compatible with the CMB observations) and then at higher frequencies (compatible with the operational window of wide-band interferometers). The latter calculation will be performed both n the case of the $\\Lambda$CDM paradigm but also in those extensions which may lead to a large spectral energy density at the scale of the wide-band interferometers without violating any of the bounds stemming from CMB observations.  The spectral energy density of the relic gravitons will be introduced in section \\ref{sec2}. Gravity is inherently a non-Abelian gauge theory there are potential ambiguities in defining univocally an energy-momentum (pseudo)-tensor for the relic gravitons: it will be shown that different choices of the energy-momentum pseudo-tensor lead to the same spectral energy density of the relic gravitons. The punch line of section \\ref{sec2} will be that the spectral energy density can be more accurately performed with numerical methods rather than resorting, as often done, to semi-analytical estimates which amount to estimate first the power spectrum and then the spectral energy density.  In section \\ref{sec3}, using the numerical techniques introduced in section \\ref{sec2}, the spectral energy density of the relic gravitons will be computed in the case of the standard $\\Lambda$CDM paradigm which leads to nearly scale-invariant spectra. The signal arising in the context of the $\\Lambda$CDM paradigm will be confronted with the sensitivity of wide-band interferometers.  After showing the compatibility of the new methods with the results of the nearly scale-invariant spectra, possible scaling violations in the spectral energy density will be discussed in section \\ref{sec4}. While in the $\\Lambda$CDM paradigm the spectral energy density of the relic gravitons is nearly scale invariant it is plausible to construct a class of models where the spectral energy density is fully compatible with the CMB and with the large-scale data at low frequencies while it is potentially detectable by wide-band interferometers. When we say that it is plausible this simply means that it is not forbidden by any of the current observational data. The proposed extensions of the $\\Lambda$CDM paradigm have also a physical interpretation (see section \\ref{sec4}) since they naturally arise when the thermal history of the Universe deviates, for sufficiently early times, from the usual assumptions of the $\\Lambda$CDM scenario which stipulates that, right after inflation, the Universe suddenly becomes dominated by radiation.  The minimal realization of the ideas pursued in section \\ref{sec4} is scrutinized in section \\ref{sec5} and it is dubbed T$\\Lambda$CDM scenario (for tensor-$\\Lambda$CDM). The T$\\Lambda$CDM paradigm consists of two supplementary parameters, i.e., in broad terms, a new pivotal frequency and a new spectral slope. The new frequency marks the onset of the high-frequency branch of the spectral energy density of the relic gravitons.  In section \\ref{sec5} the T$\\Lambda$CDM paradigm is compatible with the current bounds stemming not only from CMB and large-scale structure. It will also be required that the big-bang nucleosynthesis (BBN) as well as pulsar timing constraints are satisfied and this will allow to spell out quantitatively the restrictions on the two supplementary parameters characterizing the T$\\Lambda$CDM scenario. \\renewcommand{\\theequation}{2.\\arabic{equation}} \\setcounter{equation}{0} ",
        "conclusions": ""
    },
    "0807/0807.1532_arXiv.txt": {
        "abstract": "We present a study of the magnetic field of the Small Magellanic Cloud (SMC), carried out using radio Faraday rotation and optical starlight polarization data.  Consistent negative rotation measures (RMs) across the SMC indicate that the line-of-sight magnetic field is directed uniformly away from us with a strength $0.19 \\pm 0.06$~$\\mu$G. Applying the Chandrasekhar-Fermi method to starlight polarization data yields an ordered magnetic field in the plane of the sky of strength $1.6 \\pm 0.4$~$\\mu$G oriented at a position angle $4^\\circ \\pm 12^\\circ$, measured counter-clockwise from the great circle on the sky joining the SMC to the Large Magellanic Cloud (LMC).  We construct a three-dimensional magnetic field model of the SMC, under the assumption that the RMs and starlight polarization probe the same underlying large-scale field. The vector defining the overall orientation of the SMC magnetic field shows a potential alignment with the vector joining the center of the SMC to the center of the LMC, suggesting the possibility of a ``pan-Magellanic'' magnetic field. A cosmic-ray driven dynamo is the most viable explanation of the observed field geometry, but has difficulties accounting for the observed uni-directional field lines. A study of Faraday rotation through the Magellanic Bridge is needed to further test the pan-Magellanic field hypothesis. ",
        "introduction": "\\label{section:introduction} Magnetic fields play key roles in many astrophysical processes in the interstellar medium (ISM) --- they accelerate and confine cosmic rays, trigger star formation and exert pressure to balance against gravity \\citep{beck2007}. Therefore, to better understand galaxy evolution, investigating the structure, origin and evolution of galactic magnetic fields is necessary.  It is useful to picture the total magnetic field at any location in a galaxy as a superposition of an ordered large-scale component and a random small-scale component. An ordered (or uniform) field can either be coherent or incoherent: a coherent field has unidirectional field lines, whereas an incoherent field has field lines of the same orientation but has frequent field reversals. A large-scale dynamo is the only known mechanism that can generate large-scale coherent fields \\citep{beck2004a}.  Coherent magnetic fields have been observed in normal spiral galaxies such as the Milky Way and M31. The fields are typically in spiral-like configurations with field strengths of a few $\\mu$G \\citep{beck2007}. Because these galaxies have significant differential rotation, such observations can be explained by the standard $\\alpha$-$\\omega$ dynamo which amplifies and orders the field by small scale turbulent motion (the $\\alpha$-effect) as well as differential rotation (the $\\omega$-effect) in the galactic disk on a global e-folding time of  $\\sim$ 10$^9$ yrs \\citep{shukurov2004}. Despite its success in accounting for the large scale coherent field seen in spiral galaxies, the standard dynamo theory fails to explain the presence of coherent magnetic fields discovered in several irregular galaxies such as NGC 4449 and the Large Magellanic Cloud (LMC), due to its long amplification time scale \\citep{klein1996,chyzy2000,gaensler2005}.  The polarized radio continuum emission of NGC 4449, a dwarf irregular galaxy, at 4.9 and 8.6 GHz reveals large scale spiral like structure in the magnetic field. Moreover, this slowly rotating galaxy shows regions of coherent magnetic field from Faraday rotation studies \\citep{klein1996}. The long field amplification time scale of the classical mean field dynamo argues against this being the underlying mechanism that produces the magnetic field observed in NGC 4449. A Faraday rotation measure study of extragalactic polarized sources behind the LMC carried out by \\cite{gaensler2005} suggests that the LMC hosts a coherent axisymmetric magnetic field of strength $\\sim$ 1$\\mu$G. The random component dominates over the ordered component with a strength of $\\sim$ 3 $\\mu$G. It is believed that close encounters between the Magellanic Clouds and the Milky Way have triggered episodes of star formation in the LMC over the past 4 billion years \\citep{bekki2005}. Any coherent field built up by the standard dynamo would have been disrupted by the outflow from active star forming regions. Hence, the existence of a coherent field in the LMC suggests that a field generation mechanism with a much faster amplification time scale is at work.  A much more efficient process, the \\cite{parker1992} dynamo, could account for the large-scale magnetic fields detected in irregular galaxies such as NGC 4449 and the LMC \\citep{hanasz2004}. Vertical pressure from cosmic rays can force magnetic field lines into the galactic halo and form loops which reconnect and then are amplified by the $\\omega$ effect. This process can significantly increase the $\\alpha$ effect and can operate over a much shorter amplification time scale than the $\\alpha-\\omega$ dynamo \\citep{hanasz2004}.  As a close neighbor of the Milky Way, the large angular extent of the Small Magellanic Cloud (SMC) on the sky allows us to determine RMs of  polarized background radio sources whose projections lie behind it. Since RM is an integral of the line of sight magnetic field strength weighted by the thermal electron density, only a coherent field can produce a consistent sign in RM as a function of position on the sky. An RM study can distinguish between coherent and incoherent fields and can indicate the geometry of any coherent field, and therefore can potentially reveal the field generation mechanism.  The alignment of non-spherical dust grains with the magnetic field in the ISM linearly polarizes optical radiation that travels through it. Therefore, measuring the optical polarization of stars in the SMC enables us to estimate the orientation of the ordered magnetic field in the plane of the sky. The spread in polarization position angle of an ensemble of starlight polarization measurements allows one to estimate the mean strength of the ordered component of the magnetic field using the Chandrasekhar-Fermi (\\citeyear{chandrasekhar1953}), or C-F method. Assuming that the field is unidirectional and knowing both the line-of-sight and the plane-of-the-sky magnetic field strength and orientation, one can construct a three dimensional magnetic field vector for the SMC, which can further constrain the field generation mechanism.  In this paper, we present the results of a radio and optical polarization study of the magnetic field in the SMC. We use RMs of polarized extragalactic background radio sources to determine the magnetic field strength and direction along the line of sight. The orientation and strength of the plane-of-the-sky magnetic field is studied using optical polarization of stars in the SMC. We start in \\S~\\ref{subsection:smc} by reviewing the properties of the SMC and we summarize previous studies on SMC's magnetism in \\S~\\ref{subsection:previousstudies}. In \\S~\\ref{subsection:faradayrotation} and \\S~\\ref{subsection:osp}, we summarize the physics behind the RM method and the C-F method respectively. We then describe the observation, data reduction procedures and present results in \\S~\\ref{section:obsdataresults}. We derive the line-of-sight magnetic field of the SMC in \\S~\\ref{section:wimbfield}. In \\S~\\ref{section:posb}, we estimate the plane-of-the-sky magnetic field and the random field strength in the SMC. A 3D magnetic field vector of the SMC is constructed in \\S~\\ref{section:3Dfield}. A discussion of possible field generation mechanisms is provided in \\S~\\ref{section:discussion}.  Throughout this paper, we represent physical quantities in the plane of the sky by the subscript $\\perp$ and those along the line of sight by the subscript $\\parallel$. We denote the average of a quantity $x$ over the plane of the sky and that averaged along the line of sight by $\\langle x \\rangle$ and $\\bar{x}$ respectively. Table~\\ref{table:glossary} contains a glossary of the variables used in this paper.   \\subsection{The Small Magellanic Cloud} \\label{subsection:smc}    The Small Magellanic Cloud is a nearby gas rich dwarf irregular galaxy. Recent precise measurements of apparent magnitudes of stars at the tip of the red giant branch in the SMC yield a distance modulus of 18.99$\\pm$ 0.03(formal)$\\pm$0.08(systematic)  \\citep{cioni2000}, which corresponds to a distance of roughly 63 $\\pm$ 1 kpc. In this paper, we adopt a distance to the SMC of 60 kpc. Basic parameters of the SMC are listed in Table~\\ref{table:smcproperties}. Both the SMC and the LMC are thought to be satellite galaxies of the Milky Way. However, a recent study by \\cite{besla2007} suggests that the Clouds are not bound to the Milky Way but are on their first passage about the Galaxy. It is still of great debate as to whether the Magellanic Clouds formed as a binary, or whether they became dynamically coupled to each other $\\sim$ 4 Gyrs ago \\citep{bekki2005}. The most recent proper motion measurements of the Clouds suggest that both scenarios are equally probable \\citep{kallivayalil2006,piatek2007}. It is believed that the last close encounter of the Magellanic Clouds $\\sim$ 0.2 Gyrs ago triggered star formation in the SMC and created the morphological and kinematic features seen in the present day SMC \\citep{yoshizawa2003}.   \\cite{stanimirovic2004} found that the depth of the SMC is within its tidal radius ($\\sim$ 4 $-$ 9 kpc). \\cite{lah2005} have measured distances to pulsating red giants in the SMC and found a distance scatter of 3.2 $\\pm$ 1.6 kpc, which agrees with the results of \\cite{stanimirovic2004}. N-body simulations of the gravitational interaction between the LMC, SMC and the Milky Way have been able to reproduce the large line-of-sight extent of the SMC and its two tidal arms \\citep{gardiner1994,yoshizawa2003}.  The gas component in the SMC shows signs of rotation whereas the old stellar component does not \\citep{hatzidimitriou1997}. The gas kinematics of the SMC were investigated by \\cite{stanimirovic2004}, who found a strong velocity gradient in HI across the SMC from the southwest to the northeast. \\cite{stanimirovic2004} constructed a rotation curve of the gas disk, and derived a maximum rotation velocity of  50 km s$^{-1}$.  \\subsection{Previous Studies of Magnetism in the SMC} \\label{subsection:previousstudies} The most common way to study magnetic fields in external galaxies is by observing synchrotron emission at radio wavelengths. \\cite{haynes1986} examined linear polarization maps of the SMC at 1.4 GHz. They found, without any Faraday rotation correction, an ordered magnetic field directed along the SMC's bar in the plane of the sky. \\cite{loiseau1987} analyzed radio continuum maps of the SMC at  408 MHz, 1.4 GHz and 2.3 GHz and obtained a total equipartition field strength of $\\sim$ 5 $\\mu$G by using an average non-thermal spectral index $\\alpha$ of 0.87 (specific intensity of synchrotron emission  $I_\\nu\\propto\\nu^{-\\alpha}$ ) and a depth of the synchrotron emitting region of 6 kpc. \\cite{haynes1990} observed the SMC at 2.5, 4.8 and 8.6 GHz and concluded that the SMC has a large-scale magnetic field, since weak polarized emission is detected across the whole SMC body.  \\cite{chi1993} measured the $\\gamma$-ray flux from the SMC to determine the field strength from radio synchrotron emission without needing to invoke the equipartition assumption. They obtained an estimate which exceeded the equipartition value and concluded that energy equipartition is not valid in the SMC. \\cite{pohl1993}, however, took the energy density in cosmic ray electrons into account, and demonstrated that energy equipartition between magnetic field and cosmic rays is not necessarily violated.  One should note that calculating the total equipartition field requires knowledge of the depth of the synchrotron emitting layer of the SMC and the inclination of magnetic field with respect to the plane of the sky, which are both poorly constrained in the SMC. Also, as explained by \\cite{beck2005}, the classical equipartition energy formula underestimates the true equipartition field strength, since the former involves  integrating the radio spectrum with a fixed frequency interval instead of a fixed energy interval, and with insufficient knowledge of the ratio of the total energy density of cosmic ray nuclei to that of the electrons and positrons. We will further explore this issue in \\S~\\ref{subsection:equipartition}.    Optical polarization from stars in the SMC can be used to map the geometry of the plane-of-the-sky component of the magnetic field, assuming that the observed polarization is due to scattering by non-spherical foreground dust grains aligned by the local magnetic field. Polarization measurements for 147 SMC stars have been made by \\cite{mathewson1970a,mathewson1970b}, \\cite{schmidt1970,schmidt1976}, and \\cite{magalhaes1990}. Since the polarization ``vectors\" \\footnotemark[1] \\footnotetext[1]{Position angles provide information on the orientation of the polarization plane, but not the direction. Hence, there is a 180$^\\circ$ direction ambiguity.}, after removal of Galactic foreground  polarization,  appear to run parallel to the direction connecting the Magellanic Clouds (Figure~\\ref{fig:starlight}), the ``Pan-Magellanic\" magnetic field hypothesis emerged, which suggests the existence of a large scale magnetic field associated with the entire Magellanic system. \\citeauthor{wayte1990} 's  (\\citeyear{wayte1990}) reanalysis of previously obtained starlight polarization data sets appeared to support the idea of this Pan-Magellanic magnetic field. However, as the Galactic foreground polarization also runs along the projection of the line joining the Magellanic Clouds \\citep{schmidt1970}, any contribution from Galactic foreground that has not been correctly subtracted could be misinterpreted as an intrinsic magnetic field connecting the Magellanic Clouds. In addition, anisotropic scattering in the ISM may also polarize starlight \\citep{widrow2002}. Therefore, one has to be cautious interpreting these results. In \\S~\\ref{subsection:starlightcf}, we will further analyze the optical starlight polarization data using the C-F method to derive the ordered magnetic field strength of the SMC in the plane of the sky.   \\subsection{Faraday Rotation } \\label{subsection:faradayrotation} When linearly polarized light travels through a magnetized plasma, the plane of polarization rotates due to birefringence. The change in the polarization position angle $\\Delta$$\\phi$ in radians is given by \\begin{equation} \\Delta\\phi = {\\rm RM}  \\lambda^{2} \\end{equation} where $\\lambda$ is the wavelength of the radiation measured in meters and RM is the rotation measure, defined by \\begin{equation} {\\rm RM} =0.812 \\int ^{observer}_{source} {n_{e}(l){B_{\\parallel} (l)}} dl~~~\\rm{rad~m^{-2}} \\label{eq:rmdef} \\end{equation} In the above equation, $n_e(l)$ (in cm$^{-3}$) is the thermal electron density, $B_{\\parallel}(l)$ (in $\\mu$G) is the line of sight magnetic field strength and d$\\it{l}$ (in pc) is a line element along the line of sight. The sign of the RM gives the direction of the line of sight component of the average field. For example, a negative RM represents a field whose line of sight component is directed away from us.  RMs for extragalactic radio sources behind the SMC can be decomposed into various contributions along the line of sight: the intrinsic RM of the source, the RM through the intergalactic medium (IGM), the RM through the SMC, and the foreground Milky Way RM. \\begin{equation} \\rm{RM_{observed}}=\\rm{RM_{Intrinsic}+RM_{IGM}+RM_{SMC}+RM_{Milky Way}} \\label{eq:rmsum} \\end{equation}    $\\rm{RM_{Milky Way}}$ can be estimated by observing RMs of extragalactic sources whose projections on the sky lie outside, but close to the SMC. RMs of extragalactic sources at Galactic latitudes $|b| > 30^\\circ$ have a standard deviation $\\sim$ 10 rad m$^{-2}$  \\citep{johnstonhollitt2004b}. In addition, \\cite{broten1988}  showed that the extragalactic RMs in the neighborhood of the SMC have ${\\rm |RM_{Intrinsic}+RM_{IGM}+RM_{Milky Way}|}$ $\\le$ 25 rad m$^{-2}$. This implies that the intrinsic RM and the RM through the IGM are both small compared to the statistical errors of our RM measurements (See Table~\\ref{table:rawrms}). After the removal of the Galactic foreground, the observed RM should adequately represent the RM through the SMC.  Ionospheric Faraday rotation may also contaminate our data. However, since the magnitude of RM induced by the ionosphere is typically only  $\\sim$ 1 rad m$^{-2}$ \\citep{tinbergen1996}, this is not of great concern in our experiment as the statistical errors of our RM measurements greatly exceed this value (See Table~\\ref{table:rawrms}).   Faraday rotation is complementary to other measurement techniques such as equipartition, synchrotron intensity and starlight polarization since RMs provide the direction of the magnetic field (and hence the field coherency), while other techniques only provide the field orientation and estimates of the field strength, but not its direction. With independent knowledge of the thermal electron density and the line of sight depth of the SMC, one can estimate the average line-of-sight magnetic field strength using Equation~(\\ref{eq:rmdef}) assuming that there is no correlation between electron density and magnetic field on small scales. If such correlation or anti-correlation exists, it will result in either underestimation or overestimation of the field strength by a factor of up to two to three \\citep{beck2003}.  \\subsection{Optical Starlight Polarization and the Chandrasekhar-Fermi Method} \\label{subsection:osp}  Starlight polarization alone does not directly give the magnetic field strength. However, measuring the spread in polarization position angles for an ensemble of stars allows one to estimate the mean strength of the ordered component of the magnetic field \\citep{chandrasekhar1953}. This technique assumes that the magnetic field is frozen into the gas and that turbulence leads to isotropic fluctuation of the magnetic field around the mean field direction \\citep{heitsch2001,sandstrom2001}. Assuming equipartition between the turbulent kinetic and the magnetic energy, the ordered magnetic field strength averaged over the plane of the sky, $\\langle B_{o,\\perp}  \\rangle$, is given by \\citep{heitsch2001}: \\begin{equation} \\label{eq:cfmethod} \\langle B_{o,\\perp} \\rangle ^2  = 4 \\pi \\rho \\frac { \\sigma_{v_{los}}^2} { \\sigma (\\rm {tan}~ \\delta_p )^2 } \\end{equation} where $\\rho$ is the density of the medium, $\\theta_{p}$ is the measured polarization position angle, $\\langle \\theta_{p} \\rangle$ is the weighted mean of the measured position angles, $\\delta_{p} \\equiv {\\theta_{p}} - \\langle \\theta_{p} \\rangle$, and $\\sigma_{v_{los}}$ is the dispersion of the line-of-sight velocity in the medium. ",
        "conclusions": "\\label{section:conclusions} We have measured the Faraday rotation of extragalactic polarized sources behind the Small Magellanic Cloud to determine the SMC's magnetic field strength and geometry. Our study reveals that the SMC has a galactic-scale field of  0.19 $\\pm$ 0.06 $\\mu$G directed coherently away from us along the line of sight. Optical polarization data on stars in the SMC are re-analyzed using the Chandradsekhar-Fermi method to give an ordered component of the magnetic field in the plane of the sky, of strength 1.6 $\\pm$ 0.4 $\\mu$G. Under the assumption that the Faraday rotation measures and optical starlight polarization probe the same underlying large scale field in the SMC, we have constructed a 3D magnetic field vector of the SMC. It is found that this magnetic field vector possibly aligns with the Magellanic Bridge. This potential alignment needs to be verified by future studies of RMs towards extragalactic sources behind the Magellanic Bridge. The random magnetic field strength in the SMC derived from RM data alone and that derived by combining the results of the C-F method with equipartition were found to be in agreement ($\\sim$ 3 $\\mu$G). This implies that our underlying assumption, that these 3 independent methods probe different components of the same large scale field, is self-consistent.   The SMC is a slowly rotating galaxy, for which the standard mean-field dynamo is not expected to be at work because of the subcritical dynamo number. The cosmic-ray driven dynamo has a short enough amplification time scale to explain the observed coherent field. With modifications by tidal interactions, the field generated by the cosmic-ray driven dynamo could potentially be aligned with the Magellanic Bridge. However, this model faces difficulties in explaining the observed uni-directional field lines. Therefore, the origin of the magnetic field in the SMC is still an open question which needs to be followed up with more observations. The relatively small number of background rotation measures makes it difficult to interpret the observed RMs in detail. Future observations of the SMC with the Square Kilometre Array will provide  $\\sim$ 10$^5$ RMs in a field of 40 square degrees surrounding the SMC \\citep{beck2004b}; with which different possible origins of the magnetic field in the SMC can be fully evaluated \\citep{stepanov2008}.  \\textbf{Acknowledgements} We thank Joseph Gelfand for carrying out the ATCA observations, Anvar Shukurov, Erik Muller, Alyssa Goodman and Douglas Finkbeiner for useful discussions, and Rainer Beck, Marita Krause and Ellie Berkhuijsen for their help and hospitality during S.~A.~M. 's visit to the Max Planck Institute for Radio Astronomy. M.~H. acknowledges support from the National Radio Astronomy Observatory (NRAO), which is operated by Associated Universities Inc., under cooperative agreement with the National Science Foundation. This research was supported by the National Science Foundation through grant AST-0307358 to Harvard College Observatory. The Australia Telescope Compact Array is part of the Australian Telescope, which is funded by the Commonwealth of Australia for operation as a National Facility managed by CSIRO. The Southern H-Alpha Sky Survey Atlas (SHASSA) is supported by the NSF.      {\\it Facilities:} ATCA  \\begin{appendix}"
    },
    "0807/0807.1018_arXiv.txt": {
        "abstract": "We report on \\integral\\ observations of the bright black-hole transient \\mysou\\ performed during the period August-September 2004, including the fast transition (10 hr) observed simultaneously with \\integral\\ and \\rxte\\ on August 15th and previously reported. Our data cover three different spectral states, namely Hard/Intermediate State, Soft/Intermediate State and High/Soft State. We investigate the spectral variability of the source across the different spectral states. The hard X-ray spectrum becomes softer during the HIMS-to-SIMS transition, but it hardens when reaching the HSS state. A principal component analysis demonstrates that most of the variability occurs through two independent modes: a pivoting of the spectrum around 6 keV (responsible for 75\\% of the variance) and an intensity variation of the hard component (responsible for 21\\%). The pivoting is interpreted as due to changes in the soft cooling photon flux entering the corona, the second mode as fluctuations of the heating rate in the corona. These results are very similar to those previously obtained for Cygnus X-1. Our spectral analysis  of the spectra of \\mysou\\  shows a high energy excess with respect to pure thermal Comptonisation models in the HIMS: a non-thermal power-law component seems to be requested by data. In all spectral states joint IBIS, SPI and JEM-X data are well represented by hybrid thermal/non-thermal Comptonisation ({\\sc eqpair}). These fits allow us to track the evolution of each spectral components during the spectral transition. The spectral evolution seems to be predominantly driven by a reduction of  the ratio of the electron heating rate to the soft cooling photon flux in the corona, $l_h/l_s$. The inferred accretion disc soft thermal emission increases by about two orders of magnitude, while the Comptonised luminosity decreases by at most a factor of 3. This confirms that the softening we observed is due to a major increase in the flux of soft cooling photons in the corona associated with a modest reduction of the electron heating rate. ",
        "introduction": "Black hole candidates (BHCs) are known to show different spectral states in the X and $\\gamma$-ray domain. The two main spectral states are the Low/Hard State (LHS), with the high energy spectrum described by a cut-off power-law (typically $\\Gamma$ $\\sim$1.5 and $E_{cut} \\sim$100 keV), and the High Soft State (HSS) with a thermal component peaking at few keV and the high energy power-law much softer ($\\Gamma > 2.2$) \\cite{zdz00}. Usually, this spectral variability is interpreted as due to changes in the geometry of the central parts of the accretion flow (see Done, Gierli\\'nski \\& Kubota 2007 for a recent review). In the LHS the standard geometrically thin and optically thick disc \\cite{ss73} would be truncated far away from the last stable orbit. In the innermost parts, a hot accretion flow is responsible for the high energy emission, via thermal Comptonisation of the soft photons coming from the truncated disc. In the HSS, the optically thick disc would extend close to the minimum stable orbit producing the dominant thermal component. The weak non-thermal emission at higher energy is believed to be due to up-scattering of the soft thermal disc emission in active coronal regions above the disc \\cite{zdzg04}.  However, recently long \\xmm\\ observations of \\mysou\\ have shown that a standard thin accretion disc may remain at the innermost stable circular orbit around the black hole also during the LHS (Miller et al. 2006, Reis et al. 2008)  Since its discovery \\cite{markert73}, the X-ray binary \\mysou\\ has been thoroughly  studied at all wavelength from radio to gamma-rays. Its star companion is still unknown, even though upper limit on the optical luminosity allowed to classify the source as Low Mass X-ray Binary (LMXB; Shahbaz, Fender, \\& Charles 2001). Classified as BHC \\cite{zdz98}, \\mysou\\ is a transient source spending long periods in outburst. Hynes et al. (2003) estimated a mass function of 5.8 $\\pm$ 0.5 $M_{\\odot}$ and Zdziarski et al. (2004) derived a lower limit on the distance of 7 kpc. Before the launch of \\rxte\\ the source had been observed mostly in the Low/Hard spectral state (LHS), even though few spectral transitions to the Very High state (VHS) were reported before \\cite{miya91}. Thereafter \\mysou\\ remained bright and mostly in the LHS until 1999 when it went into quiescence. After the quiescent state observed by \\beppo\\ (Kong et al. 2000, Corbel et al. 2003), \\mysou\\ showed two new outbursts: in 2002/2003 (Smith et al. 2002a,  Belloni et al. 2005) and in 2004 after one year in quiescence (Buxton et al. 2004, Belloni et al. 2004). The long-term variability of \\mysou\\ (1987-2004) is extensively presented in Zdziarski et al. (2004).  The spectral evolution of the outburst of a black hole in low-mass system is often investigated by plotting the flux of the source as a function of the X-ray hardness. In this Hardness Intensity Diagram (HID), \\mysou\\ usually follows a {\\it q}-like pattern throughout the outburst (Homan \\& Belloni 2005; Belloni 2005). The HSS and LHS correspond respectively to the left and right hand side vertical branches of the {\\it q}. State transitions occur when the source crosses the upper or lower horizontal branches, those constitute intermediate states between the LHS and HSS. Depending on the location of the source on a horizontal branch, Belloni et al. (2005) have identified two different flavours of intermediate state: the Hard Intermediate State (HIMS) and Soft Intermediate State (SIMS). An alternative states classification defines the upper horizontal branch as VHS (see Tanaka \\& Lewin 1995; van der Klis 1995; McClintock \\& Remillard 2006).  During a typical outburst the source starts in a faint LHS, move upward along the LHS branch, then move leftward towards the HSS along the VHS branch, and when luminosity decreases the source moves down along the HSS before transiting right back to the LHS along the lower horizontal branch, and then back to quiescence in the LHS. A puzzling property of this evolution is that the spectral changes in \\mysou\\  lag the variations of the luminosity: hard-to-soft state transitions during the rising phase occur at higher luminosities than the soft-to-hard ones during the declining phase. The so-called ``hysteresis'' \\cite{zdzg04} is also observed in other LMXBs such as GRS 1758--258 and 1E 1740.7--2942 (Smith et al. 2002b; Del Santo et al. 2005).   The aim of this work is to use \\integral\\ \\cite{winkler03} data collected during the softening of  the 2004 outburst to study the broad-band spectral evolution of \\mysou\\ during a hard-to-soft state transition. Some of these data obtained on 2004 August 14$^{th}$-16$^{th}$, right across the HIMS-SIMS  fast transition, are presented in Belloni et al. (2006). This transition, marked by the disappearing of type-C QPO and by the appearance of a type-B QPO (Nespoli et al. 2003), is also associated to the ejection of fast relativistic jets, which has led to the identification of a 'jet-line' in the HID (Fender, Belloni \\& Gallo 2004). Moreover, focussing on the high energy part of the spectrum, Belloni et al. (2006) reported on the disappearing of the high energy cut-off (measured at 70 keV in the HIMS) in the SIMS. Recently, a similar evolution of such cut-off has been found for the BHC GRO J1655--40 \\cite{joinet08}. These authors propose that the cut-off increased significantly or vanished completely as the radio-jet turned off in the SIMS.  Here, we extend the work preliminarily presented in Del Santo et al. (2006) by studying the spectral evolution of \\mysou\\ over a longer period of time covering almost the whole span of the LHS-HSS transition (from 2004 August 9 to 2004 September 11). In this work, different Comptonisation models for spectral fitting have been used; furthermore, a principal component analysis provided the variability modes of \\mysou. ",
        "conclusions": "We have presented \\integral~ observations of \\mysou\\ which covered part of the 2004 outburst. We followed the source spectral evolution through the Hard-Intermediate and Soft-Intermediate states, until the High-Soft state.  During the HIMS-to-SIMS transition we found that the high energy cut-off disappears, in agreement with Belloni et al. (2006). Furthermore, the hard X-ray spectrum appears softer during the transition states than in the HSS.  Our detailed spectral analysis of time averaged spectra at different stages of the transition confirms that the spectral transition is driven by changes in the soft cooling photon flux in the corona associated with an increase of disc temperatures (leading to dramatic increase of disc luminosity). The measured disc temperature versus luminosity relation suggests that the internal disc radius decreases. In contrasts, the heating rate of the electrons in the corona appears to remain nearly constant. Although other models such as dynamic accretion disc corona models cannot be ruled out, these results are consistent with the so-called truncated disc model (Done et al. 2007).  In all spectra, including those in of HIMS, we found a significant contribution from a non-thermal component. This component appears as a high energy excess above the pure thermal Comptonisation spectrum. We associated this component with the presence of a non-thermal tail in the distribution of the Comptonising electron.  Finally we analysed the variability of the source during the state transition with a PCA. The  PCA shows that the variability of \\mysou~ can be described as produced through two independent modes: \\begin{itemize} \\item{spectral evolution with the spectrum pivoting around 6 keV (PC1) contributing for 75\\% of the observed variance. The pivoting mode may be interpreted as caused by changes in the soft cooling photons flux in the hot Comptonising plasma associated with an increase of temperature of the accretion disc. These changes are well in agreement with the variation in the flux of the soft cooling photons found by fitting averaged spectra with {\\sc eqpair}.} \\item{intensity variation of the hard power-law component (almost constant slope) on top of constant soft component (PC2) contributing at 21\\%, also consistent with the variations of either the compactness ratio $l_{h}/l_{s}$ or the hard X-ray luminosity, both found by the spectral fitting. This variability mode could be due to magnetic flares or local instabilities in the corona.} \\end{itemize}"
    },
    "0807/0807.4121_arXiv.txt": {
        "abstract": "We present interferometric and single-dish mid-infrared observations of a sample of massive young stellar objects (BN-type objects), using VLTI-MIDI (10\\,\\micron) and Subaru-COMICS (24.5\\,\\micron). We discuss the regions S140, Mon R2, M8E-IR, and W33A. The observations probe the inner regions of the dusty envelope at scales of 50 milli arcsecond and 0.6\\arcsec ($\\sim$100-1000\\,AU), respectively. Simultaneous model fits to spectral energy distributions and spatial data are achieved using self-consistent spherical envelope modelling. We conclude that those MYSO envelopes that are best described by a spherical geometry, the commensurate density distribution is a powerlaw with index $-1.0$. Such a powerlaw is predicted if the envelope is supported by turbulence on the 100-1000\\,AU scales probed with MIDI and COMICS, but the r\\^{o}le of rotation at these spatial scales need testing. ",
        "introduction": "Massive stars ($>8 M_{\\odot}$) play a major role in the dynamical, thermal and chemical evolution of galaxies. The question of how massive stars form is therefore central for a proper understanding of these processes. Currently there is no consensus regarding the answer to this question, and two diverging viewpoints regarding the initial formation stages are generally referred to. On the one hand, a high-mass star is formed from the monolithic gravitational collapse of a massive molecular condensation (core). The principal stellar product is a single or binary massive star that is formed from gravitationally bound core material \\cite{2003ApJ...585..850M} and accreted onto the central star through a disk. The competing viewpoint is due to \\cite{1997MNRAS.285..201B} in which a massive core fragments into a collection of equal mass stellar seeds that compete via Bondi-Hoyle accretion for inflowing material. The seed that ends up being a massive star happened to be located in a fortuitous spatial position (near the centre of the potential well) where most of the inflowing material gathered. The amount of observational evidence \\cite{2007ApJ...654..304K,2007A&A...476.1243M,2007A&A...462L..17A,2006MNRAS.367..737B} and theoretical considerations \\cite{2004MNRAS.349..678E,2006ApJ...641L..45K} that favour the monolithic collapse scenario is growing. Recent developments in massive star formation are reviewed in \\cite{2007ARA&A..45..481Z,2007ARA&A..45..565M}.   Whichever scenario dictates massive star formation (SF), they both start out with a massive and cold molecular core that is gravitationally unstable. Hidden from direct view by the surrounding dusty envelope, a massive star forms at the centre of the cold core and its radiation is reprocessed to longer wavelengths before it escapes the structure. Observations in the infrared (IR) and (sub)mm of star forming regions in the solar neighbourhood have identified cold objects with bolometric luminosities equalling that of a (single) massive star, i.e. a few times $\\rm 10^{4}\\,L_{\\odot}$. This luminosity is such that the central object has the ability to ionise its surroundings if due to a single main sequence star, yet only little (if any) recombination emission is observed from these sources. These objects are identified with the massive young stellar object (MYSO) evolutionary phase of SF in which the formation of a central massive star is nearing its completion but the central object still experiences mass accretion. The latter is evidenced by the ubiquitous bipolar outflow signatures associated with these star forming regions \\cite{2005ccsf.conf..105B}. The star itself, although unmistakenly present, is not observable as the surrounding dusty envelope is highly optically thick ($A_{v}$ up to 50 to 100).  However crucial clues to the formation process can be conveyed by the radial structure of the dust envelope. It is determined by the forces that operate during the onset and the subsequent evolution of the initial molecular core. For example, the radial density profile is predicted to have a powerlaw  with a value for the power index that depends on the dominant physics. The envelope emission is a direct consequence of the dust radial density distribution and the exact power index can be extracted from the observables by the use of radiative transfer models.  Spherical envelope geometries can be assumed for the dust emission dominating at wavelengths larger than $30\\mu$m, where the SEDs of MYSOs are observed to be remarkably similar, arguing for little deviation from spherical symmetry \\cite{1986A&A...167..315C,1990ApJ...354..247C,1991MNRAS.251..584H,1991A&A...252..801G,1994ApJ...427..889W,1998ApJ...500..280F,2000A&A...357..637H,2000ApJ...537..283V,2002ApJS..143..469M,2002ApJ...566..945B}. Emission of MYSOs at shorter wavelengths poses more serious interpretation problems.  Here we present an investigation into the structure of dusty MYSO envelopes at intermediate and high angular resolution in the mid-IR. We discuss Subaru-COMICS observations at 24.5\\,\\micron~with an angular resolution of 0.6\\arcsec~of a sample of 14 MYSOs, and VLTI-MIDI observation at 10\\,\\micron~of 1 MYSO with a resolution of 50 milli arcseconds. ",
        "conclusions": "High angular resolution mid-IR observations are starting to give a clear and detailed view of how the circumstellar material of a MYSO is organised. In this contribution we have presented an analysis of circumstellar emission of MYSOs in the mid-IR at two different wavelengths (24.5\\,\\micron~and 10\\,\\micron) and two different resolutions (0.6\\arcsec~and 0.05\\arcsec). We have found that when the emission at these wavelengths is dominated by dust emission from the envelope the radial density profiles can be described by a powerlaw of the form $\\rho=\\rho_{0}\\,r^{-n}$ with a power index of approximately $n=1.0$. Steep powerlaws $n\\geq1.5$ are excluded as they require high optical depth denied by the silicate absorption profile. Steep powerlaws are found but only in those cases where their is evidence for a contribution by the outflow cavity to the mid-IR emission. We note that our results for the radial density distribution does not negate the findings of previous studies conducted in the (sub)mm wavelength region. For example \\cite{2002ApJS..143..469M} claim a value of of $1.8\\pm0.4$ from analysis for 350\\,\\micron~intensity profiles. However our results seem to be more in line with lower values found by \\cite{2000ApJ...537..283V}. The powerlaw dependence of density with radius is predicted to have a powerindex of $n=1.0$ when the envelope is supported by turbulence \\cite{1989ApJ...342..834L}, although on the scales probed with MIDI rotational flattening may start to play a significant role."
    },
    "0807/0807.4590_arXiv.txt": {
        "abstract": "{ We study the global dynamics of the universe within the framework of the Interacting Dark Matter (IDM) scenario. Assuming that the dark matter obeys the collisional Boltzmann equation, we can derive analytical solutions of the global density evolution, which can accommodate an accelerated expansion, equivalent to either the {\\em quintessence} or the standard $\\Lambda$ models, with the present time located after the inflection point. This is possible if there is a disequilibrium between the DM particle creation and annihilation processes with the former process dominating, which creates an effective source term with negative pressure. Comparing the predicted Hubble expansion of one of the IDM models (the simplest) with observational data we find that the effective annihilation term is quite small, as suggested by a variety of other recent experiments. ",
        "introduction": "Over the past decade the analysis of high quality cosmological data (supernovae type Ia, CMB, galaxy clustering, etc.) have suggested that we live in a flat and accelerating universe, which contains cold dark matter to explain clustering and an extra component with negative pressure, the vacuum energy (or in a more general setting the {\\em dark energy}), to explain the observed accelerated cosmic expansion (Spergel et al. 2007; Davis et al. 2007; Kowalski et al. 2008; Komatsu et al. 2009 and references therein). Due to the absence of a physically well-motivated fundamental theory, there have been many theoretical speculations regarding the nature of the above exotic dark energy (DE) among which a cosmological constant, scalar or vector fields (see Weinberg 1989; Wetterich 1995; Caldwell, Dave \\& Steinhardt 1998; Brax \\& Martin 1999; Peebles \\& Ratra 2003; Perivolaropoulos 2003; Brookfield et al. 2006; Boehmer \\& Harko 2007 and references therein).  Most of the recent papers in this kind of studies are based on the assumption that the DE evolves independently of the dark matter (DM). The unknown nature of both DM and DE implies that we can not preclude future surprises regarding the interactions in the dark sector. This is very important because interactions between the DM and {\\em quintessence} could provide possible solutions to the cosmological coincidence problem (Grande, Pelinson \\& Sol\\'a 2009). Recently, several papers have been published in this area (eg., Amendola et al. 2003; Cai \\& Wang 2005; Binder \\& Kremer 2006; Campo et al. 2006; Wang, Lin \\& Abdalla 2006; Das, Corasaniti, \\& Khoury 2006; Olivares, Atrio-Barandela \\& Pavon; He \\& Wang 2008, and references therein) proposing that the DE and DM could be coupled, assuming also that there is only one type of non-interacting DM.  However, there are other possibilities: (a) It is plausible that the dark matter is self-interacting (IDM) [Spergel \\& Steinhardt 2000], a possibility that has been proposed to solve discrepancies between theoretical predictions and astrophysical observations, among which the less cuspy halo profiles, predicted by the IDM model, allowing for the observed gamma-ray and microwave emission from the center of our galaxy (Flores \\& Primack 1994; Moore et al. 1999; Hooper, Finkbeiner \\& Dobler 2007; Regis \\& Ullio 2008 and references therein) and the discrepancy between the predicted optical depth, $\\tau$, from the Gunn-Peterson test in the spectra of high-z QSOs and the WMAP-based value (eg., Mapelli, Ferrara \\& Pierpaoli 2006; Belikov \\& Hooper 2009; Cirelli, Iocco \\& Panci 2009 and references therein). It has also been shown that some dark matter interactions could provide an accelerated expansion phase of the Universe (Zimdahl et al. 2001; Balakin et al. 2003; Lima, Silva \\& Santos 2008), (b) The DM could potentially contain more than one particle species, for example a mixture of cold and warm or hot dark matter (Farrar \\& Peebles 2004; Gubser \\& Peebles 2004), with or without inter-component interactions.  In this work we are not concerned with the viability of the different such possibilities, nor with the properties of interacting DM models. The single aim of this work is to investigate whether there are repercussions of DM self-interactions for the global dynamics of the universe and specifically whether such models can yield an accelerated phase of the cosmic expansion, without the need of dark energy. Note that we do not ``design'' the fluid interactions to produce the desired accelerated cosmic evolution, as in some previous works (eg., Balakin et al. 2003), but rather we investigate under which circumstances the analytical solution space of the collisional Boltzmann equation, in the expanding Universe, allows for a late accelerated phase of the Universe. ",
        "conclusions": "In this work we investigate the evolution of the global density of the universe in the framework of an interacting DM scenario by solving analytically the collisional Boltzmann equation in an expanding Universe. A disequilibrium between the DM particle creation and annihilation processes, whichever its cause and in favor of the particle creation term, can create an effective source term with negative pressure which, acting like dark energy, provides an accelerated expansion phase of the universe. There are also solutions for which the present time is located after the inflection point. Finally, comparing the observed Hubble function of a few high-redshift elliptical galaxies with that predicted by our simplest IDM model ($m=0$), we find that the effective annihilation term is quite small. In a forthcoming paper we will use a multitude of cosmologically relevant observations to jointly fit the predicted, by our generic IDM model, Hubble relation and thus possibly provide more stringent constraints to the free parameters of the models. We plan also to derive the perturbation growth factor in order to study structure formation within the IDM model."
    },
    "0807/0807.3653_arXiv.txt": {
        "abstract": "We analyse mid-infrared (MIR) spectroscopic properties for 19 ultra-luminous infrared quasars (IR~QSOs) in the local universe based on the spectra from the Infrared Spectrograph on board the {\\it Spitzer Space Telescope}. The MIR properties of IR~QSOs are compared with those of optically-selected Palomar-Green QSOs (PG~QSOs) and ultra-luminous infrared galaxies (ULIRGs). The average MIR spectral features from $\\sim$5 to 30$\\mu$m, including the spectral slopes, 6.2$\\mu$m PAH emission strengths and [NeII] 12.81$\\mu$m luminosities of IR~QSOs, differ from those of PG~QSOs. In contrast, IR~QSOs and ULIRGs have comparable PAH and [NeII] luminosities.  These results are consistent with IR~QSOs being at a transitional stage from ULIRGs to classical QSOs. We also find that the colour index $\\alpha$(30, 15) is a good indicator of the relative contribution of starbursts to AGNs for all QSOs. Correlations between the [NeII] 12.81$\\mu$m and PAH 6.2$\\mu$m luminosities and those between the [NeII], PAH with 60$\\mu$m luminosities for ULIRGs and IR~QSOs indicate that both [NeII] and PAH luminosities are approximate star formation rate indicators for IR~QSOs and starburst-dominated galaxies; the scatters are, however, quite large ($\\sim$ 0.7 to 0.8~dex). Finally the correlation between the EW (PAH 6.2$\\mu$m) and outflow velocities suggests that star formation activities are suppressed by feedback from AGNs and/or supernovae. ",
        "introduction": "Since the discovery of ultra-luminous infrared galaxies (ULIRGs, L$_{\\rm IR}$$>$10$^{12}L_{\\odot}$) by the {\\it Infrared Astronomical Satellite} ({\\it IRAS}) in the 1980's \\citep[e.g.,][]{Houck85}, it is widely accepted that ULIRGs result from strong interactions/mergers between gas-rich disk galaxies. These mergers form elliptical galaxies and ULIRGs are an important intermediate stage in the process during which at least a fraction of ULIRGs manifest as dust-enshrouded QSOs \\citep[see, e.g.,][]{Sanders96, Lonsdale06}. In addition, active galactic nuclei (AGNs) triggered by mergers tend to appear at the final merging stages (e.g., \\citealt{Sanders88}, \\citealt{Zheng99}, \\citealt{Cui01}, \\citealt{Veilleux02} and reference therein).  There is mounting evidence that QSOs with far-infrared (FIR) excess have massive starbursts in their host galaxies. For example, \\citet{CS01} investigated 9 QSOs with FIR excess and found that their host galaxies are tidally interacting or major merger systems with obvious recent star-forming activities. From the detections of mid-infrared (MIR)/FIR\\footnote{In this paper, MIR refers to 5-35 $\\mu$m, FIR 35-350 $\\mu$m, and IR 8-1000$\\mu$m.} emissions for FeLoBALs (Broad Absorption Line QSOs with low-ionisation and iron absorption lines) by Multiband Imaging Photometer on {\\it Spitzer} \\citep[MIPS;][]{Rieke04} on board the {\\it Spitzer Space Telescope} \\citep{Werner04}, \\citet{Farrah07a} find that all their 9 FeLoBALs are extremely infrared (IR) bright, and concluded that these QSOs are in transition from ULIRGs to classical QSOs with ongoing or recent starbursts, because the iron absorption may be from iron ejected during starbursts. \\citet{Hao05} studied 31 QSOs/Seyfert'1s selected from the local ULIRG samples (termed as IR~QSOs for simplicity). By comparing the FIR spectral index of IR~QSOs with those of optically selected Palomar-Green QSOs (PG~QSOs; \\citealt{Schmidt83}), they argued that the FIR excess of IR~QSOs relative to PG~QSOs is from massive starbursts and inferred star formation rate (SFR) in the host galaxies of IR~QSOs. Recently, from studies of z$\\sim$6 QSOs with strong sub-mm emissions, \\citet{Carilli07}, \\citet{WangR07} and \\citet{WangR08} concluded that massive starbursts also exist in their host galaxies. The conclusion is consistent with the results of \\citet{Hao08} that high redshift (sub)mm-loud QSOs follow the same trend for FIR to bolometric luminosities established by low redshift IR~QSOs \\citep{Hao05}. All these studies suggest that there exists a transitional stage, during which both the central black hole and the spheroidal component of QSO hosts grow rapidly in a coeval fashion.  However, there is still a debate about the origin of FIR emission from QSOs, because one cannot firmly exclude the possibility that FIR emissions are from dust tori heated by central AGNs \\citep[for detailed discussions see][]{Haas03}. Moreover, from the molecular gas properties of PG~QSOs with IR excess and comparisons with ULIRGs, \\citet{Evans06} find that the $L_{\\rm IR} / L'_{\\rm CO}$ and $L_{\\rm IR} / L'_{\\rm HCN}$ ratios for PG~QSOs are higher than those of ULIRGs, implying that AGNs contribute significantly to the dust heating and hence to the FIR emission. Therefore, other SFR indicators besides the FIR emission for QSOs are important for further understanding the coeval growth of supermassive black holes and their host galaxies.  Recently, the QUEST (Quasar and ULIRG Evolution Study) group \\citep[see][]{Schweitzer06, Netzer07} reported the detection of polycyclic aromatic hydrocarbon (PAH) emission features in PG~QSOs using the Infrared Spectrograph (IRS) on {\\it Spitzer} \\citep{Houck04}. For 11 out of 26 PG~QSOs PAHs have been clearly detected. Furthermore, the average spectrum of the undetected 15 PG~QSOs also shows PAH features. Since the PAH emissions are closely related to star formation, not to AGNs \\citep[see, e.g.,][]{Shi07}, such detections strongly suggest that star formation occurs widely in QSOs. Their analysis shows that 30\\% or more FIR emission in these PG~QSOs is from starbursts. Furthermore, given that the low-excitation fine-structure emission line [NeII] 12.81$\\mu$m is one of the dominant emission lines of HII regions and that the PAH molecules are easily destroyed by high energy photons from AGNs \\citep[e.g.,][]{Wu07}, [NeII] emission may be an alternative, perhaps even better, tracer of star formation for QSOs (see \\S4).  In this paper we study the MIR spectroscopic properties of IR~QSOs based on {\\it Spitzer} IRS observations, and examine their connections and evolutionary relations to ULIRGs and PG~QSOs. The sample selection, data acquisition and reduction are described in Sect. 2 and 3. The major results and discussions are given in Sect. 4 and 5. Finally we summarise our results in Sect. 6. We adopt cosmological parameters {\\it H}$_{\\rm 0}$=70 km s$^{-1}$ Mpc$^{-1}$, $\\Omega$$_{\\rm m}$=0.3, and $\\Omega$$_{\\rm \\Lambda}$=0.7 throughout this paper. ",
        "conclusions": "By comparing the MIR spectroscopic properties of IR~QSOs, ULIRGs and PG~QSOs, we found that the indicators of relative contributions of starbursts to AGNs (such as colour index $\\alpha$(30, 15) and EW [PAH 6.2$\\mu$m]) for IR QSOs are between those of ULIRGs and PG QSOs. These results are consistent with the findings of \\citet{CS01} and \\citet{Hao05} that (at least some) infrared luminous QSOs (IR~QSOs) are at a transitional stage from ULIRGs to classical QSOs. Below we consider the star formation rates and AGN/star formation feedback in more detail.  \\subsection{SFR determined by the [NeII]12.81$\\mu$m and PAH luminosities}  As we argued, the AGN contribution to [NeII] emission for ULIRGs and IR~QSOs is probably very small, and there is a tight correlation between [NeII]12.81$\\mu$m and PAH 6.2$\\mu$m luminosities (see \\S\\ref{sec:statis}). We examine in more detail how they can be used as approximate SFR indicators for ULIRGs and IR~QSOs.  \\begin{figure*} \\includegraphics[angle=0,scale=.7]{f7.ps} \\caption{Top panel: Luminosities L(60$\\mu$m) vs. L([NeII]12.81$\\mu$m) for IR~QSOs, PG~QSOs (from \\citealt{Schweitzer06}), and ULIRGs with detectable [NeII] and 60$\\mu$m emissions. Bottom panel: Luminosities L(60$\\mu$m) vs. L(PAH 6.2$\\mu$m) for IR~QSOs, PG~QSOs in our sample (green) and from \\citet{Schweitzer06} (black), and ULIRGs with detectable PAH and 60$\\mu$m emissions. The 6.2$\\mu$m PAH fluxes for PG~QSOs from \\citet{Schweitzer06} were estimated by assuming a 7.7/6.2 flux ratio of 4.7 derived from the measurement of the ULIRG NGC~6240 with a buried AGN \\citep{Armus06}. The number of each galaxy type is indicated in the bracket, while the dashed lines represent the linear regressions for the above two correlations. The horizontal bar on the bottom right of each panel indicates the mean error of the 60$\\mu$m luminosity.} \\label{fig_sfr} \\end{figure*}  Fig.~\\ref{fig_sfr} shows the [NeII]12.81$\\mu$m (top panel) and PAH 6.2$\\mu$m (bottom panel) luminosities versus the 60$\\mu$m luminosity for IR~QSOs, PG~QSOs and ULIRGs with firmly detected [NeII]12.81$\\mu$m and PAH 6.2$\\mu$m emissions. A Spearman Rank-order analysis show that both correlations are significant at $>$ 99.99$\\%$ level. The dashed lines in Fig.~\\ref{fig_sfr} represent the least-squares regression fits: \\begin{equation} \\log L_{\\rm NeII} = (0.90\\pm0.06)\\log L_{\\rm 60\\mu m} - (2.22\\pm0.69), \\end{equation} \\begin{equation} \\log L_{\\rm PAH} = (0.78\\pm0.06)\\log L_{\\rm 60\\mu m} - (0.31\\pm0.74). \\end{equation} The fitting formula (2) is consistent with that of \\citet{Ho07} for star-forming galaxies (within the large errors). Note that our sample objects have much higher $60\\mu{\\rm m}$ and [NeII] luminosities than their star-forming galaxies. Thus both [NeII] 12.81$\\mu$m and PAH 6.2$\\mu$m luminosities can be used as approximate SFR indicators not only for normal star-forming galaxies, but also for galaxies with high infrared luminosities, such as ULIRGs and IR~QSOs (see \\citealt{Brandl06} and \\citealt{Farrah07b}).  However, the mean scatters (about 0.7 to 0.8 dex) in the relation of [NeII]12.81$\\mu$m, PAH 6.2$\\mu$m with $60\\mu{\\rm m}$ luminosities are larger than that (about 0.6 dex) of star-forming galaxies with lower infrared luminosity (\\citealt{Ho07}). Comparing Fig.~\\ref{fig_sfr} with Fig.~\\ref{fig_ne2pah}, one can see that the scatter in the relation of [NeII]12.81$\\mu$m vs. PAH 6.2$\\mu$m luminosities (about 0.6 dex) is smaller than that in the relations of [NeII]12.81$\\mu$m, PAH 6.2$\\mu$m luminosities with L(60$\\mu$m). It is also clear that most large scatters are from ULIRGs. This is perhaps not surprising since the range in the 9.7$\\mu$m silicate absorption depth among ULIRGs is quite large (see \\citealt{Armus07}, \\citealt{Spoon07}). In short, the large scatters for ULIRGs seen in the relations may be due to complicated, patchy extinctions among these galaxies in the MIR band. A detailed discussion on extinction for ULIRGs can be found in \\citet{Farrah07b}.  \\subsection{AGN/Star formation feedback in the transitional stage}  \\begin{figure*} \\includegraphics[angle=90,scale=.8]{f8.ps} \\caption{The velocity of outflows vs. the equivalent width of the 6.2$\\mu$m PAH emission feature (EW (PAH 6.2$\\mu$m)) for eleven IR~QSOs (black filled circle) which show obvious H$\\beta$ blueshifts. The Kendalls ($\\tau$) rank correlation analysis gives coefficient of linear regression as $-0.66$ with a significance of 99.5$\\%$. The outflow velocities were derived from the H$\\beta$ emission line profiles \\citep{Zheng02, Lipari05}.} \\label{fig_outflow} \\end{figure*}  One explanation for the observed correlation between spheroidal and black hole mass \\citep[e.g.,][]{Magorrian98, Ferrarese05} is that star formation and the growth of central black holes may be self-regulated: AGNs/star formation can drive nuclear outflows which in turn suppress further cooling and star formation (\\citealt{Silk98}). While the detailed processes are still to be understood, it is now increasingly clear that feedback and outflows play an important role in galaxy formation and evolution.  So far most observational evidence for AGN feedback is from radio observations at the centre of clusters or groups of galaxies \\citep{Batcheldor07}. On the galaxy scale, evidence is still limited. As discussed above, IR~QSOs have high SFRs and accretion rates \\citep{Hao05}, outflow properties in these objects may thus provide hints on the feedback processes on galaxy or group scale.  It is well known that low-ionisation broad absorption line QSOs (loBAL~QSOs) comprise about 15\\% of BAL~QSO population. They are defined as a subclass of BAL~QSOs with an obvious blueshifted absorption in Mg II $\\lambda$$\\lambda$2795,2802 and other low-ionisation species \\citep{Weymann91}. Such absorption troughs arise from resonance absorption by outflowing gas and dust \\citep{Voit93}. In addition, there is a rare class of loBAL~QSOs, showing absorption features from excited iron (termed as FeLoBAL~QSOs). The outflow velocities for most BAL~QSOs span a large range, up to a few times $10^{4}$ \\kms, which may be formed on a scale of $<1\\,{\\rm pc}$ and directly associated with the wind from an accretion disk or molecular torus \\citep{Weymann85}. However, recent spectral analyses based on Keck observations for LoBAL~QSOs or FeLoBAL~QSOs reveal that the outflow velocities of some LoBAL~QSOs range from several hundred to several thousand \\,\\kms\\, and the wind is from regions of a few hundred pc \\citep[e.g.,][]{Ganguly08, Hamann00, deKool02}, which is much larger than the central engine of AGNs, but similar to the size of the nuclear starburst region of ULIRGs \\citep{Downes98}.  On the other hand, \\citet{CS02} studied four loBAL~QSOs at z$<$0.4 (Mrk~231, IRAS~14026$+$4341, IRAS~F07599+6508, PG~1700+518; all four are in our sample) and found that all are ULIRGs with merging signatures. They argued that loBAL~QSOs cannot simply be explained by orientation effects, rather, they are directly related to young systems, still surrounded by gas and dust. It represents a short stage in the early life of a large fraction of QSOs. Moreover, \\citet{Farrah07a} reported the detections of mid/far-infrared emission from 9 FeLoBAL QSOs by MIPS. They found that all of their objects are IR bright with infrared luminosities as high as ULIRGs. Thus all loBAL~QSOs and FeLoBAL QSOs with MIR to FIR information are infrared ultra-luminous, suggesting a link among loBAL~QSOs, FeLoBAL QSOs and IR~QSOs.  For loBAL~QSOs or FeLoBAL QSOs, the velocities of outflows can be measured from blueshifted absorption lines. Unfortunately it is still unclear whether most IR~QSOs are loBAL~QSOs or FeLoBAL QSOs, because not many IR QSOs have been observed in the UV. However, there are measurements for the blueshifts in the permitted optical emission lines for our IR~QSO sample \\citep{Zheng02}. The line profiles often contain two Gaussian components, one broad and one narrow. The outflow velocities are determined by the blueshift of the broad Gaussian component relative to the narrow component of permitted emission lines (H$\\beta$) for IR~QSOs \\citep[see][]{Zheng02}. We assume the blueshifted broad Gaussian component is related to the outflow of clouds \\citep{Leighly01, Batcheldor07}.  Fig.~\\ref{fig_outflow} shows the correlation between EW (PAH 6.2$\\mu$m) and outflow velocities for IR~QSOs. As the outflow velocity increases, the EW (PAH 6.2$\\mu$m) decreases. Since the EW (PAH 6.2$\\mu$m) expresses the relative contribution of star formation to AGN in QSOs (see $\\S$4.2), the correlation shown in Fig.~\\ref{fig_outflow} implies that QSOs with higher outflow velocities have a lower ratio of SFR to accretion rate. It may be because more energetic AGNs and central massive starbursts can drive faster outflows which then suppress or even quench star formation by heating up or expelling the cold gas and dust in the QSO hosts, leading to a lower ratio of SFR to accretion rate. Another possibility is that the higher velocity outflows are more efficient in ejecting dust cocoons of AGNs, leading to AGNs becoming the dominant source compared with starbursts. We should caution, however, that the sample shown in Fig.~\\ref{fig_outflow} is small; a larger sample would allow us to better understand the relative importance of star formation and AGN in feedback and driving outflows."
    },
    "0807/0807.3379_arXiv.txt": {
        "abstract": "\\begin{center}{\\bf Abstract}\\end{center} A new kind of accelerating flat model with no dark energy that is fully dominated by cold dark matter (CDM) is investigated. The number of CDM particles is not conserved and the present accelerating stage is a consequence of the negative pressure describing the irreversible process of gravitational particle creation. A related work involving  accelerating CDM cosmology  has been  discussed before the SNe observations [Lima, Abramo \\& Germano, Phys. Rev. D53, 4287 (1996)]. However, in order to have a transition from a decelerating to an accelerating regime at low redshifts, the matter creation rate proposed here includes a constant term of the order of the Hubble parameter.  In this case, $H_0$ does not need to be small in order to solve the age problem and the transition happens even if the matter creation is negligible during the radiation and part of the matter dominated phase. Therefore, instead of the vacuum dominance at redshifts of the order of a few, the present accelerating stage in this sort of Einstein-de Sitter CDM cosmology is a consequence of the gravitational particle creation process.  As an extra bonus, in the present scenario does not exist the coincidence problem that plagues models with dominance of dark energy.  The model is able to harmonize a CDM picture with  the present age of the universe,  the latest measurements of the Hubble parameter and the Supernovae observations. ",
        "introduction": "A large amount of data relevant to cosmology (involving Supernovae type Ia and cosmic background radiation probes) have provided strong evidence that the observed universe is well described by an accelerating,  flat Friedmann-Robertson-Walker (FRW) model \\cite{Riess,Riess07,CMB,Ages}. However, the substance or mechanism behind the current cosmic acceleration remains unknown and constitutes a challenging problem of modern cosmology.  In relativistic cosmology, an accelerating regime is obtained by assuming the existence of a dark energy component (in addition to cold dark matter), an exotic fluid endowed with negative pressure in order to violate the strong energy condition \\cite{review}. The simplest theoretical representation of dark energy is by means of a cosmological constant $\\Lambda$, which acts on the Einstein field equations (EFE) as an isotropic and homogeneous source with constant equation of state (EoS) $w \\equiv p/\\rho = -1$.  All observational data available so far seems to be in good agreement with the  cosmic concordance model, i.e., a vacuum energy plus cold dark matter ($\\Lambda$CDM) scenario. Nevertheless, $\\Lambda$CDM models are plagued with several problems. For instance, it is very difficult to reconcile the small value required by observations ($\\simeq 10^{-10} \\rm{erg/cm^{3}}$) with estimates from quantum field theories ranging from 50-120 orders of magnitude larger \\cite{weinberg}.   Such problem has inspired many authors to propose alternative candidates in the literature \\cite{Sfield,decaying, decaying1,list1,XM1}, among them: (i) a relic scalar field slowly rolling down its potential, (ii) a $\\Lambda (t)$-term or a decaying vacuum energy density, (iii) the ``X-matter'',  an extra component characterized by equation of state $p_x=\\omega \\rho_x$, where $\\omega$ may be constant or a redshift dependent function, (iv) a Chaplygin-type gas whose equation of state is $p=-A/\\rho^{\\alpha}$, where A and $\\alpha$ are positive parameters.  More recently, some attention has also been paid to a possible interaction between  the dark sector components \\cite{DSI}.  The space parameter of such  models are usually highly degenerated and some of them contain the $\\Lambda$CDM scenario as a particular case. In point of fact, the plethora of possible candidates does  not help to identify  the nature of this mysterious component since there is no compelling  direct evidence yet for dark energy (or its dynamical effects). In other words, the evidence supporting its existence is not strong enough to be considered established beyond doubt (see \\cite{BLB} for a critical discussion).  Roughly speaking, a realistic cosmological scenario should be in agreement with at least four well established observational results, namely: (i) the existence of a dark non-baryonic component as required by the dynamics of galaxies and clusters, the matter power spectrum and other independent probes like the temperature anisotropies of the cosmic microwave background (CMB) from last scattering surface (ii) the late time cosmic acceleration, (iii) the (nearly) flatness of the Universe, and, finally, (iv) a Hubble parameter $H_0 \\approx 72\\;\\mbox{km/s.Mpc}$ with the Universe being older than 12 Gyrs in order to accommodate the oldest observed structures (globular clusters). When confronted with this simple requirements, we see that the CDM or Einstein-de Sitter cosmology is in  clear contradiction with results (ii) and (iv).  Therefore, if one assumes that the dark energy does not exist, the first task is to explain how a flat CDM dominated Universe can accelerate at late times because, potentially, accelerating cosmologies solve the age problem.  In this concern, we recall that the presence of a negative pressure is the key ingredient required to accelerate the expansion. This kind of stress  occurs naturally in many different contexts when  the physical systems depart from a thermodynamic equilibrium states \\cite{LL}. In general, such states are connected with phase transitions (for example, in an overheated van der Waals liquid), and for some systems the existence of negative pressure seems to be inevitable \\cite{sakharov}. In this connection, as first pointed out by Zeldovich \\cite{zeld}, the process of  cosmological particle  creation at the expenses of the gravitational field can phenomenologically be described by a  negative pressure and the associated entropy production. In principle, such an approach is completely different from the one developed  by  Hoyle and Narlikar \\cite{narlikar} adding extra terms to the Einstein-Hilbert action describing the so-called C-field.  In the latter case, the creation phenomenon is explained trough  a process of interchange of energy and momentum between matter itself and the C-field as happens, for instance, in vacuum decaying cosmologies \\cite{decaying}.  The gravitational matter creation processes was investigated from a microscopic viewpoint by Parker and collaborators \\cite{Parker} by considering the Bogoliubov mode-mixing technique in the context of quantum field theory in curved space-time \\cite{BirrellD}. Despite being rigorous and well-motivated, those models were never fully realized,  probably due to the lack of a well-defined  prescription of how matter creation is to be incorporated in the classical EFE.  The consequences of gravitational matter creation have also been macroscopically investigated mainly as a byproduct of bulk viscosity processes near the Planck era as well as during the reheating of  the  inflationary scenarios \\cite{Murphy,Hu}. However, the first self-consistent macroscopic formulation of the matter creation process was  put forward by Prigogine and coworkers \\cite{Prigogine} and somewhat clarified by Calv\\~{a}o, Lima and Waga \\cite{LCW} through a manifestly covariant formulation. It was also shown that matter creation, at the expenses of the gravitational field, can effectively be discussed in the realm of the relativistic nonequilibrium thermodynamics. Later on, it was also demonstrated that the matter creation is an irreversible process completely  different from bulk viscosity mechanism \\cite{LG92} (see also \\cite{SLC02} for a more complete discussion). Several interesting  features of cosmologies with creation of matter and radiation have been investigated by many authors \\cite{ZP2,LGA96,LA99,Susmann94,ZSBP01,freaza02} (see also \\cite{Makler07} for recent studies on this subject).  In comparison to the standard equilibrium equations, the irreversible creation process is described by two new ingredients: a balance equation for the particle number density and a negative pressure term in the stress tensor. Such quantities are related to each other in a very definite way by the second law of thermodynamics \\cite{Prigogine,LCW}.  The {\\it leitmotiv} of this approach is that the matter creation process, at the expense of the gravitational field, can happen only as an irreversible process constrained by the usual requirements of non-equilibrium thermodynamics.  In this context, we are proposing here a new flat cosmological scenario where the cosmic acceleration is powered uniquely by the  creation of cold dark matter particles. As we shall see, the model is consistent with the supernovae type Ia data, and a transition redshift of the order of a few is also obtained. In this extended CDM model, the Hubble parameter does not need to be small in order to solve the age problem and the transition happens even if the matter creation is negligible during the radiation and considerable part of the matter dominated phase. Moreover, the so-called coincidence problem of dark energy models is replaced here by a  gravitational particle creation process at low redshifts. ",
        "conclusions": "In this paper we have proposed a flat cold dark matter cosmology whose late time acceleration is powered by an irreversible creation of CDM particles. In our scenario there is no dark energy, and, as such, the so-called coincidence problem is also absent. It should be stressed that $H_0$ does not need to be small in order to solve the age problem. Further, the transition from a decelerating to an accelerating regime at late times happens even if the matter creation is negligible during the radiation and considerable part of the matter dominated phase (this is equivalent to take $\\beta = 0$ in all the expressions). Therefore, like in flat $\\Lambda$CDM scenarios, there is just one free parameter, and the resulting model provides an excellent fit to the observed dimming of distant type Ia supernovae (see Figs. 4a and 4b).  Note also that the flat model ($\\Omega_m=1$) with creation of CDM particles proposed here can easily be extended  to include negative ($\\Omega_m <1$) and positive ($\\Omega_m >1$) spatial curvatures. The same happens with the inclusion of a small (conserved) baryonic component whose density parameter today is severely constrained by the primordial nucleosynthesis and WMAP results. In this case, the value of the transition redshift as derived in section IV will be slightly modified.  On other hand, the existence of such a model also means that the accelerating expansion does not represent a direct evidence for a non-zero cosmological constant or, more generally, to the existence of dark energy as usually assumed by many authors. Naturally, new constraints on the relevant parameters ($\\gamma$ and $\\beta$) from complementary observations need to be investigated in order to see whether the matter creation model proposed here provides a realistic description of the observed Universe. New bounds on these parameters coming from the background and perturbed equations in the presence of a conserved baryonic component will be discussed in a forthcoming communication.  \\appendix"
    },
    "0807/0807.2239_arXiv.txt": {
        "abstract": "Over the last 15 years, around a hundred very young stars have been observed in the central parsec of our Galaxy. While the presence of young stars forming one or two stellar disks at $\\approx 0.1\\,$pc from the supermassive black hole (SMBH) can be understood through star formation in accretion disks, the origin of the S stars observed a factor of 10 closer to the SMBH has remained a major puzzle. Here we show the S stars to be a natural consequence of dynamical interaction of two stellar disks at larger radii. Due to precession and Kozai interaction, individual stars achieve extremely high eccentricities at random orientation. Stellar binaries on such eccentric orbits are disrupted due to close passages near the SMBH, leaving behind a single S star on a much tighter orbit. The remaining star may be ejected from the vicinity of the SMBH, thus simultaneously providing an explanation for the observed hypervelocity stars in the Milky Way halo. ",
        "introduction": "The central parsec of our Galaxy harbors different groups of young stars orbiting the central SMBH. These include one or two stellar disks at $\\approx 0.1\\,$pc from Sgr A$^*$ made up of 5--6\\,Myr old stars \\citep{pau06}, and a number of stars on highly eccentric orbits even closer, the so-called S stars \\citep{eg97}. With orbital periods as small as 15 years, the S stars provide powerful constraints on the mass and size of the central dark object \\citep{ghe05}. While the origin of the stellar disk(s) observed in the Galactic center can be understood \\citep{lb03,gs03,nc05,ger01,mp03}, the S stars show a puzzling combination of interesting features \\citep{eg97}: They are very young ($< 10$\\,Myr), with less than 0.01\\,pc a factor of 10 closer to the SMBH than the closest young population known (the stellar disks), and move around Sgr A$^*$ on randomly oriented eccentric orbits.   A number of mechanisms to create the S stars have been proposed, but so far none of them was able to account for all their properties \\citep[see also review by][]{a05}: It has been suggested that the S stars are not young but rejuvenated stars or exotic objects, but this is in conflict with their spectral properties, showing they are ordinary main-sequence B-type stars \\citep{ghe05,eis05,mge08}. In situ formation is excluded by the inexplicably high pressure required to overcome the strong tidal forces, and will also lead to a stellar disk \\citep{ghe05}, which cannot account for the random orientation and high eccentricity of the orbits within the stars' lifetime \\citep{gh07}. The latter argument also argues against the suggested scenario of an infalling cluster of young stars \\citep{ger01}, which would moreover require a very high mass \\citep{pmg03}, or an intermediate-mass black hole in its center for it to survive the tidal force of Sgr A$^*$ \\citep{hm03}. Exchange capture with compact remnants \\citep{al04} only works with a large cluster of black holes and fails to explain the youth and high initial eccentricities. The Lense-Thirring effect may account for the innermost star \\citep{lb03}, S2, being dragged out of the disk plane, but cannot explain the orbital orientation of the longer-period S stars, nor the proximity of the orbit itself. Interaction of stars in a central cluster with a fossil accretion disk can lead to high eccentricities and orbital decay \\citep{sk05} if the disk mass is sufficiently high, but fails to explain the observed bias toward young B-type stars. Furthermore, such an accretion disk is expected to fragment and form stars and should thus have been observed.  Tidal disruption of binaries falling into the Galactic center on highly eccentric orbits has been shown to be an effective mechanism to create S stars, but until now required very massive star clusters passing close to Sgr A$^*$ \\citep{gq03} or the presence of a large enough population of compact massive perturbers in combination with many young stars on eccentric orbits \\citep{pha06}.  In the following we show how the dynamical interaction of two stellar disks at larger radii, as they are observed in the Galactic center, naturally leads to the creation of S stars. Section \\ref{sec:model} describes the setup of our numerical simulations. In \\S\\ \\ref{sec:high-ecc}, we present the outcome of these simulations, showing how the disk stars can eventually gain very high eccentricities. Section \\ref{sec:disrupt} explains how binary stars on such eccentric orbits around the SMBH get disrupted, leaving behind an S star and possibly a hypervelocity star. In \\S\\ \\ref{sec:discuss} we discuss our results, showing how the mechanism suggested here explains all the observed properties of the S stars. ",
        "conclusions": "\\label{sec:discuss}  The mechanism described above naturally explains the reported mysteries of the S stars: The disruption of stellar binaries as tight as 0.1--1\\,AU can leave a single star orbiting the central SMBH a factor of 10 closer than the original binary. This is just the distance ratio of S stars and stellar disks observed in the center of our Galaxy.  The orientation of the disks rapidly changes due to precession. Furthermore, the change of orientation of a binary's orbit is a natural consequence of the Kozai effect. Once a binary is disrupted and leaves a single star on a much closer orbit, this orbit is decoupled from the originating stellar disk's motion, as both the precession frequency and the Kozai oscillation period are functions of central distance, explaining the widely varying orbital orientations observed.  The eccentricities created for the tightly bound stars after tidal disruption are as high as 0.99. Resonant relaxation \\citep{rt96} and Kozai interactions between the S stars are expected to turn these extremely high eccentricities into an isotropic distribution on a short timescale, just as is observed for the S stars.  While all of the S stars observed so far are main-sequence B-type stars, observations of the stellar disks in the Galactic center mostly revealed OB giants and Wolf-Rayet stars. But even for the unusually flat mass function suggested \\citep{pau06}, one would expect a factor of 5 more B-type main-sequence stars in the disks than O stars. Most of these stars have not been observed yet, mainly due to crowding of the observed region, limited spatial resolution, and strong nebular emission \\citep{pau06}. Our computation assuming equal-mass binaries predicts only one S star above 20\\,$\\msun$ per 18 B-type S stars. The lack of observed O-type S stars is therefore not statistically significant. The number of expected massive S stars would drop further for random pairing of binary component masses, since disruption of a high mass ratio binary will only moderately change the orbit of the massive member. Furthermore, the massive O stars might not form in as tight binaries as are required to produce S stars. Hence, it is legitimate to assume that both S stars and disk stars formed in the same environment, thus providing a solution to the ``paradox of youth''.  In this Letter, we studied the evolution of two isolated stellar disks orbiting a SMBH. In a forthcoming publication, we will analyze the effects of other contributions of a perturbing potential, such as a stellar cusp \\citep[e.g.,][]{sea+07}, a massive circumnuclear disk \\citep{cssy05}, an intermediate-mass black hole and/or a star cluster \\citep[probably IRS 13;][]{mpsr04}. Preliminary results indicate, e.g., that the presence of a cusp of stellar black holes \\citep[e.g.,][]{fak06} does not prevent the dynamical formation of S stars as described in this Letter."
    },
    "0807/0807.2525_arXiv.txt": {
        "abstract": "Our previous analysis indicates that small-scale fluctuations in the intracluster medium (ICM) from cosmological hydrodynamic simulations follow the lognormal probability density function. In order to test the lognormal nature of the ICM directly against X-ray observations of galaxy clusters, we develop a method of extracting statistical information about the three-dimensional properties of the fluctuations from the two-dimensional X-ray surface brightness.  We first create a set of synthetic clusters with lognormal fluctuations around their mean profile given by spherical isothermal $\\beta$ models, later considering polytropic temperature profiles as well.  Performing mock observations of these synthetic clusters, we find that the resulting X-ray surface brightness fluctuations also follow the lognormal distribution fairly well. Systematic analysis of the synthetic clusters provides an empirical relation between the three-dimensional density fluctuations and the two-dimensional X-ray surface brightness.  We analyze \\chandra observations of the galaxy cluster Abell 3667, and find that its X-ray surface brightness fluctuations follow the lognormal distribution.  While the lognormal model was originally motivated by cosmological hydrodynamic simulations, this is the first observational confirmation of the lognormal signature in a real cluster.  Finally we check the synthetic cluster results against clusters from cosmological hydrodynamic simulations.  As a result of the complex structure exhibited by simulated clusters, the empirical relation between the two- and three-dimensional fluctuation properties calibrated with synthetic clusters when applied to simulated clusters shows large scatter.  Nevertheless we are able to reproduce the true value of the fluctuation amplitude of simulated clusters within a factor of two from their two-dimensional X-ray surface brightness alone.  Our current methodology combined with existing observational data is useful in describing and inferring the statistical properties of the three dimensional inhomogeneity in galaxy clusters. ",
        "introduction": "\\label{sec:intro}  Galaxy clusters have been one of the most important probes of cosmology (e.g., Bartlett \\& Silk 1994; Eke, Cole, \\& Frenk 1996; Viana \\& Liddle 1996; Kitayama \\& Suto 1996, 1997; Kitayama, Sasaki, \\& Suto 1998; Holder et al. 2000; Haiman, Mohr, \\& Holder 2001; \\,\\, Majumdar \\& Mohr 2004).  In the context of dark energy surveys, which attracts much of the attention of the cosmology and particle physics communities, galaxy cluster surveys are also unique in that they most directly probe the growth of structure rather than relying solely on distance measurements \\citep[e.g.,][]{albrecht2006}.  In order to capitalize on this, galaxy cluster surveys, in particular those utilizing the Sunyaev-Zel'dovich effect \\citep[for reviews see, for example,][]{carlstrom02, birkinshaw99, rephaeli95, sunyaev80}, are currently operating and many more are planned in the near future.  However, one of the biggest challenges in interpreting these surveys is relating physical quantities of galaxy clusters, namely mass, to observable ones.  In particular, these mass-observable relations may be sensitive to the inherent complex structure of clusters.  Therefore we must better understand galaxy clusters to utilize fully the potential of galaxy cluster surveys in constraining cosmological parameters.   Recent observations of galaxy clusters have revealed a rich variety of structural complexity.  Recent X-ray satellites with their improved angular resolution, collecting area, and simultaneous spectral measurement capabilities have unveiled complex temperature structure \\citep[e.g.,][]{markevitch2000, furusho01}, shock fronts \\citep[e.g.,][]{jones2002}, cold fronts \\citep[e.g.,][]{markevitch2000}, and X-ray holes \\citep[e.g.,][]{fabian2002}.  Improved observational strategies and analysis methods of lensing observations of galaxy clusters show that the mass distribution, as opposed to just the gas, is often complicated as well \\citep[e.g.,][]{bradac06}.  Both X-ray and lensing observations of galaxy clusters reveal that clusters are frequently undergoing mergers \\citep[e.g.,][]{briel04}. With such various and sundry structural complexities may galaxy clusters reliably be used as cosmological probes ?  The complex structure seen in galaxy clusters motivates our investigation of the intracluster medium (ICM) inhomogeneity.  We note, however, that we take a statistical approach to modeling the inhomogeneity rather than directly modeling such complex phenomena as shocks, cold fronts, etc.  Motivated by results from cosmological hydrodynamic simulations we explore the ramifications of a lognormal model of the inhomogeneity of the ICM.  This model was first proposed (Kawahara et al.\\ 2007, hereafter Paper I; Kawahara et al.\\ 2008), in this context, to explain the discrepancies between emission weighted and spectroscopic temperature estimates from galaxy clusters \\citep[][]{mazzotta04, rasia05, vikhlinin06}. They found that local inhomogeneities of the ICM play an essential role in producing the systematic bias between spectroscopic and emission weighted temperatures.  Thus far, the lognormal model has been motivated by and applied only to clusters from cosmological hydrodynamic simulations. Therefore it is crucial to see if inhomogeneities in real galaxy clusters also show the lognormal signature. In reality, this is not a straightforward task since one can observe clusters in X-rays only through their projection over the line of sight. Thus we develop a method of extracting statistical information of the three-dimensional properties of fluctuations from the two-dimensional X-ray surface brightness.  The rest of the paper is organized as follows. We first summarize the log-normal model in \\S\\ref{sec:model}.  We create synthetic clusters to explore the relationship between the intrinsic cluster inhomogeneity and X-ray observables in \\S\\ref{sec:synthetic}.  In \\S \\ref{sec:obs} we apply our methodology to \\chandra observations of the galaxy cluster Abell 3667, and then attempt to quantify the nature of cluster inhomogeneity.  We also compare our synthetic cluster results with cosmological hydrodynamic simulations in \\S\\ref{sec:con}. Finally, we summarize our results in \\S\\ref{sec:sum}. Throughout the paper the Hubble constant is parameterized by $h$ in the usual way, $H_0 = 100\\, h$ km s$^{-1}$ Mpc$^{-1}$. ",
        "conclusions": "\\label{sec:sum}  We have developed a method of extracting statistical information on the ICM inhomogeneity from X-ray observations of galaxy clusters.  With a lognormal model for the fluctuations motivated by cosmological hydrodynamic simulations, we have created synthetic clusters, and have found that their X-ray surface brightness fluctuations retain the lognormal nature.  In addition, the result that $\\SigSx$ and $\\Sign$ are linearly related implies that one can, in principle, estimate the statistical properties of the three dimensional density inhomogeneity ($\\Sign$) from X-ray observations of galaxy clusters ($\\SigSx$ and $\\alphaSx$).  We have compared the predictions of our model to \\chandra X-ray observations of the galaxy cluster A3667.  For the first time in a real galaxy cluster we were able to detect the lognormal signature of X-ray surface brightness fluctuations, which was originally motivated by simulations. Based on the synthetic cluster results, this enabled an estimate of the statistical properties of the inhomogeneity of the ICM of A3667.  In the context of lognormally distributed inhomogeneity, we obtain $\\Sign \\approx 0.4$ for the gas density fluctuations of A3667.  It is encouraging that the value of the fluctuation amplitude for Abell 3667 is in reasonable agreement with typical values from the simulated clusters.  Finally we check the validity and limitation of our method using several clusters from cosmological hydrodynamic simulations. Unlike the fairly idealized synthetic clusters, simulated clusters exhibit complex structure more akin to real galaxy clusters.  As a result, the empirical relation between the two- and three-dimensional fluctuation properties calibrated with synthetic clusters when applied to simulated clusters shows large scatter.  Nevertheless we are able to reproduce the true value of the fluctuation amplitude of simulated clusters within a factor of two from their two-dimensional X-ray surface brightness alone.  Our current methodology combined with existing observational data is useful in describing and inferring the statistical properties of the three dimensional inhomogeneity in galaxy clusters.  The fluctuations in the ICM have several implications in properly interpreting galaxy cluster data. In particular, our current model may be useful in interpreting data from current and future galaxy cluster surveys using the Sunyaev-Zel'dovich effect, which have the potential to provide tight constraints on cosmology."
    },
    "0807/0807.3135_arXiv.txt": {
        "abstract": "In this paper, we study the photon-photon pair production optical depth for gamma-rays with energies from 10 to 200 GeV emitted by powerful blazars due to the diffuse radiation field of broad line region (BLR). There are four key parameters in the BLR model employed to determine the $\\gamma-\\gamma$ attenuation optical depth of these gamma-rays. They are the gamma-ray emitting radius $R_{\\gamma}$, the BLR luminosity $L_{\\rm{BLR}}$, the BLR half thickness $h$ and the ratio $\\tau_{\\rm{BLR}}/f_{\\rm{cov}}$ of the Thomson optical depth to the covering factor of BLR. For FSRQs, on average, it is impossible for gamma-rays with energies from 10 to 200 GeV to escape from the diffuse radiation field of the BLR. If $\\it GLAST$ could detect these gamma-rays for most of FSRQs, the gamma-ray emitting region is likely to be outside the cavity formed by the BLR. Otherwise, the emitting region is likely to be inside the BLR cavity. As examples, we estimate the photon-photon absorption optical depth of gamma-rays with energies from 10 to 200 GeV for two powerful blazars, HFSRQ PKS 0405$-$123 and FSRQ 3C 279. Comparing our results with $\\it GLAST$ observations in the future could test whether the model employed and the relevant assumptions in this paper are reliable and reasonable, and then limit constraints on the position of the gamma-ray emitting region relative to the BLR and the properties of the BLR. ",
        "introduction": "Thanks to the development of gamma-ray astronomy, the sample of GeV or TeV gamma-ray emitters is increasing. With the detection of high energy gamma-rays in over 60 blazars, such as 3C 279 and PKS 1510-089, in the GeV energy range by the EGRET experiment aboard the Compton Gamma Ray Observatory (Catanese et al. 1997; Fichtel et al. 1994; Lin et al. 1997; Mukherjee et al. 1997; Thompson et al. 1995, 1996; Villata et al. 1997), an exceptional opportunity presents itself for the understanding of the central engine operating in blazars. Some of blazars have been also firmly detected by atmospheric Cherenkov telescopes (ACTs) at energies above 1 TeV, such as Mrk 421 (Punch et al. 1992) and Mrk 501 (Quinn et al. 1996). The Large Area Telescope instrument on the Gamma-Ray Large Area Space Telescope ($\\it GLAST$), the next-generation high energy gamma-ray telescope which will be launched next year, will observe gamma-rays with energies from 20 MeV to greater than 300 GeV, and have the unique capability to detect thousands of gamma-ray blazars to redshifts of at least $z=4$, with sufficient angular resolution to allow identification of a large fraction of their optical counterparts (see e.g. Chen, Reyes, \\& Ritz 2004). The High Energy Stereoscopic System (HESS) is an imaging atmospheric Cherenkov detector dedicated to the ground based observation of gamma-rays at energies above 100 GeV (Funk et al. 2004; Hinton 2004; Hofmann 2003), and successfully detected TeV gamma-rays from many sources (see e.g. Aharonian et al. 2005). The Major Atmospheric Gamma Imaging Cherenkov telescope (MAGIC) is, currently the largest single-dish Imaging Air Cherenkov Telescope in operation, the one designed to have the lowest energy threshold of $\\sim$ 30 GeV among the new Cherenkov telescopes (see e.g. Baixeras et al. 2004). The MAGIC telescope detected TeV gamma-rays from HESS J1813+178 (Albert et al. 2006). $\\it GLAST$, combined with new-generation TeV instruments such as VERITAS, HESS and MAGIC, will tremendously improve blazar spectra studies, filling in the band from 20 MeV to 10 TeV with high significance data for hundreds of active galactic nuclei (AGNs) (see Gehrels \\& Michelson 1999). Future measurements of the gamma-ray spectra shape and its variability of blazars would tremendously improve our understanding to blazars.  Blazars are radio-loud AGNs characterized by luminous nonthermal continuum emission extending from radio up to GeV/TeV energies, very rapid variability, high and variable polarization, flat radio spectrum, high brightness temperatures and supperluminal motion of compact radio cores (e.g. Urry \\& Padovani 1995). These unusual characteristics of blazars are generally believed to be attributed by emission from a relativistic jet oriented at a small angle to the line of sight (Blandford \\& Rees 1978). The spectral energy distributions of blazars consist of two broad bumps. The first component is from the synchrotron radiation processes. The second component is contributed by inverse Compton emission of the same electron population, including synchrotron self-Compton (SSC) scattering synchrotron seed photons, and external Compton (EC) scattering seed photons from sources outside the jet (see e.g. B\\\"ottcher 1999). The external soft photon fields not only provide target photons for EC processes to produce the second component, but also absorb these gamma-rays from the EC processes, because gamma-rays with energies above 10 GeV interact with optical--UV photons and then could be attenuated by photon-photon pair production collisions with the external soft photons. The positions of gamma-ray emitting regions are still an open and controversial issue in the researches on blazars. It is suggested that gamma-rays are produced within broad line region (BLR) and that the gamma-ray emitting radius $R_{\\gamma}$ ranges roughly between 0.03 and 0.3 pc (Ghisellini \\& Madau 1996). It is argued by Georganopoulos et al. (2001) that the radiative plasma in relativistic jets of powerful blazars are within cavities formed by BLRs. However, it is also argued by other researchers that the gamma-ray emitting regions are outside BLRs (Lindfors, Valtaoja, \\& Turler 2005; Sokolov \\& Marscher 2005).  Observations found a large population of X-ray strong FSRQs, labeled HFSRQs by Perlman et al. (1998). HFSRQs have strong broad emission lines, flat hard X-ray spectra similar to classical FSRQs, but high synchrotron peak frequencies and steep soft X-ray spectra similar to intermediate or high synchrotron peak frequency BL Lac objects (e.g., Georganopoulos 2000; Padovani et al. 2003). Therefore, it is expected that HFSRQs are likely to be GeV--TeV gamma-ray emitters. For example, RGB J1629+4008 has a synchrotron peak $\\sim 2\\times 10^{16}\\/\\ \\rm{Hz}$ and a predicted Compton flux $\\nu F_{\\rm{\\nu}}$ of nearly $10^{-10}\\/\\ \\rm{erg\\/\\ s^{-1}\\/\\ cm^{-2}}$ at $10^{25}\\/\\ \\rm{Hz}$ (Padovani et al. 2002). However, the BLRs can reprocess and scatter the central UV radiation to produce diffuse line emission and continuum radiation, respectively. The two external diffuse radiation could absorb gamma-rays with energies from 10 to 200 GeV. Then, it is unknown whether these gamma-rays could be detected by $\\it GLAST$ even if blazars intrinsically produce gamma-rays with energies from 10 to 200 GeV in the observer's frame. In this paper, we focus our efforts on the photon-photon absorption for these gamma-rays by the two diffuse radiation fields of the BLRs in the powerful blazars, and then hope to limit constraints on the positions of gamma-ray emitting regions by comparing our results with $\\it GLAST$ observations in the near future.  The structure of this paper is as follows. Section 2 consists of four subsections, in which subsection 2.1 presents the cross section of pair production, subsection 2.2 the photon-photon attenuation optical depth, subsection 2.3 the diffuse radiation field of the BLR, and subsection 2.4 the dependence of $\\tau_{\\gamma\\gamma}$ on parameters. Section 3 presents the photon-photon absorption optical depth for FSRQs, section 4 presents applications to individual sources, and the last section is for discussions and conclusions. Throughout this paper, we use a flat cosmology with a deceleration factor $q_0=0.5$ and a Hubble constant $H_0=75 \\/\\ \\rm{km \\/\\ s^{-1} \\/\\ Mpc^{-1}}$. ",
        "conclusions": "It can be seen from equation (12) that the emissivity of emission lines is unrelated to the covering factor, and then the relevant absorption for gamma-rays is also unrelated with the value of $f_{\\rm{cov}}$. The gamma-ray absorption contributed by the diluted blackbody radiation is in proportion to the ratio $\\tau_{\\rm{BLR}}/f_{\\rm{cov}}$ (see eq. 13). There is a large uncertainty in the BLR covering factor, and the uncertainty could affect our results. Early estimates of this quantity indicated a covering factor $f_{\\rm{cov}}\\sim 5-10\\%$, while recent observations indicated instead $f_{\\rm{cov}}\\sim 30\\%$ (e.g., Maiolino et al. 2001). Estimates of the BLR covering factors need to measure UV continuum emission from accretion disks. In the case of blazars, these UV continuum emission are swamped by the beaming components from relativistic jets, and this makes it difficult to calculate the BLR covering factors for blazars. The way out to roughly estimate $f_{\\rm{cov}}$ is to know the BLR luminosity $L_{\\rm{BLR}}$ of blazars, in which the thermal emission is clearly dominating the synchrotron ones. D'Elia et al. (2003) found only two of these blazars in literatures: 3C 273 and PKS 2149-306. They obtained $f_{\\rm{cov}}\\sim 7\\%$ for the first source, while $f_{\\rm{cov}}\\sim 10\\%$ for the second one. D'Elia et al. (2003) adopted the BLR covering factor $f_{\\rm{cov}}\\sim 10\\%$ as the typical value of $f_{\\rm{cov}}$ for their blazar sample. Donea \\& Protheroe (2003) suggested $f_{\\rm{cov}}\\sim 3\\%$ for the spherical distribution of the BLR clouds in quasars. The issue of the Thomson optical depth of BLRs has been rarely studied and has gotten nowhere for blazars. Blandford \\& Levison (1995) adopted the Thomson optical depth of BLR as $\\tau_{\\rm{BLR}}=0.01$. Then the ratio $\\tau_{\\rm{BLR}}/f_{\\rm{cov}}$ of the Thomson optical depth to the covering factor is likely to be typically of $\\tau_{\\rm{BLR}}/f_{\\rm{cov}}\\sim 1/3$ for the spherical distribution of the BLR clouds in blazars. This value of $\\tau_{\\rm{BLR}}/f_{\\rm{cov}}\\sim 1/3$ is on the same order of magnitude as $\\tau_{\\rm{BLR}}/f_{\\rm{cov}}=1$, assumed in the previous sections.  As stated in section 3, the external radiation energy densities of FSRQs are on the order of $U_{\\rm{ph}}\\sim 10^{-2.1} \\/\\ \\rm{erg\\/\\ cm^{-3}}$, and this typical value limits constraints on the probable positions of inner and outer radii of BLR or the half thickness of spherical shell of BLR. The probable positions of the inner and outer radii are around $r_{\\rm{BLR,in}}\\sim 0.2 \\/\\ \\rm{pc}$ and $r_{\\rm{BLR,out}}\\sim 0.6 \\/\\ \\rm{pc}$, respectively. The median of this probable inner and outer radii is around $0.4 \\/\\ \\rm{pc}$, which is consistent with the radius $r_{\\rm{BLR}}\\simeq 0.43 \\/\\ \\rm{pc}$ derived from the reverberation mapping. This agreement implies that the radius $r_{\\rm{BLR}}$ derived from the reverberation mapping probably denotes the median of $r_{\\rm{BLR,in}}$ and $r_{\\rm{BLR,out}}$, and our assumption of $r_{\\rm{BLR}}$ in the previous sections is appropriate. The energy density distribution for the BLR with $r_{\\rm{BLR,in}}\\sim 0.2 \\/\\ \\rm{pc}$, $r_{\\rm{BLR,out}}\\sim 0.6 \\/\\ \\rm{pc}$ and $\\tau_{\\rm{BLR}}/f_{\\rm{cov}}\\sim 1/3$ has an average of energy densities within the outer radius on the order of $U_{\\rm{ph}}\\sim 10^{-2.1} \\/\\ \\rm{erg \\/\\ cm^{-3}}$. The corresponding absorption optical depth for gamma-rays is presented by the curve III in Figure 7, from which it can be seen that the two diffuse radiation fields of the BLRs are not transparent to these gamma-rays with energies from 10 to 200 GeV.  Because of the relativistic beaming effect, most of gamma-rays for FSRQs are contained within a radiation cone with a half open angle of $\\Delta \\varphi\\sim 1/ \\Gamma\\sim 3.8^{\\circ}$. If the central UV photons directly from accretion disks travel through the radiation cone, the UV photons can have photon-photon pair creation processes with gamma-rays within the radiation cone. Assuming the radius of the UV radiation region is around $R_{D}\\sim 100 R_{\\rm{s}}\\sim 0.001 \\/\\ \\rm{pc}$ ($R_{\\rm{s}}=2GM/c^2$ and $M=10^8 M_{\\rm{\\odot}}$), and the gamma-ray emitting region $R_{\\rm{\\gamma}}\\sim 0.1 \\/\\ \\rm{pc}$, the angle between the jet direction and the travelling direction of UV photons at $R_{\\rm{\\gamma}}$ is $\\sim \\arcsin R_{\\rm{D}}/R_{\\rm{\\gamma}}\\sim 0.57^{\\circ}$. Then the photons within the radiation cone have the collision angles of $\\theta \\lesssim 4.4^{\\circ}$. For the UV photons at the frequency $\\nu \\simeq10^{16.5} \\/\\ \\rm{Hz}$ with energies of $\\varepsilon_2\\simeq 2.56\\times 10^{-4}$ and these gamma-rays of $\\varepsilon_1\\simeq 3.91\\times 10^{5}$ corresponding to energies around 200 GeV, the left of equation (3) has a upper limit of $\\lesssim 0.15$, which is less than the value at the right of equation (3), and then the two kinds of photons cannot be absorbed by the photon-photon pair creation processes. Therefore, the central UV radiation have negligible contributions to the absorption for gamma-rays relative to the diffuse radiation from the BLRs.  The energies carried by the annihilated gamma-rays would be deposited in the lower frequency radiation field by the synchrotron and inverse Compton processes of the electron-positron pairs and the pair cascade and pair annihilation processes (see Protheroe \\& Stanev 1993; Saug\\'e \\& Henri 2004; Zdziarski \\& Coppi 1991). The intense creation of pairs would produce a strong radiation at low energy X-rays or at GeV energy (Protheroe \\& Stanev 1993; Saug\\'e \\& Henri 2004; Zdziarski \\& Coppi 1991). The electron-positron pair cascade can cause the soft X-ray excesses (Zdziarski \\& Coppi 1991). The produced electron-positron pairs could make a difference at lower energies, around 1 GeV, where the pileup of photons below the absorption feature results in a significant flattening in the observed spectrum relative to the emission spectrum (Protheroe \\& Stanev 1993). Then the electron-positron pairs produced by the photon-photon annihilation processes in FSRQs could result in the similar observational effects mentioned above at around 1 GeV gamma-rays or low energy X-rays.  In summary, we used a BLR model to study the photon-photon absorption optical depth for gamma-rays with energies from 10 to 200 GeV in the observed spectrum. There are four key parameters for the BLR model: the gamma-ray emitting radius $R_{\\gamma}$, the BLR luminosity $L_{\\rm{BLR}}$, the half thickness of BLR $h$ and the ratio $\\tau_{\\rm{BLR}}/f_{\\rm{cov}}$. The BLR luminosity $L_{\\rm{BLR}}$ determines the distance $r_{\\rm{BLR}}$, which limits the BLR inner and outer radii of $r_{\\rm{BLR,in}}$ and $r_{\\rm{BLR,out}}$ if the parameter $h$ is given. We discussed the dependence of the optical depth $\\tau_{\\gamma\\gamma}$ on the parameters $R_{\\rm{\\gamma}}$, $h$ and $r_{\\rm{BLR}}$. The diffuse external radiation energy density in the gamma-ray emitting region is typically on the order of $U_{\\rm{ph}}\\sim 10^{-2} \\/\\ \\rm{erg \\/\\ cm^{-3}}$ for FSRQs, and the ratio $\\tau_{\\rm{BLR}}/f_{\\rm{cov}}$ might be typically on the order of $\\tau_{\\rm{BLR}}/f_{\\rm{cov}}\\sim 1$ for blazars. Then the inner and outer radii of BLR are required to be around $r_{\\rm{BLR,in}}\\sim 0.2 \\/\\ \\rm{pc}$ and $r_{\\rm{BLR,out}}\\sim 0.6 \\/\\ \\rm{pc}$, respectively. It is impossible for gamma-rays from 10 to 200 GeV emitted by FSRQs to escape from the diffuse radiation fields of the BLRs if $R_{\\gamma}\\sim r_{\\rm{BLR,in}}$ (see Fig. 7). If $\\it GLAST$ could detect gamma-rays with energies from 10 to 200 GeV for most of FSRQs, there would be $R_{\\rm{\\gamma}}>r_{\\rm{BLR,in}}$ for FSRQs on average, otherwise, $R_{\\rm{\\gamma}}\\lesssim r_{\\rm{BLR,in}}$. If $\\tau_{\\rm{BLR}}/f_{\\rm{cov}}\\sim 1/3$ and $R_{\\rm{\\gamma}}\\sim r_{\\rm{BLR,in}}$ for the classical FSRQ 3C 279, the diffuse soft photon field is not transparent to gamma-rays with energies from multi-ten to 200 GeV and is transparent to gamma-rays around 10 GeV. $\\it GLAST$ observations in the energy domain of multi-ten to 200 GeV could limit constraints on the position of the gamma-ray emitting region relative to the inner radius of BLR and the BLR properties for 3C 279. For the powerful HFSRQ PKS 0405$-$123, the diffuse photon field is not transparent to gamma-rays with energies from 10 to 200 GeV in the observed spectrum. $\\it GLAST$ observations in the future could limit constraints on the gamma-ray emitting region and the BLR in this HFSRQ. Our results are model dependent, especially dependent on the gamma-ray emitting radius $R_{\\rm{\\gamma}}$ and the inner radius $r_{\\rm{BLR,in}}$. To determine the relationship of the inner radius $r_{\\rm{BLR,in}}$ with the radius $r_{\\rm{BLR}}$ derived from the reverberation mapping needs more theoretical and observational researches. $\\it GLAST$ observations in the future can give more observational constraints on the gamma-ray emitting regions and the BLRs for the powerful blazars."
    },
    "0807/0807.4186_arXiv.txt": {
        "abstract": "Gravitational scattering between massive planets has been invoked to explain the eccentricity distribution of extrasolar planets. For scattering to occur, the planets must either form in -- or migrate into -- an unstable configuration. In either case, it is likely that a residual gas disk, with a mass comparable to that of the planets, will be present when scattering occurs. Using explicit hydrodynamic simulations, we study the impact of gas disks on the outcome of two-planet scattering. We assume a specific model in which the planets are driven toward instability by gravitational torques from an outer low mass disk. We find that the accretion of mass and angular momentum that occurs when a scattered planet impacts the disk can strongly influence the subsequent dynamics by reducing the number of close encounters. The eccentricity of the innermost surviving planet at the epoch when the system becomes Hill stable is not substantially altered from the gas-free case, but the outer planet is circularized by its interaction with the disk. The signature of scattering initiated by gas disk migration is thus a high fraction of low eccentricity planets at larger radii accompanying known eccentric planets. Subsequent secular evolution of the two planets in the presence of damping can further damp both eccentricities, and tends to push systems away from apsidal alignment and toward anti-alignment. We note that the late burst of accretion when the outer planet impacts the disk is in principle observable, probably via detection of a strong near-IR excess in systems with otherwise weak disk and stellar accretion signatures. ",
        "introduction": "The eccentricity distribution of massive extrasolar planets beyond the tidal circularization radius is broad, with a median eccentricity $<e> \\simeq 0.25$ \\citep{marcy05}. Since giant planet formation models predict that the envelope is accreted when the planet has an almost circular orbit, the observed eccentricities are generally interpreted as evidence of evolutionary processes that occur subsequently. One possibility is that the same gravitational torques from the gas disk that drive orbital migration also excite eccentricity \\citep{artymowicz91,ogilvie03,goldreich03,moorhead08}. Alternatively, the eccentricity may arise from gravitational instability within an unstable multiple system \\citep{rasio96,weidenschilling96,lin97}. {\\it N}-body experiments show that a wide variety of unstable initial conditions -- which can include systems with two, three or many giant planets -- evolve under scattering to yield an eccentricity distribution similar to that observed \\citep{chatterjee07,juric07}. The agreement between theory and observations supports the scattering scenario, though other possibilities are not ruled out.  To date, most studies of planet-planet scattering have considered purely gravitational evolution from gas-free initial conditions. Gas has been ignored primarily for reasons of computational necessity, since fairly robust arguments suggest that at least {\\em small} amounts of gas are likely to survive to the epoch when scattering occurs (planetesimal disks will also be present, though these will likely be dynamically important only in systems similar to the Solar System, in which at least some giant planets at large radii have low envelope masses). In particular, for two-planet systems most initial separations within the boundary that is Hill stable for all time \\citep{gladman93} are unstable on very short time scales.  Realistically, an unstable two planet system is therefore likely to arise only if the planets are driven toward such a state, for example due to gas disk migration. Efficient migration requires a disk mass that is of the same order as the planet mass. Multiple planet systems with $N_p > 2$ exhibit broader instability regions, but even in this case the most likely scenario is that instability sets in once eccentricity damping due to the presence of gas becomes small enough. Again, this will occur once (to order of magnitude) the gas disk mass falls to the level of the planet mass, and scattering is likely to occur in an environment that is gas poor but not gas free.  In this paper we use hydrodynamic simulations to study the influence of low mass residual gas disks on the outcome of gravitational scattering. As noted above, there are many possible scenarios that yield different initial conditions for scattering experiments -- here we adopt a simple model that is computationally tractable. We assume that two Jupiter-mass planets form on well-separated circular orbits. Models of giant planet formation suggest that the onset of runaway accretion occurs when the cores attain a mass of 10~$M_\\oplus$ \\citep{hubickyj05}, and, since neighboring cores are likely to form with a separation of $\\sim 10$ mutual Hill radii \\citep{kokubo98}, this suggests that plausible initial separation ratios are $\\simeq 1.3$ or greater. If the gas disk interior to the outer planet is relatively weak migration driven by the residual outer gas results in convergence of the orbits. The initial separation alluded to above lies close to the 3:2 resonance, so depending upon the precise initial separation convergent migration might result in resonance crossing. If the resonance {\\em is} encountered, the result can either be trapping into resonance \\citep{lee02,snellgrove01,kley05}\\footnote{The existence of a number of known resonant pairs of planets supports the idea that trapping occurs in at least some cases.} or, if the disk is turbulent \\citep{adams08}, avoidance of trappping and ongoing convergence of the orbits. In the latter case the outer planet continues to migrate inward on a viscous time scale until the separation of the planets is driven past the Hill stability criterion, and a scattering event occurs. The goal of this paper is to answer the question: what effect does the remnant disk now have on the scattering planetary system? ",
        "conclusions": "The expense of the hydrodynamic simulations means that we have been able to run only a handful of realizations of scattering in the presence of gas, and as a consequence it is unwise to draw definitive conclusions. Nonetheless we find, as expected, that remnant disks that have masses comparable to the masses of the giant planets {\\em do} influence the outcome of gravitational scattering, primarily via eccentricity damping that occurs due to accretion when planets intrude upon the gas disk. If, as we have assumed, the remnant gas lies exterior to the orbit of the outermost planet, damping when the planet enters the disk tends to increase the planetary separation and stabilize the system against close encounters and possible ejection of one of the planets. The stabilization occurs rapidly enough that we were able to evolve the systems with gas up to the point where they would plausibly be Hill stable.  The observed eccentricity distribution of extrasolar planets\\footnote{A detailed study of possible selection biases in the Keck planet search \\citep{cumming08} shows that the survey is complete for velocity amplitudes $K > 20 \\ {\\rm ms}^{-1}$ and $e < 0.6$. The ``observed\" distribution is therefore a reasonable proxy for the true distribution in the sense that neither low nor moderately high $e$ objects are systematically missed. There may, however, be some upward bias to the derived eccentricities of detected planets in cases where the measurement signal to noise is relatively low \\citep{shen08}.} matches the theoretically predicted one {\\em from simulations that do not include gas} \\citep{chatterjee07,juric07}. If gas is implicated in driving planetary systems toward instability, its effect on the subsequent dynamics cannot be so strong as to ruin this agreement. We find that this is plausible. Specifically, models computed with an exterior disk whose mass equals that of the planets yield inner planet eccentricities that are comparable to those achieved in the absence of gas. Since radial velocity surveys favor the detection of planets at smaller orbital radii, the observed sample today is likely dominated by planets which would have been unaffected by gas during the scattering process. Conversely, we find that the eccentricity of the outer planet is significantly damped by accretion of gas once it enter the disk. The dynamical dominance of accretion over other damping mechanisms is likely the most robust result of this paper.  Although we have only considered one specific scattering scenario, it seems plausible that the generic effect of residual gas disks will be to increase the fraction of retained rather than ejected planets following scattering. If so, analysis of the abundance of multiple planets ought to provide the strongest constraints on models of this type. If the specific model for migration and scattering we have used were actually commonplace in extrasolar planetary systems, we would expect to find a high fraction of multiple planet systems in which the outer planet had a lower eccentricity than the inner planet. Observationally, it is known that there are hints for additional companions in many systems that host known planets \\citep{wright07}, but the quantitative statistical analysis necessary to place limits on their masses and periods (analogous to the work of \\citealt{cumming08} for single planets) has not, to our knowledge, been attempted. Our impression is that while our toy model -- in which the hypothetical outer planet has the same mass as the detected inner one -- could already be ruled out observationally, a more realistic scenario in which the outer planet was modestly less massive probably remains viable. More extensive theoretical work will be needed to quantify the extent to which residual gas disks increase the predicted fraction of multiple systems after scattering.  On Myr timescales, the inner planet's eccentricity can be damped indirectly by the remnant gaseous disk via excitation of the outer planet and subsequent dissipation.  The lifetime of the remnant disk is uncertain, but is probably $\\sim 10^5$ years \\citep{simon95,wolk96,clarke01,alexander06}.  Given the large eccentricity of the inner planet and the relative proximity of the outer planet, there is strong secular coupling and eccentricity oscillations are induced.  The damping of the outer planet by the remnant gas disk tends to push the system away from apsidal alignment and toward anti-alignment; the amount of damping depends on the disk lifetime.  Depending on the starting configuration, strong damping can produce systems that are near the apsidal separatrix \\citep{barnes06}, circulating, or librating about anti-alignment.  However, strong damping should not produce systems that librate about alignment.  In cases with relatively little damping, the apsidal alignment is weakly affected, although eccentricities can be modestly reduced (see figure \\ref{longterm_var}).  Finally, we note that the accretion rate onto the planet when the planet first enters the disk is rather high. This gas will, in reality, join a circumplanetary disk, and how quickly it is accreted by the planet will depend upon the evolution of that disk.  However, it is clear that the luminosity of the planet could be large due to the implied late burst of accretion. Naively estimating the accretion luminosity as $L_{acc} \\sim G M_p \\dot{M} / R_p$, we find that for Jovian parameters the luminosity over the first 500~yr is of the order of $10^{32}$-$10^{33}$ erg s$^{-1}$. Since the accreted gas has angular momentum, the mass inflow onto the planet would be buffered by a circumplanetary disk, though only modestly since the viscous timescale of such disks is typically of the order of $10^3 \\ {\\rm yr}$ or less \\citep{canup02}. Unlike the main envelope accretion event that occurs in core accretion models for giant planet formation, this late burst of accretion is in principle observable (as a near-IR excess without matching stellar accretion signatures) since the accretion luminosity of the circumstellar disk at this epoch is small. If we assume that late accretion is detectable for $10^3 \\ {\\rm yr}$, and occurs in 10-100\\% of young stars, then we would need to survey $10^4$ to $10^5$ systems to have a chance of directly observing the last stage of giant planet formation. Surveys of this magnitude are well beyond current capabilities, but are not inconceivable in the future."
    },
    "0807/0807.3073_arXiv.txt": {
        "abstract": "The origin of dramatically different electron distributions responsible for Comptonization in black hole X-ray binaries (BHBs) in their various states is discussed. We solve the coupled kinetic equations for photons and electrons without approximations on the relevant cross-sections accounting for Compton scattering, synchrotron radiation, and Coulomb collisions. In the absence of external soft photons,  the electrons are efficiently thermalized  by synchrotron self-absorption and Coulomb scattering  even for pure nonthermal electron injection. The resulting quasi-thermal synchrotron self-Compton spectra have very  stable slopes and electron temperatures similar to the hard states of BHBs.  The hard spectral slopes observed in the X-rays, the cutoff at 100 keV, and the MeV tail  together require low magnetic fields ruling out the magnetic dissipation mechanism. The motion of the accretion disk toward the black hole results in larger Compton cooling and lower equilibrium electron temperature.  Our self-consistent simulations show that in this case both electron and photon distributions attain a power-law-dominated shape similar to what is observed in the soft state. The electron distribution in the Cyg X-1 soft state  might require a strong magnetic field, being consistent with the magnetically dominated  corona. ",
        "introduction": "The physical processes giving rise to the X-ray/gamma-ray emission of accreting black-hole binaries (BHBs) have been a matter of debates over the last four decades. The hard-state spectra, showing a  strong cut-off  around 100 keV, are well described by thermal Comptonization \\citep[e.g.][]{P98,ZG04}, while a weak MeV tail  requires the presence of nonthermal particles \\citep{McConnell94,Ling97}. The origin of seed soft photons for Comptonization is, however, much less clear. An apparent correlation between the spectral slope and the amount of Compton reflection \\citep{ZLS99} argues in favor of the accretion disk, while the observed optical/X-ray correlation \\citep{MIC82,Kanbach01} leans toward  the synchrotron hypothesis \\citep[e.g.][]{FGM82,WZ01}. Interesting questions are then: what stabilizes the X-ray spectral slope at $\\alphax\\sim 0.6$--0.8, and what  fixes the temperature of the emitting plasma at $k\\Te\\sim$50--100 keV \\citep{ZJP97,P98,ZG04}? Do the feedback from the cool accretion disk and the thermostatic properties of electron-positron pairs \\citep{HM93,HMG94,SPS95,MBP01}  play a role here? Or does the cooling by synchrotron radiation \\citep{NY95} act as a stabilizer?  In the soft state, BHB spectra are dominated by thermal disk emission of temperature $k\\Tbb\\sim$0.4--1.5 keV. At higher energies the spectrum is power-law-like and shows no signatures of the cut-off  \\citep{Grove98} extending possibly up to 10 MeV \\citep{McConnell02}. This emission is well described by Comptonization in almost purely nonthermal plasmas \\citep{PC98,G99,coppi99, ZGP01,ZG04}.  We can then ask why the electrons are nearly thermal in the hard state, and what causes such a dramatic change in the electron distribution when transition to the soft state happens. \\citet{PC98} proposed that the two states are distinguished by the way the energy is supplied to the electrons: by thermal heating, dominating during the hard state, and by nonthermal acceleration, operating in the soft state. However, their treatment of Coulomb collisions \\citep[using {\\sc eqpair} code by][]{coppi92,coppi99} was approximate, and they have neglected the  effect of synchrotron boiler, involving the emission and absorption of synchrotron photons, which can act as an efficient particle thermalizer \\citep{GGS88}.    \\citet{GHS98} studied for the first time the combined effect of the synchrotron boiler and Compton cooling on the electron distribution and photon spectra (but neglected Coulomb scattering). They considered a two-phase corona model \\citep{HM93,HMG94,SPS95}, where half of the high-energy radiation was assumed to be reprocessed by the disk to soft photons. As the actual geometry of the emitting region is not known, we start from pure synchrotron self-Compton models (i.e. with no external soft photons) and compute self-consistently the electron (positron) and photon distributions.  We then investigate  how the additional soft photons (e.g., associated with the inner radius of  the  cool accretion disk) affect the equilibrium distributions and compare the results of simulations with the data on Cyg X-1. The preliminary results of this study were presented by \\citet{VP08a}. ",
        "conclusions": "The hard state of BHBs can well be described by the quasi-thermal SSC mechanism. The feedback from the cool disk is not needed to stabilize the spectral slope and the electron temperature. Electrons can be injected to the active region with the power-law spectrum, but Coulomb scattering and synchrotron self-absorption thermalize them efficiently. This reduces the need  for mysterious 'thermal heating' that was invoked previously to explain thermal Comptonization spectra of BHBs. The MeV tail together with the hard X-ray spectra of BHBs with photon indices $\\alphax\\lesssim$0.7 and a cutoff at 100 keV require rather low magnetization $\\etab<0.1$ and a large size of $R\\gtrsim 60\\Rs$.   In that case, magnetic reconnection can be ruled out as a source of energy.  At high optical depth of the emitting region $\\taup\\gtrsim10$, in the absence of disk radiation, the spectrum is close to saturated Comptonization, peaking at a few keV. This Wien-type spectrum might be associated with the ultrasoft  state of BHBs. At low $\\taup$, the electrons are hotter and the spectra are softer due to the efficient synchrotron cooling.  A behavior similar to what is observed during the spectral state transitions in BHBs can be reproduced by varying the ratio of injected soft luminosity and the power dissipated in the hot phase, which could be caused by varying the radius of the inner cool disk. The increasing Compton cooling causes dramatic changes in the electron distribution from almost purely thermal to nearly nonthermal.  The photon distribution also changes from quasi-thermal SSC  to the nonthermal Comptonization of the disk photons. In the soft state of Cyg X-1,  a strong magnetic field can thermalize electrons at sufficiently high temperature, which is consistent with a magnetically dominated corona being responsible for the high-energy emission."
    },
    "0807/0807.1076_arXiv.txt": {
        "abstract": "Several accurate analyses of the CMB temperature maps from the Wilkinson Microwave Anisotropy Probe (WMAP) have revealed a set of anomalous results, at large angular scales, that appears inconsistent with the statistical isotropy expected in the concordance cosmological model $\\Lambda$CDM. Because these anomalies seem to indicate a preferred direction in the space, here we investigate the signatures that a primordial magnetic f\\/ield, possibly present in the photon-baryon f\\/luid during the decoupling era, could have produced in the large-angle modes of the observed CMB temperature f\\/luctuations maps. To study these imprints we simulate Monte Carlo CMB maps, which are statistical\\/ly anisotropic due to the correlations between CMB multipoles induced by the magnetic f\\/ield. Our analyses reveal the presence of the North-South angular correlations asymmetry phenomenon in these Monte Carlo maps, and we use these information to establish the statistical signif\\/icance of such phenomenon observed in WMAP maps. Moreover, because a magnetic f\\/ield produces planarity in the low-order CMB multipoles, where the planes are perpendicular to the preferred direction def\\/ined by the magnetic f\\/ield, we investigate the possibility that two CMB anomalous phenomena, namely the North-South asymmetry and the quadrupole-octopole planes alignment, could have a common origin. Our results, for large-angles, show that the correlations between low-order CMB multipoles introduced by a suf\\/f\\/iciently intense magnetic f\\/ield, can reproduce some of the large-angle anisotropic features mapped in WMAP data. We also reconf\\/irm, at more than 95\\% CL, the existence of a North-South power asymmetry in the WMAP f\\/ive-year data. ",
        "introduction": "The f\\/ive-year data from the Wilkinson Microwave Anisotropy Probe (WMAP)~\\citep{b1,b2,b3,b4,b5} contain the most valuable cosmological information to study the large-scale properties of the universe (see~\\citealt{b6,b7,b8,b9,b10,b11,b12}) for previous releases of WMAP data). One of these features concerns the hypothesis that the set of temperature f\\/luctuations of the cosmic microwave background radiation (CMB) is a stochastic realization of a random f\\/ield, meaning that its angular distribution on the celestial sphere is statistical\\/ly isotropic at all angular scales. Examination of the three-year WMAP data conf\\/irms highly significant departures from statistical isotropy at large angular scales~\\citep{b22,b21,b19,b14,b15,b16,b17,% b20,b13,b18}, previously found also in first-year WMAP data~\\citep{b33,b34,b23,b24,b31,b25,b26,b27,b28,b29,b32,b30}. Evidences for such large-angle anisotropy come from the asymmetry of the CMB angular correlations between the northern and southern ecliptic hemispheres (hereafter NS-asymmetry), with indications of a preferred axis of maximum hemispherical asymmetry~\\citep{b23,b24,b25,b26,b27,b28,b29,b16,b30,b13,b18}. Furthermore, other manifestations of large-angle anisotropy include the unlikely quadrupole-octopole planes alignment (refering both to the strong planarity of these multipoles as well as to the alignment between such planes, see, e.g.,~\\citealt{b33,b34,b35,b32,b19,b22,b20,b17}). Actually, several anomalous results were already reported after analyses of the WMAP data~\\citep{b37,b38,b31,b36,b39,b40,b41,b107,b42}. For a dif\\/ferent point of view see, e.g.,~\\citet{b43,b44,b45,b46,b47,b48,b49,b50}.  Possible sources of these anomalies include non-CMB contaminants, like residual foregrounds~\\citep{b51,b52,b53,b54,b55,b56,b57}, incorrectly  subtracted dipole and/or dynamic quadrupole terms~\\citep{b32,b58,b20}, or systematic errors~\\citep{b59,b60}. For this, particular ef\\/forts have been done by the WMAP team in last releases in order to improve the data processing by minimizing the ef\\/fects of foregrounds (mainly coming from dif\\/fuse Galactic emission and astrophysical point-sources), artifacts (in the mapmaking process, in the instrument characterization, etc.), and systematic errors~\\citep{b11,b1,b2}. As a result, the data released by the WMAP team include the Internal Linear Combination (hereafter ILC-5yr) full-sky CMB map suitable for large-angle temperature f\\/luctuations studies~\\citep{b10,b1,b3}. Here we investigate the ILC-5yr map, and for completeness, we also study the other full-sky cleaned CMB maps that were dif\\/ferently processed from WMAP f\\/ive- and three-year data releases in order to account for foregrounds and systematics. Thus, we also consider the Kim-Naselsky-Christensen~\\citep{b94}, WMAP-3yr ILC~\\citep{b10}, de Oliveira-Tegmark~\\citep{b53}, and Park-Park-Gott~\\citep{b17} CMB maps, hereafter termed the HILC-5yr, ILC-3yr, OT-3yr, and PPG-3yr, respectively.  A number of studies has been done looking for physical explanations of the above mentioned anomalies, specially searching for a unifying mechanism relating both the NS-asymmetry and the CMB lower multipoles alignment (see, e.g.,~\\citealt{b19,b61,b62}). Additionally, some processes that breaks down statistical isotropy during the inf\\/lationary epoch have been suggested~\\citep{b63,b64,b65,b66}. Nonetheless, one can also interpret such CMB large-angle anisotropy as being of cosmological origin, in this sense globally axisymmetric space-times have been proposed to account for the mapped preferred axis in WMAP maps~\\citep{b67,b70,b68,b69,b71,b72}. Previous works investigated the ef\\/fects of a primordial homogeneous magnetic f\\/ield on the CMB temperature f\\/luctuations on all angular scales (see, e.g.,~\\citealt{b90,b105}). Here we study such primordial f\\/ield as a possible physical mechanism to produce two large-angle phenomena, that is, the CMB NS-asymmetry and the lower CMB multipoles alignment.  In section~\\ref{magnetic fields} we present this primordial magnetic f\\/ield scenario, and show the ef\\/fect of such a f\\/ield on the CMB temperature f\\/luctuations. Then, in order to reveal such ef\\/fects on simulated CMB maps we develop a geometrical-statistical method, which is presented in section~\\ref{method}. After that we use our anisotropic indicator to perform, in section~\\ref{data analyses}, the analyses of both sets of data, the Monte Carlo simulated CMB maps as well as the WMAP maps. At the end, in section~\\ref{conclusions}, we discuss our results and formulate our conclusions.   \\section[]{Primordial magnetic f\\/ields scenario}\\label{magnetic fields}  There is strong observational evidence for the presence of large-scale intergalactic magnetic f\\/ields of fews $\\mu G$, and magnetic f\\/ields of similar strength within clusters of galaxies (see, e.g.,~\\citealt{b73,b74,b75,b76,b77}). Nowadays it is believed that these magnetic f\\/ields are amplif\\/ications of small primordial magnetic f\\/ields of the order of few nanoGauss, that would have occured due to dif\\/ferent processes, like galactical dynamo (see, e.g.,~\\citealt{b78,b79,b80}), during anisotropic protogalactic collapses, or due to dif\\/ferential rotation in galaxies (see, e.g.,~\\citealt{b81,b82}). In turn, such primordial seeds of magnetic f\\/ields would have several origins, for instance an electroweak phase transition~\\citep{b84,b83,b85} or quark-hadron phase transition (see, e.g.,~\\citealt{b86,b87}). Here, we assume the existence of a primordial homogeneus magnetic f\\/ield and investigate their ef\\/fects on the CMB temperature f\\/luctuations at large-angles.  According to the Einstein equations for linearized metric perturbations, in absence of a magnetic f\\/ield, the vector metric perturbations go like $a^{-2}$ and the velocity induced by these perturbations goes like $a^{-1}$ (where $a$ is the scale factor that accounts for the expansion of the universe). Therefore, they decay very fast with the expansion of the universe~\\citep{b88} and the velocities produced by vector metric perturbations do not contribute signif\\/icatively to the CMB temperature f\\/luctuations.  The presence of a homogeneus magnetic f\\/ield in the early universe change this scenario because such f\\/ields modify the behavior of charged particles in the primordial plasma via Lorentz forces producing additional velocity gradients in the f\\/luid. In this way, a magnetic f\\/ield induces Alfv\\'en waves in the primordial plasma that propagate at velocity $v_{\\mbox{\\rm\\footnotesize A}}$, changing the speed of sound in the photon-baryon f\\/luid as $c^2_s \\longrightarrow c^2_s + v^2_{\\mbox{\\rm\\footnotesize A}} \\cos^2\\theta$, where~\\citep{b89} \\begin{eqnarray} v_{\\mbox{\\rm\\footnotesize A}}^2 = \\frac{B^2_0}{4\\pi(\\rho + p)} \\, , \\end{eqnarray} with $\\rho$ and $p$ being the density and pressure in the radiation dominated era, respectively, $B_0$ is the strength of the magnetic f\\/ield $\\textbf{B}$, $\\theta$ is the angle between $\\textbf{B}$ and the $\\vec{k}$-mode of the Fourier expansion of Alfv\\'en velocity $v_{\\mbox{\\rm\\footnotesize A}}$. These Alfv\\'en wave modes induce small rotational velocity perturbations which, for the scales of our interest, have the form $\\vec{v} \\, \\approx \\, \\vec{v}_0 \\, v_{\\mbox{\\rm\\footnotesize A}}\\, k\\, t\\, \\cos \\theta$, where $t$ is the cosmic time, $k \\equiv |\\vec{k}|$, and $\\vec{v}_0$ is the initial velocity, which we assume to have a power spectrum of the form~\\citep{b90,b91} \\begin{eqnarray} \\label{espectrov} \\langle v_{0i} \\, v_{0j} \\rangle \\propto k^n \\, (\\delta_{ij} - k_i \\, k_j) \\, . \\end{eqnarray} Vectorial contributions to CMB temperature f\\/luctuations are present via Doppler and Sachs-Wolfe ef\\/fects~\\citep{b90} \\begin{eqnarray} \\frac{\\delta T}{T}(\\hat{n})^{\\mbox{(\\footnotesize vec)}} \\, =\\, -\\vec{V}_T \\cdot \\hat{n}|^{t_0}_{t_{\\rm\\small dec}} \\, +\\, \\int^{t_0}_{t_{\\rm\\small dec}} \\dot{\\vec{\\nu}} \\cdot \\hat{n} \\, dt \\, , \\end{eqnarray} where $\\hat{n}$ indicates any direction in the sky, $\\vec{V}_T = \\vec{V} + \\vec{v}$, $\\vec{\\nu}$ is the vectorial metric perturbation, $\\vec{V}$ is the velocity produced by $\\vec{\\nu}$, and the integration is between the actual time $t_0$ and the decoupling time $t_{\\rm\\small dec}$. Thus, only the rotational velocity perturbations $\\vec{v}$ contributes ef\\/fectively to temperature f\\/luctuations because $\\vec{\\nu}$ and $\\vec{V}$ decay quickly with the expansion of the universe~\\citep{b88}. Therefore, the signatures of a homogeneus primordial magnetic f\\/ield $\\textbf{B}$ on the CMB temperature f\\/luctuations are~\\citep{b90} \\begin{eqnarray} \\label{Flutuacoes} \\frac{\\delta T}{T}^{\\textbf{B}}(\\hat{n},k) \\,\\approx\\, \\hat{n} \\cdot \\vec{v}_{0} \\, v_{A} \\, k \\, t_{\\rm\\small dec} \\, \\cos \\theta \\, . \\end{eqnarray} As we can see, the existence of a preferred direction associated to $\\textbf{B}$ af\\/fects the CMB temperature f\\/luctuations through the angle $\\theta$, which implies, a straight dependence on the orientation of the vector ${\\textbf{B}}$, consequently, a break down of the statistical isotropy in the CMB sky.  The most suitable quantity to simulate statistical\\/ly anisotropic CMB skies is the correlation matrix of multipole coef\\/f\\/icients. For this, we expand the sky temperature f\\/luctuations in spherical harmonics \\begin{eqnarray}\\label{Temp} \\frac{\\delta T}{T}(\\hat{n}) &=& \\sum_{\\ell,m} a_{\\ell m} Y_{\\ell m}(\\hat{n})  \\, , \\\\ a_{\\ell m} &=& \\int \\frac{\\delta T}{T}(\\hat{n}) \\, Y^{\\ast}_{\\ell m}(\\hat{n}) \\, d\\Omega \\, , \\end{eqnarray} and then use eqs.~(\\ref{espectrov}),~(\\ref{Flutuacoes}) and~(\\ref{Temp}) to calculate the correlation matrix $\\langle a_{\\ell m} a^{\\ast}_{\\ell' m'} \\rangle$. In this work we assume a scale-invariant power spectrum for the magnetic f\\/ield, i.e., $n=-5$ in the eq.~(\\ref{espectrov}).  The explicit calculation of the correlation matrix elements, for a Harrison-Zel'dovich scale-invariant power spectrum for the magnetic f\\/ield, gives~\\citep{b90,b91,b104} \\begin{eqnarray}\\label{CorrelMat1} \\langle a_{\\ell m} a^{\\ast}_{\\ell' m'} \\rangle^{\\textbf{B}} \\!=\\! \\delta_{m m'} [ \\delta_{\\ell \\ell'} C^{\\textbf{B}}_{\\ell m} + (\\delta_{\\ell+1,\\ell'-1}+\\delta_{\\ell-1,\\ell'+1}) D^{\\textbf{B}}_{\\ell m} ],\\!\\!\\! \\end{eqnarray} where \\begin{eqnarray}\\label{CorrelMat2} C^{\\textbf{B}}_{\\ell m} \\,=\\, 27.12 \\times 10^{-16} \\left( \\frac{B_0}{1\\,nG} \\right)^4 \\times \\nonumber \\\\ \\frac{2\\ell^4 + 4\\ell^3-\\ell^2-3\\ell+6m^2-2\\ell m^2-2 \\ell^2 m^2} {(\\ell+2)(\\ell+1)\\ell(\\ell-1)(2\\ell-1)(2\\ell+3)} \\, , \\end{eqnarray} and \\begin{eqnarray}\\label{CorrelMat3} \\frac{D^{\\textbf{B}}_{\\ell m}}{C^{\\textbf{B}}_{\\ell m}} = \\frac{9\\pi}{32} \\frac{\\sqrt{(\\ell+m+1)(\\ell-m+1)(\\ell+m)(\\ell-m)}} {2\\ell^4+4\\ell^3-\\ell^2-3\\ell+6m^2-2\\ell m^2-2 \\ell^2 m^2} \\times \\nonumber \\\\ \\frac{\\sqrt{(2\\ell-1)(2\\ell+3)}(\\ell-1)(\\ell+2)}{(2\\ell+1)} \\, , \\hspace{1cm} \\end{eqnarray} noticing that $B_0$ is given in nanoGauss ($nG$). Thus, the computation of the correlation matrix leads to non-zero elements of the type $(\\ell, m)=(\\ell', m')$ and $(\\ell, m)=(\\ell' \\pm 2, m')$, while all the other elements are zero. In other words, the multipole matrix correlation is non-diagonal which is an inheritance of its angular dependence on $\\theta$ and $\\vec{v}_0$, shown in the eq.~(\\ref{Flutuacoes}), and this fact being a consequence of the presence of the magnetic f\\/ield at early times. In this scenario, correlations appear between dif\\/ferent scales, like $\\ell$ and $\\ell \\pm 2$,  and then the multipole coef\\/f\\/icients $a_{\\ell m}$ at such scales are correlated. For this reason, one concludes that large-angle anisotropy features should appear in the CMB maps produced by this primordial magnetic f\\/ield. If we consider a magnetic f\\/ield in a $\\Lambda$CDM universe the correlation matrix will be \\begin{eqnarray} \\label{Matrixtotal} \\langle a_{\\ell m} a^{\\ast}_{\\ell' m'} \\rangle \\,=\\, \\langle a_{\\ell m} a^{\\ast}_{\\ell' m'} \\rangle^{\\mbox{\\footnotesize $\\Lambda$CDM}} + \\langle a_{\\ell m} a^{\\ast}_{\\ell' m'} \\rangle^{\\textbf{B}} \\, , \\end{eqnarray} where \\begin{eqnarray}\\label{Clsss} \\langle a_{\\ell m} a^{\\ast}_{\\ell' m'} \\rangle^{\\mbox{\\footnotesize $\\Lambda$CDM}} \\,=\\, C^{\\mbox{\\footnotesize $\\Lambda$CDM}}_{\\ell} \\,\\delta_{m m'} \\,\\delta_{\\ell \\ell'} \\, . \\end{eqnarray} We observe that correlations between temperature f\\/luctuations do not only depend on the angular separation between two points, but also on their orientation with respect to the magnetic f\\/ield, that is, $m$-dependence. This $m$-dependence, caused by the fact that $B_0 \\ne 0$, means that now the CMB temperature f\\/luctuations are actually statistical\\/ly anisotropic.  Our purpose here is just to illustrate the ef\\/fect induced by these large-angle anisotropies in CMB maps, and compare them with those features found in the CMB maps from WMAP. To study the relationship between low-$\\ell$ (i.e., from $\\ell = 2$ to $\\ell = 10$) anomalies in such anisotropic magnetic f\\/ield scenario, we produce, according to eqs.~(\\ref{CorrelMat1})--(\\ref{Clsss}), f\\/ive sets (with dif\\/ferent strengths $B_0$) of Monte Carlo simulated CMB maps. For the statistical\\/ly isotropic part (that is, the f\\/irst term in eq.~(\\ref{Matrixtotal})) we use the CMBFAST tool to calculate the angular power spectrum $C^{\\Lambda \\mbox{\\sc cdm}}_{\\ell}$, and after that we use Cholesky decomposition of the matrix (\\ref{Matrixtotal}) in order to deal with the non-diagonality. Then we randomically simulate the $a_{\\ell m}$ coef\\/f\\/icient sets and from them we generate the CMB temperature f\\/luctuations maps. The simulations were performed for magnetic f\\/ields intensities $B_0 = (20 \\pm 10)\\, nG$, to be in agreement with known limits at cosmological scales for the $n=-5$ power spectrum (see, e.g.,~\\citealt{b101,b91,b104}, and~\\citealt{b106} for other power spectrum indices). ",
        "conclusions": "We investigated a plausible primordial scenario where a homogeneous magnetic f\\/ield acting on the photon-baryon f\\/luid at the recombination era, introduces a preferred direction that establish a statistical isotropy breaking. This setting of\\/fers a possible explanation for those anomalies found in WMAP data that seems to be associated to a preferred axis in the space, like the hemispherical NS-asymmetry  and the planarity of some CMB low-order multipoles (where such a plane is  perpendicular to that axis). It is found that this magnetic f\\/ield induces correlations between the CMB  multipoles, manifested through non-diagonal terms in the multipole correlation matrix (see eqs.~(\\ref{CorrelMat1})--(\\ref{CorrelMat3})). Accordingly, we have simulated f\\/ive sets of CMB skies considering an equal number of  strengths $B_0$ of the magnetic f\\/ield. These MC CMB maps, with multipoles in the range $\\ell = 2 - 10$, are used to investigate their large-angle correlations using our sigma-map method. With these data sets we performed a quantitative analysis of their large-angle signatures and perform a comparison with similar features computed from WMAP maps. For this, we consider a set of full-sky cleaned CMB maps, obtained from f\\/ive-year and three-year WMAP data releases, namely the ILC-5yr, KCN-5yr, ILC-3yr, OT-3yr, and PPG-3yr maps, which were dif\\/ferently processed in order to account for foregrounds and systematic errors.  Our results can be summarized as follows. F\\/irst, our analyses corroborate that an uneven hemispherical distribution in the power of the large-angular correlations, best known as NS-asymmetry phenomenon, is present in all these WMAP maps at more than 95\\% CL as compared with statistical\\/ly isotropic CMB maps produced according to the $\\Lambda$CDM cosmological model (i.e., $B_0=0$). Second, we found that this hemispherical asymmetry phenomenon, is such that higher is the value of the f\\/ield strenght $B_0$, greater is the probability that such asymmetry be present in the MC CMB maps, and be revealed with a high dipole value $S_{1}^{\\mbox{\\sc mc}}$ in its corresponding sigma-map-MC. In other words, our results show that the correlations introduced by a magnetic f\\/ield mechanism in low-order CMB multipoles, for a suf\\/f\\/iciently intense f\\/ield, reproduce the NS-asymmetry present in WMAP data as a common phenomena and not as an anomalous one. Third, as expected, the MC CMB maps with $B_0 \\ne 0$ exhibit the planarity ef\\/fect in their low-order multipoles (see, for illustration, f\\/igure~\\ref{fig5}) as quantif\\/ied in the Table~\\ref{table1}, where the $\\ell$-multipole power is concentrated in the $m = \\pm \\ell$ modes, and as evidenced by the intense red spots representative of strong angular correlations appearing in the equatorial region of the sigma-maps-MC (see, for il\\/lustration, f\\/igure~\\ref{fig2}). Fourth, as shown in the Table~\\ref{table1}, we found a significant correlation between the NS-asymmetry and quadrupole-octopole (planarity and) alignment phenomena. For completeness, we have also verif\\/ied that our results are robust under dif\\/ferent sets of parameters, as mentioned in the previous section, involved in the sigma-map method calculations.  The large angular scales CMB anomalies challenge the statistical isotropy expected in the $\\Lambda$CDM concordance model. Our results suggests that perhaps some of them could be suitable explained by just one physical phenomenon. As a matter of fact, there are several attempts to f\\/ind the origin of large-angle CMB anomalies, but further investigations are needed to fully comprehend if they have one or more causes. In this sense, statistical\\/ly anisotropic models explaining several anomalies with a minimum set of hipothesis should be explored.  We believe that our results shall motivate the study of other statistical\\/ly anisotropic scenarios, in particular, those where correlations between CMB anomalies appear. The magnetic f\\/ield scenario analysed here is just the simplest one, and considering more realistic f\\/ields new ef\\/fects could be present on the CMB temperature and polarization data (see, e.g.,~\\citealt{b95,b96,b97,b102,b100,b98,b99}, and references therein). In the same form, as well, possible relationships between these new ef\\/fects and statistical\\/ly anisotropic phenomena found in WMAP data could be established, all this deserving a better investigation."
    },
    "0807/0807.4653_arXiv.txt": {
        "abstract": "We give an overview of ISINA: INTEGRAL Source Identification Network Algorithm. This machine learning algorithm, using Random Forests, is applied to the IBIS/ISGRI dataset in order to ease the production of unbiased future soft gamma-ray source catalogues. First we introduce the dataset and the problems encountered when dealing with images obtained using the coded mask technique. The initial step of source candidate searching is introduced and an initial candidate list is created. A description of the feature extraction on the initial candidate list is then performed together with feature merging for these candidates. Three training and testing sets are created in order to deal with the diverse timescales encountered when dealing with the gamma-ray sky. Three independent Random Forest are built: one dealing with faint persistent source recognition, one dealing with strong persistent sources and a final one dealing with transients. For the latter, a new transient detection technique is introduced and described: the Transient Matrix. Finally the performance of the network is assessed and discussed using the testing set and some illustrative source examples. ",
        "introduction": "Since its launch in 2002, the INTEGRAL (International Gamma-Ray Astrophysics Laboratory) satellite has carried out more than 5 years of observations in the energy range from 5 keV to 10 MeV. In particular the IBIS (Imager on Board INTEGRAL spacecraft) imaging instrument (\\citealt{ubertini03}) has been optimised for survey work, with a large ($30^o$) field of view with excellent imaging and spectroscopic capability, and has formed the basis of several previous INTEGRAL surveys. Because of the success of INTEGRAL, with longer exposure times and larger sky coverage, the amount of data accumulated over the mission lifetime has dramatically increased, and with it the number of detected objects (see Fig.\\ref{fig:source_dist}). Moreover different types of objects are also being detected, with a high ratio of unknowns.  Conventional methods for the production of source catalogues are becoming less and less adequate, requiring the use of novel techniques for the reduction of visual inspection required for the compilation of future catalogues. Moreover novel algorithms for classification and identification will produce more objective and homogeneous sets compared to the subjective samples obtained through visual inspection.  In this paper we will describe the implementation of a supervised machine learning algorithm which has been developed for the purpose of source identification on the IBIS/ISGRI images based on Random Forests (\\citealt{randfo}). In machine learning, a Random Forest is a classifier that builds many single decision trees, and uses these to output a predicted class for a given input. Each single decision tree produces a predicted class for a given input, whilst the Random Forest determines the mode of all these small predictors to produce a more accurate classifier. The algorithm will be described in more detail, but in order to fully appreciate it the data being classified must be thoroughly understood.  \\begin{figure} \\centering \\includegraphics[width=0.5\\textwidth]{INTRO_PLOT} \\caption{Numbers of sources in the first, second, and third IBIS/ISGRI catalogues, classified by type: Cataclysmic Variables (CV), Active Galactic Nuclei (AGN), High Mass X-ray Binaries (HMXB), Low Mass X-ray Binaries (LMXB), Other and Unknowns (?)} \\label{fig:source_dist} \\end{figure}  In Fig.\\ref{fig:section_flow} we illustrate how the paper is divided into sections following the procedural flow of the algorithm. In section 2 we will give an overview of the dataset used in this work, and in section 3 we describe the methods employed in order to obtain an initial candidate list to be used by the classification network. Section 4 will describe the process of extracting features for the candidates from the IBIS/ISGRI images and the different methods to merge the features for them to be useful in a classification network. Section 5 will deal with constructing a reliable training set for our classification network to train on and a corresponding test set in order to asses the recognition performance of the network. We will describe the main Random Forest algorithm in section 6 in the context of our training and testing sets. The assessment and discussion, together with some examples of how the algorithm performs on specific objects is presented in sections 7, 8 and 9.  \\begin{figure*} \\centering \\includegraphics[width=0.75\\textwidth]{section_flow_3} \\caption{Flow diagram showing the levels involved in the classification procedure. The sections in which the methods are described are indicated in red.} \\label{fig:section_flow} \\end{figure*} ",
        "conclusions": "We have developed a reliable algorithm to aid the production of future IBIS/ISGRI gamma-ray survey catalogues. The algorithm will help produce less subjective catalogues, unbiased by human intervention. Meant for source identification, ISINA has also turned out useful in discriminating source types. We have shown how to automate the task of selecting and reducing a set of candidates from IBIS/ISGRI images.  The distribution of recovered objects, sorted by type, together with the objects published in the $3^{rd}$ catalogue present in the testing set are shown in Fig.\\ref{fig:types}. It is clear that the majority of objects are recovered correctly with a $50\\%$ global vote threshold. It is interesting to note that the only populations to suffer from a substantial decrease in recovered objects are the CVs and the unknown source types. The drop in the number of CVs can easily be explained by the fact that most of the non-recovered ones lie in crowded regions, where the systematic noise is greatest. However, we point out that missing objects in these regions does not cause a big problem for the creation of future catalogues. This is because crowded regions will be the most inspected ones, so that if the network does not recover certain objects, human intervention will. The other population to have a significant drop in the number of recovered objects is the unknown category which is a bit less trivial to assess. This is because, by definition, the only real way to determine their nature is to have longer exposure times for the regions where these are present. We also point out that both the CV population and the unknown one was not part of our training set, which might also explain the relatively low recovered rate for these. This may also have an impact in our final result as the three networks have now specialised in recovering their subclass of objects. One last observation of the general behavior of the network on the testing set is that despite the fact that the remaining classes perform well within the network, it has to be pointed out that our training and testing sets might have misclassified objects within them (false positives and true negatives). Given the nature of the classification task, the training set will always be biased towards this. However given the extremely fast data growth the problem can only get better, and these small discrepancies will systematically reduce.  We would also like to point out the potential of such a network for exploratory data analysis in other wavelengths. The networks described here can be easily tuned to deal with different images, taken from different observatories. The features defined are quite generic, and anyway may be adjusted according to the new dataset, probably suffering from different systematic effects than the ones presented here.  Finally we would like to stress the need for such algorithms in contemporary astronomy. The network presented here has been constructed to deal with a particular dataset: namely the IBIS/ISGRI one, however multi-wavelength information might well prove useful, not only in discriminating real/fake excess but also in identifying source types. Moreover as the astronomical data size is increasing exponentially with the advent of more survey projects, larger detector areas, finer resolutions and larger fields of view the need for such algorithms is ever more pressing as human intervention for such tasks is becoming less and less feasible.  \\begin{figure} \\centering \\includegraphics[width=0.5\\textwidth]{CONCLUSION_FIN} \\caption{Number of sources in the testing set classified by type (red) and recovered objects using a $50\\%$ global vote cut (blue). Note that objects in the transient category are also present in there respective class types.} \\label{fig:types} \\end{figure}"
    },
    "0807/0807.1240_arXiv.txt": {
        "abstract": "% We report a coordinated multi-band photometry of the RS Oph 2006 outburst and highlight the emission line free $y$-band photometry that shows a mid-plateau phase at $y \\sim 10.2$ mag from day 40 to day 75 after the discovery followed by a sharp drop of the final decline. Such mid-plateau phases are observed in other two recurrent novae, U Sco and CI Aql, and are interpreted as a bright disk irradiated by the white dwarf.  We have calculated theoretical light curves based on the optically thick wind theory and have reproduced the early decline, mid-plateau phase, and final decline.  The final decline is identified with the end of steady hydrogen shell-burning, which turned out at about day 80. This turnoff date is consistent with the end of a supersoft X-ray phase observed with {\\it Swift}. Our model suggests a white dwarf mass of $1.35 \\pm 0.01 ~M_\\odot$, which indicates that RS Oph is a progenitor of Type Ia supernovae. We strongly recommend the $y$-filter observation of novae to detect both the presence of a disk and the hydrogen burning turnoff. Observational data of $y$ magnitudes are provided here together with other multi-wavelength light curve data. ",
        "introduction": "It is well known that the X-ray turnoff time is a good indicator of the white dwarf (WD) mass \\citep[e.g.,][]{hac06ka, hac06kb}. When the hydrogen shell-burning atop the WD extinguishes, a supersoft X-ray phase ends \\citep[e.g.,][]{kra96}.  In a visual light curve, however, this turnoff is not clear because many strong emission lines contribute to it.  To avoid such contamination to the continuum flux, we have observed RS Oph with the Str\\\"omgren $y$-band filter.  The $y$-filter is an intermediate bandpass filter designed to cut the strong emission lines in the wide $V$ bandpass filter, so that its light curve represents the continuum flux of novae \\citep[e.g.,][]{hac06kb}.  We have further modeled the light curve of RS Oph and have estimated the WD mass by fitting our modeled light curve with the observation.  Optical observations were started just after the discovery of the 2006 outburst \\citep{nar06}.  Each observation are listed in Tables 1 and 2.  We have put a special emphasis on  the Str\\\"omgren $y$ filter to avoid contamination by the strong emission lines.  These $y$ filters were made by Custom Scientific, Inc.,\\footnote{See http://www.customscientific.com.} and distributed to each observer by one of the authors (M. Kato). Kiyota, Kubotera, Maehara, and Nakajima (VSOLJ members) started observation on 2006 February 13 and obtained data for $y$ magnitudes from 2006 February 17 until 2007 June 20. The Osaka Kyoiku University (OKU) team obtained $V$ and $y$ magnitudes of 42 nights starting from 2006 February 17 until 2006 November 17. The magnitudes of this object were measured by using the local standard star, TYC2 5094.92.1 (Kiyota) or TYC2 5094.283.1 (the other observers). We adapted the brightness and color of ($y= V= 9.57$, $B-V= 0.56$) for TYC2 5094.92.1 and ($y= V= 9.35$, $B-V=1.23$) for TYC2 5094.283.1 from the Tycho2 catalog.  \\begin{figure}[!h] \\plotone{fig1.epsi} \\caption{ $y$ ({\\it open circles})and $Ic$ ({\\it open squares}) light curves for the RS Oph (2006).  Thick solid lines are combined free-free and blackbody. } \\end{figure}   \\begin{table}[!ht] \\label{oku-data} \\caption{Optical Data of Osaka Kyoiku Univ.} \\smallskip \\begin{center} {\\tiny \\begin{tabular}{llllllll} \\tableline \\noalign{\\smallskip} Days & JD & $y$ & err. & obs. & $V$ & err. & obs.\\\\ \\noalign{\\smallskip} \\tableline \\noalign{\\smallskip} 6.33 & 3784.33 & 6.98 & 0.02 & 11 & 6.47 &  0.01 &  6 \\\\ 16.30 & 3794.30 & 8.26 & 0.01 & 14 & 7.61 &  0.01 &  5 \\\\ 24.25 & 3802.25 & 8.83 & 0.01 & 11 & 8.02 &  0.02 & 10 \\\\ 31.30 & 3809.30 & 9.48 & 0.01 & 10 & 8.76 &  0.01 & 10 \\\\ 34.27 & 3812.27 & 9.65 & 0.01 & 10 & 8.90 &  0.02 & 10 \\\\ 37.24 & 3815.24 & 9.69 & 0.02 & 10 & 8.94 &  0.03 & 10 \\\\ 41.25 & 3819.25 & 9.92 & 0.01 & 10 & 9.22 &  0.03 & 10 \\\\ 48.22 & 3826.22 & 10.08 & 0.01 & 10 & 9.39 &  0.02 & 10 \\\\ 51.21 & 3829.21 & 10.07 & 0.01 & 11 & 9.45 &  0.01 & 10 \\\\ 55.22 & 3833.22 & 10.11 & 0.01 & 10 & 9.54 &  0.01 & 10 \\\\ 72.19 & 3850.19 & 10.20 & 0.01 & 10 & 9.80 &  0.02 & 10 \\\\ 73.19 & 3851.19 & 10.28 & 0.02 & 10 & 9.82 &  0.01 & 10 \\\\ 78.14 & 3856.14 & 10.36 & 0.02 & 10 & 9.98 &  0.02 & 10 \\\\ 80.13 & 3858.13 & 10.54 & 0.01 & 10 & 10.08 &  0.01 & 10 \\\\ 83.13 & 3861.13 & 10.68 & 0.03 & 10 & 10.23 &  0.02 & 10 \\\\ 98.09 & 3876.09 & 11.47 & 0.03 & 10 & 11.00 &  0.01 & 10 \\\\ 99.13 & 3877.13 & 11.47 & 0.01 & 10 & 11.03 &  0.01 & 10 \\\\ 102.03 & 3880.03 & 11.52 & 0.02 & 11 & 11.10 &  0.03 & 10 \\\\ 108.08 & 3886.08 & 11.88 & 0.07 & 10 & 11.40 &  0.02 & 10 \\\\ 112.17 & 3890.17 & 11.93 & 0.04 & 10 & 11.45 &  0.02 & 10 \\\\ 115.17 & 3893.17 & 11.85 & 0.04 & 10 & 11.38 &  0.02 & 10 \\\\ 122.15 & 3900.15 & 11.80 & 0.03 & 10 & 11.41 &  0.02 & 10 \\\\ 128.16 & 3906.16 & 11.99 & 0.06 &  5 & 11.56 &  0.11 &  4 \\\\ 152.04 & 3930.04 & 12.01 & 0.06 & 10 & 11.71 &  0.03 &  9 \\\\ 153.15 & 3931.15 & 11.97 & 0.03 & 10 & 11.69 &  0.04 & 10 \\\\ 177.01 & 3955.01 & 11.68 & 0.02 & 10 & 11.47 &  0.03 & 10 \\\\ 184.04 & 3962.04 & 11.83 & 0.02 & 10 & 11.62 &  0.02 & 10 \\\\ 192.96 & 3970.96 & 11.67 & 0.02 & 10 & 11.48 &  0.02 & 10 \\\\ 220.93 & 3998.93 & 11.53 & 0.02 & 10 & 11.38 &  0.01 & 10 \\\\ 227.97 & 4005.97 & 11.70 & 0.04 & 10 & 11.51 &  0.02 & 10 \\\\ 238.95 & 4016.95 & 11.68 & 0.03 & 10 & 11.49 &  0.03 & 10 \\\\ 247.90 & 4025.90 & 11.69 & 0.03 & 10 & 11.60 &  0.03 & 10 \\\\ 248.90 & 4026.90 & 11.76 & 0.04 & 10 & 11.61 &  0.02 & 10 \\\\ 257.93 & 4035.93 & 12.00 & 0.04 & 48 & --  & --  & -- \\\\ 261.90 & 4039.90 & 11.90 & 0.03 & 10 & 11.70 &  0.05 & 10 \\\\ 264.88 & 4042.88 & 11.67 & 0.02 & 10 & -- & -- & -- \\\\ 266.94 & 4044.94 & 11.54 & 0.07 & 7  & -- & -- & -- \\\\ 267.91 & 4045.91 & 11.61 & 0.16 & 6  & -- & -- & -- \\\\ 268.90 & 4046.90 & 11.50 & 0.03 & 10 & -- & -- & -- \\\\ 269.88 & 4047.88 & 11.33 & 0.03 & 10 & -- & -- & -- \\\\ 270.88 & 4048.88 & 11.30 & 0.02 & 10 & -- & -- & -- \\\\ 278.86 & 4056.86 & 11.42 & 0.05 & 10 & -- & -- & -- \\\\ \\noalign{\\smallskip} \\tableline\\\\ \\end{tabular} } \\end{center} \\end{table}    \\begin{table}[!ht] \\label{vsolj-data} \\caption{Optical Data of VSOLJ} \\smallskip \\begin{center} {\\tiny \\begin{tabular}{llllllll} \\tableline \\noalign{\\smallskip} Days & JD & $B$ & $V$ & $y$ & vis. & $R$ & $I_c$\\\\ \\noalign{\\smallskip} \\tableline \\noalign{\\smallskip} 1.36 & 779.36 & - & - & - & 4.46 & - & - \\\\ 2.33 & 780.33 & - & 5.23 & - & 5.11 & 4.46 & 3.74 \\\\ 3.36 & 781.36 & - & 5.73 & - & 5.61 & - & - \\\\ 4.33 & 782.33 & - & - & - & 5.70 & - & - \\\\ 5.36 & 783.36 & - & 6.43 & - & 6.32 & - & - \\\\ 6.35 & 784.35 & 7.53 & 6.80 & 6.89 & 6.70 & 6.06 & - \\\\ 7.33 & 785.33 & - & - & - & 6.93 & - & - \\\\ 9.33 & 787.33 & - & - & - & 7.10 & - & - \\\\ 10.35 & 788.35 & - & - & - & 7.10 & - & - \\\\ 11.32 & 789.32 & 8.18 & 7.46 & 7.78 & 7.16 & 6.04 & 5.86 \\\\ 12.34 & 790.34 & 8.33 & 7.68 & 7.93 & - & - & 6.00 \\\\ 13.31 & 791.31 & 8.54 & 7.66 & 7.99 & 7.43 & 6.05 & 5.99 \\\\ 15.33 & 793.33 & - & - & - & 8.00 & - & - \\\\ 16.33 & 794.33 & 8.85 & 8.00 & 8.25 & - & 6.11 & - \\\\ 19.29 & 797.29 & 8.99 & 8.29 & 8.41 & - & 6.24 & - \\\\ 20.30 & 798.30 & 9.08 & 8.37 & 8.71 & 8.10 & 6.55 & 6.71 \\\\ 21.30 & 799.30 & 8.97 & 8.38 & 8.63 & 8.10 & 6.61 & 6.78 \\\\ 22.29 & 800.29 & 9.22 & 8.39 & 8.69 & - & 6.97 & 6.84 \\\\ 23.26 & 801.26 & - & - & - & 8.30 & - & - \\\\ 24.27 & 802.27 & 9.21 & 8.53 & 8.88 & - & 6.75 & 6.96 \\\\ 27.30 & 805.30 & - & 8.82 & 9.14 & - & - & - \\\\ 29.31 & 807.31 & 9.68 & 9.03 & 9.25 & - & 6.91 & - \\\\ 30.28 & 808.28 & 9.70 & 8.97 & 9.29 & 8.10 & 7.15 & 7.36 \\\\ 31.26 & 809.26 & 9.91 & 9.18 & 9.51 & - & 7.31 & 7.54 \\\\ 32.29 & 810.29 & 9.89 & 9.08 & 9.44 & - & 7.72 & - \\\\ 34.25 & 812.25 & 10.14 & 9.22 & 9.63 & - & 7.48 & 7.69 \\\\ 35.32 & 813.32 & 10.14 & 9.46 & 9.62 & - & 7.33 & - \\\\ 36.24 & 814.24 & 10.13 & 9.26 & 9.65 & 9.60 & 7.58 & 7.74 \\\\ 37.26 & 815.26 & 10.11 & 9.29 & 9.68 & - & 7.53 & 7.74 \\\\ 38.25 & 816.25 & 10.10 & 9.41 & 9.72 & 9.50 & 7.78 & 7.75 \\\\ 40.27 & 818.27 & 10.30 & 9.56 & 9.84 & - & 7.72 & 7.84 \\\\ 41.28 & 819.28 & 10.44 & 9.80 & 9.95 & 9.40 & 7.64 & - \\\\ 42.27 & 820.27 & - & - & - & 9.50 & - & - \\\\ 46.26 & 824.26 & 10.44 & 9.60 & 9.95 & - & 8.00 & 8.07 \\\\ 47.25 & 825.25 & 10.60 & 9.93 & 10.06 & 9.90 & 7.84 & - \\\\ 48.28 & 826.28 & 10.53 & 9.86 & 10.03 & - & 7.99 & 8.06 \\\\ 50.26 & 828.26 & 10.56 & 9.88 & 10.06 & 9.90 & 8.05 & 8.10 \\\\ 51.24 & 829.24 & 10.64 & 9.93 & 10.09 & 10.10 & 7.93 & - \\\\ 53.23 & 831.23 & 10.61 & 9.84 & 10.08 & - & 8.20 & 8.27 \\\\ 54.22 & 832.22 & 10.70 & 9.95 & 10.08 & - & 8.02 & - \\\\ 55.22 & 833.22 & 10.71 & 9.98 & 10.13 & - & 8.04 & - \\\\ 56.27 & 834.27 & 10.69 & 9.95 & 10.10 & - & 8.21 & 8.14 \\\\ 61.21 & 839.21 & 10.87 & 10.06 & 10.23 & - & 8.23 & - \\\\ 63.18 & 841.18 & 10.72 & 9.89 & 10.17 & - & 8.58 & 8.32 \\\\ 64.22 & 842.22 & 10.86 & 10.07 & 10.19 & 10.20 & 8.23 & - \\\\ 65.26 & 843.26 & 10.77 & 10.00 & - & - & 8.49 & 8.25 \\\\ 68.22 & 846.22 & 10.99 & 10.16 & 10.23 & - & 8.38 & - \\\\ 69.25 & 847.25 & 10.92 & 10.10 & 10.23 & - & 8.46 & 8.30 \\\\ 72.18 & 850.18 & 11.07 & 10.21 & 10.33 & - & 8.44 & - \\\\ 73.26 & 851.26 & 10.93 & 10.08 & 10.23 & - & 8.63 & 8.33 \\\\ 76.27 & 854.27 & 11.04 & 10.17 & 10.27 & - & 8.67 & 8.42 \\\\ 78.24 & 856.24 & 11.13 & 10.26 & 10.44 & - & 8.78 & 8.48 \\\\ 80.17 & 858.17 & 10.97 & 10.33 & 10.52 & - & 8.85 & 8.65 \\\\ 82.25 & 860.25 & 11.36 & 10.49 & 10.62 & 10.80 & 8.89 & 8.66 \\\\ 89.16 & 867.16 & 11.81 & 11.01 & 11.16 & - & 9.29 & - \\\\ 92.16 & 870.16 & 11.99 & 11.13 & 11.34 & - & 9.51 & - \\\\ 98.15 & 876.15 & 12.15 & 11.20 & 11.45 & - & 9.88 & 9.39 \\\\ 99.19 & 877.19 & 12.30 & 11.34 & 11.57 & 11.20 & 9.83 & - \\\\ \\noalign{\\smallskip} \\tableline\\\\ \\end{tabular} } \\end{center} \\end{table}   \\setcounter{table}{1} \\begin{table}[!ht] \\caption{Optical Data of VSOLJ (continued)} \\smallskip \\begin{center} {\\tiny \\begin{tabular}{llllllll} \\tableline \\noalign{\\smallskip} Days & JD & $B$ & $V$ & $y$ & vis. & $R$ & $I_c$\\\\ \\noalign{\\smallskip} \\tableline \\noalign{\\smallskip} 100.08 & 878.08 & 12.24 & 11.14 & 11.49 & - & 10.02 & 9.44 \\\\ 102.13 & 880.13 & 12.40 & 11.43 & 11.67 & 11.10 & 9.93 & - \\\\ 106.12 & 884.12 & 12.64 & 11.69 & 11.92 & - & 10.22 & - \\\\ 108.18 & 886.18 & 12.69 & 11.65 & 11.96 & - & 10.33 & 9.69 \\\\ 109.17 & 887.17 & 12.68 & 11.65 & 11.99 & - & 10.35 & 9.66 \\\\ 110.09 & 888.09 & - & - & - & 11.80 & - & - \\\\ 111.14 & 889.14 & - & - & - & 11.40 & - & - \\\\ 114.09 & 892.09 & 12.79 & 11.48 & 12.47 & - & 10.45 & 9.65 \\\\ 115.12 & 893.12 & - & - & - & 11.70 & - & - \\\\ 127.18 & 905.18 & - & - & - & 11.90 & - & - \\\\ 128.07 & 906.07 & 12.83 & 11.79 & 11.99 & - & 10.41 & - \\\\ 129.06 & 907.06 & 12.99 & 11.86 & 12.06 & - & 10.49 & - \\\\ 133.05 & 911.05 & 13.09 & 11.76 & 12.25 & - & 10.67 & 9.78 \\\\ 136.13 & 914.13 & 13.05 & 12.12 & 12.33 & - & 10.71 & 9.65 \\\\ 137.99 & 915.99 & 13.22 & 12.18 & 12.31 & - & 10.78 & - \\\\ 140.04 & 918.04 & 13.22 & 12.17 & 12.37 & - & 10.77 & - \\\\ 141.99 & 919.99 & 13.22 & 12.21 & 12.39 & - & 10.77 & - \\\\ 152.09 & 930.09 & 13.12 & 11.86 & 12.07 & 12.20 & 10.60 & 9.60 \\\\ 153.18 & 931.18 & 13.17 & 11.76 & 11.97 & - & 10.48 & 9.58 \\\\ 155.00 & 933.00 & 13.13 & - & 12.07 & - & 10.57 & - \\\\ 164.00 & 942.00 & 13.24 & 12.14 & 12.31 & - & 10.83 & - \\\\ 165.00 & 943.00 & 13.28 & 12.01 & 12.19 & - & 10.85 & 9.71 \\\\ 166.00 & 944.00 & 13.24 & 12.08 & 12.26 & - & 10.82 & - \\\\ 167.07 & 945.07 & 13.18 & 11.91 & 12.02 & - & 10.71 & 9.64 \\\\ 168.02 & 946.02 & 13.18 & 12.01 & 12.19 & 11.90 & 10.74 & - \\\\ 172.00 & 950.00 & 13.05 & 11.79 & 11.92 & - & 10.57 & - \\\\ 173.05 & 951.05 & 13.02 & 11.70 & 11.80 & - & 10.49 & 9.44 \\\\ 174.00 & 952.00 & 12.97 & 11.69 & 11.90 & - & 10.54 & 9.47 \\\\ 175.04 & 953.04 & 12.91 & 11.59 & 11.69 & - & 10.44 & 9.39 \\\\ 176.10 & 954.10 & 12.98 & 11.53 & 11.59 & - & 10.41 & 9.37 \\\\ 178.99 & 956.99 & 12.95 & 11.66 & 11.76 & - & 10.49 & - \\\\ 180.03 & 958.03 & 12.94 & 11.56 & 11.74 & - & 10.52 & 9.46 \\\\ 182.15 & 960.15 & - & 11.55 & - & - & 10.49 & 9.44 \\\\ 182.99 & 960.99 & 13.00 & 11.77 & 11.85 & - & 10.58 & - \\\\ 184.00 & 962.00 & 12.97 & 11.71 & 11.81 & - & 10.62 & 9.50 \\\\ 188.11 & 966.11 & 13.00 & 11.71 & 11.86 & - & 10.60 & 9.51 \\\\ 190.00 & 968.00 & 13.04 & 11.79 & 11.87 & - & 10.59 & - \\\\ 192.98 & 970.98 & 12.86 & 11.64 & 11.76 & - & 10.49 & - \\\\ 193.99 & 971.99 & 12.82 & 11.56 & 11.70 & - & 10.43 & - \\\\ 199.95 & 977.95 & 12.88 & 11.68 & 11.79 & - & 10.53 & - \\\\ 201.03 & 979.03 & - & 11.43 & - & - & 10.43 & 9.55 \\\\ 202.02 & 980.02 & 12.92 & 11.85 & 11.88 & - & 10.62 & 9.51 \\\\ 203.06 & 981.06 & 12.95 & 11.77 & 11.92 & 11.90 & 10.64 & 9.51 \\\\ 204.96 & 982.96 & 13.08 & 11.97 & 12.02 & - & 10.75 & - \\\\ 205.94 & 983.94 & 13.18 & 12.02 & 12.13 & - & 10.81 & - \\\\ 210.95 & 988.95 & 13.39 & 12.20 & 12.28 & - & 10.89 & 9.74 \\\\ 216.00 & 994.00 & - & 11.51 & - & - & - & 9.41 \\\\ 218.96 & 996.96 & 12.97 & 11.49 & 11.55 & - & 10.35 & 9.36 \\\\ 219.95 & 997.95 & 12.93 & 11.49 & 11.60 & - & 10.36 & - \\\\ 220.93 & 998.93 & 12.97 & 11.47 & 11.61 & - & 10.35 & - \\\\ 223.92 & 1001.92 & 12.91 & 11.44 & 11.46 & - & 10.33 & 9.31 \\\\ 224.92 & 1002.92 & 12.97 & 11.52 & 11.64 & - & 10.42 & - \\\\ 227.93 & 1005.93 & 12.95 & 11.46 & 11.61 & 11.70 & 10.46 & 9.40 \\\\ 230.94 & 1008.94 & - & - & - & 11.80 & - & - \\\\ 236.90 & 1014.90 & 13.03 & 11.66 & 11.81 & - & 10.51 & - \\\\ 237.94 & 1015.94 & 12.96 & 11.54 & 11.62 & 11.80 & 10.44 & 9.40 \\\\ 238.89 & 1016.89 & - & - & - & 11.90 & - & - \\\\ 239.91 & 1017.91 & 12.89 & 11.50 & 11.53 & - & 10.43 & 9.36 \\\\ 244.93 & 1022.93 & 12.91 & 11.56 & 11.55 & - & 10.44 & 9.37 \\\\ 245.91 & 1023.91 & 12.92 & 11.59 & 11.61 & - & 10.48 & 9.41 \\\\ \\noalign{\\smallskip} \\tableline\\\\ \\end{tabular} } \\end{center} \\end{table}   \\setcounter{table}{1} \\begin{table}[!ht] \\caption{Optical Data of VSOLJ (continued)} \\smallskip \\begin{center} {\\tiny \\begin{tabular}{llllllll} \\tableline \\noalign{\\smallskip} Days & JD & $B$ & $V$ & $y$ & vis. & $R$ & $I_c$\\\\ \\noalign{\\smallskip} \\tableline \\noalign{\\smallskip} 251.90 & 1029.90 & - & - & - & 11.90 & - & - \\\\ 254.88 & 1032.88 & - & - & - & 11.90 & - & - \\\\ 255.88 & 1033.88 & - & - & - & 11.80 & - & - \\\\ 259.90 & 1037.90 & - & - & - & 11.90 & - & - \\\\ 265.88 & 1043.88 & 12.84 & 11.62 & 11.60 & - & 10.59 & - \\\\ 268.90 & 1046.90 & 12.70 & 11.45 & 11.42 & - & 10.42 & 9.44 \\\\ 269.87 & 1047.87 & 12.77 & 11.35 & 11.42 & - & 10.31 & - \\\\ 270.87 & 1048.87 & 12.70 & 11.29 & 11.36 & - & 10.24 & - \\\\ 271.85 & 1049.85 & 12.72 & 11.31 & 11.37 & - & 10.27 & - \\\\ 273.89 & 1051.89 & 12.66 & 11.26 & 11.26 & - & 10.21 & 9.25 \\\\ 274.87 & 1052.87 & 12.78 & 11.29 & 11.35 & - & 10.26 & - \\\\ 276.87 & 1054.87 & 12.88 & 11.38 & 11.51 & - & 10.36 & - \\\\ 330.36 & 1108.36 & - & 11.42 & 11.50 & - & 10.49 & 9.37 \\\\ 336.34 & 1114.34 & - & 11.36 & 11.45 & - & 10.49 & 9.40 \\\\ 337.32 & 1115.32 & - & 11.11 & - & - & 10.35 & 9.29 \\\\ 341.33 & 1119.33 & - & 11.21 & 11.29 & - & 10.34 & 9.30 \\\\ 345.33 & 1123.33 & - & 11.44 & - & - & 10.52 & 9.47 \\\\ 347.32 & 1125.32 & - & 11.61 & 11.70 & - & 10.65 & 9.52 \\\\ 355.35 & 1133.35 & 13.15 & 11.96 & 12.08 & - & 11.01 & 9.85 \\\\ 357.32 & 1135.32 & - & 12.02 & 12.11 & - & 11.15 & 9.96 \\\\ 358.32 & 1136.32 & - & 11.79 & 11.81 & - & 11.01 & 9.81 \\\\ 360.33 & 1138.33 & 12.62 & 11.80 & 11.91 & - & 10.96 & 9.83 \\\\ 361.32 & 1139.32 & - & 11.81 & 11.81 & 12.00 & 11.03 & 9.89 \\\\ 365.32 & 1143.32 & - & 11.36 & 11.42 & - & 10.64 & 9.57 \\\\ 366.35 & 1144.35 & - & - & - & 12.00 & - & - \\\\ 368.35 & 1146.35 & 12.74 & 11.54 & 11.56 & 12.00 & 10.58 & - \\\\ 369.31 & 1147.31 & - & 11.32 & 11.38 & - & 10.54 & 9.43 \\\\ 370.32 & 1148.32 & 12.14 & 11.40 & 11.32 & - & 10.57 & 9.45 \\\\ 372.35 & 1150.35 & 12.60 & 11.34 & 11.38 & - & 10.33 & - \\\\ 374.35 & 1152.35 & 12.74 & 11.44 & 11.45 & - & 10.43 & - \\\\ 375.29 & 1153.29 & - & 11.30 & 11.34 & - & 10.44 & 9.36 \\\\ 377.36 & 1155.36 & - & - & - & 11.70 & - & - \\\\ 378.32 & 1156.32 & 12.30 & 11.15 & 11.17 & - & 10.30 & 9.24 \\\\ 380.30 & 1158.30 & - & 11.12 & 11.24 & - & 10.26 & 9.23 \\\\ 381.30 & 1159.30 & 12.53 & 11.26 & - & - & 10.27 & - \\\\ 382.31 & 1160.31 & 12.67 & 11.46 & 11.49 & 11.80 & 10.53 & 9.38 \\\\ 383.32 & 1161.32 & 12.37 & 11.26 & 11.26 & - & 10.29 & 9.21 \\\\ 388.27 & 1166.27 & 12.74 & 11.51 & 11.58 & - & 10.55 & - \\\\ 389.28 & 1167.28 & 12.69 & 11.52 & 11.56 & - & 10.58 & - \\\\ 390.26 & 1168.26 & 12.71 & 11.50 & 11.55 & - & 10.55 & - \\\\ 391.27 & 1169.27 & 12.94 & 11.68 & 11.59 & - & 10.71 & - \\\\ 394.20 & 1172.20 & 12.65 & 11.75 & 11.54 & - & 10.56 & - \\\\ 400.28 & 1178.28 & 12.63 & 11.45 & 11.47 & - & 10.43 & - \\\\ 409.29 & 1187.29 & - & - & - & 11.60 & - & - \\\\ 416.22 & 1194.22 & 12.63 & 11.45 & 11.49 & - & 10.47 & - \\\\ 418.23 & 1196.23 & 12.66 & 11.35 & 11.37 & - & 10.37 & - \\\\ 422.20 & 1200.20 & 12.45 & 11.24 & 11.20 & - & 10.23 & - \\\\ 424.21 & 1202.21 & 12.53 & 11.25 & 11.26 & - & 10.27 & - \\\\ 432.30 & 1210.30 & 12.56 & 11.20 & 11.17 & - & 10.12 & 9.10 \\\\ 442.25 & 1220.25 & 12.26 & 10.95 & 11.05 & - & 10.02 & - \\\\ 450.14 & 1228.14 & 12.33 & 11.16 & 11.15 & - & 10.22 & - \\\\ 454.28 & 1232.28 & 12.33 & 11.14 & 11.17 & - & 10.23 & - \\\\ 460.24 & 1238.24 & 12.34 & 11.22 & 11.13 & 11.30 & 10.23 & - \\\\ 462.25 & 1240.25 & - & 10.95 & - & - & - & - \\\\ 463.26 & 1241.26 & 12.19 & 10.98 & 10.95 & - & 10.06 & - \\\\ 474.08 & 1252.08 & 12.23 & 10.99 & 10.99 & - & 10.09 & - \\\\ 479.12 & 1257.12 & - & - & - & 11.20 & - & - \\\\ 480.12 & 1258.12 & 12.49 & 11.15 & 11.21 & - & 10.21 & - \\\\ 483.15 & 1261.15 & - & - & - & 10.60 & - & - \\\\ 485.19 & 1263.19 & 12.11 & 10.91 & - & - & - & - \\\\ 489.24 & 1267.24 & - & 11.32 & - & - & - & - \\\\ 494.17 & 1272.17 & 12.57 & 11.40 & - & - & 10.43 & - \\\\ \\noalign{\\smallskip} \\tableline\\\\ \\end{tabular} } \\end{center} \\end{table} ",
        "conclusions": ""
    },
    "0807/0807.2243_arXiv.txt": {
        "abstract": "We present high resolution submillimeter interferometric imaging of two of the brightest high--redshift submillimeter galaxies known: GN20 and AzTEC1 at 0.8 and 0.3 arcsec resolution respectively.  Our  data -- the highest resolution submillimeter imaging of high redshift sources accomplished to date -- was collected in three different array configurations: compact, extended, and very extended.We derive angular sizes of 0.6 and 1.0 arcsec for GN20 and 0.3 and 0.4 arcsec for AzTEC1 from modeling their visibility functions as a Gaussian and elliptical disk respectively.  Because both sources are B--band dropouts, they likely lie within a relatively narrow redshift window around $z\\sim 4$, which indicates their angular extent corresponds to physical scales of 4-8 and 1.5-3 kpc respectively for the starburst region.  By way of a series of simple assumptions, we find preliminary evidence that these hyperluminous starbursts -- with star formation rates $>1000$ $M_\\odot$ yr$^{-1}$ -- are radiating at or close to their Eddington limit.  Should future high resolution observations indicate that these two objects are typical of a population of high redshift Eddington--limited starbursts, this could have important consequences for models of star formation and feedback in extreme environments. ",
        "introduction": "\\label{sec:intro}  \\begin{deluxetable*}{cccccccc} \\tablewidth{0pt} \\tablecaption{Track Details} \\tablehead{ \\colhead{Target} & \\colhead{Configuration} & \\colhead{$u-v$ Coverage} & \\colhead{Beam Size} & Date & \\colhead{$<\\tau_{\\rm 225GHz}>$}  & Obs. Time$^a$ & \\colhead{Reference$^b$} \\\\ \\colhead{} & \\colhead{} & \\colhead{[k$\\lambda$]} & \\colhead{[arcsec]} & \\colhead{[dd.mm.yy]} & \\colhead{} & \\colhead{[hrs]} & \\colhead{} } \\startdata GN20 & COM & 15-75 & $2.99\\times 2.26$ & 20.02.05, 05.03.05 & 0.04, 0.06 & 10.4 & I06 \\\\ & EXT & 40-200 & $0.81\\times0.75$ & 10.02.08, 11.02.08 & 0.04, 0.04 & 5.3 & This work \\\\ \\hline AzTEC1 & COM & 20-75 & $2.69\\times 2.19$ & 17.01.07 & 0.05 & 5.6 & Y07 \\\\ & EXT & 50-250 & $0.86\\times 0.55$ & 16.01.08 & 0.03 & 4.0 & This work \\\\ & VEX & 60-550 &  $0.25\\times 0.35$ & 05.04.08 & 0.03 & 3.4 & This work \\enddata \\tablenotetext{a}{Total on--source integration time in that configuration.} \\tablenotetext{b}{I06: \\citet{iono2006}; Y07: \\citet{younger2007}} \\label{tab:tracks} \\end{deluxetable*}  \\begin{figure*} \\plottwo{f1a.eps}{f1b.eps} \\caption{The $u-v$ coverage for our high resolution interferometric imaging of GN20 (left) and AzTEC1 (right).  Included are all tracks in three different SMA configurations: compact (COM: black), extended (EXT: blue), and very extended (VEX: red).  For further details, including weather conditions and on-source integration times, see Table~\\ref{tab:tracks}.} \\label{fig:uv} \\end{figure*}  Wide area surveys at millimeter \\citep[e.g.,][]{greve2004,bertoldi2007,scott2008} and submillimeter \\citep[e.g.,][]{smail1997,hughes1998,barger1998,pope2006,coppin2006} wavelengths have revealed a large population of ultra-- and hyperluminous infrared galaxies -- ULIRGs and HyLIRGs -- at high redshift \\citep[median $z\\sim 2$ for a radio--selected sample][]{chapman2005}.  Since their initial discovery it has become clear that these ``submillimeter galaxies\" (SMGs) are likely massive, gas--rich merging systems \\citep{frayer1998,frayer1999,chapman2003b,greve2005,tacconi2006,tacconi2008} that represent massive galaxies in formation \\citep{scott2002,blain2004}.  Extremely luminous infrared objects take on increasing cosmological importance at $z\\gsim 1$ \\citep{sanders1996,lefloch2005}, and may dominate cosmic star formation for up to the first half of the lifetime of the universe \\citep{blain1999,blain2002}.  Despite significant progress over the past decade, a more complete understanding of SMGs has been hampered in part by the relatively poor resolution of submillimeter cameras ($\\sim 10-18$ arcsec FWHM).  In particular, the size scale of the starburst region -- traced by the rest frame far infrared (IR) -- potentially provides important insights into the nature of the engine driving the tremendous luminosity of these systems.  If they are scaled up versions of local ULIRGs, we would expect far--IR emission on scales of $\\sim 5-10$ kpc \\citep[e.g.,][see also Iono et al. 2008, in preparation]{downes1998,iono2007}.  Hydrodynamic modeling of merger driven nuclear starbursts \\citep[e.g.,][]{mihos1994} of the kind thought to drive many SMGs \\citep{chapman2003b,tacconi2006,tacconi2008} can be somewhat more compact, a result that could have important physical consequences: Eddington arguments suggest a minimum size scale for such regions \\citep{murray2005,thompson2005}.  Unfortunately, at typical SMG redshifts, all of these size scales are far smaller than the typical angular resolution of submillimeter cameras on single--dish instruments.  The first breakthrough came with deep radio continuum surveys, which leveraged the local far--IR/radio correlation \\citep[][]{condon1992} in combination with statistical arguments \\citep{ivison2002,ivison2007} to associate faint radio counterparts within the submillimeter beam with SMGs.  Higher resolution radio imaging of these sources \\citep{chapman2004,biggs2008} found a range of source structures on physical scales of $\\sim 1-8$ kpc with a median of 5 kpc.  While promising, this technique assumes a spatially resolved far--IR/radio correlation which is not particularly well understood locally \\citep[e.g.,][]{hippelein2003,murphy2006,tabatabaei2007}.  As a consequence, these results are not straightforward to interpret.  This motivates high--resolution imaging of the rest frame far--IR directly, via submillimeter interferometry.  The vast majority of previous work was done at resolutions of $\\sim 1-2$ arcsec, and found that the far--IR continuum in SMGs originates at physical scales of $\\lsim 4-8$ kpc \\citep{neri2003,greve2005,tacconi2006,wang2007,younger2007,younger2008,dannerbauer2008}.  More recently, very high resolution CO imaging by \\citet{tacconi2008} showed that gas motions in typical SMGs are disordered on scales of $\\sim 1-2$ kpc, suggesting that they are ongoing major mergers.  In this paper, we present high resolution (beam size $\\lsim 1$ arcsec) 890\\micron continuum imaging of two of the brightest SMGs known -- GN20 \\citep{pope2006,iono2006} and AzTEC1 \\citep{younger2007,scott2008} -- with the Submillimeter Array \\citep[SMA:][]{ho2004}.  By targeting the brightest -- and therefore likely the most luminous \\citep{blain1993} -- objects, we constrain the physical scale of the far--IR in extreme conditions. In addition, since these objects are thought to lie at higher redshift than radio--selected samples \\citep{younger2007}, they offer an intriguing probe of the nature of star formation at earlier epochs.  Throughout this work we assume a concordance cosmology with $\\Omega_m = 0.3$, $\\Omega_\\Lambda = 0.7$, and $h=0.7$. ",
        "conclusions": "We present high resolution interferometric submillimeter imaging of two of the brightest -- and therefore likely most luminous \\citep{blain1993} -- high redshift starburst galaxies known -- GN20 and AzTEC1.  The visibility functions for these sources indicate characteristic angular sizes of $\\sim 0.5-1.2$ and $\\sim 0.2-0.4$  arcsec respectively.  Both are B--band dropout optical sources, which indicates a redshift of $3.5\\lsim z\\lsim4.5$, and thus these angular size measurement correspond to physical scales of $4-8$ and $1.5-3$ kpc.  Assuming a simple morphology and a dynamical state typical of high--redshift SMGs, we find preliminary evidence that GN20 and AzTEC1 are both close to the limiting luminosity derived via Eddington arguments.  If these two sources are indicative of a large population of hyperluminous starbursts at high redshift, this may have important consequences for models of star formation and feedback in extreme environments."
    },
    "0807/0807.1731_arXiv.txt": {
        "abstract": "Accretion discs at sub-pc distances around supermassive black holes are likely to cool rapidly enough that self-gravity results in fragmentation. Here, we use high-resolution hydrodynamic simulations of a simplified disc model to study how the outcome of fragmentation depends upon numerical resolution and cooling time, and to investigate the incidence of binary formation within fragmenting discs.  We investigate a range of cooling times, from the relatively long cooling time-scales that are marginally unstable to fragmentation down to highly unstable cooling on a time-scale that is shorter than the local dynamical time. The characteristic mass of fragments decreases with reduced cooling time, though the effect is modest and dependent upon details of how rapidly bound clumps radiate. We observe a high incidence of capture binaries, though we are unable to determine their final orbits or probability of survival. The results suggest that faster cooling in the parent disc results in an increased binary fraction, and that a high primordial binary fraction may result from disc fragmentation.  We discuss our results in terms of the young massive stars close to the Galactic Centre, and suggest that observations of some stellar binaries close to the Galactic Centre remain consistent with formation in a fragmenting accretion disc. ",
        "introduction": "\\label{sec:intro} In recent years the topic of gravitational instabilities in accretion discs has been the subject of a great deal of theoretical and numerical research.  Such gravitational instabilities have applications both to accretion disc theory, in terms of the transport of angular momentum, and also to theories of star and planet formation, as under certain circumstances the instability can lead to fragmentation \\citep[e.g.][]{durisen_ppv,lodato08}.  In particular, the fragmentation of accretion discs has been invoked to explain both the formation of gas giant planets \\citep[e.g.][]{boss97,mayer02}, and also to explain the formation of young, massive stars observed close to super-massive black holes at the centres of galaxies \\citep[e.g.,][]{ps77,lb03}.  In order for a gravitationally unstable disc to fragment, two conditions must be met.  The first is that the disc must be sufficiently cold, or sufficiently massive, to satisfy the \\citet{toomre64} criterion \\begin{equation}\\label{eq:Q} Q = \\frac{c_{\\mathrm s} \\kappa}{\\pi G \\Sigma} \\lesssim 1 \\end{equation} Here $c_{\\mathrm s}$ is the local sound speed, $\\kappa$ the epicyclic frequency (which in a thin disc is approximately equal to the local orbital frequency), and $\\Sigma$ the disc surface density.  In addition, in order to fragment the disc must cool sufficiently rapidly to overcome compressional heating during collapse.  \\citet{gammie01} used a local analysis, and found the fragmentation boundary to be \\begin{equation} t_{\\mathrm {cool}} < \\beta_{\\mathrm {crit}} \\Omega^{-1} \\, , \\end{equation} where $t_{\\mathrm {cool}}$ is the local cooling time-scale, $\\Omega$ is the orbital frequency, and $\\beta_{\\mathrm {crit}} \\simeq 3$.  Subsequent studies using global simulations have verified the validity of this criterion \\citep{rice03,rice05,meija05,boley07}, although the exact value of $\\beta_{\\mathrm {crit}}$ required for fragmentation can be larger by a factor of 2--3 depending on the details of the thermal physics adopted (cooling law, equation of state, etc.).  To date, however, most work has either made use of simple thermal physics \\citep[e.g.][]{rice03,rice05,nayak07}, or used more realistic thermal physics in the specific case of self-gravitating protoplanetary discs \\citep[e.g.][]{cai08,sw08}.  More generally, most studies of fragmentation have tended to look at cases of critical cooling (with $\\Omega t_{\\mathrm {cool}} \\simeq \\beta_{\\mathrm {crit}}$), despite the fact that some discs, notably those around black holes, are thought to cool much more rapidly \\citep[e.g.][]{bl01,goodman03}.  In this paper we use numerical hydrodynamics to study the influence of rapid cooling on the fragmentation properties of gravitationally unstable accretion discs.  In Section \\ref{sec:method} we present the details of our numerical simulations, in which we adopt a cooling law consistent with radiative cooling in an optically thick disc.  We use this model to study the non-linear evolution of the instability over a wide range in cooling times.  We also consider in detail the effects of numerical resolution on our calculations, varying the number of fluid elements by a factor of 64 between our lowest- and highest-resolution calculations.  We present our results in Section \\ref{sec:results}, first considering issues of numerical convergence and then addressing the effects of varying the cooling time.  In Section \\ref{sec:dis} we discuss the limitations of our analysis, and the implications of our results for the formation of stars near to the Galactic Centre, before summarizing our conclusions in Section \\ref{sec:summary}. ",
        "conclusions": "\\label{sec:dis} Our simulations are robust and well-understood from a numerical perspective, but concerns remain when we consider the physical interpretation of our results.  The most obviously unphysical aspect of our simulations is our ``thin disc approximation'', using an artificial cooling law to allow fragmentation characteristic of a thin disc to be modelled at moderate numerical resolution.  It would be preferable to construct global models of thin discs, but this is not computationally feasible with current techniques, especially in the case of short cooling times.  Our disc model is effectively a local approximation, so the primary concern regards the boundary conditions.  In our models the real physical boundaries (at $r=1$ and $r=3$) are sufficiently far from the region of interest that they have no influence on the results, and the radial variation of the cooling time is negligible in the fragmentation region.  However, in a real thin disc there are many such filaments, and the typical radial separation between filaments can be as little as a few times $H$.  Dynamical interactions between adjacent filaments seem likely, especially once bound, point-mass-like objects have formed, and the likely outcome of this increased rate of clump-clump interactions is the formation of even more capture binaries than seen in our simulations.  In some of our simulations (those with slower cooling) some of the bound clumps undergo significant radial scattering even in the absence of adjacent filaments (moving as far as $\\Delta r\\simeq 0.3$ from their initial location), and consequently their cooling times change in an unphysical way.  However, this only occurs at late times in the simulations (to which we do not attach great significance), and moreover in such dense clumps the dominant term in the cooling time is the density dependence.  Consequently we are satisfied that our artificial cooling model is robust, and although our simulations should be regarded as numerical experiments rather than physical simulations, we are confident attaching limited physical significance to our results.  An additional concern regards the density dependence of our cooling law.  Fixing the cooling time to scale as $\\rho^{1/3}$ is correct only if the opacity is constant and the disc is optically thick, but it is not clear that these assumptions hold in reality.  In real self-gravitating discs the temperatures are expected to be $\\sim100$K, so dust grains dominate the disc opacity \\citep[e.g.][]{bl94,semenov03}.  In the case of a fragmenting disc with $Q=1$, the disc is optically thick if the orbital period is $\\lesssim 1000$yr \\citep{levin07}.  Real discs are generally only thought to be self-gravitating in their outer regions, at radii $\\gtrsim50$AU in the case of protostellar discs \\citep[e.g.][]{sw08} or at radii $\\gtrsim0.1$pc in the case of a disc around a $\\sim10^6$M$_{\\odot}$ black hole \\citep[e.g.][]{kp07}, and in both cases this suggests that fragmenting discs are marginally optically thick.   Individual clumps (with higher densities) thus have optical depths much greater than unity, so the optically thick approximation appears valid across a wide range of disc scales.  It is clear, however, that the opacity is not strictly constant in real discs.  Dust opacity is essentially independent of density, but varies strongly as a function of temperature: the opacity scales as $T^2$ for $T\\lesssim160$K, and is approximately constant for $160$K$\\lesssim T\\lesssim1000$K \\citep{bl94}.  It is reasonable to assume that the temperature will increase somewhat during collapse, leading to an increase in the dust opacity, so it seems likely that our adopted function $D(\\rho)$ somewhat over-estimates the efficiency of radiative cooling.  Thus cooling is likely slower than predicted by our simulations in real fragmenting clumps.  Models of core-collapse during star formation predict only moderate increases in temperature, however, so the uncertainly due to the temperature dependence of dust opacity is likely small.  We note in passing that we also neglect further complexities of the collapse process, such as changes in gas composition with increasing temperature (due to dust sublimation, H$_2$ dissociation, etc.): such issues are well beyond the scope of this work.  An additional complication is that is still not clear whether or not it is possible to form a disc with $\\beta<1$.  A coherent disc structure cannot be assembled in the presence of such rapid cooling, as the growing disc would simply fragment on a dynamical time-scale.  Various authors have argued in favour of rapid cooling in self-gravitating discs around massive black holes \\citep[e.g.][]{bl01,goodman03}, but these arguments generally rely on the assumption of roughly steady-state conditions in the disc, rather than considering disc formation.  \\citet{rafikov07} argued that convective cooling results in an cooling time of $\\beta \\simeq 0.1$, but again does not consider disc formation in detail.  By contrast, \\citet[][see also \\citealt{levin07}]{nayak06} argues against the presence of rapid cooling in a fragmenting, star-forming disc around a massive black hole, instead arguing that the disc mass is accreted gradually.  The cooling time therefore decreases slowly as the disc mass is accumulated, so fragmentation occurs at $\\beta\\simeq\\beta_{\\mathrm {crit}}$.  If, however, the disc was formed in a rapid ``accretion event'', it seems that fragmentation could occur with $\\beta<1$.  In addition the thermodynamics of the disc formation process are not fully understood, and one can conceive of a scenario where the disc is assembled with a longer cooling time and then begins to cool faster as it evolves \\citep[e.g.][]{clarke07}.  One possible mechanism appeals to the ``opacity gap'' at $10^3$--$10^4$K: if a disc was assembled from warm, $\\gtrsim 1500$K gas then the condensation of dust at $\\sim1000$--1500K would result in a dramatic increase in the cooling rate as the temperature dropped, potentially resulting in a cooling time-scale shorter than the dynamical time-scale.  Our simple model is not able to address this issue in detail, but this problem provides fertile ground for future research.  Despite the fact that our simulations show good numerical convergence with increasing particle number, numerical resolution is still an issue that must be considered with care.  In particular, as noted in Section \\ref{sec:res_conv}, the dynamical interactions between  different clumps are under-resolved even in our highest-resolution simulations.  This is a common feature of hydrodynamic simulations which seek to model large dynamic ranges in length and density, and in our simulations this results in unphysically rapid mergers of binary and multiple systems (as seen, for example, in Fig.\\ref{fig:panels}).  We treat this problem by simply stopping our calculations before such under-resolved mergers begin to dominate the simulations, and attach only limited significance to our results beyond the point at which the first such mergers begin to happen.  [Note, however, that real mergers between fragments are expected to occur in such systems, and indeed have been invoked to explain the large discrepancy between the typical fragment mass and the observed stellar masses near the Galactic Center \\citep{levin07}.]  A more sophisticated approach could, for example, make use of ``sink particles'' \\citep[e.g.][]{bate95}, modified thermal physics \\citep[e.g.][]{jappsen05}, or sub-resolution physics \\citep[e.g.][]{nayak07} in order to follow the calculations further forward in time, but in our models (and in particular in the simulations with fast cooling) it is not clear that such approaches would yield more accurate results.  Here, a better approach may be to ``re-simulate'' a small region of our computational domain at much higher resolution in order to study the inter-clump dynamics in more detail, but this is beyond the scope of this paper.   \\subsection{Application to the Galactic Centre} We now consider the implications of our results for the formation of stars near the Galactic Centre (GC).  Many young, massive stars are known to exist close to the black hole at the centre of the Galaxy \\citep{ghez98,genzel03,ghez05,paumard06}, and many of the stars at $\\simeq0.1$pc are known to form a coherent ring or disc \\citep[e.g.][]{genzel03}\\footnote{Note that there is some controversy as to whether one or two discs of young stars exist at the GC \\citep{paumard06,lu06}.  For the purposes of this discussion this is unimportant: the existence of one disc of young stars is sufficient to motivate our arguments, and they are not significantly altered by the existence of a second disc.}.  A popular theory for the formation of these stars is fragmentation of an accretion disc due to gravitational instabilities \\citep[e.g.][]{lb03,gt04,nc05}, and recent hydrodynamic simulations have found that the properties of the stars formed in this manner are broadly consistent with the observed GC stellar disc \\citep[e.g.][]{nayak07,rda08}.  However, to date most of these simulations have used greatly simplified thermal physics, adopting scale-free cooling and/or simplified equations of state, and given the critical role of cooling in the development of gravitational instabilities this treatment is less than ideal.  Additionally, for numerical reasons most previous simulations of the formation of the GC stellar disc have made use of relatively slow cooling time-scales (with $\\beta\\simeq3$), although \\citet{nayak07} did run a single model with $\\beta=0.3$ (as discussed in Section \\ref{sec:res_cool}).  However, analytic estimates of the cooling times in black hole discs generally find that the cooling time should be short compared to the dynamical time \\citep[e.g.][]{goodman03,rafikov07} although, as discussed above, how such systems form remains uncertain.  Estimates of the cooling parameter in a fragmenting disc at the GC range from $\\beta\\sim0.1$ \\citep{rafikov07} to $\\beta=\\beta_{\\mathrm {crit}} \\simeq 3$ \\citep{nayak06}, which is approximately the range spanned by our simulations.  As discussed above, we find that a dramatic decrease in the cooling time leads to only a modest decrease in the characteristic stellar mass, and given the dependence of the characteristic mass on other disc properties (notably the disc aspect ratio $H/R$), this is unlikely to have a significant effect on the MF of the GC stellar disc.  However, our simulations indicate that fast cooling results in a very high primordial binary/multiple fraction (of order unity), and although binaries do form when cooling is slower, they are less common.  We therefore suggest that, if the GC stellar disc formed by fragmentation of a gaseous accretion disc, a high binary fraction may be indicative of rapid cooling in the parent accretion disc.  It is not clear, however, if this prediction provides a means of distinguishing between different proposed scenarios for the formation of the GC stellar disc.  The massive star population in the field is known to have a high multiplicity fraction, and the binary fraction for Wolf-Rayet stars in the solar neighbourhood is $\\sim 40$\\% \\citep{vdhucht01}.  Little is currently known about binary statistics in the GC stellar cluster, with neither observations nor dynamical arguments \\citep[e.g.][]{cuadra08} providing real constraints.  It is entirely possible (and in fact likely) that many unresolved close binaries near the GC may be masquerading as single objects in current data, but some binary systems have been observed.  In particular IRS 16SW, an eclipsing binary with a period of $\\simeq$19.5 days, appears to be a contact binary consisting of two approximately equal-mass Wolf-Rayet stars with masses $\\gtrsim 50$M$_{\\odot}$ \\citep{martins06,peeples07,rafelski07}.  Systems of this type appear to be a plausible outcome of fast cooling in a fragmenting, self-gravitating accretion disc, and we suggest that such systems should be common if the stellar disc at the GC formed via fragmentation of a gaseous accretion disc.   We have used numerical hydrodynamic simulations to investigate the role of the cooling rate in the fragmentation of gravitationally unstable accretion discs.  Our simulations make use of a simple cooling law, designed to mimic the behaviour of radiative cooling in an optically thick, geometrically thin disc.  We performed detailed numerical convergence tests, and are satisfied that our simulations are well-resolved.  We find that decreasing the cooling time results in the formation of a larger number of fragments, but with a correspondingly smaller characteristic mass.  We also observed a high incidence of binaries, all of which form by dynamical capture.  Rapid cooling, on a time-scale shorter than the local dynamical time, results in a dramatic increase in the incidence of binary and multiple systems, although we are not able to determine the final properties of such systems.  We have considered our results in the context of the young stellar disc(s) at the Galactic Centre, and find that - at the very least - observations of some stellar binaries close to the Galactic Center do not rule out formation in a fragmenting accretion disc."
    },
    "0807/0807.4608_arXiv.txt": {
        "abstract": "{} {We examine the constraints that can currently be obtained on standard cold dark matter models from the most currently used data set: CMB anisotropies, type Ia supernovae and the SDSS luminous red galaxies. We also examine how these constraints are widened when the equation of state parameter $w$ and the curvature parameter $\\Omega_k$ are left as free parameters. Finally, we investigate the impact on these constraints of a possible form of evolution in SNIa intrinsic luminosity. } {We obtained our results from MCMC analysis using the full likelihood of each data set.} { For the $\\Lambda$CDM model, our 'vanilla' model, cosmological parameters are tightly constrained and consistent with current estimates from various methods. When the dark energy parameter $w$ is free we find that the constraints remain mostly unchanged, i.e. changes are smaller than the 1 sigma uncertainties. Similarly, relaxing the assumption of a flat universe leads to nearly identical constraints on the dark energy density parameter of the universe $\\Omega_\\Lambda $ , baryon density of the universe $\\Omega_b $, the optical depth $\\tau$,  the index of the power spectrum of primordial fluctuations $n_S$,  with most one sigma uncertainties better than 5\\%. More significant changes appear on other parameters: while preferred values are almost unchanged, uncertainties for the physical dark matter density $\\Omega_ch^2$, Hubble constant $H_0$ and $\\sigma_8$ are  typically twice as large.  The constraint on the age of the Universe, which is very accurate for the vanilla model, is the most degraded. We found that  different methodological approaches on large scale structure estimates lead to appreciable differences in preferred values  and uncertainty widths. We found that possible evolution in SNIa intrinsic luminosity does not alter these constraints by much, except for $w$, for which the uncertainty is twice as large. At the same time, this possible evolution is severely constrained. } {  We conclude that  systematic uncertainties for some estimated quantities are similar or larger than statistical ones. } ",
        "introduction": "The field of cosmology has made impressive developments in the past decade. It is now widely accepted that the present universe is well represented by a dark energy dominated universe with dark matter being in the form of a cold component and matter fluctuations arising from primordial adiabatic Gaussian fluctuations generated in the early universe, a picture consistent with simple inflationary models.  Such a paradigm was obtained after a remarkable advance in the quality of observational data. The Hubble diagram derived from type Ia supernovae allowed to obtain the first evidence for dark energy and provided measurements on the properties of this component (\\cite{Perlmutter}; \\cite{Riess98}). Another important contribution came from the results of the Wilkinson Microwave Anisotropy Probe (WMAP) (\\cite{WMAP1}; \\cite{WMAP3}), which have provided strong constraints on many cosmological parameters and allowed for precise estimations of cosmological parameters% . Those constraints are consistent with a wide set of constraints coming from different cosmological tests and have considerably reinforced early claims of  flat cosmology (\\cite{1997AA...322..365L}; \\cite{1998ApJ...509L..65W}). The use of combined constraints is therefore a way to increase the accuracy of cosmological parameter estimations and reduce degeneracies.\\\\ \\indent Among the many different cosmological tests available beyond the CMB, two are particularly robust and based on geometrical considerations: the already mentioned Hubble diagram of type Ia supernovae,  and the recent large galaxy surveys that provided a first detection of the predicted baryonic acoustic oscillations (\\cite {eisenstein}, E05 hereafter; \\cite{Tegmark}, T06 hereafter; \\cite{Percival}). All these observations reinforce the case of the  concordance model in cosmology, in which  the Universe appears to be in an accelerated expansion due to some ``dark energy'' whose true nature remains one of cosmology's most compelling mysteries. The ability of the concordance model to accommodate recent high precision data is indeed impressive. Although alternative views exist, it is fair to say that the so called concordance cosmology has led to specific {\\em predictions} that were verified a posteriori. The most common interpretations for this ``dark energy'' are an unknown energy component of the Universe or a modification in the equations of general relativity (for some reviews on this subject, see Frieman et al. 2008, Padmanabhan 2006). In order to distinguish between  models beyond the $\\Lambda$CDM model, one needs to have very good constraints on the equation of state of dark energy by combining reliable results from quality observations. Some multi-observational constraints have been published recently (\\cite{Tegmark}; \\cite{WMAP5}; Kowalski et al. 2008; \\cite{wright2007}; \\cite{xia2008}), but do not  include a detailed likelihood  analysis  of the three most recent observational data sets (i.e. information from one or several data sets was reduced to a single number).\\\\  In this paper, we propose to constrain several cosmological parameters, including the $w$ parameter in the equation of state of dark energy (supposed constant), using the recent results on the matter power spectrum of LRG galaxies, SNIa and WMAP-5 year data. Although there are other promising methods to test cosmological models and dark energy properties, such as galaxy clusters (\\cite{vikhlinin}, \\cite{Allen}) or weak lensing (\\cite{benjamin}), we have not included these in our analysis due to the level of incertitude that may still be present in these methods (\\cite{Ferramacho}; \\cite{Vauclair}; \\cite{Nusa}; \\cite{2005AA...436..411B} ). Given the high number of parameters to be constrained, we adopt the Markov chain Monte Carlo (MCMC) method to perform our analysis. The paper is structured as follows: In section 2 we present and discuss the quality of the data. In Section 3 we describe the method and the cosmological models to be constrained and in Section 4 we present the obtained results. Our conclusions are given in Section 5.   \\begin{figure*}[t!] \\centering \\hspace*{-17mm} \\includegraphics[width=1.1\\textwidth]{plots_1D_WMAP5_unionsn_2.eps} \\vspace*{-35mm} \\caption[]{Posterior distributions for the parameters constrained in our MCMC analysis. Doted curves represent the mean likelihood of the samples, while solid lines show the fully marginalized posterior. The four parameter sets constrained are represented by different colours: Vanilla (black), Vanilla + $\\Omega_k$(green), Vanilla + $w$ (blue), Vanilla + $\\Omega_k + w$ (red). \\label{1Dcontour} } \\end{figure*}   \\begin{figure*}[tcb!] \\includegraphics[width=0.93\\textwidth]{2D_contours_vanilla_omk_w_2D.eps} \\vspace*{-8mm} \\caption[]{2D constraints on most correlated cosmological parameters when Vanilla+$\\Omega_k+w$ is considered. The curves show the marginalised 1$\\sigma$ and 2$\\sigma$ confidence regions} \\label{2Dcontour} \\end{figure*}   \\begin{table*}[t!]  \\centering \\begin{tabular}{c c c c c l} \\hline $\\mathbf{Parameter}$\t\t&$\\mathbf{Vanilla}$\t &$\\mathbf{Vanilla + \\Omega_k}$\t\t\t&$\\mathbf{Vanilla+w}$ & $\\mathbf{Vanilla+\\Omega_k+w}$  & \\bf{Parameter Definition}\\\\ \\hline $        \\Ob h^2$      &$ 0.0227\\pm 0.0005$                            &   $0.0227\\pm0.0006$           & $0.0228\\pm0.0006$      & $0.0227\\pm0.0005$ \t        &Normalised baryon matter density times $h^2$\\\\ $           \\Omega_c h^2$      &$0.112\\pm0.003 $                     &  $0.109\\pm0.005$            & $0.109\\pm0.005$      & $0.109\\pm0.005$\t        &Normalised CDM density times $h^2$\\\\ $                 \\theta$      &$ 1.042\\pm0.003$                       &  $1.042\\pm0.003$            & $1.042\\pm0.003$\t    & $1.042\\pm0.003$\t \t&Ratio of sound horizon to angular diameter distance\\\\ $                 \\tau$      &$ 0.085\\pm0.017$                       &  $0.088\\pm0.017$            & $0.087\\pm0.017$\t    & $0.088\\pm0.017$\t        &Reionization optical depth\\\\ $             n_s$      &$0.963\\pm0.012$                             &  $0.964\\pm0.013$            & $0.967\\pm0.014$\t    & $0.964\\pm0.014$\t        &Primordial spectral index at $k=0.05 \\rm Mpc^{-1}$\\\\ $                        log(10^{10} A_s)$      &$3.07\\pm0.04$        &  $3.06\\pm0.04$            & $3.06\\pm0.04$\t    & $3.06\\pm0.04$\t        &$A_s$ is the primordial scalar power at $k=0.05 \\rm Mpc^{-1}$ \\\\ $                       \\Omega_k$      &$ 0$                         &  $-0.005\\pm0.007$            &  $ 0 $\t    & $-0.005\\pm0.0121$\t        &Spatial curvature   \\\\ $                        w$      &$ -1 $                             &  $ -1$            &  $-0.965\\pm0.056$\t    & $-1.003\\pm0.102$\t        &Dark energy EoS ($w=p_\\Lambda/\\rho_\\Lambda$)\\\\ \\hline $                      \\Ol$      &$ 0.738\\pm0.015                $   &  $0.735\\pm0.016$           & $0.739\\pm0.014$      & $0.733\\pm0.020$\t\t&Dark energy density ($h^{-2}\\rho_{\\Lambda}$)\\\\ $                      Age       $      &$  13.7\\pm0.1$              &  $13.9\\pm0.4$            & $13.7\\pm0.1$\t    & $13.9\\pm0.6$\t        &Age of the universe in Gyr\\\\ $                      \\Om$      &$ 0.262\\pm0.015$                   &  $0.270\\pm0.019$            & $0.261\\pm0.020$\t    & $0.272\\pm0.029$\t\t&Normalised matter density\\\\ $                    \\sigma_8$      &$0.806\\pm0.023$                 &  $0.791\\pm0.030$            & $0.816\\pm0.014$      & $0.788\\pm0.042$\t\t&Matter density fluctuation amplitude \\\\ $                   z_{re}$      &$10.9\\pm1.4$                       &  $11.0\\pm1.5$            & $11.0\\pm1.5$\t    & $11.0\\pm1.4$\t\t&Reionization redshift\\\\ $                    h $      &$ 0.716\\pm0.014$                      &  $0.699\\pm0.028$            & $0.713\\pm0.015$\t    & $0.698\\pm0.037$\t\t&Hubble parameter\\\\   \\hline \\end{tabular} \\caption{Summary of the posterior distribution mean values for the different sets of parameters constrained, with the corresponding 68\\% confidence intervals.} \\end{table*} ",
        "conclusions": "In this paper, we have used three of the most solid sets of observational data in cosmology to perform an up to date constraint on the $\\Lambda$CDM model: CMB, the Hubble diagram of distant SN and the shape of the power spectrum of galaxy distribution on large scales. We have chosen to use the power spectrum shape because the analysis using the correlation function was appreciably longer.  We have confirmed that 6 parameters only are enough to reproduce the three currently used data sets. Using appropriate likelihoods for each data set,  we obtained very good constraints with typical  uncertainties of the order of a few percent. The comparison between the values obtained in this work and those found  using only WMAP3+SDSS (\\cite{Tegmark})  data do not show significant changes in most of the parameters, which demonstrates the great constraining power of these large scale observations. Indeed, the introduction of the supernova data produces changes mostly on the parameters that govern the geometry of the Universe, namely $\\Ol$ and $h$, which have their mean values reduced by around 0.02/0.03. As for the dark energy EoS parameter $w$, we do not obtain a better constraint by adding the supernova data, and the results point to a value highly compatible with a cosmological constant although the 1$\\sigma$ error bars are appreciably larger than on other parameters (of the order of 5 to 10 \\%) .  We found that adding extra freedom to the vanilla model does not decrease the quality of the constraints in an appreciable way for several parameters:  as can be seen from figure \\ref{1Dcontour}, constraints on the baryonic content of the universe $\\Omega_b$, the dark energy content $\\Ol$, the matter content of the universe $\\Omega_c$ and the primordial index of scalar fluctuations $n_S$ are almost unchanged.  We get a slightly wider distribution in other  parameters, with preferred values unchanged but with uncertainties increased by a factor of up to two. The most drastic change happens for the age of the universe, with an uncertainty increased by a factor of more than five, although the uncertainty remains still only of the order of 5\\%.  Our constraints have been established within some assumptions. For instance we do not investigate the possible contribution of tensor modes (Zaldarriaga \\& Seljak 1997; Liddle \\& Lyth 2000) or neutrino mass (\\cite{2006JCAP...06..019G}) . Obviously, the addition of new freedom could modify some of our constraints. There is no general method to avoid this type of limitation; for instance, one of the most critical assumptions on constraints obtained from the CMB is the power law shape of the spectrum of primordial fluctuations. Relaxing this assumption could lead to dramatic modification of the constraints (\\cite{2003AA...412...35B}, \\cite{2007PhRvD..76l3504H}).  While preparing this paper, we benefited from the release of the 5-year WMAP results. The WMAP team presents MCMC constrains on the $\\Lambda CDM $ model (\\cite{WMAP5}), using WMAP+SN+BAO. They used the BAO measurements on the angular diameter scale $D_A$ derived from the SDSS LRG correlation function (\\cite{eisenstein}), while we preferentially used the full power spectrum shape data. By comparing the two approaches we find very similar constraints on all parameters and similar $1\\sigma$ confidence intervals, which indicates that the use of the distance parameter $A$ in implementing the LSS constraint is essentially sufficient. Also, we find that the use of the power spectrum instead of the correlation function  leads to appreciable difference of up to two sigma. Finally our investigation on the consequence of a specific model for the evolution of SNIa luminosities has been interesting in several respects: first, although evolution does not allow any conclusion on cosmological parameters from the Hubble diagram alone, we found that most of the constraints from combinations remained unchanged, and at the same time the amount of evolution is severely constrained to be less than 5 to 20\\% for redshifts between 0.5 to 1.  The current status of cosmology seems then to be highly compatible with the $\\Lambda$CDM. Present-day existing data allows to estimate cosmological parameters to within a few percent whithin a specific model. This raises the fundamental question of the role of astrophysical sources of uncertainties. Clearly the next step is to understand in more detail the nature of dark energy (cosmological constant, dynamic scalar field or modified gravity) and this is yet to be achieved. The ever increasing quality of data, anticipated from future large observational programs, could be used to rule out or validate some of theses alternatives. However, future  projects would need to keep systematic effects, both from methodological and astrophysical origin, controlled to achieve a typical anticipated precision of one percent or better."
    },
    "0807/0807.1507_arXiv.txt": {
        "abstract": "We present an optical/near-infrared search for a counterpart to the perplexing radio transient \\gcrt, a source located $\\sim1\\degr$ from the Galactic Center.  Motivated by some similarities to radio bursts from nearby ultracool dwarfs, and by a distance upper limit of $70\\,$pc for the emission to not violate the $10^{12}\\,$K brightness temperature limit for incoherent radiation, we searched for a nearby star at the position of \\gcrt.  We found only a single marginal candidate, limiting the presence of any late-type star to $>1\\,$kpc (spectral types earlier than M9), $>200\\,$pc (spectral types L and T0--T4), and $>100\\,$pc (spectral types T4--T7), thus severely restricting the possible local counterparts to \\gcrt.  We also exclude any white dwarf within $1\\,$kpc or a supergiant star out to the distance of the Galactic Center as possible counterparts.  This implies that \\gcrt\\ likely requires a coherent emission process, although whether or not it reflects a new class of sources is unclear. ",
        "introduction": "A blind search program using the Very Large Array (VLA) at 330$\\,$MHz ($90\\,$cm), made possible by new wide-field radio imaging techniques, has resulted in the discovery of two radio transients near the Galactic Center. One of these, \\gcrt, exhibits dramatic 1-Jy outbursts that last about 10$\\,$min and have a recurrence interval of about 77$\\,$minutes \\citep{hlk+05}.   \\gcrt\\ exhibited 5 bursts at 330$\\,$MHz, each lasting about 10$\\,$minutes, at apparently regular intervals of 77.1$\\,$min during a 7-hour VLA observation in September 2002.  After several non-detections during the summer of 2003, a single burst was detected during a short Giant Metrewave Radio Telescope (GMRT) observation in September 2003 \\citep{hlr+06}, confirming that it is  recurrent.  It has since also been faintly detected in an additional GMRT observations from 2004, where it had a very steep nonthermal radio spectrum across the $32\\,$MHz bandpass \\citep{hrp+07}.  It is at Galactic coordinates $\\ell = 358.89\\degr$, $b=-0.54\\degr$, so it may be as close as 180$\\,$pc to the Galactic Center (GC), though the distance of \\gcrt\\ from the Earth is completely unknown.  If it is at a distance $d>70\\,$pc from Earth, then the radio flux density combined with the rapid ($\\sim 2\\,$min) decay of the bursts constrains its brightness temperature to exceed the $10^{12}\\,$K limit for incoherent synchrotron radiation \\citep{readhead94} and thus is very likely a coherent emitter.  The positional coincidence with the GC, along with the greatly increasing source density near the Center, suggests it is at a distance of about 8$\\,$kpc, but this could be a selection effect because all the fields searched for transients in this program were in that direction.  Consequently, it is very important to rule out local classes of sources as possible explanations.    Among the many possibilities discussed for the origin of \\gcrt, several involved low-mass stars or substellar objects (brown dwarfs, or BDs), which we generically call ultracool dwarfs.  These were initially rejected as being unlikely or inconsistent with the observations \\citep{hlk+05}.  However, additional observations are required to strengthen the case against the ultracool dwarfs and to confirm that the radio source is indeed a new type of source.  Alternative models have also been proposed, ranging from a nulling pulsar \\citep{kp05}, a double pulsar \\citep*{tpt05}, a transient white dwarf (WD) pulsar \\citep{zg05}, to a precessing radio pulsar \\citep{zx06}, but these will largely be constrained through other observations.  A number of BDs and late-type stars have now been detected at GHz frequencies, in many cases despite the absence of flaring optical or X-ray emission that would accompany radio emission from normal stars (see \\citealt{berger06} for a review).  These objects exhibit a range of nonthermal flaring behaviors, from strong, narrowband, fully polarized, bursting emission with frequency drifts \\citep{bp05}, to periodic bursts suggestive of pulsar-like beaming \\citep{hbl+07}, to order-of-magnitude variations spanning years \\citep{adh+07}.  In particular, the behavior found by \\citet{hbl+07} for the M9 dwarf TVLM~513$-$46546 is reminiscent of \\gcrt: bright bursts of coherent emission that follow the 2~hr rotation of the star.  Although the bursts are orders of magnitude weaker than those seen for \\gcrt, the observations differ by an order of magnitude in wavelength, and there are no current limits on the radio behavior of ultracool dwarfs at meter wavelengths.  We then have two independent but complementary motivations for searching for a counterpart to \\gcrt\\ in the near-infrared (near-IR) and optical bands.  First, we wish to see if there is {any} possible counterpart closer than $70\\,$pc: otherwise, we must conclude that the emission from \\gcrt\\ is coherent.  Second, given the similarity between it and TVLM~513$-$46546, we wish to see if \\gcrt\\ is associated with a nearby ultracool dwarf, although such an object could be more distant than $70\\,$pc as they can emit coherently.  A single ultracool dwarf in the error circle of \\gcrt\\ would be strong evidence that it is the counterpart, even in the absence of a precise position, as one expects $<1$ such object in a $2\\arcmin$ field (\\citealt{rkl+99}; \\citealt*{cbk08}).  We concentrated on the near-IR for several reasons: (1) if the source is at the distance of the Galactic Center, the average $K$-band extinction, $A_K$, at this position is approximately 2--3 magnitudes, corresponding to a visual extinction of $\\geq$20 magnitudes (see e.g., \\citealt{dsbb03}); (2) if the source is a relatively nearby late-type star, it would be intrinsically red; and (3) the field towards the GC is very crowded, and near-IR observations often have better seeing than bluer bands. We augmented the traditional near-IR bands ($JH\\ks$) with $I$-band to aid in color discrimination (see below).   The structure of this paper is as follows.  First, to aid the analysis that follows, we update our determination of the radio position of \\gcrt\\ in \\S~\\ref{sec:position}.  We then describe our optical and near-infrared data in \\S~\\ref{sec:obs}.  The primary analysis is in \\S~\\ref{sec:bd}, where we attempt to constrain any ultracool dwarf counterpart to \\gcrt.  Following this, in \\S~\\ref{sec:other} we also examine what constraints on other types of counterparts we can determine from our data.  Finally, we give our conclusions in \\S~\\ref{sec:conc}. ",
        "conclusions": "\\label{sec:conc} We have examined relatively deep images of the field of \\gcrt\\ taken under good conditions (seeing better than $0\\farcs5$) meant to search for a possible nearby counterpart to the radio transient.  Such an object would have to be at a distance $<70\\,$pc in order to not violate the $10^{12}\\,$K brightness temperature limit on incoherent sources.  Beyond that, we searched for nearby ultracool dwarfs at a range of distances.  Examining data at wavelengths from $0.8\\,\\um$ to $2.1\\,\\um$, we find that the field changes dramatically.  Only three objects are detected at the shortest wavelength, compared to more than 20 at the longest wavelength.  Two of the three sources detected at $0.8\\,\\um$ are background giant stars.  The third (object C) is most likely a late-type K star at a moderate distance ($\\sim 6\\,$kpc).  We cannot entirely rule out the possibility that it is an object of type $\\sim$L5, but this is not  consistent with the $I$-band photometry and would still be at a distance of $\\sim 200\\,$pc: further than the $100\\,$pc limit discussed above.  If it is of a late K type, object C is unlikely to be the source of the radio emission. While isolated K stars typically do not show strong radio flares, it is possible but generally at fainter flux levels \\citep{gudel02}.  Given the radio luminosities typically seen ($\\sim10^{23}\\,{\\rm erg}\\,{\\rm s}^{-1}$; \\citealt{gudel02}), it would be far too faint to be the source of the radio emission from \\gcrt.  However, the radio luminosities seen to date for K stars may not include the brightest transient events.  Some K stars in binaries (e.g., RS~CVn binaries; \\citealt{gudel02}) also show coherent radio flares.  Such objects have peak fluxes of $<1\\,$Jy, even at relatively low radio frequencies \\citep[e.g.,][]{vdodb94}, and for the distances implied for object C (possibly underestimated if it is a binary, especially if it is evolved like the K star in the prototypical RC~CVn binary HR~1099) it would again be far too faint.  If we have not detected the counterpart, we use constraints from the shorter wavelengths, where the field is much less crowded, to limit the presence of any low-mass object.  For distances of $< 100\\,$pc, we can exclude stars earlier than T7 (with related limits of stars earlier than T4 at $200\\,$pc, M9 at $1.3\\,$kpc, M5 at $2\\,$kpc, etc.) or so based on the $I$- and $J$-band photometry.  Our search is limited by the sample of comparison objects that we used: nearby ultracool dwarfs with good photometry and (often) astrometry.  We therefore cannot exclude the possibility that object C is some peculiar type of ultracool dwarf, where the mis-match in the $I$-band reflects some intrinsic difference of this object, but we require additional constraints to be able to decide conclusively one way or the other.  While the source is bright enough at $\\ks$-band for spectroscopy with large telescopes, $I$-band spectroscopy is likely to prove most productive, as the source is much easier to isolate from the nearby objects.  Additional photometry will suffer from the near degeneracy between effective temperature and reddening seen in Figure~\\ref{fig:objC} and the difficulty of calibrating data where the spectrum is dropping so sharply, but spectra should be able to distinguish cleanly between the largely featureless spectrum of a late-K/early-M star and the deep metal hydride absorption of an L dwarf.  Our conclusions are also dependent on a lack of significant variability: fluctuations of $<0.5\\,$mag would not change things, but larger variations might.  We see good agreement between the two sets of near-IR data (from PANIC and NIRI).  However, those data were taken within two nights of each other.  Variations over longer timescales, such as the $\\sim1\\,$yr between the near-IR and optical observations, or even short flares, could still be present, although we note that we compared our data with photometry of the 3 sources (A, B, and D) within the error circle that are visible in the UKIRT Infrared Deep Sky Survey (UKIDSS) Galactic Plane Survey Data Release 3 (see \\citealt{lwa+07,lhl+07}) taken 2 years after our near-IR data and see no sign of variability at levels $\\gsim 0.3\\,$mag; below that, it is difficult to compare given the differences in calibration, seeing, and photometric technique.  Additional monitoring could help address this issue.  With the same data, we have attempted to constrain other source types as possible counterparts to \\gcrt.  We find that no white dwarf nearer than $\\approx 1.5\\,$kpc could be present in the error circle, nor could there be a supergiant star at the distance of the Galactic Center.  However, both of those conclusions have the same limitations as that discussed above.  Namely, we searched for objects that resemble other, known objects.  If the radio emission from \\gcrt\\ really marks it as a unique object than the counterpart that we seek could have dramatically different properties.  The presence of a binary companion, accretion disk, or other anomaly could also complicate things \\citep[e.g.,][]{hws+07}, although most such additions to the system would make it brighter and hence would still largely be ruled out.  With all of the analysis above, we have excluded virtually any white-dwarf or non-degenerate star as counterpart to \\gcrt\\ within $100\\,$pc, and many sources are excluded within $1\\,$kpc.  We are therefore left with the conclusion that \\gcrt\\ likely emits via a coherent process.  What this means, though, is unclear.  Coherent emission with brightness temperatures as high as $10^{15}\\,$K (e.g., \\citealt*{swr08}; \\citealt{ob08}) and sometimes higher have been seen from isolated stars and binaries, either relating to magnetic activity from the stars or from interaction between the members of the binary \\citep{gudel02,osten08}.  The emission seen from \\gcrt\\ does not really resemble the known radio emission from such sources, although this may reflect the limitations of our observations.  As discussed above, most sources have been observed to be significantly fainter than \\gcrt, but the comparison may be unfair since the behavior across such a wide range of frequencies and timescales is not known.  If we remove the flux/luminosity constraints, the emission from \\gcrt\\ is actually not dissimilar to some ultracool dwarfs (such as the results of \\citealt{brr+05} and \\citealt{had+06}).  Time-resolved analysis at frequencies of $<500\\,$MHz ($>60\\,$cm) have been rare especially in recent years \\citep[e.g.,][]{vdodb94}, and we really do not know how known radio sources behave there.  It seems likely that either \\gcrt\\ represents new low-frequency behavior from a known class of sources, or that it is indeed the first member of a new class. Searches for counterparts to \\gcrt\\ at other wavelengths (especially X-rays, where the searches so far have not been very constraining; \\citealt{hlk+05}), along with further characterization of the radio properties of both \\gcrt\\ and other transient sources, are required to help elucidate its nature."
    },
    "0807/0807.3391_arXiv.txt": {
        "abstract": "s{ Motivated by recent Pierre Auger result on the correlation of the highest-energy cosmic rays with the nearby active galactic nuclei, we explore possible ultrahigh energy cosmic ray (UHECR) source distributions and their effects on GZK horizons. Effects on GZK horizons by local over-density of UHECR sources are examined carefully with constraints on the degree of local over-density inferred from the measured UHECR spectrum. We include the energy calibration effect on the Pierre Auger data in our studies. We propose possible local over-densities of UHECR sources which are testable in the future cosmic ray astronomy.} ",
        "introduction": "Recently, Pierre Auger observatory published results on correlation of the highest-energy cosmic rays with the positions of nearby active galactic nuclei (AGN) \\cite{Abraham:2007bb}. Such a correlation is confirmed by the data of Yakutsk \\cite{Ivanov:2008it} while it is not found in the analysis by HiRes \\cite{Abbasi:2008md}. In the Auger result, the correlation is maximal for the threshold energy of cosmic rays at $5.7\\times 10^{19}$ eV, the maximal distance of AGN at $71$ Mpc and the maximal angular separation of cosmic ray events at $\\psi=3.2^{\\circ}$. Due to increasing efforts on verifying the Auger result, it is worthwhile to examine the above correlation from a phenomenological point of view.  Since the angular scale of the observed correlation is a few degrees, one expects that these cosmic ray particles are predominantly light nuclei. The effect of GZK attenuations on these cosmic ray particles \\cite{Greisen,ZK} can be described by a distance scale referred to as ``GZK horizon\". By definition, the GZK horizon associated with a threshold energy $E_{\\rm th}$ is the radius of a spherical region which is centered at the Earth and produce $90\\%$ of UHECR events arriving on Earth with energies above $E_{\\rm th}$.  Assuming a uniform distribution of UHECR sources with identical cosmic ray luminosity and spectral index \\cite{Harari:2006uy}, the GZK horizon for protons with $E_{\\rm th}=57$ EeV is about $200$ Mpc while the V-C catalog\\cite{VC06} used by Pierre Auger for the correlation study is complete only up to $100$ Mpc. Such a deviation may arise from non-uniformities of spatial distribution, intrinsic luminosity and spectral index of local AGN as mentioned in \\cite{Abraham:2007bb}. In addition, the energy calibration also plays a crucial role since the GZK horizon is highly sensitive to the threshold energy $E_{\\rm th}$. Energy values corresponding to the dip and the GZK cutoff of UHECR spectrum were used to calibrate energy scales of different cosmic ray experiments \\cite{Berezinsky:2008qh,Kampert:2008pr}. It has been shown that all measured UHECR energy spectra can be brought into good agreements by suitably adjusting the energy scale of each experiment \\cite{Berezinsky:2008qh}. Furthermore, it has been shown that a different shower energy reconstruction method infers a $30\\%$ higher UHECR energy than that determined by Auger's fluorescence detector-based shower reconstruction \\cite{Engel:2007cm}.  In this presentation, we report our results \\cite{Lu:2008wj} on examining the consistency between Auger's UHECR correlation study and its spectrum measurement. The impact by the local over-density of UHECR sources is studied. We also study the energy calibration effect on the estimation of GZK horizon and the spectrum of UHECR. Certainly a $20\\%-30\\%$ upward shift on UHECR energies reduces the departure of theoretically calculated GZK horizon to the maximum valid distance of V-C catalog \\cite{Abraham:2007bb}. The further implications of this shift will be studied in fittings to the shifted Auger spectrum. ",
        "conclusions": "We have discussed the effect of local over-density of UHECR sources on shortening the GZK horizon. The result is summarized in Table I. It is seen that such an effect is far from sufficient to shorten the GZK horizon at $E_{\\rm th}=57$ EeV to $\\sim 100$ Mpc for a local over-density consistent with the measured UHECR spectrum. With a $30\\%$ energy shift, each cosmic ray event in Auger's correlation study would have an energy above $74$ EeV instead of $57$ EeV. GZK horizons corresponding to $E_{\\rm th}=74$ EeV then match well with the maximum valid distance of V-C catalog. Fittings to the shifted Auger spectrum indicate a possibility for the local over-density of UHECR sources.  We point out that the local over-density of UHECR sources is testable in the future cosmic ray astronomy where directions and distances of UHECR sources can be determined. Table IV shows percentages of cosmic ray events that come from sources within $30$ Mpc for different values of $E_{\\rm th}$ and $n(l<30 {\\rm Mpc})/n_0$. Although these percentages are calculated with $\\gamma=2.4$, $m=3$ and $E_{\\rm cut}=1000$ EeV, they are however not sensitive to these parameters. \\begin{table}[hbt] \\begin{center} \\caption{ Percentages of cosmic ray events originated from sources within 30 Mpc for different values of $E_{\\rm th}$ and $n(l<30 {\\rm Mpc})/n_0$.} \\begin{tabular}{|c|c|c|c|c|}\\hline $n(l<30 {\\rm Mpc})/n_0$ & $E_{\\rm th}=57 \\, {\\rm EeV}$ & $E_{\\rm th}=70 \\, {\\rm EeV}$ & $E_{\\rm th}=80 \\, {\\rm EeV}$ & $E_{\\rm th}=90 \\, {\\rm EeV}$ \\\\ \\hline $1$ & $0.17$ & $0.27$ & $0.36$ & $0.46$\\\\ \\hline $2$ & $0.30$ & $0.43$ & $0.53$ & $0.63$\\\\ \\hline $4$ & $0.46$ & $0.60$ & $0.70$ & $0.77$\\\\ \\hline $10$ & $0.68$ & $0.79$ & $0.85$ & $0.89$\\\\ \\hline \\end{tabular} \\end{center} \\label{event_dist} \\end{table} For $E_{\\rm th}=57$ EeV and $n(l<30 {\\rm Mpc})/n_0=1$, only $17\\%$ of cosmic ray events come from sources within $30$ Mpc. For $n(l<30 {\\rm Mpc})/n_0=2$ and the same $E_{\\rm th}$, $30\\%$ of cosmic ray events are originated from sources in the same region.  In conclusion, we have shown that the deviation of theoretically calculated GZK horizon to the maximum valid distance of V-C catalog can not be resolved by merely introducing the local over-density of UHECR sources. On the other hand, if Auger's energy calibration indeed underestimates the UHECR energy, such a discrepancy can be reduced. More importantly, fittings to the shifted Auger spectrum indicate a possible local over-density of UHECR sources, which is testable in the future cosmic ray astronomy.   \\noindent{\\bf Acknowledgements}  We thank A. Huang and K. Reil for helpful discussions. We also thank F.-Y. Chang, T.-C. Liu and Y.-S. Yeh for assistances in computing. This work is supported by National Science Council of Taiwan under the grant number 96-2112-M-009-023-MY3."
    },
    "0807/0807.4264_arXiv.txt": {
        "abstract": "Hydrogen-helium mixtures at conditions of Jupiter's interior are studied with first-principles computer simulations. The resulting equation of state (EOS) implies that Jupiter possesses a central core of 14 -- 18 Earth masses of heavier elements, a result that supports core accretion as standard model for the formation of hydrogen-rich giant planets. Our nominal model has about 2 Earth masses of planetary ices in the H-He-rich mantle, a result that is, within modeling errors, consistent with abundances measured by the 1995 Galileo Entry Probe mission (equivalent to about 5 Earth masses of planetary ices when extrapolated to the mantle), suggesting that the composition found by the probe may be representative of the entire planet. Interior models derived from this first-principles EOS do not give a match to Jupiter's gravity moment $J_4$ unless one invokes interior differential rotation, implying that jovian interior dynamics has an observable effect on the measured gravity field. ",
        "introduction": "The discovery of over two hundred extrasolar planets resembling the giant planet Jupiter in mass\\footnote{Schneider, J., the Extrasolar Planets Encyclopaedia http://exoplanet.eu} and composition~\\citep{citation2} has raised fundamental questions about the origin and inner structure of these bodies.   Establishing the existence of a dense core in Jupiter is vital for understanding the key processes of planetary formation. An observed correlation between the metallicity of the parent star and the likelihood of giant planets orbiting it~\\citep{citation3} may indicate that an initial core aggregated from solid planetesimals triggers the gravitational collapse of nebular gases to form giant planets~\\citep{citation4}, or alternatively, increased metallicity enhances the direct collapse of giant planets from the nebula~\\citep{citationX}. Current planetary models~\\citep{SC95,citation6} predict only a very small core between 0 and 7 Earth masses ($M_{\\rm E}$) for Jupiter, lending support to core-accretion theories with comparatively small cores~\\citep{Po96}, late-stage core erosion scenarios~\\citep{Gu04}, or suggesting that jovian planets are able to form directly from gases without a triggering core~\\citep{citation7}. Using first-principles simulations for hydrogen-helium mixtures at high pressure, we show here that Jupiter possesses a significant central core of 16$\\pm$2 $M_{\\rm E}$ of heavier elements, a result that supports core accretion as the standard model for formation of hydrogen-rich giant planets. In our model, Jupiter's mantle, defined as the H-He-rich outer layers, shows no enrichment of heavier elements, indicating that compositional measurements from the 1995 Galileo Entry Probe mission may be representative of most of the planet. The derived interior models also provide the first evidence that Jupiter's interior does not rotate as a solid body because no match to Jupiter's gravity moment $J_4$ is obtained unless one invokes interior differential rotation, implying that jovian interior dynamics have an observable effect on the measured gravity field.  While laboratory techniques cannot yet probe deep into Jupiter's interior (reaching pressures P $\\sim$ 100--1000 GPa and temperatures T $\\sim$ 10000 K), advances in first-principles computer simulation techniques have made it possible to characterize the equation of state (EOS) of hydrogen-helium mixtures for the entire planet. The determination of the EOS, in combination with the modeling reported here, enables us to more precisely specify Jupiter's mantle metallicity and core mass. ",
        "conclusions": ""
    },
    "0807/0807.2959_arXiv.txt": {
        "abstract": "We introduce a model which uses semi-analytic techniques to trace formation and evolution of galaxy disks in their cosmological context. For the first time we model the growth of gas and stellar disks separately. In contrast to previous work we follow in detail the angular momentum accumulation history through the gas cooling, merging and star formation processes. Our model successfully reproduces the stellar mass--radius distribution and gas-to-stellar disk size ratio distribution observed locally. We also investigate the dependence of clustering on galaxy size and find qualitative agreement with observation. There is still some discrepancy at small scale for less massive galaxies, indicating that our treatment of satellite galaxies needs to be improved. ",
        "introduction": "Understanding the origin of galaxy disks is an important aspect of the study of galaxy formation and evolution. Almost all observed star formation occurs either in galaxy disks or in material which came from galaxy disks (e.g. in starbursts). The size of disks is closely related to their gas surface density which in turn is critical in setting the star formation rate. Galaxy sizes also correlate with many other physical properties, stellar mass, luminosity, circular velocity, etc. These relations, as well as their evolution, pose strong constraints on galaxy formation models.   A long-standing problem for cosmological simulations has been to reproduce galaxies with the proper size distribution. The high efficiency of gas cooling at early times, causes the later assembly of disks to proceed by coalescence of cold gas clumps. As these clumps merge onto the main galaxy, they lose much of their initial angular momentum through dynamical friction. The resulting disks are then substantially smaller than observed and contain relatively little stellar masss. Much effort has been directed to solving this problem, most of it invoking some form of feedback to delay collapse (\\cite[Sommer-Larsen et al. 1999]{Sommer1999}, \\cite[Thacker and Couchman 2001]{Thacker2001}), or adopting an alternate initial fluctuation spectrum with reduced small-scale power (\\cite[Sommer-Larsen et al. 2001]{Sommer2001}). Some authors (\\cite[Governato et al. 2004]{Governato2004}) claim the problem can at be significantly reduced by improving the numerical resolution, but there is no consensus yet on the true solution.  In this work we implement a new treatment of disk formation within the Munich galaxy formation model used in \\cite[De lucia \\& Blaizot 2007]{delucia2007}, hereafter DLB07. Our model distinguishes between gas and stellar disks and treats their growth separately. We model the accumulation of mass and angular momentum in a gas disk by conserving the specific angular momentum of the cooling gas and the cold gas component of accreted satellite galaxies. When star formation occurs in the disk, we further assume that the cold gas which is converted into stars has the typical specific angular momentum of the current gas disk. We apply our model to the Millennium Simulation and we compare the stellar mass--size relation of disks and the distribution of the ration of stellar and gas disk sizes to observation.  The dependence of galaxy clustering on galaxy size is also investigated. ",
        "conclusions": "We have modeled the growth of gas and stellar disks in galaxies separately using a semi-analytic approach which tracks the accumulation of their angular momentum.  The differing angular momentum accumulation histories of galaxies and of dark matter halos make it possible for a galaxy to have higher or lower specific angular momentum than, and to be misaligned with, its surrounding halo. The gas and stellar disks can also be misaligned, allowing the modelling of galaxy warps.  We show that our model can reproduce many observed size-related relations including the stellar mass vs. size relation and the gas-to-stellar disk size ratio distribution. Both the mean values and the scatter are well reproduced, especially for galaxies less massive than $10^{10.5}M_{sun}$. For the dependence of clustering on galaxy size, we qualitatively reproduce observed trends, but some discrepancy remains on the small scale for less massive galaxies. Further improvement of our treatment of satellite galaxies, including tidal stripping, may help to clarify the reasons behind this."
    },
    "0807/0807.0501_arXiv.txt": {
        "abstract": "MicroLux is a GPS-based high precision and high speed timing add-on to the Calar Alto Lucky Imaging camera AstraLux. It allows timestamping of individual CCD exposures at frame rates of more than 1 kHz with an accuracy better than one microsecond with respect to the UTC timeframe. The system was successfully used for high speed observations of the optical pulse profile of the Crab pulsar in January and November 2007. I present the technical design concept of MicroLux as well as first results from these observations, in particular the reconstructed pulse profile of the pulsar\\footnote{Based on observations collected at the Centro Astron\\'omico Hispano Alem\\'an (CAHA) at Calar Alto, operated jointly by the Max-Planck Institut f\\\"ur Astronomie and the Instituto de Astrof\\'isica de Andaluc\\'ia (CSIC).}. ",
        "introduction": "High spatial or spectral resolution observing techniques are nowadays a common standard at large modern optical telescopes. However, only few instruments exist that explore the temporal dimension with resolution higher than a few 10\\,Hz. On one hand, classical CCD imaging cameras have typical readout times of several ten seconds, resulting in increasingly small duty cycles if high temporal resolution is demanded. Cameras with fast readout capabilities on the other hand, e.g. frame transfer CCDs, suffer from increased readout noise, limiting the minimum brightness or exposure time significantly.  Nevertheless, high temporal resolution at optical wavelengths has its applications, especially in the near infrared (NIR), where sky brightness dominates the noise budget even at short exposure times. High speed photometry of stellar occultations by the moon, planets, or planetary rings and satellites, is a standard observing technique since several decades now. Typical instruments built for this purpose at visible wavelengths are based on photon counting Avalanche photodiodes\\cite{Kanbach:2003,Stefanescu:2008} or conventional CCDs using high readout clocks\\cite{Dhillon:2007}.  With the advent of electron multiplying CCDs (EMCCD) with virtually zero readout noise and single-photon detection capabilities at wavelengths below 1\\,$\\mu$m, high speed photometry of faint targets in the visible is now technically feasible even with medium sized telescopes, and with minimum technical effort. The Lucky Imaging camera AstraLux\\cite{Hormuth:2008} at the Calar Alto 2.2-m telescope is such a system. Originally designed for high spatial resolution observations, the frame transfer EMCCD can be read out with frame rates of up 1.5\\,kHz, depending on binning and windowing settings. If all 512$\\times$512\\,pixels of the 24$\\times$24\\,arcsec large field of view are read unbinned, the resulting frame rate is still $\\approx$34\\,Hz.  Unfortunately, high temporal resolution alone without precise knowledge of the absolute times of the individual measurements is not always sufficient. While the fringe pattern observed during a stellar occultation can be used to determine the star's diameter on the basis of the known frame rate only, absolute timing is needed for precision measurements of e.g. eclipsing binary star parameters or the radio-optical delay of a pulsar. Decades ago, the signals of terrestrial navigation systems or television stations were used as time and frequency standards\\cite{Papaliolios:1970,Lohsen:1981}. Today, the Global Positioning System (GPS) satellite network is a powerful and reliable tool for absolute time measurements, and in fact GPS based timing solutions are employed in modern high speed photometry instruments.  \\begin{figure}[htb!] \\begin{center} \\begin{tabular}{c} \\includegraphics[width=20.5cm,angle=90]{gpslayout.eps} \\end{tabular} \\end{center} \\caption[example] { \\label{fig:layout} Block layout of the MicroLux timing system. See text for further description of the individual components. } \\end{figure}  \\begin{figure}[htb!] \\begin{center} \\begin{tabular}{c} \\includegraphics[width=12cm,angle=0]{gpsocxo.eps} \\end{tabular} \\end{center} \\caption[example] { \\label{fig:gpsocxo} The Jupiter Pico T GPS receiver on its adapter board and the Isotemp OCXO 143-55 module used in MicroLux.\t} \\end{figure}   \\begin{figure} \\begin{center} \\begin{tabular}{c} \\includegraphics[width=5.7cm,angle=270]{vco.ps} \\includegraphics[width=5.7cm,angle=270]{pd.ps} \\end{tabular} \\end{center} \\caption[example] { \\label{fig:vco} Measurements of the VCO voltage and the phase difference between the PLL signals after disconnection of the GPS receiver's antenna at 20\\,s after measurement start. While the PLL maintains perfect synchronisation for more than a minute, the loss of satellite signal is clearly indicated by oscillations of the VCO voltage after only few seconds. Monitoring the VCO voltage and its first derivative is thus an extremely sensitive tool for the detection of GPS signal loss conditions. } \\end{figure}  Commercially available devices for GPS time tagging often allow for only low frame rates, not suitable for very high speed photometry, or were specifically designed for complex synchronisation tasks, and are quite expensive. In order to enable photometry at more than 1\\,kHz frame rate with sub-microsecond accuracy and at an affordable price, the MicroLux system was designed to enable precision time-tagging of individual frames taken with the AstraLux Lucky Imaging camera. ",
        "conclusions": "At total hardware costs of less than 1000\\,Euro and within less than two months of development time it was possible to design and built MicroLux, a high-speed, high-precision time-tagging system. Using a GPS disciplined ovenised crystal oscillator, MicroLux can determine the start times of individual CCD exposures with an accuracy better than 1\\,$\\mu$s at frame rates of more than 1000\\,Hz. Originally built for the Lucky Imaging camera AstraLux at the Calar Alto 2.2-m telescope, MicroLux could be used with other image acquisition systems as well thanks to its universal and modularised design.  MicroLux has been used mainly for observations of the optical pulse profile of the Crab pulsar up to now, demonstrating that the intended design goals could be fully reached. Further analysis will make use of the absolute timing capabilities of MicroLux with respect to the UTC timeframe to determine the pulsar's radio-optical delay and its dependency on wavelength.\\\\ Other applications include precise timing of binary star eclipses or stellar occultations by the moon, planets, planetary satellites, or asteroids.  Further planned improvements of MicroLux aim at increased versatility, higher maximum supported frame rates, and miniaturisation of the system, hopefully leading to a final design suitable for field-operation, e.g. in the case of stellar occultations by asteroids, where multiple observers with timing capabilities are needed to gather full information. The relatively simple design, based mainly on commercially available parts, and the low price are certainly great benefits in this respect."
    },
    "0807/0807.2504_arXiv.txt": {
        "abstract": "We present optical and IR spectra of a set of low-mass stars and brown dwarfs in the Coronet cluster (aged $\\sim$1 Myr), obtained with the multifiber spectrograph FLAMES on the VLT and with the IRS instrument on Spitzer. Most of the objects had been selected via their X-ray emission in a deep Chandra survey. The optical spectra reveal spectral types between M1 and M7.5, confirm the youth of the objects (via Li 6708\\,\\AA\\ absorption), and show the presence of accretion (via H$\\alpha$) and shocks (via forbidden line emission). The IRS spectra, together with IR photometry from the IRAC/MIPS instruments on Spitzer and 2MASS, confirm the presence of IR excesses characteristic of disks around $\\sim$70\\% of the objects. Half of the disks do not exhibit any silicate emission, or present flat features characteristic of large grains ($>$6\\,$\\mu$m). The rest of the disks show silicate emission with indications of amorphous and crystalline silicate grains a few microns in size. About 50\\% of the objects with disks do not show near-IR excess emission, corresponding to the presence of ``transitional'' disks, according to their classical definition. This is a very high fraction for such a young cluster. The large number of ``transitional'' disks suggests lifetimes comparable to the lifetimes of typical optically thick disks. Therefore, these disks may not be in a \\textit{short-lived} phase, intermediate between Class II and Class III objects. The median spectral energy distribution of the disks in the Coronet cluster is also closer to a flat disk than observed for the disks around solar-type stars in regions with similar age.  The differences in the disk morphology and evolution in the Coronet cluster could be related to fact that these objects have very late spectral types compared to the solar-type stars in other cluster studies. Finally, the optical spectroscopy reveals that one of the X-ray sources is produced by a Herbig Haro object in the cloud. ",
        "introduction": "}  The formation of very low-mass stars and brown dwarfs (BD) has been a matter of discussion for a long time (Luhman et al. 2006). The difficulties in fragmenting a cloud or stopping the envelope mass accretion at or near the hydrogen burning limit suggested that very low-mass objects may not form like solar-type stars, but require other mechanisms like ejection in multiple systems (Bate et al. 2005; Umbreit et al. 2005) or photoerosion by nearby high-mass stars that can remove part of the initial core material (Withworth \\& Zinnecker 2004). At the same time, multiple studies have focused on the detection of accretion disks around very low-mass young stars and BD (Fern\\'{a}ndez \\& Comer\\'{o}n 2001; Natta \\& Testi 2001; Jayawardhana et al. 2002, 2003; Pascucci et al. 2003; Klein et al. 2003; Luhman et al. 2005; Mohanty 2004, 2005). H$\\alpha$ spectroscopy reveals active accretion, and the excesses in the near- and mid-IR are consistent with disks similar to the disks around classical T Tauri stars (CTTS; Muench et al. 2001; Apai et al. 2002, 2005; Pascucci et al. 2003; Mohanty et al. 2004; Jayawardhana et al. 2005; Scholz \\& Jayawardhana 2008). IR spectroscopy finds further similarities: dust evolution, grain growth/settling and disk flattening, formation of crystalline silicates, and inner holes (Furlan et al. 2005;  Apai et al. 2005; Gizis et al. 2005; Sterzik et al. 2004; Sargent et al. 2006; Natta et al. 2007; Morrow et al. 2008; Henning 2008), which are thought to be the first steps in the formation of planetary systems, occur in these disks. Therefore, objects of all masses may be formed by cloud fragmentation and subsequent disk formation and accretion (Luhman et al. 2007).  Nevertheless, our understanding of disk evolution and accretion mechanisms in very low-mass stars and BD is still far from complete. The timescales for very low-mass disk evolution, expected to be longer than those of solar-type T Tauri stars (TTS) and Herbig Ae/Be stars (HAeBe; Sterzik et al. 2004; Dahm \\& Hillenbrand 2007; Bouy et al. 2007), are still very uncertain (Lada et al. 2006), due to the lack of statistically significant samples of coeval very low-mass stars and BD. IR observations of M-type stars reveal differences in innermost disk evolution (Kessler-Silacci et al. 2007; Sicilia-Aguilar et al. 2007), which may be the result of differences in the disk structure, presence of dead zones, and accretion mechanisms (Hartmann et al. 2006) and could affect the formation of planetary systems in such disks. As it happens with the solar-type TTS, deeper and more complete multiwavelength observations suggest that there is an important fraction of harder-to-detect, flattened and transitional disks with inner holes around the lower-mass objects. For these very low-mass objects, with masses $<$0.2 M$_\\odot$ down to the BD regime, the morphology of the disks could be different (Morrow et al. 2008), but the difficulties in observing faint objects and determining the low-mass initial mass function in the regions translate into very small samples and large uncertainties. If, like we observe in solar-type TTS, individual differences in disk evolution are remarkable even for objects of similar ages and spectral types, the knowledge of disks and accretion around very low-mass stars and BD and their evolution is highly limited at the present time, and may be biased toward ``special'' systems (with large IR excess and strong accretion rates).  The Coronet cluster (also known as CrA) is an obscured star-forming region, located at a distance of $\\sim$150 pc (d$\\sim$ 130 pc according to Marraco \\& Rydgren 1981 and De Zeeuw et al. 1999; d$\\sim$ 170 pc according to Knude \\& Hog 1998), associated with the Herbig Ae star R CrA (Taylor \\& Storey 1984) and a dense molecular core (Loren 1979). The Coronet cluster is relatively isolated at the edge of the Gould Belt. It is located out of the galactic plane, and it is probably part of the Sco OB2 Association (De Zeeuw et al. 1999) and of the large Sco-Cen region (De Geus et al. 1992). Cambr\\'{e}sy (1999) estimated a total mass of 1600 M$_\\odot$ for the whole cloud, using a distance of 170 pc. Casey et al. (1998) obtained an age of 3 Myr for the TY CrA multiple system, considering a distance of 130 pc; the cluster would be younger if the distance is larger. An age around 1 Myr is more consistent with the presence of nebular material and Class 0 objects (Chen et al. 1997). Nisini et al. (2005) estimated an age of 0.5-1 Myr based on 4 embedded objects, being therefore one of the closest and youngest clusters. Submillimeter, millimeter, IR, and optical surveys (Henning et al. 1994; de Muizon et al. 1980; Wilking et al. 1985; Marraco \\& Rydgren 1981) found a few HAeBe and TTS, suggesting a modest stellar population. The most massive stars in the region are the Herbig Ae star R CrA and the B9 star TY CrA. X-ray surveys using Einstein, XMM-Newton, ROSAT, and Chandra  (Walter 1986; Walter et al. 1997; Neuh\\\"{a}user et al. 1997, 2000; Hamaguchi et al. 2005a,b) identified several low-mass members. Deeper X-ray observations (Garmire \\& Garmire 2003) found 118 potential (1$\\sigma$) sources, many of which were consistent with young stars, some of them deeply embedded, which have not been described in detail yet.  Here we present the results of optical and IR spectroscopy on a collection of low- and very low-mass members of the Coronet Cluster. The new and archive data, including the Chandra survey, optical spectroscopy with FLAMES/VLT, IR spectroscopy with IRS/Spitzer, and IRAC and MIPS/Spitzer photometry are introduced in Section \\ref{data}, where we also describe the selection methods of the sample of very low-mass objects. The derived stellar and disk properties are presented in Section \\ref{analysis}, and their implications for star formation and disk evolution are discussed in Section \\ref{discussion}. Finally, the results are summarized in Section \\ref{conclusions}. ",
        "conclusions": "}  The study of the disks around very low-mass objects in the Coronet cluster reveals significant differences with other regions. The disk fraction is smaller ($\\sim$70\\%) than in clusters with a similar age, and more striking, about $\\sim$50\\% of the disks show ``transitional'' SEDs with inner holes (TO). The SEDs of the disks without inner holes are significantly flatter than observed in similarly young clusters and even in regions a few Myr older (Sicilia-Aguilar et al. 2006a), and the median SED is close to that of a flat disk ($\\lambda$F$_\\lambda$ $\\sim$ $\\lambda^{-4/3}$). The disks with no inner gaps show typically evidence of accretion, while only some of the TO are actively accreting. In any case, the H$\\alpha$ EW and profiles suggest that accretion rates are probably low to moderate ($<$10$^{-8}$-10$^{-9}$\\,M$_\\odot$/yr), as it is common for very low-mass objects.  Since the members of the Coronet cluster have later spectral types than typical disks studies (M0-M7.5 instead of G-early M), our work suggests that M-type stars and BD tend to have flatter disks and inner holes with a significantly higher frequency than solar-type stars. Partial evidence of this had been suggested in other studies, based on smaller samples of stars (Megeath et al. 2005; McCabe et al. 2006; Lada et al. 2006; Allers et al. 2006; Morrow et al. 2007).  The fact that most disks around the very low-mass members of the Coronet cluster have inner holes or are strongly flattened may be responsible for the controversy about disk fractions for very low-mass stars and BD. It also reveals that observations at mid-IR wavelengths are required in order to determine accurately the disk fractions for M-type stars and BD.  The large fraction of ``transitional'' disks suggest a lifetime comparable to the typical disk lifetimes, which is inconsistent with the rapid timescales due to disk photoevaporation (Alexander et al. 2006a,b), and with the relatively short lifetimes inferred for solar-type stars in other regions (Hartmann et al. 2005; Sicilia-Aguilar et al. 2006a). Therefore, at least some of the ``transitional'' disks may not be in a short-lived, intermediate evolutive stage.  The youth of the Coronet cluster members suggests that part of these disks could have been formed already as flattened or ``transition-like'' structures. Whether this effect is due to the formation conditions of these low-mass objects in a low-mass star-forming region like the Coronet cluster (angular momentum of the collapsing cores, presence of binary companions), or is a consequence of a different evolutive path of disks around the less massive sources, needs to be explored by targeting large samples of very low-mass objects in older regions and in massive star-formation environments.  Finally, the optical spectroscopy confirms that one of the X-ray sources in the Coronet cluster corresponds to a newly detected HH object."
    },
    "0807/0807.3114_arXiv.txt": {
        "abstract": "GH\\,10 is a broad-lined active galactic nucleus (AGN) energized by a black hole of mass 800,000~$M_\\odot$.  It was the only object detected by Greene et al.\\ in their Very Large Array (VLA) survey of 19 low-mass AGNs discovered by Greene \\& Ho.  New VLA imaging at 1.4, 4.9, and 8.5~GHz reveals that GH\\,10's emission has an extent of less than 320~pc, has an optically-thin synchrotron spectrum with a spectral index $\\alpha = -0.76\\pm0.05$ ($S_\\nu \\propto \\nu^{+\\alpha}$), is less than 11\\% linearly polarized, and is steady - although poorly sampled - on timescales of weeks and years. Circumnuclear star formation cannot dominate the radio emission, because the high inferred star formation rate, 18~$M_\\odot$~yr$^{-1}$, is inconsistent with the rate of less than 2~$M_\\odot$~yr$^{-1}$ derived from narrow H$\\alpha$ and [O~II] $\\lambda$3727 emission. Instead, the radio emission must be mainly energized by the low-mass black hole.  GH\\,10's radio properties match those of the steep-spectrum cores of Palomar Seyfert galaxies, suggesting that, like those Seyferts, the emission is outflow-driven.  Because GH\\,10 is radiating close to its Eddington limit, it may be a local analog of the starting conditions, or seeds, for supermassive black holes. Future imaging of GH\\,10 at higher linear resolution thus offers an opportunity to study the relative roles of radiative versus kinetic feedback during black-hole growth. ",
        "introduction": "Although little is known about the mass function of nuclear black holes with mass below about $10^6~M_\\odot$, this is an important regime for several reasons.  For example, the remarkable correlation between host bulge velocity dispersion and black hole mass strongly hints that black holes play an essential role in galaxy evolution \\citep{kor04}, yet there are few observational constraints on the starting conditions, or so-called seeds, for the supermassive black holes with masses above about $10^6~M_\\odot$.  Also, until the {\\em Laser Interferometer Space Antenna\\/} launches, an understanding of the low-mass end of the black-hole mass function may provide one of the few constraints on these seed black holes \\citep{hug02}. Moreover, it remains to be established whether small galaxies can commonly host nuclear black holes and, if so, whether the same mass versus dispersion relation mentioned above is obeyed.  Motivated by these and similar issues, \\citet{gre04} conducted the first systematic search for a population of black holes with mass below about $10^6~M_\\odot$.  They used SDSS DR1 to select a sample of broad-lined active galactic nuclei (AGNs) with virial mass estimates $10^{5-6}~M_\\odot$.  This sample of 19 objects was the first uniformly selected sample of low-mass AGNs.  The sample is characterized by Eddington ratios, $L_{\\rm bol}/L_{\\rm Edd}$, that are close to unity, implying radiatively efficient accretion onto these black holes.  With its low-mass black holes and high Eddington ratios, this sample probes a poorly explored region of parameter space, potentially providing new insights into the physical drivers of radio properties in AGNs.  For this reason, \\citet{gre06} observed the sample at 4.9 GHz with the Very Large Array (VLA).  A single object, \\object{GH 10}, was detected with a 4.9-GHz--to--optical flux ratio $R \\sim 2.8$, formally radio-quiet \\citep{kel89}.  A stacked image of the remaining 18 objects yielded an upper limit of $R \\le 0.27$, emphasizing the atypical traits of GH\\,10.  To help understand the physical origin of these atypical traits, new multi-frequency VLA data were obtained on GH\\,10 and are presented in Section 2.  As summarized in Section 3.1, the radio emssion from GH\\,10 is found to be compact, steep-spectrum, unpolarized, and steady in time.  A star-formation origin for this emission is considered, and dismissed, in Section 3.2, while Section 3.3 explores the implications of a black-hole origin for the radio emission.  The paper ends, in Section 4, with a summary of the key findings and suggestions for future directions.  The following cosmological parameters are assumed: $H_0 = 71$~km~s$^{-1}$~Mpc$^{-1}$, $\\Omega_m = 0.27$, and $\\Omega_\\Lambda = 0.75$ \\citep{spe03}, leading to a luminosity distance of 363~Mpc \\citep{gre04,gre06}.  Also, the sign convention adopted for the spectral index $\\alpha$ near a frequency $\\nu$ is the flux density $S_\\nu \\propto \\nu^{+\\alpha}$. ",
        "conclusions": "GH\\,10 was the only object detected by \\citet{gre06} in their VLA survey of the 19 low-mass AGNs discovered by \\citet{gre04}.  New VLA imaging reveals that the emssion from GH\\,10 has an extent of less than 320~pc, with an optically thin synchrotron spectrum that is less than 11\\% linearly polarized.  The radio emission is steady, although poorly sampled, on timescales of weeks and years.  Circumnuclear star formation cannot dominate the radio emission, because the inferred star formation rate, 18~$M_\\odot$~yr$^{-1}$, is inconsistent with the upper limit of 2~$M_\\odot$~yr$^{-1}$ derived from the narrow optical emission lines.  Instead, the radio emission must be dominantly energized by the low-mass black hole, which is thought to be radiating near its Eddington limit.  The radio properties of GH\\,10 strongly resemble those of the steep-spectrum cores of Palomar Seyfert galaxies, which feature outflow-driven emission unmodified by the effects of internal and/or external opacity.  This suggest that, like those Palomar Seyferts, the radio emission from GH\\,10 is outflow-driven, a suggestion that can be tested with future VLBI imaging.  Such imaging could also offer an opportunity to study the relative roles of radiative versus kinetic feedback during black-hole growth, given that GH\\,10 may be a local analog of the starting conditions, or seeds, for supermassive black holes.  Finally, the compact, steep-spectrum, unpolarized, and steady nature of the emission from GH\\,10 will guide future observations of the new 1.4-GHz detections of low-mass AGNs reported by \\citet{gre07b}."
    },
    "0807/0807.1057_arXiv.txt": {
        "abstract": "{Understanding the gas abundance distribution is essential when tracing star formation using molecular line observations. Changing density and temperature conditions cause gas to freeze-out onto dust grains, and this needs to be taken into account when modeling a collapsing molecular cloud.} {This study aims to provide a realistic estimate of the CO abundance distribution throughout the collapse of a molecular cloud. We provide abundance profiles and synthetic spectral lines which can be compared to observations.} {We use a 2D hydrodynamical simulation of a collapsing cloud and subsequent formation of a protoplanetary disk as input for the chemical calculations. From the resulting abundances, synthetic spectra are calculated using a molecular excitation and radiation transfer code.} {We compare three different methods to calculate the abundance of CO. Our models also consider cosmic ray desorption and the effects of an increased CO binding energy. The resulting abundance profiles are compared to observations from the literature and are found to agree well.} {The resulting abundance profiles agree well with analytic approximations, and the corresponding line fluxes match observational data. Our developed method to calculate abundances in hydrodynamical simulations should greatly aid in comparing these to observations, and can easily be generalized to include gas-phase reaction networks.} ",
        "introduction": "The study of low-mass Young Stellar Objects (YSOs), from early protostellar cores to T Tauri stars surrounded by disks, involves observations of either the thermal emission from dust grains or molecular line emission. While dust emission yields information about the temperature profile \\citep[e.g.,][]{draine1984}, and, through measurements of the shape of the Spectral Energy Distribution, the evolutionary stage~\\citep{lada1984, adams1987}, spectral lines are the only way to constrain the kinematical properties of YSOs. The interpretation of molecular lines are crucial for the understanding of protoplanetary disk formation and the distribution of angular momentum during these stages of star formation~\\citep{evans1999}. However, deriving the gas motion from spectral line profiles is complicated by degeneracies between the velocity field topology, inclination of the angular momentum axis, optical depth, and geometrical effects~\\citep{brinch2007}.  In addition to these effects, molecular abundances may vary greatly. The gas abundance is determined by the ongoing chemistry driven by the co-evolving density and temperature distributions, and to some extent also the velocity field. The variation in molecular abundances throughout a YSO may have a huge impact on line profiles and therefore it is important to understand the chemistry in order to interpret these line profiles. The effect of the chemistry is enhanced when observations at higher resolution are done. If the observations are done at a relatively low resolution ($\\gtrsim$10$''$), the emission is smeared out over a large portion of the object and therefore an average constant abundance is mostly adequate. However, when going to a higher resolution, using sub-millimeter interferometry ($\\sim$1$''$), the emission is averaged over a smaller part of the object and thus variations in the abundance need to be taken into account.  The problem of determining the abundance distribution has been addressed before in the literature. \\citet{lee2004} took a semi-analytical approach using a self-similar collapse of an isothermal sphere in which they followed ``fluid-elements''. In that paper, the chemistry was followed in a series of Bonnor-Ebert spheres during the pre-stellar phase and in a self-similar inside-out collapse during the protostellar phase. This model provides a good analytical description of the early stages of star formation, but, lacking any rotational velocity, no disk is formed in this scenario. \\citet{aikawa2001,aikawa2003, aikawa2005} followed the chemical evolution of contracting pre-stellar clouds using an approach similar to \\citet{lee2004} but using an isothermal cloud collapse only. Their calculations consider only the very early stages of the star formation process. \\citet{jorgensen2002,jorgensen2005} introduced the ``drop model'', where chemical depletion occurs as a step function when certain temperature and density conditions of any underlying model are met. While this approach is simple yet quite successful, there may be cases where the abundance changes only very gradually so that a step function is no longer a good approximation. More recently, \\citet{Tsamis2008} calculated HCO$^+$, CS, and N$_2$H$^+$ abundances and spectra from a simulation of an inside-out collapsing core.   In this paper we investigate detailed evolution of the gas-phase and solid state abundance of CO in a collapsing, rotating core. We base our study on a hydrodynamical simulation, where we follow the CO freeze-out and evaporation for a number of test particles that flow with the gas. The lay-out of the paper is as follows: In Section~\\ref{model} we present our numerical simulations and the equations used to solve the CO freeze-out chemistry. In Section \\ref{results} we show three different chemical models and how our results compare to observations as well as the results of previous papers. Sections \\ref{discussion} and~\\ref{summary} hold a discussion and a summary respectively. ",
        "conclusions": "\\label{summary} In this paper we have investigated different scenarios for the depletion of CO in a collapsing cloud. We have  used a hydrodynamical code to describe the dynamical evolution of the cloud and the chemistry has been solved using a trace particle approach.  Of the various models shown in Sect.~\\ref{results} the one which exclude cosmic ray desorption and the ones using higher binding energies differ the most. The drop abundance model provides a good approximation to Model 2 and 3. Therefore, if it is possible to fit a drop profile to observations, as was done by~\\citet{jorgensen2005}, it should be possible to run a hydrodynamical simulation with customized initial conditions, so that it reproduces this profile and thereby putting an observational constraint on the hydrodynamics. Also in cases where an analytical model is used, using the drop abundance approximation is quite good.  We furthermore conclude that the freeze-out chemistry has little effect on optically thick spectral lines, especially when observed in low resolution. It is therefore not entirely unreasonable to model such lines using a constant abundance model disregarding the chemistry. However, one should be careful when using optically thin lines to determine which average gas phase abundance to use for a constant abundance model, since the average abundances derived from such lines are indeed depending on the freeze-out chemistry. The maximum discrepancy between the flat abundance profile model (Fig.~\\ref{spec3}) and the freeze-out models shown in Fig.~\\ref{spec2} is a factor of three in the integrated intensity. Such a change in the average abundance does not affect the optically thick lines a lot, so depending on the accuracy needed, using a flat abundance distribution to model data in which freeze-out is present may not be entirely unreasonable.  Finally, the effect of the chemical depletion is considerably enhanced when the resolution of the observations is increased, making it more important to get the abundance profiles right. This will especially be important for future ALMA observations (with a resolution of $\\sim$ 0.01$''$) where the level of detail will be so great that getting the exact abundances will be crucial for interpreting the data.\\\\  \\noindent \\emph"
    },
    "0807/0807.4444_arXiv.txt": {
        "abstract": "The transient accreting millisecond pulsar SAX J1808.4--3658 has shown several outbursts to date but the transition from outburst to quiescence has never been investigated in detail. Thanks to the Swift observing flexibility, we monitored for the first time the decay to quiescence during the 2005 outburst. At variance with other transients, wide luminosity variations are observed. In addition, close to quiescence, SAX J1808.4--3658 seems to switch between two different states. We interpret them in terms of the accretion states accessible to a magnetized, fast rotating neutron star. ",
        "introduction": "SAX J1808.4--3658 (SAX J1808 in the following) was the first neutron star transient during whose outbursts coherent pulsations were detected (Wijnands \\& van der Klis 1998). This source, together with six other low mass transients (see Wijnands 2005 for a review), form a separate (sub-)class within neutron star transients (see Campana et al. 1998a for a review). Distinctive properties, besides coherent pulsations, are weaker outburst peak luminosities ($\\sim (1-5)\\times 10^{36}\\ergs$), very low mass companions (mass functions $<2\\times 10^{-3}\\msole$), short orbital periods ($P_{\\rm orb}\\lsim 4.3$ hr), very faint quiescent luminosities ($\\lsim 5\\times 10^{31}\\ergs$) and absence of a soft X--ray spectral component in quiescence.  A further difference with respect to `classical' neutron star transients concerns the return to quiescence following an outburst. The best light curve to date of `classical' neutron star transients is represented by BeppoSAX observations of Aql X-1 (Campana et al. 1998b). Below a luminosity level of $\\sim 10^{36}\\ergs$ Aql X-1 turned off with an exponential decay with an $e-$folding time of $\\sim 1$ d. The source then remained quiescent for at least one month following the 1998 outburst (Campana et al. 1998b) and for five months following the 2000 outburst (Rutledge et al. 2001; Campana \\& Stella 2003). The return to quiescence for transient pulsating neutron stars has been monitored for SAX J1808 (Wijnands et al. 2003; Wijnands 2005) and for XTE J1751--305 (Markwardt et al. 2002) and Swift J1756.9--2508 (Krimm et al. 2007). The behaviour of the latter two sources was  somewhat similar to Aql X-1 with a fast decay, even though in the case of XTE J1751--305 there was still activity 15 d after the transition to quiescence had occurred (according to RXTE/PCA, corresponding to a 0.5--10 keV unabsorbed luminosity of $\\sim 10^{35}\\ergs$ for a source distance of 7 kpc). In case of SAX J1808 the decay was completely different from that observed in Aql X-1. The most striking feature is the strong erratic variability by a factor of $\\gsim 30$. This erratic behaviour was not concentrated in the first few weeks after turn off but lasted for months (e.g. Wjinands et al. 2004). The most likely explanation for this behaviour comes from the similarities with dwarf novae where a similar behaviour has been observed (Kato et al. 2004; Patterson et al. 2002) and to a lesser extent in black hole transients (Kuulkers, Howell \\& van Paradijs 1996). In the case of dwarf novae this has been interpreted as due to a combination of short orbital period and low mass ratio (Osaki \\& Meyer 2002; Truss et al. 2002; Campana et al. 2008, in preparation).  Here we report on a monitoring campaign carried out with Swift, following the 2005 outburst. At variance with previous monitoring with RXTE/PCA we are able to follow the source down to very low levels, a factor of $\\sim 50$ fainter than before, close to quiescence. In Sect. 2 we discuss the data and their analysis. Sect. 3 contains our discussion and conclusions. ",
        "conclusions": "SAX J1808 is the prototype of the accreting millisecond X--ray pulsar (AMP) class. The source underwent several outbursts to date which have been monitored in great details thanks to RXTE observations (Wijnands 2005; Wijnands et al. 2003). These observations however were limited to the brightest part of the outburst, being the PCA on RXTE a collimated instrument (and therefore heavily background limited). Here we report on the first campaign aimed at studying the faintest tail of the outburst and return to quiescence. Thanks to the fast repointing and flexible scheduling capabilities of the Swift satellite, we monitored the return to quiescence of SAX J1808 during the 2006 outburst. The bright phase of the outburst is similar to what has been observed in the past: following the peak and the smooth decay a flaring behaviour sets in with at least three rebrightening episodes. After this the source turns to quiescence even if same low level activity (factor of $\\sim 5$) is still present.  The puzzling result of this observational campaign on SAX J1808 is that out of 13 observations with luminosity below $\\sim 10^{34}\\ergs$, 8 times the source was found with a luminosity around $\\sim 5\\times 10^{32}\\ergs$ (see Fig. 1), i.e. at a luminosity level $\\sim 10$ times higher than the true quiescent level (Campana et al. 2002; Heinke et al. 2007). Given the wild luminosity variations this is somewhat strange and calls for the presence of a `metastable' state. This luminosity level is difficult to interpret in case of standard disk accretion but finds a natural explanation if the accretion process is mediated by the magnetic field of the fast spinning neutron star. When the magnetospheric radius (i.e. the radius at which the neutron star magnetic field starts controlling the accretion flow) is larger than the corotation radius, the system is in the (so-called) propeller regime. In this case matter is halted from the rotating magnetosphere and does not accrete onto the neutron star surface (see e.g. Campana \\& Stella 2000 and references therein). When the magnetospheric radius becomes larger than the light cylinder radius ($r_{\\rm lc}=c\\,P/2\\,\\pi$, where $c$ is the velocity of light and $P$ the spin period), the magnetic field cannot corotate any more with the neutron star and a dipole losses will take place. The luminosity range spanned in the propeller regime is $\\sim 440\\,P_{\\rm 2.5}^{3/2}\\,M_{1.4}^{-3/2}$ (where $P_{\\rm 2.5}$ is the spin period in units of 2.5 ms and $M_{1.4}$ the neutron star mass in units of $1.4\\msole$; see Campana \\& Stella 2000). In order to have the `metastable' state to coincide with the lowest luminosity in the propeller regime (i.e. just before the reactivation of the pulsar), we have to require a magnetic field of $\\sim 7\\times 10^7$ G, in line with previous estimates: $\\sim (1-10)\\times 10^8$ G using disk-magnetosphere interaction models (Psaltis \\& Chakrabarty 1999), $\\sim (1-6)\\times 10^8$ G using simple considerations on the position of the magnetospheric radius during quiescent periods (Di Salvo \\& Burderi 2003); $(0.4-1.5)\\times 10^8$ G modelling period changes during outbursts with a magnetic torque model (Hartman et al. 2008).  The lowest luminosity level in the proposed interpretation is ascribed to the interaction of a turned-on pulsar with the interbinary and circumstellar environment. Variability is expected at this stage due to the rapidly changing environment as observed, e.g., in the millisecond radio pulsar PSR J1740--5340 (D'Amico et al. 2001; Ferrario et al. 2001) which gets eclipsed over a range of orbital phases for different orbital cycles."
    },
    "0807/0807.2262_arXiv.txt": {
        "abstract": "{Galactic synchrotron emission, rotation measure (RM) and the deflection of ultra-high-energy-cosmic-rays (UHECR) permit detailed studies of the Galactic magnetic field (GMF). The synchrotron emission has also to be measured in order to be separated from other astrophysicaly interesting signals like the CMB.} {We present a publicly available code called \\textsc{hammurabi} for generating mock polarized observations of Galactic synchrotron emission for telescopes like LOFAR, SKA, Planck and WMAP, based on model inputs for the GMF, the cosmic-ray density distribution and the thermal electron density. We also present mock UHECR deflection measure (UDM) maps based on model inputs for the GMF. In future, when UHECR sources are identified, this will allow us to define UDM as a GMF probe in a similar way as polarized radio sources permit us to define rotation measures.} {To demonstrate the code's abilities mock observations are compared to real data as a means to constrain the input parameters of our simulations with a focus on large-scale magnetic field properties.} {The Galactic magnetic field models in the literature seem to fail to reproduce any additional observational data which was not included in their design.} {As expected, attempts at trying to model the synchrotron, UHECR deflection and RM input parameters, show that any additional observational data set greatly increases the constraints on the models. The \\textsc{hammurabi} code addresses this by allowing to perform simulations of several different data sets simultaneously, providing the means for a more reliable constraint of the magnetized inter-stellar-medium.} ",
        "introduction": "There are several different observations which probe the Galactic magnetic field (GMF): synchrotron radiation, rotation measure, UHECR deflection, dust related observations (e.g. starlight polarization \\citet{Heiles2000}, polarized dust emission), and more recently atom alignment spectroscopic observations were proposed \\citep[see][]{2006ApJ...653.1292Y}.  A better understanding of the inter-stellar magnetized plasma and the radio emission processes related to it is also paramount for current and future CMB experiments \\citep[see, e.g.,][]{MAMD2008, Page2007}, where in particular measurements of the polarized CMB signal will be limited by our knowledge of foreground emission \\citep{Tucci2005}.  Observational studies of the GMF are usually based on one single sort of observable. However, since the different observables and measurement methods provide complementary information, constraining the GMF is logically better if all possible observations are considered simultaneously. We address the need to confront those observations with models by presenting a publicly available software\\footnote{\\tt Software available at \\\\ http://www.mpa-garching.mpg.de/$\\sim$waelkens/hammurabi} capable of generating mock synchrotron and Faraday rotation observations as well as mock UHECR deflection maps.  The code is constructed in such a fashion that it should be feasible to adapt it to do any sort of all or partial-sky line-of-sight integral observations, (see, e.g., \\citealt{Waelkens2008} and \\citealt{Sun2008}).  We will give first a brief account of previous work on GMF modeling in Sect. \\ref{sec::previous} and then we will give a description of the involved Faraday rotation, UHECR deflection and synchrotron physics, and the simplifying assumptions used in our computation (Sect. \\ref{sec::physics}). The implementation of a line-of-sight (hereafter LOS) integration scheme which mimics the radiative transfer is discussed in Sect. \\ref{sec::Implementation}. We present some test output using simple standard magnetic field, thermal- and cosmic-ray electron models. We discuss the role of the turbulent field and sub-grid modeling in Sect. \\ref{sec::SSGMF}, and the one of helical magnetic fields in Sect. \\ref{sec::HGMF}. Finally we conclude in Sect. \\ref{sec::Conclusion}. ",
        "conclusions": "\\label{sec::Conclusion} We have presented the \\textsc{hammurabi} code, a tool for simulating rotation measure maps, total and polarized Galactic synchrotron emission maps taking into account Faraday rotational-depolarization effects, as well as UHECR deflection maps.  Drawing a set of well known input models from the literature (NE2001 for the thermal electron distribution, \\citet{Sun2008} and \\citet{Page2007} for the cosmic-ray electron distribution and the Galactic magnetic field), we show example outputs comparing them to corresponding observations to illustrate the code's abilities as a scientific tool for charting the Galactic magnetized plasma and the cosmic-ray electron distribution.  Full galaxy simulations are currently limited by the finite HEALPix-grid resolution, thus being unable to probe fluctuations with characteristic lengths of less than $4 {\\rm pc}$. A second caveat is that the widely used approximation of a power-law energy distribution of the cosmic-ray electrons has been shown to be questionable by sophisticated cosmic ray propagation simulations \\citep{Strong2007}. The degree to which this affects the precision of our simulations has yet to be assessed. Furthermore, our non-ray-tracing calculations of UHECR deflections might not fulfill strict precision requirements as soon as the UHECR's trajectory through the Galaxy is significantly bent; they serve rather as a first approximation.  It was shown that, unsurprisingly, models designed to fit only a fraction of the available observational data on GMF might fail to reproduce the remaining observational information not taken into account in their construction.  This is an indication that the constraints to the GMF might be highly degenerate. Hence confronting the models with the broadest possible range of observations on the Galactic magnetic field is paramount to achieving any useful statement about the field.  Furthermore we have displayed \\textsc{hammurabi}'s capabilities of generating low-frequency mock observations, where Faraday effects play a significant role. This can be used for feasibility studies and the analysis of the actual observations of forth coming low-frequency radio telescopes like LOFAR or the SKA (note that the latter can also observe at high-frequencies).  \\textsc{hammurabi} can be applied for constraining the Galactic magnetized plasma and increasing our understanding on radio Galactic emission and UHECR source locations. Already a worthy scientific goal by itself, it will also have implications on our understanding on foreground emission subtraction for experiments like Planck or WMAP."
    },
    "0807/0807.2106_arXiv.txt": {
        "abstract": "We present a population synthesis study of the observed properties of the magnetars investigating the hypothesis that they are drawn from a population of progenitors that are more massive than those of the normal radio pulsars. We assume that the anomalous X-ray emission is caused by the decay of a toroidal or tangled up field that does not partake in the spin down of the star. Our model assumes that the magnetic flux of the neutron star is distributed as a Gaussian in the logarithm about a mean value that is described by a power law $\\Phi=\\Phi_{0}\\left(\\displaystyle{\\frac{M_{\\rm p}}{9\\Msun}}\\right)^{\\gamma}$~G~cm$^2$ ($8 \\Msun \\le M_{\\rm p}\\le 45\\Msun$) where $M_{\\rm p}$ is the mass of the progenitor.  We find that we can explain the observed properties of the magnetars for a model with $\\Phi_{0}=2 \\times 10^{25}$~G~cm$^2$ and $\\gamma=5$ if we suitably parametrise the time evolution of the anomalous X-ray luminosity as an exponentially decaying function of time. Our modelling suggests that magnetars arise from stars in the high mass end ($20\\Msun \\le M_{\\rm p} \\le 45\\,\\Msun$) of this distribution. The lower mass progenitors are assumed to give rise to the radio pulsars.  The high value of $\\gamma$ can be interpreted in one of two ways.  It may indicate that the magnetic flux distribution on the main sequence is a strong function of mass and that this is reflected in the magnetic fluxes of the neutron stars that form from this mass range (the fossil field hypothesis). The recent evidence for magnetic fluxes similar to those of the magnetars in a high fraction ($\\sim 25$\\%) of massive O-type stars lends support to such a hypothesis.  Another possibility is that the spin of the neutron star is a strong function of the progenitor mass, and it is only for stars that are more massive than $\\sim 20\\Msun$ that magnetar-type fields can be generated by the $\\alpha-\\omega$ dynamo mechanism (the convective dynamo hypothesis).  In either interpretation, it has to be assumed that all or a subset of stars in the mass range $\\sim 20\\Msun - 45\\Msun$, which on standard stellar evolution models lead to black holes via the formation of a fall-back disc, must give rise to magnetars.  Unlike with the radio pulsars, the magnetars only weakly constrain the birth spin period, due to their rapid spin-down.  Our model predicts a birthrate of $\\sim 1.5-3\\times 10^{-3}$~yr$^{-1}$ for the magnetars. ",
        "introduction": "The anomalous X-ray pulsars (AXPs) and Soft Gamma Repeaters (SGRs) have spin down properties indicative of magnetic fields of $\\sim 10^{14-15}\\,$G which places them at the high end of the neutron star field distribution.  The source of the quiescent X-ray luminosity in both the SGRs and the AXPs is unclear, but it is generally believed that it is associated with the decay of some component of the magnetic field.  The standard model of \\citet{Duncan92} assumes that the fields are generated at the time of formation of the neutron star by an efficient $\\alpha- \\omega$ dynamo that operates at low Rossby numbers and requires millisecond birth periods.  Neutron stars born with initial periods $P_i$ are predicted to generate large scale magnetic fields of $3\\times 10^{17}\\,{\\rm G}\\left ({1\\,\\rm ms/P_i}\\right)$ under optimum conditions, which is more than adequate to explain the fields in magnetars.  A consequence of the rapid initial spin predicted for these models is that the supernova explosions that create the magnetars are expected to be an order of magnitude more energetic than ordinary core-collapse supernovae, if one makes the standard assumption that angular momentum is lost by magnetic braking and not by gravitational radiation or due to emission in a jet. However, the energetics of some well studied supernova remnants associated with magnetars, appear to suggest that their formation may not always be accompanied by a hypernova explosion \\citep{Vink06}. The origin of magnetic fields in magnetars therefore remains unclear.  Recent observations have provided two additional clues on the origin of magnetars. Firstly, several magnetars have been associated with massive star clusters. The magnetar candidate CXO J164710.2-455216 has been associated with the open cluster Westerlund~1 \\citep{Muno06} and SGR 1806-20 with the giant Galactic HII complex W31 \\citep{Figer05}, both suggesting progenitors with $M\\gsimeq 40\\Msun$. Massive ($\\gsimeq 20\\Msun$) progenitors appear also to be indicated for SGR 1627-41 \\citep{Corbel99} and SGR 0526-66 \\citep{Klose04}. Interestingly, the progenitors have masses in the range $\\sim 20 -45\\Msun$ where interaction with a fall-back disc is expected to play an important role in determining the ultimate outcome of stellar evolution - namely whether a neutron star or a black hole forms \\citep{Heger03}.  A second clue comes from the discovery that more massive stars tend to have higher magnetic fluxes than less massive stars, and that the incidence of magnetism on the main sequence increases with mass.  Some $25$\\% of O and early B stars in the Orion cluster are strongly magnetic with fields in the range $600 - 1500$~G (\\citet{Petit08}, \\citet{Donati02}) so that the progenitors of the magnetars are also likely to be strongly magnetic.  Fossil magnetic fluxes similar to those observed in the magnetars may already be present in stellar cores prior to collapse and could give rise to global magnetic fields of $\\sim 10^{15}$~G either directly through magnetic flux conservation, or indirectly by acting as seed fields for a dynamo. In this context, we note that while the dynamo mechanism is effective in generating weak fields of the order of the equi-partition value in the convective cores of stars with radiative envelopes \\citep{Brun05}, much stronger fields can be generated by the dynamo mechanism if a fossil field is present in the radiative region \\citep{Featherstone06}.  In this paper, we carry out a population synthesis study to analyse the observed numbers of magnetars and their magnetic field and period distributions allowing for X-ray selection effects. Our basic premise is that the birth magnetic field of a neutron star is related primarily to the mass of the progenitor star $M_{\\rm p}$. Our aim is to determine the form of this relationship in the hope that it may provide some insights into the origin of magnetars. ",
        "conclusions": "\\label{results}  We have computed a series of models for different values $M^*$, $\\gamma$, $\\sigma_{\\log\\Phi0}$, $P_{0}$, $\\sigma_{P0}$, $L_0$, $\\tau_{d0}$ and $\\delta$ until a model is obtained that satisfies the following criteria: (a) the model must envelope and provide a reasonable representation of the observed period and magnetic field distributions and (b) the total number of observed magnetars must be less than the theoretically predicted number after X-ray selection effects are taken into consideration. Given the small number of magnetars and the unknown number of active, but currently dormant magnetars (e.g. objects such as SGR~1627-41 that  has only recently re-activated after a nearly 10 year long quiescent period \\citep{Palmer08}), the real Galactic distributions is not well enough sampled to justify a more detailed modelling.  \\begin{figure*} \\label{best_fit} \\begin{center} \\hspace{0.1in} \\epsfxsize=0.7\\textwidth \\epsfbox[35 286 586 703]{best_fit.ps} \\caption{Our model (solid line) overlapped to the magnetars data sample (dashed histograms) taken from http://www.physics.mcgill.ca/~pulsar/magnetar/main.html and with the addition of the recently confirmed member of the AXP class, 1E1547.0-5408 (=PSR~J1550-5418) \\citep{Camilo07}.} \\end{center} \\end{figure*}  The total number of active magnetars in the Galaxy depends strongly on $M^*$ because of the rapidly declining initial mass function. Our calculations show that in addition to the initial mass function, the parameters $\\gamma$, $\\Phi_{0}$, $\\sigma_{\\log\\Phi_{0}}$, $L_0$, $\\tau_{d0}$ and $\\delta$ together determine the magnetic field, period, and period derivative distributions.  Our results are rather insensitive to $P_{0}$ and $\\sigma_{P_0}$ because of the rapid initial spin down of magnetars. Hence, this study cannot exclude the possibility that magnetars are born at millisecond periods as required by some recent models \\citep{Geppert06}, or that the birth spin period may be a function of magnetic field as assumed in FW.  For simplicity, in our models we have adopted $P_0=3$ seconds and $\\sigma_{P_0}=1$ second, although the FW prescription with an upper initial period cutoff at $8$ seconds is also acceptable.  We present in Figure~1 our favoured model with $M^*=20\\Msun$, $\\log\\Phi_{0} = 25.2$, $\\sigma_{\\log\\Phi_{0}} = 0.4$, $L_0=10^{34}$, $\\tau_{d0}=5\\times 10^5$~yr, and $\\delta=1.3$. The corresponding birthrate is $3\\times 10^{-3}$~yr$^{-1}$. This model best envelopes the field, period, period derivative, characteristic age, luminosity and temperature distributions and predicts 26 active magnetars after applying the {\\it ROSAT} detectability criteria.  We exclude the possibility of a significantly lower $M^*$ (e.g. $\\le 18\\Msun$) on the grounds that it will lead to a conflict with the birth rate of radio pulsars which we assume arise from stars with $M\\le M^*$. On the other hand, the use of a significantly higher $M^*$ ($> 22 \\Msun$) leads to too few magnetars and inconsistencies with the observed distributions.  We have also constructed a $P-\\dot {P}$ diagram which is shown in Figure 2 where our {\\it ROSAT} selected objects are compared with the observations of all magnetars with known $\\dot P$. Thus, among the observed objects we have also included SGR~0526-66, which is located in the LMC, and CXOU~J010043-721134 which is in the SMC, since their position in the $P-\\dot {P}$ diagram is unlikely to depend on which galaxy they reside.  The predicted number of active magnetars from our favoured model can be reconciled with the observed number of 5 {\\it ROSAT} detected sources if we assume that only $\\sim 20$\\% of magnetars were in an active state which brought them to a high enough luminosity to be detected by this satellite. In this case, the actual birthrate of magnetars would be $\\sim 3\\times 10^{-3}\\,$yr$^{-1}$, which of course depends on our basic assumption that all stars in the mass range $20-45\\Msun$ produce magnetars. However, we note that the total number of observed magnetars in the Galaxy is only 14. If no additional magnetars exist in our Galaxy, that is if all magnetars have already been discovered, this means that we have produced too many magnetars from progenitors in the mass range $20-45\\Msun$. This may suggest that some 50\\% of stars in this mass range will produce black holes. If this is the case, the birth rate of magnetars would reduce to $\\sim 1.5\\times 10^{-3}\\,$yr$^{-1}$.  Studies of SGR bursts conducted by \\citet{Kouveliotou94} indicated that there cannot be more than 7 active SGRs in the Galaxy and \\citet{Kouveliotou98} suggested that magnetars are born at a rate of about $0.1$ per century.  The recent study carried out by \\citet{Gill07} was based on a study of the five AXPs detected in the {\\it ROSAT} All-Sky Survey, and indicated a birthrate of 0.22 per century with their progenitors being massive main sequence stars. Both these estimates are generally consistent with the birth rate that we deduce. Another recent study carried out by \\citet{Muno08} who searched for magnetars in archival Chandra and XMM-Newton observations of the Galactic plane, yielded a birthrate in the range $3\\times 10^{-3}-6\\times 10^{-2}\\,$ yr$^{-1}$. The upper end of this range is excluded by our calculations.  On the fossil field hypothesis, the high value indicated for the magnetic flux index $\\gamma$ may simply reflect the intrinsic magnetic flux distribution on the main sequence.  Observations of main sequence stars show that the maximum magnetic fluxes observed on the main sequence map on to the magnetic fluxes of the highest field magnetic white dwarfs ($10^8-10^9$~G ) and neutron stars ($10^{14}-10^{15}\\,$G) rather well. However, due to sensitivity limitations of polarimetric observations, only the upper end of the magnetic field distribution of main sequence stars (the strongly magnetic stars) can be observed, so we only have a partial picture of magnetism on the main sequence. Unfortunately, most stars in the mass range $8-20\\Msun$ that give rise to radio pulsars have magnetic fluxes well below the currently observed range on the Main Sequence so that the index $\\gamma$ cannot be empirically estimated. We note however that magnetic fluxes similar to those inferred in magnetars occur in $\\sim 25$\\% of massive B and O-type stars.  If we somehow dismiss the close similarities between the magnetic fluxes of massive main sequence stars and magnetars by putting them down to mere coincidence, alternatives to the fossil field hypothesis need to be explored. It is possible that the mass of the progenitor determines the spin of the nascent neutron star and thereby the strength of a dynamo generated field. Support for this hypothesis comes from the calculations of \\cite{Heger05} which allow for angular momentum transport by magnetic fields generated by differential rotation during stellar evolution. These show that more massive stars tend to produce more rapidly spinning neutron stars.  For a $35\\Msun$ progenitor (the highest mass they considered), they predict that the neutron star will have a spin period of $3$~ms which is rapid enough for the generation of a a magnetar-type field by the $\\alpha-\\omega$ dynamo mechanism of \\citet{Duncan92}. However, for this to be a viable explanation of the magnetars, the $\\alpha-\\omega$ dynamo mechanism would need to be effective for neutron stars that arise from progenitors with significantly lower masses ($\\sim 20-22\\Msun$) or the derived birthrate of magnetars would be too low to explain the observations. Furthermore, since the calculations of \\citet{Heger05} have been seriously challenged by \\citet{Zahn07}, it appears that we are still waiting for more stellar evolution calculations, that allow for both magnetic field and rotation, to show us whether or not millisecond rotation periods can result from massive stars. If such periods can be achieved, this would give support to the idea that the dynamo model is a viable alternative to the fossil field hypothesis, although it is likely that a fossil field could still play the role of a seed field.  \\begin{figure} \\begin{center} \\epsfxsize=0.4\\textwidth \\epsfbox[32 168 524 646]{ppdot_diag.ps} \\caption{Filled squares: observed magnetars (http://www.physics.mcgill.ca/~pulsar/magnetar/main.html) with the addition of 1E1547.0-5408 \\citep{Camilo07}; filled circles: magnetars as derived from our modelling (see text). The solid line is an empirically determined boundary (a ``magnetar death line'') given by $\\log(\\dot P)=8.4\\log(P)-20$. } \\end{center} \\label{ppdot_diag} \\end{figure}  It is relevant that we comment on the alternative to our basic hypothesis on spin evolution, namely that the field that decays is also the field that drives spin evolution during the magnetar phase. We have considered this possibility using ``Avenue C'' of \\citet{Colpi00}, but excluded it from our present considerations because it failed to populate the high field end of the observed magnetar field distribution. It is conceivable that with the use of different field decay parameters for avenue C we may also be able to model the data, but we expect that such a model will need to have field characteristics during the magnetar phase that are so similar to those adopted in the present calculations that our major conclusions will remain largely unchanged.  In conclusion, we note that the origin of the crustal field component in magnetars remains unresolved.  \\citet{Zahn07} have shown that the shearing of a poloidal field of fossil origin in a differentially rotating radiative region can lead to the generation of a toroidal field so that field complexity of the type that appears to be required in magnetars may arise naturally from stellar evolution, when the complex interplay between rotation and fossil magnetic fields is taken into consideration. On the other hand, since fall-back discs are expected to play a role in the evolution of stars in the mass range $\\sim 20 - 45\\Msun$, it is tempting to speculate that the crustal field component that characterises the magnetars may have its origin in the interaction of the fall-back disc with the magnetic field of the stellar core, in cases where a magnetar is the outcome of stellar evolution. If this is the case, one would expect that most of the mass in the fall-back disc would be magnetically ejected."
    },
    "0807/0807.1796_arXiv.txt": {
        "abstract": "{} {We have studied the characteristic timescales of 80 AGNs at 22, 37 and 90\\,GHz examining the properties of the wavelet method and comparing them to traditional Fourier-based methods commonly used in astronomy.} {We used the continuous wavelet transform with the Morlet wavelet to study the characteristic timescales. We also gain information when the timescale is present in the flux curve and if it is persistent or not.} {Our results show that the sources are not periodic and changes in the timescales over a long time are common. The property of wavelets to be able to distinguish when the timescale has been present is superior to the Fourier-based methods. Therefore we consider it appropriate to use wavelets when the quasi-periodicities in AGNs are studied.} {} ",
        "introduction": "Wavelet methods have been used for the past 10 years to study the variability timescales in different types of astronomical objects \\citep[e.g.][]{foster96, priestley97, scargle97}. Wavelets have properties superior to Fourier based methods, the most useful being the locality of the timescale. In Fourier based methods a single sinusoid is fitted to the whole light curve, whereas wavelets enable us to study temporally local timescales or periods that are present in the light curves. This is a very good property when studying flux curves of AGNs, which are almost never strictly periodic over long times.  Wavelets were used to study variability timescales of a larger sample of AGNs by \\cite{kelly03}. They used the cross-wavelet transform to study quasi-periodic variations in the Pearson-Readhead VLBI sample, monitored by the University of Michigan Radio Observatory (UMRAO). A quasi-periodic behaviour was found in a little over half of the 30 sources they studied. Individual sources have also been studied for quasi-periodicity and periodicity using wavelets \\citep[e.g.][]{hughes98, lehto99, kadler06, ciprini07}.  In \\cite{hovatta07}, hereafter Paper I, we studied the variability timescales of a large sample of AGNs using the structure function (SF), the discrete autocorrelation function (DCF) and the Lomb-Scargle periodogram. We found that larger flares in these sources happen quite rarely, only once in 4--6 years while short-term variations are continuously present in the flux curves. Also, many of the sources had changed their behaviour during the 25 years of monitoring which is one reason why we wanted to study the sample again using wavelets. In Paper I we also studied the properties of the methods and their differences by searching for the characteristic timescales. We noticed that there are several problems when using periodograms to study the timescales because none of the sources in our sample showed strict periodicities in their flux curves or were well characterised by sinusoids.  In this paper we will compare the results from this wavelet analysis to Paper I. This paper is organised as follows: In Sect. \\ref{sect:sample} we describe the source sample and data used. Section \\ref{sect:method} includes the description of the method and in Sect. \\ref{sect:results} we present the results. The discussion is in Sect. \\ref{sect:discussion} and the conclusions are drawn in Sect. \\ref{sect:conclusions}. ",
        "conclusions": "\\label{sect:conclusions} We studied a sample of 80 sources with the continuous wavelet transform using data at 22, 37 and 90\\,GHz. Our aim was to study the variability behaviour of the sources and also to better understand the method and to compare it with Fourier-based methods. We found no clear periodicities in the sources. Instead in most of the sources the timescales appeared only for a short time in the total monitoring period or changed over the years. These kind of properties are not revealed with Fourier-based methods and therefore wavelets should be used more when quasi-periodicities are studied. In our sample only 1/4 had a timescale that had been present for at least 4 cycles. In eight of these sources the timescale was present in more than one frequency band. This shows that the sources are never strictly periodic and conclusions about periodicities in the radio frequencies should be drawn with caution.  The average timescale from the wavelet analysis is 4.5 years which is the same as was obtained in the DCF analysis of Paper I. This timescale should not be affected by the length of the monitoring period because we used all the sources in which the timescale had repeated at least twice in the analysis and therefore longer timescales also are included in the calculations. The weakness of wavelets is that we cannot detect the timescale with the same accuracy as with DCF or periodograms and therefore both type of methods should be used when periodicities or quasi-periodicities are searched for. The wavelets should be used to verify whether the timescale is persistent or short-lived. If a persistent timescale is found, the Fourier methods can be used to accurately determine the timescale. In our analysis where the characteristic timescales are studied, the accuracy of wavelets is enough to determine the approximate timescale."
    },
    "0807/0807.1105_arXiv.txt": {
        "abstract": "Many recent observational studies have concluded that planetary systems commonly exist in multiple-star systems.  At least $\\sim$20\\,\\%, and presumably a larger fraction of the known extrasolar planetary systems are associated with one or more stellar companions. These stellar companions  normally exist at large distances from the planetary systems (typical projected binary separations are on the orders $10^2\\,$--\\,$10^4$AU) and are often faint (ranging from F to T spectral types). Yet, secular cyclic angular momentum exchange with these distant stellar companions can significantly alter the orbital configuration of the planets around the primaries.  One of the most interesting and fairly common outcomes seen in numerical simulations is the opening of a large mutual inclination angle between the planetary orbits, forced by  differential nodal precessions caused by the binary companion. The growth of the mutual inclination angle between planetary orbits induces additional large-amplitude eccentricity oscillations of the inner planet due to the quadrupole gravitational perturbation by the outer planet. This eccentricity oscillation may eventually result in the orbital decay of the inner planet through tidal friction, as previously proposed as Kozai migration or Kozai cycles with tidal friction (KCTF). This orbital decay mechanism induced by the binary perturbation and subsequent tidal dissipation may stand as an alternative formation channel for close-in extrasolar planets. ",
        "introduction": "As of June 2008, at least $\\sim$20\\,\\% of the $\\sim$240 stars hosting planets are components of multiple-star systems \\citep{desidera07,raghavan06}. Most of the binaries that harbor planets are very wide, with typical sky-projected separations ranging from $10^2$ to $10^4$\\,AU. Majority of the stellar companions to planetary systems are FGK-type stars, but there are also later-type companions, including 8 M-dwarfs and one T-type brown dwarf.  These distant, and often faint, companions, however, may still play a significant role in  altering planetary orbits around the primaries on a very long timescale.  As first pointed out by \\citet{zucker02}, the two populations of planets, those around single stars and those around components of binary systems, seem to exhibit different orbital properties. The effect of secular binary perturbations is particularly discernible in the distribution of planets' orbital eccentricities that are possibly excited by the Kozai mechanism (\\S~\\ref{kozai}). For example, a large fraction of planets with highly eccentric orbits ($e > 0.7$) are associated with stellar companions.  Also, the population of planets in binaries shows a higher median eccentricity (0.25) compared to that of the planets around single stars (0.20). \\footnote{observational biases seem to influence the observed eccentricity distribution only marginally; see for example \\citet{shen08}.} Another curious trend that was first noticed by \\citet{zucker02} and has remained robust still to date is the population of massive hot Jupiters ($\\mpl \\sin{i} > 2 \\MJ$ and $P < 100\\,$days) found exclusively among binary systems (there are a few recent exceptions such as XO-3\\,b or HD\\,17156\\,b, for which the presence of stellar companions has not yet been  ruled out).  Interestingly, secular eccentricity-pumping mechanisms such as the Kozai mechanism are closely related to the formation of close-in (possibly transiting) planets. In addition to the traditional planetary migration theory incorporating planet-disk interactions, there is a growing attention to an alternative mechanism involving circularization of initially eccentric orbits through tidal dissipation in the central stars.  There are two major ways a planet can gain a sufficiently large orbital eccentricity that leads to the tidal circularization and orbital decay: mutual, impulsive gravitational scattering between planets upon dissipation of the nascent gas disk \\citep{nagasawa08}, and adiabatic secular eccentricity growth induced by the Kozai mechanism \\citep{wu03,fabrycky07}.  The latter scenario, often called Kozai migration or Kozai cycles with tidal friction (KCTF), has been adopted for single-planet systems in binaries and produced models that are consistent with the observations, given certain assumptions in the initial conditions. In this study, we investigate the dynamical outcomes of multiple-planet systems secularly perturbed by stellar companions and find that Kozai migration process is significantly enhanced with the presence of additional planetary companions. ",
        "conclusions": ""
    },
    "0807/0807.0179_arXiv.txt": {
        "abstract": "% We study the super-orbital modulation present in the \\cyg\\ X-ray data, usually attributed to the precession of the accretion disk and relativistic jets. We find a new, strong, 326$\\pm$2~d period modulation starting in 2005, in \\swift\\ and \\asm\\ light curves (LCs). We also investigate \\vela\\ and \\ariel\\ archival data and confirm the previously reported $\\sim$290~d periodic modulation, and therefore confirming that the super-orbital period is not constant. Finally, we study \\asm\\ LC before 2005 and find that the previously reported $\\sim$150~d period is most likely an artifact due to the use of a Fourier-power based analysis under the assumption that the modulation has a constant period along the whole data sample. Instead, we find strong indications of several discrete changes of the precession period, happening in coincidence with soft and failed state-transition episodes. We also find a hint of correlation between the period and the amplitude of the modulation. The detection of gamma-rays above 100 GeV with MAGIC in September 2006 happened in coincidence with a maximum of the super-orbital modulation. The next maximum will happen between 2 and 14 of July 2008, when the observational conditions of \\cyg\\ with ground-based Cherenkov telescopes, such as MAGIC and VERITAS, are optimal. ",
        "introduction": "%  Cygnus X-1\\citep{Bowyer1965} is the best established candidate for a stellar mass black-hole (BH). It is composed of a $21\\pm 8$~M$_\\odot$ BH turning around an O9.7~Iab companion of $40\\pm 10$~M$_\\odot$ \\citep{Ziolkowski2005} in a circular orbit of 5.6 days \\citep{Brocksopp99a}. High resolution radio imaging has unveiled the presence of a highly colimated, relativistic \\citep{Stirling2001}, radiatively inefficient \\citep{Gallo2005} jet. The X-ray source displays soft and hard states and relatively frequent failed transitions between them.  There are strong evidences of a high energy non-thermal component extending up to soft gamma-rays \\citep{Mcconnell2002,Cadolle2006}. The steady emission of gamma-rays above 100 GeV is strongly constrained by the observations with MAGIC which has obtained, however, a very strong evidence of an intense, fast flaring episode at these energies \\citep{Albert07}.  A $\\sim$5.6~d period modulation, attributed to the orbital motion of the compact object around the companion, has been observed at various wavelengths by numerous authors \\citep[e.g.,][]{Pooley99,Brocksopp99a,Brocksopp99b,LaSala98,Lachowicz06}. On the other hand, a super-orbital $\\sim$290~d period was claimed by \\citet{Priedhorsky83} on the soft X-ray data recorded by \\vela\\ (1969-1979) and \\ariel\\ (1974-1980). Later, a $\\sim$150~d periodic variability has been reported by various authors \\citep{Pooley99,Brocksopp99b,Ozdemir01,Benlloch01,Benlloch04,Lachowicz06} using different data samples ranging between April 1991 and November 2003. It must be noted, however, that other significant modulations with periods $\\sim$200~d and $\\sim$420~d have been also found \\citep[e.g.][]{Benlloch01,Benlloch04,Lachowicz06}.   In this letter, we search the latest \\cyg\\ X-ray data for periodic modulations. We also perform a critical revision of the previous results obtained from archival X-ray data. Finally, we put our results in the context of a multiwavelength description of the source. ",
        "conclusions": "%  We find that \\cyg\\ displays a super-orbital modulation, with a period that changes, probably in a discrete way and in coincidence with soft or failed state-transition phases, over time scales ranging from a few hundred days to several years. According to our findings, the very much cited $\\sim$150~d period is most probably an artifact of applying a (sometimes biased) Fourier-transform based analysis to a data sample where more than one consecutive period modulations are present. Since 2005, \\cyg\\ shows a very powerful and stable super-orbital modulation with a period of $326\\pm2$~d.  The super-orbital modulation is usually attributed to the precession of the accretion disk~\\citep{Priedhorsky83} and relativistic jet~\\citep{Romero02}, as a result of the tidal forces exerted by the companion star on a tilted disk~\\citep{Katz73}. A mechanism for keeping the disk tilted can be provided by radiation pressure warping~\\citep{Petterson77,Pringle96,Wijers99,Ogilvie01}. In the case of tidally forced precession, the expected period $P_\\textrm{prec}$ depends on the outer radius $R_o$ and inclination of the disk $\\delta$ as $P_\\textrm{prec} \\propto R_o^{-3/2} \\cos^{-1}\\delta$ \\citep{Larwood98}. Then, the longer the period the larger the precession angle, and hence also a larger modulation amplitude is expected. This is in agreement with our results for \\asm, where we have found four different super-orbital periods, which follow this tendency (see Table~\\ref{tab:results} and Figure~\\ref{fig:asmA_lc}).  A different issue is why the precession of the disk produces a modulation in the X-ray flux. Some authors \\citep[e.g.,][]{Lachowicz06,Ibragimov07} have considered the possibility that the precession movement changes the optical thickness along the line of sight. They reject this possibility since it seems unlikely due to the fine tuning required to produce the observed modulation amplitude, which in addition should depend on the energy, which is not confirmed by our observations. The multiwavelenth data seem to support a scenario where the emission itself is anisotropic. The precession modulation is detected at similar times with identical periods in radio, soft and hard X-rays during the hard state. A unified picture, where the anisotropy is provided by the jet, has been proposed by \\citet{Brocksopp99b}. The soft X-ray emission is produced in the disk via bremsstrahlung of thermal electrons, and are then up-scattered to higher energies via Compton scattering in the hot corona or at the base of the relativistic jet (which precesses with the disk). The acceleration of electrons along magnetic field lines in the jet produces the radio emission by synchrotron emission. During the soft state, the jet and corona dissappear and no modulation is observed. According to our findings, once the source goes back to the hard state, the reconstructed disk and jet have different kinematical properties, and the modulation period changes. It seems that failed transitions produce a similar effect.  MAGIC detected a fast and intense episode of emission of gamma-rays above 100~GeV \\citep{Albert07} during MJD=54003, albeit at the limit of the telescope's sensitivity. This happened in coincidence with the soft and hard X-ray maxima (Figures~\\ref{fig:asm_swift_lc} and \\ref{fig:phaseogram}) and an unusually bright outburst detected with {\\it INTEGRAL} \\citep{Malzac08}. It is interesting to note that, according to the ephemeris shown in Table~\\ref{tab:results}, the next passage for the precession maximum will happen at MJD=54655$\\pm$2, i.e. between 6 and 10 of July 2008. The observational conditions of \\cyg\\ with ground-based Cherenkov telescopes, such as MAGIC and VERITAS, will be optimal during those days."
    },
    "0807/0807.5089_arXiv.txt": {
        "abstract": "{We investigated a set of high-resolution ($R\\sim45\\ 000$), high signal-to-noise ($S/N\\sim70$) quasar spectra to search for the signature of the so-called proximity effect in the \\ion{H}{i} Ly$\\alpha$ forest.  The sample consists of 40 bright quasars with redshifts in the range $2.1 < z < 4.7$. Using the flux transmission statistic, we determined the redshift evolution of the \\ion{H}{i} effective optical depth in the Lyman forest between $2\\la z\\la 4.5$, finding good agreement with previous measurements based on smaller samples. We also see the previously reported dip in $\\tau_\\mathrm{eff}(z)$ around redshift $z\\sim 3.3$, but as the significance of that feature is only 2.6$\\sigma$, we consider this detection tentative. Comparing the flux transmission near each quasar with what was expected from the overall trend of $\\tau_\\mathrm{eff}(z)$, we clearly detect the proximity effect not only in the combined quasar sample, but also towards each individual line of sight at high significance, albeit with varying strength. We quantify this strength using a simple prescription based on a fiducial value for the intensity of the metagalactic UV background (UVB) radiation field at 1~Ryd, multiplied by a free parameter that varies from QSO to QSO. The observed proximity effect strength distribution (PESD) is asymmetric, with an extended tail towards values corresponding to a weak effect. We demonstrate that this is not simply an effect of gravitational clustering around quasars, as the same asymmetry is already present in the PESD predicted for purely Poissonian variance in the absorption lines. We present the results of running the same analysis on simulated quasar spectra generated by a simple Monte-Carlo code. Comparing the simulated PESD with observations, we argue that the standard method of determining the UVB intensity $J_{\\nu_0}$ by averaging over several lines of sight is heavily biased towards high values of $J_{\\nu_0}$ because of the PESD asymmetry. Using instead the \\emph{mode} of the PESD provides an estimate of $J_{\\nu_0}$ that is unbiased with respect to his effect.  For our sample we get a modal value for the UVB intensity of $\\log J_{\\nu_0} = -21.51\\pm 0.15$ (in units of $\\mathrm{erg}\\,\\mathrm{cm}^{-2}\\,\\mathrm{s}^{-1}\\,\\mathrm{Hz}^{-1}\\,\\mathrm{sr}^{-1}$) for a median quasar redshift of 2.73. With $J_{\\nu_0}$ fixed we then corrected $\\tau_\\mathrm{eff}$ near each quasar for local ionisation and estimated the amount of excess \\ion{H}{i} absorption attributed to gravitational clustering. On scales of $\\sim 3$~Mpc, only a small minority of quasars show substantial overdensities of up to a factor of a few in $\\tau_\\mathrm{eff}$; these are exactly the objects with the weakest proximity effect signatures. After removing those quasars residing in overdense regions, we redetermined the UVB intensity using a hybrid approach of sample averaging and statistical correction for the PESD asymmetry bias, arriving at $\\log J_{\\nu_0} = -21.46^{+0.14}_{-0.21}$. This is the most accurate measurement of $J_{\\nu_0}$ to date. We present a new diagnostic based on the shape and width of the PESD that strongly supports our conclusion that there is no systematic overdensity bias for the proximity effect. This additional diagnostic breaks the otherwise unavoidable degeneracy of the proximity effect between UVB and overdensity.  We then applied our hybrid approach to estimate the redshift evolution of the UVB intensity and found tentative evidence of a mild decrease in $\\log J_{\\nu_0}$ with increasing redshift, by a factor of $\\sim 0.4$ from $z=2$ to $z=4$. Our results are in excellent agreement with earlier predictions for the evolving UVB intensity, and they also agree well with other methods of estimating the UVB intensity.  In particular, our measured UVB evolution is much slower than the change in quasar space densities between $z=4$ and $z=2$, supporting the notion of a substantial contribution of star-forming galaxies to the UVB at high redshift. } ",
        "introduction": "High-resolution spectra of high-redshift, quasi-stellar objects (QSOs) show a plethora of narrow absorption lines from various species along their lines of sight through the intergalactic medium (IGM). Most absorption lines blueward of an Ly$\\alpha$ emission line stem from the Ly$\\alpha$ transition of intervening neutral hydrogen, giving rise to the Ly$\\alpha$ forest \\citep{sargent80,weymann81,rauch98}. High-resolution spectra of the Ly$\\alpha$ forest enabled a statistical analysis of the absorbers, as well as their physical properties \\citep[e.g.][]{kim01,schaye03,kim04}.  The vast majority of \\ion{H}{i} Ly$\\alpha$ absorbers are optically thin to ionising photons and are thus kept highly photoionised by the metagalactic UV background (UVB) generated by the overall population of quasars and star-forming galaxies \\citep{haardt96,fardal98,haardt01}. The UVB intensity at a given frequency is determined by the evolving population of UV sources, modified by cosmological expansion and absorption \\citep[e.g.][]{dave99}. It is important to determine the amplitude of the UVB and its spectral shape as a function of redshift in order to constrain the relative contributions of quasars and star-forming galaxies to the UVB, as well as to provide a key ingredient to numerical simulations of structure formation.  There are different ways to constrain the \\ion{H}{i} UVB photoionisation rate in the IGM. Assuming its spectral energy distribution it is possible to infer the UVB intensity at the Lyman limit $J_{\\nu_0}$.  First, the UVB can be predicted by integrating the contributions of the observed source population. Given the luminosity functions of the sources, a characteristic intrinsic spectral energy distribution and the observed absorber distribution functions in the IGM, the amplitude and spectral shape of the UVB can be calculated numerically as the radiation is filtered by the IGM \\citep{bechtold87}.  In particular, comparing these models to observations provides constraints on the relative contributions of quasars \\citep{haardt96,fardal98} and star-forming galaxies \\citep{haardt01}.  Furthermore, the \\ion{H}{i} photoionisation rate at $2<z<4$ has been constrained by matching the Ly$\\alpha$ forest absorption and the IGM temperature evolution from numerical simulations of structure formation to observations \\citep{rauch97,theuns98,mcdonald00,meiksin04,tytler04,bolton05,kirkman05,jena05}.  A more direct approach to estimate the UVB photoionisation rate is based on the so-called proximity effect. In the vicinity of a luminous quasar the UV radiation field is expected to be enhanced, resulting in a statistical deficit of Ly$\\alpha$ absorbers near the quasar compared to far away from it along the line of sight \\citep{weymann81,carswell82,murdoch86}. Knowing the UV luminosity of the quasar, the deficit of lines (or generally Ly$\\alpha$ opacity) yields an estimate of the UVB intensity at the \\ion{H}{i} Lyman limit \\citep{carswell87,bajtlik88}. Traditionally, the proximity effect has been considered as a statistical effect and large samples of up to $\\sim 100$ quasars have been compiled to measure the mean UVB at $2\\la z\\la 4$ \\citep[e.g.][]{bajtlik88,lu91,giallongo96,cooke97,scott00,liske01}. Nevertheless, the proximity effect can also be detected towards individual quasars \\citep[e.g.][]{williger94,cristiani95,lu96,savaglio97,adaglio08}.  Proximity effect analyses at different redshifts can constrain the redshift evolution of the UVB, showing that it broadly peaks at $z\\sim 3$. The decrease in the UVB at $z< 2$ \\citep{kulkarni93,scott02} and $z\\ga 4$ \\citep{williger94,lu96,savaglio97} is likely due to the declining space density of quasars \\citep{haardt96,fardal98}. Whereas the UVB evolution  at low redshifts is sufficiently fast to be detectable within a single sample \\citep{scott02}, there is presently little direct evidence of evolution within the redshift range $2\\la z\\la 4$.  It is currently an open issue whether the measured UVB at $z\\la 3$ is consistent with the integrated emission of quasars, or whether a significant share of ionising photons has to be provided by star-forming galaxies \\citep[e.g.][]{madau99,sokasian03,schirber03,hunt04}.  One difficulty lies in the fact that the UVB estimates via the proximity effect are subject to several systematic uncertainties, all of which can lead to an overestimate of the inferred UV background: \\begin{enumerate} \\item \\textit{Quasar variability:} The size of the proximity effect zone scales as the average quasar luminosity over the photoionisation equilibrium timescale in the IGM ($\\sim 10^4$~yr). Quasar variability on shorter timescales will tend to overestimate the UVB due to Malmquist bias in the selected quasar samples \\citep{schirber04}. It is not known how important this effect is in practice. \\item \\textit{Underestimated quasar redshifts:} Quasar emission redshifts determined from broad high-ionisation lines are likely underestimated \\citep[e.g.][]{gaskell82,richards02}, causing the UVB to be overestimated by a factor $\\sim 3$ \\citep{espey93}. However, this effect can be largely avoided by using low-ionisation lines to determine systemic redshifts. \\item \\textit{Overdense environment:} In the standard ionisation model of the proximity effect \\citep{bajtlik88}, the Ly$\\alpha$ forest line density is extrapolated into the quasar vicinity, assuming that the matter distribution is not altered by the presence of the QSO.  If quasars typically reside in high intrinsic overdensities instead, the UVB will be overestimated by a factor of up to $\\sim 3$ \\citep{loeb95}. Conversely, the disagreement between the UVB estimates from the proximity effect and those obtained by matching simulations to the mean Ly$\\alpha$ forest absorption has recently been used as an argument that quasars are surrounded by substantial overdensities \\citep{rollinde05,guimares07,giguere08}. \\end{enumerate}  In the present paper we pursue a new approach to employ the proximity effect as a tool to investigate the UVB, based on measuring its signature towards individual QSO sight lines. We demonstrated recently \\citep[][ hereafter Paper~I]{adaglio08} that indeed essentially all quasars show this signature, even at relatively low spectral resolution. Here we follow up on this finding, now based exclusively on high resolution, high S/N quasar spectra. The plan of the paper is as follows.  We begin with a brief description of our spectroscopic data (Sect.~\\ref{txt:data}). In Sect.~\\ref{txt:sims} we explain the Monte Carlo simulations used to assess uncertainties and to interpret the observed proximity effect strength distribution. We then determine the redshift evolution of the effective optical depth in the Ly$\\alpha$ forest in Sect.~\\ref{txt:tau_evol}, followed by a comparison between the two most commonly adopted methods of revealing the proximity effect in Sect.~\\ref{txt:method}. We very briefly report on the analysis of the combined QSO sample in Sect.~\\ref{txt:comb_PE}. In Sect.~\\ref{txt:sglos} we investigate the proximity effect in individual sightlines, showing that the traditional sample-combining method is biased. We then estimate the excess Ly$\\alpha$ absorption near quasars (Sect.~\\ref{txt:data_overd}), and use this to constrain the redshift evolution of the UVB (Sect.~\\ref{txt:uvb}). We present our conclusions in Sect.~\\ref{txt:concl}.  Throughout this paper we assume a flat universe with Hubble constant $H_0=70$~$\\mathrm{km}\\,\\mathrm{s}^{-1}\\,\\mathrm{Mpc}^{-1}$ and density parameters $\\left(\\Omega_\\mathrm{m},\\Omega_\\Lambda\\right)=\\left(0.3,0.7\\right)$. All distances are expressed in physical units. ",
        "conclusions": "\\label{txt:concl}  We have analysed the largest sample of high-resolution quasar spectra presented in a single study to date. We demonstrated that the line-of-sight proximity effect is a universal phenomenon that can be found in the spectrum of essentially every individual high-redshift quasar. While in Paper~I we arrived at the same conclusion using low-resolution spectra, we now can significantly reinforce this claim. In particular, continuum placing uncertainties, which were a bit of an issue in Paper~I, play absolutely no role now, given the high spectral resolution and high S/N data exploited in this paper.  Our study rehabilitates the proximity effect as a tool to investigate the intensity at 1~Ryd (more accurately: the H$^0$ photoionisation rate) of the cosmic ultraviolet background. Previous investigations using the proximity effect nearly always produced UVB estimates that were suspiciously high and indeed at variance with other means to estimate or predict the UVB.  It has been suggested in the literature that quasars typically reside in denser than average large-scale environments, and that therefore excess absorption biases the proximity effect to appear too weak, resulting in an overestimated UVB. We argue that this is not the main reason for the discrepancies between the proximity effect and other methods.  The actually measured strength of the proximity effect in a single quasar line of sight depends on the UVB, but also on the number of absorption lines present in the `proximity effect zone'. That number is generally small, and to first approximation the presence or absence of a line at a given redshift can be described as a Poissonian random process. We have demonstrated that the distribution of the resulting `proximity effect strength' parameter (which directly relates to the UVB) is significantly skewed, even without invoking any physical overdensities. Any direct averaging over a sample of sightlines will inevitably bias the resulting UVB. Only by looking at individual lines of sight separately can this effect be detected and removed.  We proposed and used two different methods of estimating the UVB intensity in the presence of this bias. In the first method, the UVB is taken directly from the mode of the proximity effect strength distribution (PESD). This method is simple and robust against a small number of outliers; on the other hand it is limited to large samples and depends on the condition that the mode is not affected by the outliers. Our second method involves Monte Carlo simulations to calibrate the asymmetry bias of the PESD and then uses straightforward sample averaging, possibly combined with a prior rejection of outliers, and the result is then corrected for the bias. This `hybrid' method is quite stable and works also for smaller samples; it is, however, more complicated to use and depends on a successful removal of outliers.  Our best estimates of the UVB intensity using these two methods give very similar results. The modal estimate is $\\log J_{\\nu_0} = -21.51\\pm 0.15$, and the hybrid estimator gives $\\log J_{\\nu_0} = -21.46^{+0.14}_{-0.21}$, which is statistically indistinguishable. This value is in excellent agreement with other methods to estimate or predict the UVB.  If the UV background is known, the data used for the proximity effect can also be employed to reconstruct the absorption pattern without the radiative influence of the quasar. Doing this for our sample of 40 quasars we found that only 10~\\% of them showed very significant excess absorption (by more than a factor of $\\sim 3$) while most are distributed in a way more or less consistent with Poissonian fluctuations around the cosmic mean absorption expected at these redshifts. Given the widespread notion of quasars residing in massive galaxies, which in turn should sit in massive haloes with significantly clustered environments, this sounds surprising. We have no intention of questioning that picture, and we stress that our measurements constrain exclusively the neutral hydrogen distribution averaged on scales of $\\sim 3$~proper Mpc towards the quasar.  Nevertheless, our results may bear some relevance to the question of what are the typical halo masses for quasars. Most of the QSOs in our sample are optically selected and therefore presumably radio-quiet, which could mean that their haloes are not exceptionally massive. In a recent paper, \\citet{mandelbaum08} estimated halo masses of AGN in the Sloan survey from clustering and galaxy-galaxy lensing, finding that optically selected AGN follow the same relation between stellar and halo mass as normal galaxies, whereas radio-loud AGN have much higher halo masses. (Incidentally, the outstanding object in our sample in terms of overdensity is the radio quasar PKS~0237$-$233; but the next-ranked objects are all optically selected.)  Our argument against significant overdensities of quasars and in favour of the applicability of the proximity effect to measure the UV background gains weight by the excellent agreement between our estimates of the UVB intensity and those obtained by completely different methods. This notion is supported independently by considering the \\emph{shape} of the PESD. The observed distribution width clearly favours average densities close to unity; if significant overdensities were involved, a much broader distribution would be expected than the one observed. This additional constraint breaks the degeneracy between UV background and overdensities that has haunted the field for so long.  Finally, we also attempted to constrain the redshift evolution of the UVB intensity. Although the results are still quite uncertain, we could for the first time from a single quasar sample derive useful limits on the amount of evolution. We can rule out at high significance that quasar light dominates the UVB at redshifts $z\\ga 4$. Our best fit suggests a mild decrease in the UVB intensity towards higher redshifts, and the derived slope is in astonishingly good agreement with the predictions by \\citet{haardt01}. Nevertheless, an approximately constant UVB as found by \\citet{bolton05} and others is also formally consistent with our data.  We have shown that the proximity effect holds the potential to derive important cosmological quantities. While our sample is large in comparison to those of previous studies using high-resolution spectra, it still suffers in several aspects from small number statistics. As the public archives are growing, substantial progress can be expected from applying the proximity effect analysis to further spectra, in particular at higher redshift."
    },
    "0807/0807.3982_arXiv.txt": {
        "abstract": "We investigate the nature of the mid-IR excess for 31 intermediate-mass stars that exhibit an 8 $\\mu$m excess in either the Galactic Legacy Infrared Mid-Plane Survey Extraordinaire or the Mid-Course Space Experiment using high resolution optical spectra to identify stars surrounded by warm circumstellar dust. From these data we determine projected stellar rotational velocities and estimate stellar effective temperatures for the sample.  We estimate stellar ages from these temperatures, parallactic distances, and evolutionary models.  Using MIPS [24] measurements and stellar parameters we determine the nature of the infrared excess for 19 GLIMPSE stars. We find that 15 stars exhibit H$\\alpha$ emission and four exhibit H$\\alpha$ absorption.  Assuming that the mid-IR excesses arise in circumstellar disks, we use the H$\\alpha$ fluxes to model and estimate the relative contributions of dust and free-free emission.  Six stars exhibit H$\\alpha$ fluxes that imply free-free emission can plausibly explain the infrared excess at [24]. These stars are candidate classical Be stars.  Nine stars exhibit H$\\alpha$ emission, but their H$\\alpha$ fluxes are insufficient to explain the infrared excesses at [24], suggesting the presence of a circumstellar dust component. After the removal of the free-free component in these sources, we determine probable disk dust temperatures of T$_{disk}\\simeq$300--800~K and fractional infrared luminosities of $L_{IR}/L_{*}$ $\\simeq$10$^{-3}$. These nine stars may be pre-main-sequence stars with transitional disks undergoing disk clearing.  Three of the four sources showing H$\\alpha$ absorption exhibit circumstellar disk temperatures $\\simeq$300--400~K, $L_{IR}/L_{*}$ $\\simeq$10$^{-3}$, IR colors $K$-[24]$<$ 3.3, and are warm debris disk candidates. One of the four H$\\alpha$ absorption sources has $K$-[24]$>$ 3.3 implying an optically thick outer disk and is a transition disk candidate. ",
        "introduction": "With its launch in 1983, the \\textit{Infrared Astronomical Satellite} ($IRAS$) opened up a new wavelength window and enabled identification of many classes of stellar sources with circumstellar material, providing new insights into star formation and evolution.  In one of the most exciting discoveries enabled by $IRAS$, \\citet{Aumann:1984} reported the first detection of an infrared excess owing to dust in the Vega system, and this excess was later explained as a debris disk. To date, many debris disks have been found with \\textit{IRAS}, the \\textit{Infrared Space Observatory} \\textit{(ISO)}, and the \\textit{Spitzer Space Telescope} (\\textit{Spitzer}) (\\citealt{Backman:1993}; \\citealt{Spangler:2001}; \\citealt{Meyer:2004}). Most known debris disks show evidence of cool dust at wavelengths $>$ 20 $\\mu$m, and the majority of these surround intermediate-mass stars (\\citealt{Backman:1993}; \\citealt{Chen:2006}; \\citealt{Hillenbrand:2008}).  Only two $IRAS$ Point Source Catalog debris disk sources, the intermediate-mass A-type stars $\\beta$ Pictoris and $\\zeta$ Leporis, exhibited a mid-IR ($\\leq$ 12 $\\mu$m) excess detectable by $IRAS$ \\citep{Aumann:1991}.  $Spitzer$ observations of $IRAS$-discovered debris disks indicate that most are devoid of warm dust and small particles (sizes $<$ 10 $\\mu$m) \\citep{Chen:2006}.  The rarity of mid-IR excess sources may imply that dust particles emitting at terrestrial temperatures are uncommon around main-sequence stars.  Advances in infrared technologies coupled with large-aperture ground-based telescopes and $Spitzer's$ mid-IR sensitivity with IRAC enable new searches for warm/hot circumstellar dust that $IRAS$ did not detect.  Even with such improvements in sensitivity over $IRAS$, only five known intermediate-mass stars (HD 3003, HD 38678, HD 75416, HD 105234, and HD 181296) are surrounded by debris disks with temperatures analogous to our asteroid belt ($>$ 200 K) (\\citealt{Chen:2001}; \\citealt{Rhee:2007} and references therein). $Spitzer$ observations have also revealed a number of mid-IR excess objects that are primordial disks undergoing disk clearing processes similar to the solar-mass star TW Hya (\\citealt{Calvet:2005}; \\citealt{Furlan:2007}).  ``Transition disks'' may be evolved protoplanetary disks with inner holes caused by clearing of the initial gas-dominated disks from the inside out \\citep{Strom:1989}. These disks would be optically thin at near- and mid-IR wavelengths and optically thick in the far-IR.  These sources would show weak to no accretion signatures \\citep{Sicilia-Aguilar:2006b}. Intermediate-mass candidate transition disks have only recently been discovered (\\citealt{Hernandez:2006}; \\citealt{Sicilia-Aguilar:2006a}; \\citealt{Hernandez:2007}).  Such transition disks are especially interesting as possible evolutionary links between the optically thick Herbig AeBe disk systems and the older optically thin debris disk systems.  In \\citet{Uzpen:2005,Uzpen:2007} we exploited this new mid-IR capability and used $Spitzer$ to identify stars with 8 $\\mu$m infrared excesses. In \\citet{Uzpen:2007}, we identified 230 stars that exhibited 8 $\\mu$m mid-IR excesses out of 4,015 stars investigated in the \\textit{Spitzer} Galactic Legacy Infrared MidPlane Survey Extraordinaire (GLIMPSE; \\citealt{Benjamin:2003}) and MidCourse Space Experiment (\\textit{MSX}; \\citealt{Egan:2003}) catalogs. Out of the 230 stars with 8 $\\mu$m excesses, 183 are of spectral types O or early B. The excess for these stars is plausibly explained by free-free emission. Out of the 47 remaining stars of spectral types B8 or later, five stars were shown to be false excess sources, and three sources were not yet confirmed at 24 $\\mu$m \\citep{Uzpen:2007}. In systems with mid-IR excesses we modeled the 8 and 24 $\\mu$m excesses as simple blackbodies to characterize the color temperature and fractional infrared luminosity of the circumstellar disk. The results were consistent with massive debris disks and/or transitional disks with disk temperatures ranging between $\\approx$200--800~K and fractional infrared luminosities ($L_{IR}/L_{*}$) $\\approx$10$^{-4}$--10$^{-2}$. However, infrared photometry alone is not sufficient to unambiguously determine the nature of the extraphotospheric emission in these sources.  Because of their scarcity, identifying new transition disks and warm debris disks is an important step toward understanding the evolutionary path of circumstellar disks.  The nature of the mid-IR excesses emanating from circumstellar disks around intermediate-mass stars (2--8\\ \\mo) varies greatly and partially depends on age \\citep{Zuckerman:2001}. Young stars form circumstellar disks through the collapse of the initial molecular cloud as a method to conserve angular momentum \\citep{Shu:1987}. Intermediate-mass pre-main-sequence Herbig AeBe stars accrete gaseous material via their circumstellar disks as they evolve towards main-sequence stars (\\citealt{Herbig:1960}; \\citealt{Waters:1998}).  Circumstellar dust disks, presumably due to planetesimal collisions, surround some main-sequence stars (\\citealt{Backman:1993}; \\citealt{Lagrange:2000}). Circumstellar gaseous disks that emit excesses owing to free-free emission surround other main-sequence stars during late main-sequence or post-main-sequence evolution and are known as classical Be stars (\\citealt{Struve:1931}; \\citealt{Porter:2003}; \\citealt{McSwain:2005}). Circumstellar disks can surround evolved giants as well (e.g., \\citealt{Jura:1999}; \\citealt{Verhoelst:2007}; \\citealt{Deroo:2007}). By understanding the nature of the circumstellar disks, we can better understand how these disks evolve.  The two most common stellar contaminants in surveys for stars with debris disks are Herbig AeBe stars and classical Be stars.  Herbig AeBe stars are young stars accreting circumstellar material via their disks (\\citealt{Strom:1972}; \\citealt{Hillenbrand:1992}; \\citealt{Hartmann:1993}). Emission lines and near-infrared excesses evidence this accretion \\citep{Hillenbrand:1992}. \\citet{Waters:1998} suggest the following as criteria for Herbig AeBe stars: (1) an emission-line A or B star, (2) an infrared excess owing to circumstellar dust, either hot or cool, and (3) luminosity class III--V.  As the primordial disk dissipates the emission lines weaken and the infrared excess decreases, appearing only at longer wavelengths \\citep[e.g.,][]{The:1994}.  Classical Be stars are typically rapidly rotating B stars with H$\\alpha$ emission and an IR excess from free-free emission (\\citealt{Gehrz:1974}; \\citealt{Balona:2000}). These stars are variable in both emission-line strength and luminosity \\citep{Dachs:1988}. The circumstellar disks for classical Be stars are thought to be mechanisms for mass loss, unlike Herbig AeBe circumstellar disks which facilitate accretion \\citep{Balona:2000}. Classical Be stars are similar to Herbig AeBe stars in that they both exhibit emission lines due to circumstellar gas (\\citealt{Struve:1931}; \\citealt{Herbig:1960}). The greatest distinction between Herbig AeBe stars and classical Be stars is the nature of the infrared excess---dust in the former and thermal bremsstrahlung in the latter.  In this paper we present results of optical spectroscopic analysis for 30 mid-IR excess sources identified in the \\textit{Spitzer} GLIMPSE and $MSX$ catalogs to determine the nature of infrared excess \\citep{Uzpen:2007}. By investigating the nature of the infrared excess through optical spectral features, we can distinguish between various types of circumstellar disks that persist around intermediate-mass stars.  Our goal is to remove classical Be and Herbig AeBe stars from our sample to identify transition and warm debris disk candidates. In \\S 2 we describe how we obtained new optical echelle spectroscopy.  We discuss the results of our rotational velocity analysis in \\S 3.  In \\S 4 we present our methods for determining effective temperatures. In \\S 5 we discuss the spectra including the H$\\alpha$, \\ion{Fe}{2} $\\lambda$ 5317 \\AA, Paschen, \\ion{O}{1}, and \\ion{Ca}{2} IR triplet line profiles and any implications on the nature of the disk from the lines.  We describe how we estimate stellar ages through main-sequence fitting utilizing our determined fundamental parameters in \\S 6. In \\S 7 we discuss how all of the derived parameters provide insight into the nature of the infrared excess, how we use these parameters to estimate the contribution of free-free emission associated with our sources, and how we draw conclusions regarding the evolutionary state of the circumstellar disk. In \\S 8 we describe the individual sources and classify the infrared excesses. Finally we use spectral classification schemes for stars with infrared excesses, such as classical Be stars, with our stars to determine if the classifications are consistent in \\S 9. We also apply our classification schemes to well-known stellar samples to determine their classification in \\S 9. ",
        "conclusions": "In an effort to identify the nature of the infrared excesses, we obtained high-resolution optical spectroscopy for 31 luminosity class V or IV stars of spectral type B8 or later exhibiting 8 $\\mu$m mid-IR excesses in the GLIMPSE and $MSX$ catalogs. We measured stellar effective temperatures for 28 stars and projected rotational velocities and H$\\alpha$ equivalent widths for 30 stars. The projected rotational velocities of the 20 GLIMPSE stars lie between 25 and 310 km s$^{-1}$ with a mean and dispersion of 209$\\pm$76 km s$^{-1}$ consistent with the mixed nature of the sample.  We used derived stellar parameters and distances to estimate stellar ages ranging from $\\simeq$1 Myr to 100 Myr for 16 stars (9 of the 20 GLIMPSE stars). Twenty of the 31 stars (16 of the 20 GLIMPSE stars) exhibit H$\\alpha$ emission, consistent with characteristics of pre-main-sequence stars (i.e., Herbig AeBe stars) or classical Be stars.  H$\\alpha$ equivalent widths range between -39.38 \\AA\\ (characteristic of classical Be stars or Herbig AeBe stars) and 9.41 \\AA\\ (main-sequence B and A stars), suggesting a variety of origins in a sample of objects spanning a range of stellar evolutionary states.  For 16 of the 20 GLIMPSE stars showing H$\\alpha$ emission, we use the H$\\alpha$ flux to estimate the expected contribution of (free-free and free-bound) emission to the mid-IR excesses. In six out of the 16 cases, the expected IR excess can fully account for the observed infrared excess (Type I sources). We did not confirm the infrared excess at 24 $\\mu$m in one case. In the remaining nine cases (Type II sources) the expected contribution from free-free emission cannot account for all of the observed IR excess, implying that either 1) the circumstellar excess is time-variable and the non-contemporaneous nature of H$\\alpha$ and mid-IR data render this estimate invalid, or 2) there is a dust component to the IR excess. The lack of evidence for H$\\alpha$ variability for the majority of our sources, combined with the fact that in no instance does the H$\\alpha$ \\textit{over predict} the IR excess component, makes the first scenario unlikely. We conclude that the second scenario, the presence of a circumstellar dust component, is most probable. Combining correlations of spectral features for sources with known characteristics we find that the nine Type II sources with IR excesses require a dust component and are likely to be young main-sequence or near main-sequence objects harboring warm (300--800 K) circumstellar dust. The nominal fractional infrared luminosities for these nine sources after removal of the free-free component is $\\simeq$10$^{-2}$--10$^{-3}$. For the six of these nine sources with age estimates, there is a correlation between estimated age and H$\\alpha$ EW consistent with a progression of disk clearing with time. Such characteristics make these stars candidates for being ``transition disk'' or ``evolved-HAeBe'' systems containing rare dissipating primordial gas and dust disks.  One of the four systems lacking H$\\alpha$ emission has a very large $K$-[24] color, $>$ 4 magnitudes. Such a large color difference implies that the excess is optically thick at [24]. We conclude that this source is also a transition disk. The remaining three systems lacking H$\\alpha$ emission have disk temperatures of 300--400 K and fractional infrared luminosities of 10$^{-3}$.  These characteristics make them candidates for dusty debris disk systems.  We show that $K$-8, 8-24 color-color diagrams are a useful tool to distinguish cool (T$_{d}$ $<$ 300~K) dusty mid-IR excess objects from hot (T$_{d}$ $>$ 500~K) dusty or free-free mid-IR excesses.  Such diagrams reveal that additional information beyond near/mid-IR colors is necessary to distinguish hot dust from free-free excesses. This diagram also reinforces the idea that some stars cataloged as classical Be stars may in fact be late-stage pre-main-sequence stars containing circumstellar dust."
    },
    "0807/0807.4931_arXiv.txt": {
        "abstract": "The spectral--energy (and luminosity) correlations in long Gamma Ray Bursts are being hotly debated to establish, first of all, their reality against possible selection effects. These are best studied in the observer planes, namely the peak energy \\epo\\ vs the fluence $F$  or the peak flux $P$. In a recent paper (Ghirlanda et al. 2008) we started to attack this problem considering all bursts with known resdhift and spectral properties. Here we consider instead all bursts with known \\epo, irrespective of redshift, adding to those a sample of 100 faint \\ba\\ bursts representative of a larger population of 1000 objects. This allows us to construct a complete, fluence limited, sample, tailored to study the selection/instrumental effects we consider. We found that the fainter \\ba\\ bursts have smaller \\epo\\ than those of bright events. As a consequence, the \\epo\\ of these bursts is correlated with the fluence, though with a slope flatter than that defined by bursts with $z$. Selection effects, which are present, are shown not to be responsible for the existence of such a correlation. About 6\\% of these bursts are surely outliers of the \\ama\\ correlation (updated in this paper to include 83 bursts), since they are inconsistent with it for any redshift. \\epo\\ correlates also with the peak flux, with a slope similar to the \\yon\\ correlation. In this case there is only one sure outlier. The scatter of the \\yonf\\ correlation defined by the \\ba\\ bursts of our sample is significantly smaller than the \\amaf\\ correlation of the same bursts, while for the bursts with known redshift the \\ama\\ correlation is tighter than the \\yon\\ one. Once a very large number of bursts with \\epo\\ and redshift will be available, we thus expect that the \\yon\\ correlation will be similar to that currently found, whereas it is very likely that the \\ama\\ correlation will become flatter and with a larger scatter. ",
        "introduction": "The correlation between the peak spectral energy \\ep\\ and the bolometric isotropic energy \\eiso\\ emitted during the prompt (the so called \"Amati\" correlation, Amati et al. 2002) may be a key ingredient for our comprehension of the physics of Gamma Ray Bursts (GRBs).  The proposed interpretations explain the \\ama\\ correlation as either due to geometric effects (Eichler \\& Levinson 2004; Toma, Yamazaky \\& Nakamura 2005) or to the radiative process responsible for the burst prompt emission (Thompson 2006; Thompson, M\\'esz\\'aros \\& Rees 2007), though there is no unanimous consensus.  In addition, there is no agreement about the reality of the correlation itself. Indeed, Nakar \\& Piran (2005) and Band \\& Preece (2005) have pointed out the existence of outliers, while Ghirlanda et al.  (2005a) and Bosnjak et al. (2008), considering an updated Amati relation, found no new outliers besides GRB 980425 and GRB 031203.  More recently it has been argued that this correlation might be the result of selection effects related to the detection of GRBs (Butler et al.  2007 - hereafter B07).  The existence of an \\ama\\ correlation was predicted (Lloyd, Petrosian \\& Mallozzi 2000) before the finding of Amati et al. (2002). The prediction was based on the existence of a significant correlation in the observer frame between the peak energy \\epo\\ and the bolometric fluence $F$ of a sample of \\ba\\ bursts. This finding has been recently confirmed by Sakamoto et al. (2008a) using a sample of bursts detected by \\sw, \\ba\\ and \\he. In particular, they note that X-Ray Flashes and X-Ray-Rich satisfy and extend this correlation to lower fluences. Amati et al. (2002) discovered the \\ama\\ correlation with a sample of 12 bursts detected by \\sax\\ with spectroscopically measured redshifts.  Later updates (e.g. Lamb, Donaghy \\& Graziani 2005; Amati et al.  2006; Ghirlanda et al. 2008 -- hereafter G08) confirmed it with larger samples. In the most recent updates (76 bursts in G08 and 83 in this paper) the GRBs with measured redshift and \\epo\\ define a correlation $E_{\\rm peak}\\propto E_{\\rm iso}^{0.5}$, with only two outliers (GRB 980425 and GRB 031203, but see Ghisellini et al. 2006).  This last sample contains bursts detected by different instruments/satellites, i.e. \\ba, \\sax, \\he, \\ko\\ (all operative in the so--called pre--\\sw\\ era) and, since 2005 (mostly) by \\sw.  These instruments have different detection capabilities and also different operative energy ranges.  For these reasons is crucial to answer to the following question: is the \\ama\\ correlation real or is it an artifact of some selection/instrumental effect?  To investigate this issue we should move to the observer frame plane corresponding to the Amati relation (i.e. the \\epof\\ plane) where the instrumental selection effects act. There, we can place the ``cuts\" corresponding to the instrumental selection effects and see if the distribution of the data points are affected by these cuts, or if, instead, they prefer to lie in specific regions of the plane, away from these cuts.  There are two main selection effects: first, a burst must have a minimum flux to be triggered by a given instrument. This minimum flux can be converted (albeit approximately) to a fluence by adopting an average flux to fluence conversion ratio, as done in G08. This is the minimum fluence a burst should have to be detected.  We call it ``trigger threshold\", TT for short. Secondly, we need a minimum fluence also to find \\epo\\ and the spectral shape. In fact we can have bursts that, although detectable, have too few photons around \\epo\\ to reliably determine \\epo\\ itself.  Consider also that the limited energy range of any detector inhibits the determination of \\epo\\ outside that range.  We call it ``spectral threshold\" -- ST.  While both selection effects are functions of \\epo\\, we found that the latter is dominant for all the detectors.  In G08 we considered a sample of 76 GRBs with redshifts $z$ (which we will call, for simplicity, $z$GRB sample).  These bursts define an Amati correlation in the form $E_{\\rm peak}\\propto E_{\\rm iso}^{0.47\\pm 0.03}$, consistent with what found by previous works.  Moreover, G08 found that these bursts define a strong correlation also in the observer frame (\\epo$\\propto F^{0.39\\pm0.05}$).  We demonstrated that the ST truncation effect is biasing the $z$GRB sample of \\sw--bursts (i.e. bursts for which \\epo\\ has been determined from the fit of BAT spectra) while this is not the case for the no--\\sw\\ $z$GRBs (i.e. bursts for which \\epo\\ has been determined from other instruments, namely \\ko, \\ba, \\sax\\ and \\he) . This leaves open two possibilities: (i) if those described above are all the possible conceivable selection effects, the no--\\sw\\ $z$GRB sample represents an unbiased sample and, therefore, the \\epof\\ correlation defined with these bursts is real; (ii) there are other selection effects biasing the sample of $z$GRBs.  For instance, the optical afterglow luminosity might be proportional to the burst fluence, resulting in a bias in favour of $\\gamma$--ray bright bursts. Another effect concerns the \\ba\\ bursts, that had to be localized by the Wide Field Camera (WFC) of {\\it Beppo}SAX.  Although, formally, the TT for the WFC should not introduce any relevant truncation, in G08 we have shown that all bursts detected by the WFC (with and without redshift or measurable \\epo) had fluences much larger than its TT curve.  To proceed further, in this paper we consider all bursts with measured \\epo, irrespective of having or not a measured redshift.  With this enlarged sample we study if the distribution of GRBs in the \\amaf\\ plane is: (i) consistent or not with the correlation defined by the $z$GRBs; and (ii) if it is strongly biased or not by the considered selection effects.  Besides using existing samples of bursts (from \\he, Sakamoto et al. 2005, from \\sw, B07, and from \\ba, Kaneko et al. 2006, K06), we collected a \\ko\\ sample from the GCN circulars (Golenetskii et al., 2005, 2006, 2007, 2008) and we analysed a new sample of \\ba\\ bursts reaching a fluence of $10^{-6}$ [erg cm$^{-2}$].  This is the \\ba\\  limiting fluence (ST) which allows to derive a reliable \\epo\\ from the spectral analysis.  We therefore have a complete, fluence limited, sample of \\ba\\ bursts.  With this we can study if there is any \\epof\\ correlation and if it is a result of selection effects or not.   By the same token, we can study the correlation between \\ep\\ and the peak luminosity $L_{\\rm iso}$.  This correlation was first found by Yonetoku et al. (2004), with a small number (16) of bursts, and was slightly tighter than the \\ama\\ correlation, and had the same slope. It will be interesting to see if this is still the case considering our sample of 83 $z$GRBs.  We can then investigate the instrumental selection effects acting on this correlation by studying it in the \\yonf\\ plane, where $P$ is the peak flux.  \\begin{figure*} \\vskip -0.5 true cm \\centerline{\\psfig{figure=fluence_epicco_solo_pubbl.ps,width=18cm,height=14cm}} \\vskip -0.5 true cm \\caption{ Distribution in the \\amaf\\ plane of the GRBs with measured redshift (filled symbols) and bursts without measured $z$ published in the literature (open symbols). The bolometric fluence is obtained by integrating the spectrum in the range 1 keV--10 MeV.  The bright \\ba\\ sample (from K06) is shown by open circles (for well constrained \\epo) and up/down arrows (when only an upper/lower limit can be set on \\epo\\ from the spectral analysis of K06). \\he\\ (Sakamoto et al. 2005) and \\sw\\ bursts (B07 -- but see text) without redshifts are shown with open triangles and stars, respectively.  Shaded regions represent the ST curves of minimum fluence, for different instruments (see G08 for more details), down to which it is possible to fit the spectrum and constrain the spectral parameters.} \\label{fl_ep_kaneko} \\end{figure*} ",
        "conclusions": "\\begin{figure} \\hskip -0.5 true cm \\psfig{figure=fettine.ps,width=9cm,height=7.5cm} \\caption{ Graphic illustration of the \\amaf\\ correlation. We consider different fluence--bins and in each of those we count the total number $N_f$ of objects. Then we divide this fluence--bin into \\epo--bins, counting the number of objects in each small area, dividing it by $N_f$. Each small area is then characterized by a number $n_{i,f}$ (between 0 and 1), which corresponds to a different level of grey. In this way the increasing number of bursts for decreasing fluence (the ``Log$N$--Log$S$\" effect) is accounted for, and it does not influence $n_{i,f}$. Note that the data points do not fill the entire accessible region of the plane (to the right of the ST curve shown in Fig. \\ref{solobatse}), but concentrate along a stripe. } \\label{fettine} \\end{figure}    To study the role of possible instrumental selection effects on the Amati relation we have focused our attention on the observational \\amaf\\ plane. Here we can compare the distribution of different samples of GRB (for example, $z$GRBs and GRB with unknown redshift). To this aim we adopt the analysis performed by G08, referring to two different instrumental biases: the trigger threshold (TT, the minimum fluence derived considering the minimum flux required to trigger a burst) and the spectral analysis threshold (ST, the minimum fluence needed to constrain the GRB spectral properties). These curves depends on \\epo\\ and define what part of the observational plane is accessible.   First we updated the sample of bursts with redshift, adding 7 new recent GRBs, for a total of 83 objects. These GRBs define a \\ep$\\propto$\\eiso$^{0.48\\pm 0.03}$ correlation in the rest frame, very similar to that obtained with previous (and smaller) samples. In the observer plane, they define a slightly flatter correlation (\\epo$\\propto F^{0.40\\pm0.05}$). The scatter of these two correlations is the same (see Tab. \\ref{tab1}). As G08 pointed out, the \\ba\\ ST curve is not biasing the distribution of \\ba\\ bursts with redshift in the observer plane, while the \\sw\\ ST could, in the sense that the distribution of \\sw\\ bursts (with redshift) is truncated by the \\sw\\ ST curve. Then why the \\ba\\ bursts (with redshift) are not truncated by their corresponding ST? Is it because of a real, intrinsic correlation or is it due to another, hidden, selection effect? One way to answer this crucial question is to analyze all GRBs with \\epo, even without redshift. The \\ba\\ sample of GRBs is the best suited for this aim because: i) it contains a large number of bursts; ii) large sample of bright GRBs have been already analyzed, and iii) for \\ba\\ we already know the ST curve. Then we pushed the spectral analysis to the limit, deriving the spectral parameters for a representative sample of 100 \\ba\\ GRBs with  a (bolometric) fluence between $10^{-6}$ [erg cm$^{-2}$] (corresponding to the ST limit) and $2\\times 10^{-5}$ [erg cm$^{-2}$] (the limiting fluence of K06). These 100 GRBs represent a large population of 1000 GRBs, in the same fluence range. Combining our and the K06 samples we have a homogeneous and complete sample, best suited to study how \\ba\\ GRBs populate the \\amaf\\ plane. Using this complete, fluence limited, sample we find:   \\begin{itemize} \\item GRBs without redshifts, in this plane, are not spread in the region free from instrumental selection effects, but define a correlation with a flat slope ($\\sim 0.16$) and a scatter larger for smaller fluences (after accounting for the errors increasing for smaller fluences). Fig. \\ref{fettine} is a graphic illustration of this: different grey levels corresponds to different density of points in the \\amaf\\ plane, once we take out the effect of the overall increase in density going to smaller fluences (for the Log$N$--Log$S$ slope). The way we do this is the following: we consider different fluence--bins and in each of those we count the total number $N_f$ of objects. Then we divide this fluence--bin into \\epo--bins, counting the number of objects in each small area, dividing it by $N_f$. Each small area is then characterized by a number $n_{i,f}$ (between 0 and 1), which corresponds to a different level of grey. Note that the data points do not fill the entire accessible region of the plane but concentrate along a stripe.  Note that the shape of this concentration of points is not determined by the ST curve, reported in Fig. \\ref{fettine} for a typical burst lasting 20 s. The only effect of the ST curve to the found correlation is to cut it at the smallest fluences and \\epo.  The very flat slope could be due to the difficulty of having \\ba\\ GRBs with \\epo\\ smaller than $\\sim 50$ keV, whose existence is demonstrated by other instruments.  However, the paucity of the derived upper limits on \\epo\\ suggests that this effect is marginal.  \\item Formally, the scatter is not greater than the scatter of the \\ama\\ correlation (both have $\\sigma=0.23$ once the contribution to the scatter of the measurement errors are taken into account). Despite that, the entire \\ba\\ sample and the $z$GRB population define two \\amaf\\ correlations which have significantly different slopes. If their redshift distribution is similar, then they will define two different correlations also in the \\ama\\ plane: considering then the two samples together, we will define a correlation with intermediate slope and a scatter larger than the individual one.  \\item If the above point holds (i.e. if the redshift distributions of GRBs of unknown redshifts is the same of the $z$GRBs) then we can conclude that there exists an \\ama\\  correlation, not determined by selection effects, even if its slope and scatter may be different from what we know now. We should emphasize that by the term \"correlation\" we mean that GRBs will occupy a ``stripe\" in the \\ama\\ plane with a relatively large scatter (fitting it with $\\chi^2$ method one would obtain a very large reduced $\\chi^2_r$). In other words, it is very likely that there is another (third) variable responsible for the scatter. In fact one finds a tighter correlation considering, as a third variable, the jet break time (Ghirlanda, Ghisellini \\& Lazzati 2004; Laing \\& Zhang 2005) or the time of enhanced prompt emission (Firmani et al. 2005). Another cause for a large scatter is the viewing angle, if a significant number of bursts are seen slightly off--axis.   \\item In the \\ba\\ sample there are a few bursts with small or intermediate fluences but large \\epo, not present in the $z$GRB sample. Among them there are some surely outliers of the \\ama\\ correlation (as defined by the $z$GRB sample), i.e. bursts that lie at more than 3$\\sigma$  from it, for any redshift. The number of these sure outliers is however very small, amounting to the 6 per cent of the entire population.  \\item We have also investigated the \\yon\\ correlation, and its counterpart (\\yonf) in the observer plane. First, we partly confirm the original findings of Yonetoku et al. (2005, see also the update in Ghirlanda et al. 2005b): with the $z$GRB sample we find a strong \\yon\\ correlation, whose slope is flatter than originally found ($s=0.40$ instead of 0.5) and whose scatter is greater than the scatter of the \\ama\\ correlation.  \\item In the observer plane, instead, the \\yonf\\ correlation of our complete sample of \\ba\\ bursts is tighter than the the \\amaf\\ correlation ($\\sigma=0.2$ instead of $\\sigma=0.23$). Its slope is $s=0.24$, flatter than the \\yon\\ correlation ($s=0.39$), but however steeper than the \\amaf\\ slope ($s=0.15$). There is only one sure outlier.  \\item Selection effects are in this case determined by the TT curves. These effects are present, being responsible for the cutting at low peak fluxes, but they do not influence the slope and scatter for peak fluxes larger than the what defined by the TT curves.  \\item Considering the $z$GRB sample we have that the \\ama\\ correlation is tighter than the \\yon\\ one. Considering our complete \\ba\\ sample and moving to the observer plane, we have just the opposite: the \\yonf\\ correlation is tighter than the \\amaf\\ one.  \\item It is then conceivable that the \\yon\\ correlation, once a large number of burst with redshift will be available, will be stronger than the \\ama\\ one.  \\end{itemize}  The general conclusion we can draw from our study is that, although selection effects are present, they do not determine the spectral--energy and spectral--luminosity correlations. These could be characterized by a slope and scatter different from what we have determined now using heterogeneous bursts samples with measured redshift, but we found that \\ep\\ is indeed correlated with the burst energetics or peak luminosity. Therefore it is worth to investigate the physical reason for this relation."
    },
    "0807/0807.1129_arXiv.txt": {
        "abstract": "Galactic planetary nebula (PN) distances are derived, except in a small number of cases, through the calibration of statistical properties of PNe. Such calibrations are limited by the accuracy of individual PN distances which are obtained with several non-homogeneous methods, each carrying its own set of liabilities. In this paper we use the physical properties of the PNe in the Magellanic Clouds, and their accurately known distances, to recalibrate the Shklovsky/Daub distance technique. Our new calibration is very similar (within 1\\%) of the commonly used distance scale by Cahn et al. (1992), although there are important differences. We find that neither distance scale works well for PNe with classic (\"butterfly\") bipolar morphology, and while the radiation bounded PN sequences in both the Galactic and the Magellanic Cloud calibration have similar slopes, the transition from optically thick to optically thin appears to occur at higher surface brightness and smaller size than that adopted by Cahn et al. The dispersion in the determination of the scale factor suggests that PN distances derived by this method are uncertain by at least 30\\%, and that this dispersion cannot be reduced significantly by using better calibrators. We present a catalog of Galactic PN distances using our re-calibration which can be used for future applications, and compare the best individual Galactic PN distances to our new and several other distance scales, both in the literature and newly recalibrated by us, finding that our scale is the most reliable to date. ",
        "introduction": "The uncertainty associated with distances measurements of Galactic planetary nebulae (PNe) is a major obstacle to the advancement of PN research. Only $\\sim$40 Galactic PNe have distances that have been determined individually with reasonable accuracy. Distances to Galactic PNe can be determined individually in various ways, including cluster membership (Chen et al.~2003, CHW03; Alves et al.~2000, ABL00), by measuring the rate of their expansion (e.~g., Liller \\& Liller 1968, LL68; Hajian et al. 1995, HTB95), by the reddening method (e.~g., Gathier et al.~1986, GPP96; Kaler \\& Lutz 1985, KL85), and by measuring their spectroscopic parallax (Ciardullo et al. 1999, C99) or trigonometric parallax  (Harris et al. 2007, Hea07).  For the remaining $>$1800 Galactic PNe (Acker et al.~1992) one has to rely on statistical distance scales, whose calibrations are based on the reliability of the individually known PN distances, and the validity of a general correlation that links the distance-dependent to the distance-independent physical properties of PNe. The Cahn, Kaler, \\& Stanghellini (1992, CKS) distance scale of Galactic PNe is based on an attempt by Daub (1982, D82) to improve Shklovsky's distance scale (Shklovsky 1956a, 1956b) for optically thick nebulae. Shklovsky's distance scale assumes that all PNe have equal (observed) ionized mass. D82 assumed that Shklovsky's constant mass approach was still valid, but only for those PNe that are optically thin to the Lymann continuum radiation emitted by the central stars (density bounded). For the optically thick (radiation bounded) PNe, D82 based the distance scale on a calibration of an ionized mass versus surface brightness relation. CKS improved D82's calibration with the use of a larger number of calibrators (PNe with known individual distance), and calculated the statistical distances to 778 Galactic PNe.    Since its publication, distances from the CKS catalog have been used preferentially and widely in the literature. Other statistical methods that have been commonly used include those by Maciel (1984), Zhang (1995, Z95), van de Steene \\& Zijlstra (1995, vdSZ), Schneider \\& Buckley (1996, SB96) and Bensby \\& Lundstr{\\\"o}m (2001, BL01). All these distance scales rely on a set of Galactic calibrators whose distances are mostly derived from reddening or expansion properties, or from the assumption of Galactic Bulge membership, with all the consequent uncertainties. With the publication over the past decade of critical physical parameters for a large sample of Magellanic Cloud PNe (Shaw et al. 2001, 2006; Stanghellini et al. 2002, 2003), including highly accurate H$\\beta$ fluxes, physical dimensions, morphologies, and extinction constants, we have the opportunity to assess and improve the distance scale for Galactic PNe. In this paper we take advantage of the wealth of Magellanic Cloud PN data to re-calibrate the CKS distance scale, as well as other distance scales for comparison. Homogeneously determined photometric radii from {\\it HST} of a PN sample with low Galactic reddening are the best way to determine any relation that involves apparent diameters. Furthermore, the relatively recent publication of trigonometric and spectroscopic parallax and cluster membership distances to Galactic PNe allow us to test with unprecedented reliability our own and other distance scales.  The construction of any statistical distance scale for PNe is composed of three fundamental steps: the selection of a \\textit{method} that has some physical or empirical basis, the selection of a set of \\textit{calibrator PNe}, for which distances have been determined by some independent means, and an analysis of the \\textit{applicability} of the calibration to a wide variety of PNe. Until now it has been difficult to compare the viability various methods, since their calibrations and applications have varied so widely. In \\S2 we describe in detail the Shklovsky/Daub/CKS distance scale and the superiority of the Magellanic Cloud PNe as calibrators. We also derive our new calibration of this method and assess its inherent uncertainties. In \\S3 we discuss the physical underpinning of the CKS distance method in light of recent advances in modeling the evolution of PNe. In \\S4 we take a closer look at the viability of various methods for determining independent distances to Galactic PNe (i.e., the set that had previously been used as calibrators), and discuss the applicability of the Magellanic Cloud distance scale to Galactic PNe. In \\S5 we recalibrate the most used statistical distance methods with the Magellanic Cloud PNe, and then compare the accuracy of these methods to one another using the best independently determined distances to Galactic PNe. We conclude in \\S6 with our final prescription and recommendation for determining statistical distances to Galactic PNe. ",
        "conclusions": "The wealth of new data available that describe the physical parameters of Magellanic PNe has allowed us to check and re-calibrate the Shklovsky/Daub/CKS statistical distance scale, which is most commonly used in the literature, and provide distances of 645 Galactic PNe following the new distance scale calibration (Table 1). To calculate the SSV distance for other PNe, or for the same PNe but using other parameters than those in CKS, given  $\\theta$ and F, the 5 GHz flux, one can use the following equations:  $$ {\\rm log}~D_{\\rm SSV} = 3.06 + 0.37~ {\\rm log}~ \\theta - 0.68~ {\\rm log}~F, \\tau<2.1 \\eqno(8a), $$ $$ {\\rm log}~D_{\\rm SSV} = 3.79 - 0.6~ {\\rm log}~ \\theta - 0.2~ {\\rm log}~F, \\tau>2.1 \\eqno(8b). $$  If the 5 GHz flux is not available for the given PN, one can use Eq. (6) in CKS to derive the equivalent 5 GHz flux from the H$\\beta$ flux.  In this paper we have used recent data on PNe in the Magellanic Clouds to construct a set of calibrators for which the distances are known to high absolute accuracy ($\\sim10$\\%), and for which the dispersion among the distances is extraordinarily small (a few percent). Furthermore, the great distance of these nebulae allows us to establish a distance scale factor that is insensitive to uncertainties in distances to Galactic PNe that are drawn from a heterogeneous, nearby (few hundred pc) sample; a local sample has generally been necessary given the limited range over which many independent distance methods (notably trigonometric and expansion parallaxes) can provide accurate distances.  In addition, we use consistent and reliable means to determine angular sizes (from photometric radii), and the H$\\beta$ fluxes and extinction constants, derived from \\textit{HST} calibrations, are among the most reliable in the literature. There has never been a better set of calibrators for statistical distance determinations. For comparison, we selected Galactic PNe where the independent distances are the best available (including recently published data), and we evaluated the reliability of various independent distance methods by the degree to which they are consistent with our distance scale.  With this study we show that: (1) the distance scale as calibrated from the Magellanic Cloud PNe is very similar to that derived by CKS; (2) our revised distance scale agrees superbly with the most accurate distances measured for individual Galactic PNe. We also show that other methods of statistical distance determination generally do not yield results that are better than this statistical method; (3) the distance scale does not work for PNe with bipolar morphology, and we believe this is because progenitors of bipolars are often not fully ionized during the course of PN evolution (the log I -- log R$_{\\rm PN}$ relation could be used instead for these PNe); (4) with the Magellanic Cloud calibration we provide a more robust physical basis for why the Shklovsky/Daub distance scale works, despite wide variations in the expected ionized mass; we also show that the recalibration of other distance scales with Magellanic Cloud PNe might not work as well as the recalibrated Shklovsky/Daub distance scale; (5) we point out that the dispersion in the distance scale is an inherent property of the method, and cannot be reduced significantly by using better calibrators; (6) the radiation-bounded sequence for Magellanic Cloud PNe may terminate at higher surface brightness than previously derived. It seems that the new sequence and the radiation bounded to density bounded transition does not depend on metallicity very much, as it is the same for the LMC and the SMC PNe; the best available data show that the Magellanic Cloud calibration of this sequence is entirely consistent with Galactic PNe."
    },
    "0807/0807.3967_arXiv.txt": {
        "abstract": "Results from the ongoing Bulge Radial Velocity Assay ({\\it BRAVA}) are presented.  {\\it BRAVA} uses M red giant stars, selected from the {\\it 2MASS} catalog to lie within a bound of reddening corrected color and luminosity, as targets for the Cerro Tololo Inter-American Observatory 4-m Hydra multi-object spectrograph.  Three years of observations investigate the kinematics of the Galactic bulge major ($-$10$^\\circ$$<${\\it l}$<$$+$10$^\\circ$, {\\it b}$=$$-$4$^\\circ$) and minor ($-$6$^\\circ$$<${\\it b}$<$$+$5$^\\circ$, $-$0.4$^\\circ$$<${\\it l}$<$$-$0$^\\circ$) axes with $\\sim$3300 radial velocities from 32 bulge fields and one disk field.  We construct a longitude-velocity plot for the bulge stars and find that, contrary to previous studies, the bulge does not rotate as a solid body; from $-$4$^\\circ$$<${\\it l}$<$$+$4$^\\circ$ the rotation curve has a slope of roughly 100 km~s$^{-1}$~kpc$^{-1}$ and flattens considerably at greater {\\it l}, reaching a maximum rotation of 75 km~s$^{-1}$.  We compare our rotation curve and velocity dispersion profile to both the self-consistent model of Zhao (1996) and to N-body models; neither fits both our observed rotation curve and velocity dispersion profile.  We place the bulge on the plot of (V$_{max}$/$\\sigma$) vs. $\\epsilon$ and find that the bulge lies near the oblate rotator line, and very close to the parameters of NGC 4565, an edge-on spiral galaxy with a bulge similar to that of the Milky Way.  We find that our summed velocity distribution of bulge stars appears to be sampled from a Gaussian distribution, with $\\sigma=116 \\pm 2$ km~s$^{-1}$ for our summed bulge fields.   Furthermore, the high precision of our radial velocities ($\\sim$5 km~s$^{-1}$) allows us to investigate hints of cold kinematic features that were seen in a number of the line-of-sight velocity distributions from our earlier observations.  In the past, optical radial velocity studies in the bulge have not emphasized high precision, because of the large velocity dispersion and the expectation that the short orbital periods would erase any cold structures in well under a gigayear. Thus our precision is sufficient to enable a search for more cold streams analogous to those associated with the Sagittarius dwarf spheroidal galaxy; some candidate cold features were seen and follow-up observations are reported herein. ",
        "introduction": "The status of the Milky Way bulge as a distinct stellar population had historically been well established by Baade's surveys of RR Lyrae stars, and the well known concentration of red giants toward the Sagittarius region (Blanco 1965).  Furthermore, the image of the Galaxy produced by the {\\sl COBE} satellite (Weiland et al.1984; Dwek et al. 1995;  Arendt et al. 1998) depicted an unambiguous, asymmetric, peanut-shaped bulge.  Blitz \\& Spergel (1991) modeled the $2.4\\mu \\rm m$ balloon borne data of Matsumoto et al. (1982) as a bar which shows the hallmark thickening of the bulge at positive Galactic longitude that is seen in later studies (see e.g. Launhardt et al. 2002).  While the discovery of a bar in the stellar distribution was somewhat of a surprise, gas motions in the inner 2 kpc have long been modeled as due to the potential of tilted bar  (e.g. Liszt \\& Burton 1980).  Optical surveys find roughly 1/3 of spirals are barred (Sellwood \\& Wilkinson 1993); this fraction rises to 60\\% or more in the infrared (Menendez-Delmestre et al. 2007), so it is perhaps not surprising that one should find evidence of bar structure in the Milky Way.  In fact, the evidence for a triaxial (barred) bulge is convincing.  The COBE 2$\\mu$m light distribution is modeled as a bar oriented toward positive Galactic longitude (Dwek et al. 1995, Binney et al. 1997).  Star counts of  red clump stars (Stanek et al 1997; Babusiaux \\& Gilmore 2005) show a barlike structure, while microlensing events further show a central bar that is pointed roughly in the direction of the Sun (Figure \\ref{schematic}).  There is also the issue of bar survival; Sheth et al. (2007) find that the bar fraction declines dramatically with redshift to half the Hubble time, with only the most massive galaxies hosting bars by $z\\sim 0.8$.  This fact is to be contrasted with growing evidence that the bulge is old from a chemical (Rich 1990; McWilliam \\& Rich 1994; Matteucci et al. 1999; Fulbright et al. 2007; Lecureur et al. 2007) and stellar populations (Ortolani et al. 1995; Kuijken \\& Rich 2002; Zoccali et al. 2003) perspective.   There is also continuing debate over whether the Milky Way bulge/bar system is best categorized as a pseudobulge (Kormendy \\& Kennicutt 2004).  As evidence of a bar in the Milky Way has grown (see e.g. Gerhard et al. 2002) it has become clear that the Galactic bulge offers a potential laboratory to investigate the dynamics of the nearest bar in detail with radial velocities and eventually with proper motions as well as abundances.  This recent evidence has prompted renewed interest in the dynamical modeling of the Galactic bar/bulge system in order to determine, in part, the dynamical formation and evolution of galaxies e.g., the unsolved issue of whether a Cold Dark Matter halo cusp is compatible with the microlensing and rotation of the Milky Way bar (Zhao, Spergel, Rich 1995, Klypin, Zhao, Somerville 2002).  However, Zhao's (1996) dynamical model paper notes the lack of comprehensive surveys as a problem in comparing the model against data; until recently there simply has not been enough data to properly constrain the few dynamical models of the bulge available.  Indeed, many efforts have modeled the 2 micron light in projection (cf. Gerhard 2002), but surveys of the dynamics have generally either emphasized very rare evolved stars (e.g. SiO masers; Izumiura et al. 1995; OH/IR or IRAS stars; Menzies 1995) or use K/M giants. Previous studies are restricted to small isolated fields due to the problems of crowding, large and variable extinction, lack of wide field surveys in the red, and the presence of a contaminating disk population.  Minniti (1992) used K giants, but at the time, uniform selection of candidates was an issue, and concerns were raised about disk contamination, especially for the field at ({\\it l,b})=(12$^\\circ$,3$^\\circ$).  Consequently, past surveys have had insufficient statistics to properly probe the dynamics of fields at Galactic latitude $|$ {\\it l} $|$ $<$ 5$^\\circ$ (see Table 3 of Zhao 1996 for detailed summary) in order to provide a constraint for dynamical models.  Partly in response to this problem, Beaulieau et al. (2000) undertook the first dynamical survey of the bulge aimed at isolating a uniform sample of kinematic probes while at the same time covering a wide area.  Their study combines PNe from a new survey and those found in existing catalogs to span a wide range in Galactic longitude and to study both the Northern and Southern portions of the bulge.  They argue that the easily observed PNe have no metallicity bias.  However, PNe are short-lived and are therefore rare, and their population membership (and especially distances) are uncertain.  Rich et al. (2007) noted (and we confirm) that the bulge PNe have lower velocity dispersion than other populations, consistent with a partial population of non-bulge members.  The dynamical model for the bulge/bar has a number of important implications. Large samples of uniform radial velocity data are still of great value in constraining the bar versus axisymmetric models, and the nature of the orbit families supporting the bar. Further, the interpretation of the microlensing events in the bulge depends on the use of an accurate dynamical model (Han \\& Gould 2003). The recent discovery of planetary transit host stars in the bulge (Sahu et al. 2006) gives an additional incentive to improve our knowledge of the bulge/bar model, as it is debated whether stars on non-circular orbits are incapable of bearing planets, if their orbits stray too near the Galactic center.  The relatioinship to the disk, thick disk, and halo populations is of great interest.  Only with a realistic dynamical model can the issue of long term stability be explored.  There is also the issue of whether the bar formed due to the buckling of the disk (see e.g. Raha et al. 1991) or is a yet more ancient population.  The availability of bulge photometry and high precision astrometry from the Two-Micron All-sky Survey (2MASS; Skrutskie et al. 2006) along with the Schlegel et al. (1998) all sky reddening map offered, for the first time, the possibility of easily selecting a large sample of kinematic probes from the very population responsible for the 2$\\mu$m radiation comprising the projected bulge/bar (Dwek et al. 1995; Launhardt et al. 2003; Figs. 1,2).   The {\\it Bulge Radial Velocity Assay} or {\\it BRAVA} (Rich et al. 2007) was conceived as a survey of the line of sight velocity distribution of red giants across the bulge, to be compared with self-consistent dynamical models like that of Zhao (1996).  {\\it We emphasize that while we report the first and second moments of the BRAVA fields in this paper, the ultimate aim is to study the Line of Sight Velocity Distribution (LOSVD) and to compare  it with dynamical models.}   A further additional aim of {\\it BRAVA} is to explore correlations between kinematics and abundances.  Note that there is already a hint from one proper motion field study (Soto, Rich, \\& Kuijken 2007) that predominantly metal rich stars are observed to show a vertex deviation in their velocity ellipsoid and may therefore be predominantly responsible for the support of the bar.  If $\\rm [\\alpha/Fe] $ can be derived in addition to [Fe/H], it may be possible to use the $\\alpha-$enhancement as a secondary criterion to distinguish bulge from thick disk stars, via chemistry (see e.g. Fulbright et al. 2007).  Here we report results from a survey based on red giants, which comprise the bulk of the 2.4$\\mu \\rm m$ light of the bulge.  M giants offer a dynamical probe that is populous, luminous, traces the light, and is easily utilized for velocity measurements.  M giants are also extremely well studied (Frogel \\& Whitford 1987) and have been identified, and their giant branch characterized, over the whole of the bulge (Frogel et al. 1990), from $-3^\\circ$ to $-12^\\circ$ (420 to 1675 pc, adopting $R_0=8$ kpc).   The issue of the abundance \"bias\" raised by Beaulieau et al. is essentially settled based on high dispersion infrared spectra (see Rich \\& Origlia 2005, and Rich, Origlia, \\& Valenti 2007).  M giants follow roughly the abundance distribution of the K giants; given their significant contribution to the integrated light, we are confident that these stars are excellent dynamical probes.   The M giants, being luminous, are also usable as probes even in regions of high extinction, and for exploring the fringes of the bulge, as Frogel et al. (1990) as well as the 2MASS catalog, illustrate.  Early results from our survey are given in Rich et al. (2007a,b).  Here, we report on the first three-year results from {\\sl BRAVA} and also describe the sample selection and measurement uncertainties in detail.   At this time, we present $\\sim$3300 spectra for which we have obtained a radial velocity precision of $\\sim$5 km~s$^{-1}$ for our most recent data.   Future papers in this series will address the kinematics of the bulge at $-8^\\circ$ and a detailed comparison with the Zhao (1996) model (Howard et al. 2008 in prep.).  The long term plan is to make the entire survey, including the spectra, available on a public website.  In the following section we discuss our sample selection.  Section 3 discusses our observations, while Section 4 discusses our data reduction and velocity calibration.  Section 5 considers the fields and the issue of reddening.  Section 6 presents our results, while Section 7 gives our conclusions. ",
        "conclusions": "We present our methods of selecting M giant probes for the {\\sl Bulge Radial Velocity Assay}.  Our sample as reported in this paper has nearly 3300 stars, and is defined clearly to cover the full range of abundance on the first red giant branch.  We report our rotation curve and velocity dispersion profiles, which agree more closely with the self consistent model of Zhao (1996) compared to N-body models in both rotation curve and dispersion.  The reliability of our year-to-year measurements, along with our velocity precision, allow us to search for evidence of cold streams in the bulge resulting from recent merger events, analogous to the Sagittarius Dwarf spheroidal stream.  We find no evidence for hidden dynamically cold, or hot, subcomponents, either in the line of sight velocity distribution of the minor axis fields, or in the l-v plot.   Two candidate cold stream features are investigated in follow-up studies that increase the sample size by a factor of three.  The features were not confirmed, and illustrate the importance of sample size when searching for such features.  We find a departure from solid body rotation at a projected distance of $\\approx$0.6 kpc for R$_{o}$=8 kpc, in agreement with the proper motion study of Clarkson et al. (private communication).  The break from solid body rotation occurs as the density of the bulge is dropping and may signify a transition from bar to keplerian dynamics.  The bulge appears to extend to $l=\\pm 10^\\circ$, beyond which a colder, more rapidly rotating, disk-like population appears to dominate.  However, it is noted that there is a lack of data between 10$^\\circ$$<$$|$ {\\it l} $|$$<$30$^\\circ$, which prevents one from determining the relationship between our bulge sample, and the outer disk fields.  Our selection technique has also allowed us to determine that the the bulge population is normally distributed, with no significant deviation from Gaussian when considering the entire co-added sample of $\\approx3200$ stars.  We have divided our sample by magnitude and color, and find no significant differences in dynamics.   We believe that our selection technique is robust and is yielding  a bulge/bar dominated sample. The rotation curve permits us to place the Milky Way's bulge on the (V$_{max}$/$\\sigma$) vs. $\\epsilon$ plot; it is close to the peanut-shaped bulge of NGC 4565 in this diagram.  Further observations at the {\\it b}$=-8^\\circ$ latitude fields will allow us to determine if the Milky Way bulge rotates cylindrically, an indication of whether the Galactic bulge is a pseudobulge, or a more classically evolved bulge.  RMR and CDH acknowledge support from grant AST-0709479 from the National Science Foundation and from grant GO-11159 from the Space Telescope Science Institute and NASA.   CDH acknowledges support from a CTIO travel grant in support of doctoral thesis observations at CTIO.  The authors are grateful to Sylvie Beaulieau for providing her planetary nebula radial velocities for this work and Rich et al. (2007).  The authors also acknowledge useful comments by Will Clarkson, John Kormendy, and Dante Minniti.  The authors thank the support staff at CTIO.  This research has made use of the NASA/IPAC Extragalactic Database (NED) which is operated by the Jet Propulsion Laboratory, California Institute of Technology, under contract with the National Aeronautics and Space Administration.  \\clearpage"
    },
    "0807/0807.1962_arXiv.txt": {
        "abstract": "Long duration gamma ray bursts (GRBs) and X-ray flashes (XRFs) are produced by highly- relativistic jets ejected in core-collapse supernova explosions. The origin of short hard gamma-ray bursts (SHBs) has not been established. They may be produced by highly relativistic jets ejected in various processes: mergers of compact stellar objects; large-mass accretion episodes onto compact stars in close binaries or onto intermediate-mass black holes in dense stellar regions; phase transition of compact stars. Natural environments of such events are the dense cores of globular clusters, superstar clusters and young supernova remnants. We have used the cannonball model of GRBs to analyze all Swift SHBs with a well-sampled X-ray afterglow. We show that their prompt gamma-ray emission can be explained by inverse Compton scattering (ICS) of the progenitor's glory light, and their extended soft emission component by ICS of high density light or synchrotron radiation (SR) in a high density interstellar medium within the cluster. The mechanism generating the afterglow is synchrotron radiation outside the cluster. No associated supernova could be detected in the low luminosity nearby GRBs 060614 and 060505. We interpret them as SHBs seen relatively far off axis. ",
        "introduction": "Gamma ray bursts (GRBs) have been traditionally classified into short hard bursts (SHBs) and long soft bursts. Their distribution as a function of duration is bimodal with a minimum around two seconds (Mazets et al.~1981; Norris et al.~1984, Kouveliotou et al.~1993). The $\\gamma$-rays of short bursts are typically harder than those of long bursts, thus their acronym, SHBs, but their fluence is much smaller. The short hard bursts consist of a single or a complex of peaks with a total duration typically less than 2 seconds and peak widths which range between a few ms and a fraction of a second. The lag-time between their emission in soft and hard energy band is much smaller than that in long GRBs and X-ray flashes (XRFs). We shall refer to the ensemble of long GRBs and XRFs as LGRBs.   A true breakthrough in the observations of SHBs came with the launch of Swift in November 2004. Since its launch Swift has detected over 30 SHBs with X-ray and UVO afterglows, and triggered follow-up ground-based observations that for most SHBs succeeded in identifying their host galaxy and measuring its redshift. Swift has also discovered that, in more than 25\\% of these cases, perhaps in all, the SHB was followed by an extended soft emission component (ESEC), a spectrally softer component lasting tens of seconds (for recent reviews see Nakar~2007, Kann et al.~2008). The fluence of SHBs is typically much smaller than that of GRBs. The fluence of their ESEC is usually comparable to or smaller than that of the SHB. The ESECs end with a fast-decay phase, followed by an afterglow with a canonical time- dependence, similar to the one observed in a large fraction of LGRBs (Nousek et al.~2006). LGRBs are located near the center of their hosts, which are all late-type galaxies (Covino et al.~2006). Contrariwise, SHBs have an heterogeneous population of hosts, and take place both in elliptical galaxies (e.g., SHBs 050709 and 050724, Gehrels et al.~2005, Berger et al.~2005) and spiral galaxies (e.g., SHB 051221A, Soderberg et al.~2006c). Generally, the SHBs with an ESEC are centrally located, while those with no detectable ESEC are found at very large distances from their host's center (Troja et al.~2008). There is also a clear trend of their afterglow to be fainter for larger offsets (Kann et al.~2008).   Observations indicate that LGRBs are mainly produced by core-collapse supernovae (SNe) of type Ib/c. A GRB-SN association was discussed long ago (Colgate 1968, Dar et al.~1992) but was dismissed by Woosley~(1993) who suggested that only `failed supernovae' produce GRBs.  Following the discovery by Galama et al.~(1998) of SN1998bw in the error box of GRB980425, Wang \\& Wheeler~(1998), Dar \\& Plaga~(1999) and Dar \\& De R\\'ujula~(2000) suggested that perhaps most core collapse SNe produce LGRBs. This was later advocated as an observed fact, based on a comprehensive analysis of optical afterglows within the cannonball (CB) model of LGRBs (Dado, Dar \\& De R\\'ujula~2002,2003 hereafter DDD2002.DDD2003, Dar \\& De R\\'ujula~2004 hereafter DD2004,  and references therein) and on empirical grounds by Zeh, Klose \\& Hartmann~(2004). General acceptance of the GRB-SN association waited until  the spectroscopic discovery of an SN2003dh coincident with GRB030329 (Hjorth et al.~2003, Stanek et al.~2003) and b additional discoveries of spectroscopically proven GRB-SN associations such as GRB030213/SN2003lw (Malesani et al.~2003), GRB021211/SN2002lt (Della Valle et al.~2003), XRF060218/SN2006aj (Campana et al.~2006a, Pian et al.~2006, Mazzali et al.~2006) and XRF080109/SN2008D (Malesani et al.~2008, Modjaz et al.~2008).  In contrast, SHBs do not seem to be associated with any known type of supernova (e.g., Hjorth et al.~2005b). The identity of their progenitors is not established, except that giant flares from soft gamma ray repeaters (SGRs) that look like SHBs, may account for most of the observed SHBs in nearby  (redshift $z\\lsim 0.01$) galaxies (Dar~2005, Hurley et al.~2005).  The extremely bright SHBs 051103 and 070201 observed by Konus-Wind could have been such events in the nearby galaxies M81/M82 and M31 (Ofek et al.~2006, Frederiks et al.~2007). But hyperflares from SGRs cannot account for the bulk of the much more distant ones, $z \\!\\gg\\! 0.01$, unless they are highly collimated or much more energetic but less frequent.   Although the observational data on cosmological SHBs and their afterglows did not pin down their progenitors, they do provide useful clues to their progenitors, production mechanism, and production sites. Their extremely short durations and large equivalent isotropic gamma ray energies, and the absence of self absorption features in their $\\gamma$-ray spectra suggest that they are produced by highly relativistic jets ejected in violent processes involving compact stars (Shaviv \\& Dar 1995a) such as:  \\begin{itemize} \\item{} Neutron stars merger and neutron star - black hole merger in compact binaries (Blinnikov et al.~1984, Paczynski~1986, Goodman, Dar \\& Nussinov~1987, Eichler et al.~1989).  \\item{} Collapse of compact stars (neutron stars, hyper stars, quark stars) to a more com- pact star due to mass accretion, and/or loss of angular momentum and/or cooling by radiation (Dar~1998a,1999, Dar \\& De R\\'ujula~2000, Dar~2006).  \\item{} Phase transition in compact stars, such as neutron-stars, hyper-stars and quark stars (Dar~1999, Dar \\& De R\\'ujula~2000, Dar~2006).  \\item{} Accretion episodes in microblazars (Dar~1998b,1999) and on intermediate mass black holes in dense stellar regions.  \\end{itemize}  The natural environment of these progenitors are super dense stellar regions such as collapsed cores of globular clusters (GCs) and young centrally-condensed super-star cluster (SSC) (Shaviv \\& Dar 1995b) whose stellar densities and total luminosities seem to be correlated with their distance from the center of their host galaxy. If the ESEC is produced by ICS of ambient light in globular clusters with a low density interstellar medium (ISM) or by synchrotron radiation (SR) in super-star clusters with a very large ISM density, it may explain the observations that SHBs with an ESEC are centrally located, while those with no detectable ESEC are found at very large distances from their host's center (Troja et al.~2008).   The spectra and pulse-shapes of SHBs measured by the $\\gamma$-ray telescopes on board satellites, except for being harder and shorter, are very similar to those measured by the same satellites in LGRBs (Nakar~2006, Kann et al.~2008). In particular, the X-ray light curves of SHBs, which were well sampled with the Swift Burst Alert Telescope (BAT) and X-ray Telescope (XRT) show the same canonical behaviour seen in many LGRBs (Nousek et al.~2006, O'Brien et al.~2006): An early rapid temporal decay of the prompt emission during which the spectrum softens at a very fast rate. This rapid decay phase ends within a few hundred seconds when it is taken over by a plateau with a much harder power-law spectrum, typically lasting thousands to tens of thousands of seconds . Within a time of order of one day, the plateau steepens into a power-law decay, which lasts until the X-ray AG becomes too dim to be detected. Often, X-ray flares, not coinciding with any detectable $\\gamma$-ray activity, are superimposed on the X-ray light curve during the fast-decaying phase and even later. These similarities with long GRBs suggest that the initial burst and the afterglow in both SHBs and LGRBs are produced by the same radiation mechanisms: inverse Compton scattering (ICS) of glory light and synchrotron radiation (SR) from decelerating CBs.   The CB model has been successful in predicting/accommodating the observed properties of the prompt radiation of LGRBs (Dar \\& De R\\'ujula~2004 - hereafter DD2004) and their afterglows (Dado, Dar \\& De R\\'ujula 2002, 2003, 2007b, 2008a,b - hereafter DDD2002, DDD2003, DDD2007b, DDD2008a,b,c). Because of the above similarities between LGRBs and SHBs, the CB model may also reproduce correctly the main observed properties of SHBs, after only simple adjustments. Indeed, in this paper we show that using only general properties of the putative SHB progenitors and their natural environments, the CB model with the same two basic radiation mechanisms, ICS of ambient light, which dominates the prompt emission, and synchrotron radiation which dominates the afterglow emission, provides also a simple and successful description of the observed properties of SHBs and their afterglows. Moreover, a few nearby LGRBs without a detectable supernova, such as GRB 060614, which were interpreted as belonging to a new class of LGRBs (e.g., Gal-Yam et al.~2006b, Gehrels et al.~2006, Della Valle et al.~2006) can be interpreted as SHBs seen far off-axis. ",
        "conclusions": "We have demonstrated that the entire observational data on SHBs can be explained by the assumption that SHBs are produced by highly relativistic jets ejected in processes involving compact stars, such as mergers of compact stellar objects, large-mass accretion episodes onto compact stars in close binaries or onto intermediate-mass black holes in dense stellar regions, and phase transition of compact stars. Natural environments of such events are the dense cores of globular clusters or superstar clusters and young supernova remnants. We have demonstrated that the cannonball model of GRBs can reproduce the main observed properties of their prompt emission, extended soft emission component and afterglows. In particular, we have used the CB model to fit the XRT light curve of all Swift SHBs with a well sampled X-ray afterglow. We have shown that their prompt gamma-ray emission is well described by inverse Compton scattering of the progenitor's glory light, their extended soft emission component by ICS of light of the host star cluster or by synchrotron radiation in the high density interstellar medium of superstar cluster, and their afterglow by synchrotron radiation outside the cluster. We have also demonstrated that nearby GRBs of low luminosity short end of the duration distribution of LGRBs and without an associated SN, such as GRBs 060614 and 060505, may be SHBs viewed far off axis. The BATSE observations on board CGRO indicated that nearly $\\!sim\\!$25\\% of all GRBs are SHBs. The evidence from INTEGRAL that the rate of LGRBs is comparable to that of SNe of type Ib/c, which are a significant fraction of all core collapse SNe, implies that also the production rate of SHBs is large and comparable to the birth-rate of neutron stars. Hence neutron star merger, which has a rate much smaller than the birth-rate of neutron stars, cannot be the main source of SHBs. A high rate of SHBs (most of which are beamed away from us), comparable to that of SNe Ib/c, suggests that  phase transition in neutron stars or mass accretion episodes onto stellar and intermediate mass black holes are more probable sources of SHBs.  \\noindent {\\bf Acknowledgment}: We would like to thank A. De R\\'ujula for a long and fruitful collaboration and  for his contributions to this manuscript. A.D. would like to thank the Theory Division of CERN for its hospitality during  this work."
    },
    "0807/0807.1686_arXiv.txt": {
        "abstract": "While the strong anti-correlation between chromospheric activity and age has led to the common use of the Ca II H \\& K emission index ($R^\\prime_{\\rm HK} = L_{HK}/L_{bol}$) as an empirical age estimator for solar type dwarfs, existing activity-age relations produce implausible ages at both high and low activity levels. We have compiled $R^\\prime_{\\rm HK}$ data from the literature for young stellar clusters, richly populating for the first time the young end of the activity-age relation. Combining the cluster activity data with modern cluster age estimates, and analyzing the color-dependence of the chromospheric activity age index, we derive an improved activity-age calibration for F7-K2 dwarfs (0.5 $<$ B--V $<$ 0.9 mag).  We also present a more fundamentally motivated activity-age calibration that relies on conversion of $R^\\prime_{\\rm HK}$ values through the Rossby number to rotation periods, and then makes use of improved gyrochronology relations. We demonstrate that our new activity-age calibration has typical age precision of $\\sim$0.2 dex for normal solar-type dwarfs aged between the Hyades and the Sun ($\\sim$0.6-4.5 Gyr). Inferring ages through activity-rotation-age relations accounts for some color-dependent effects, and systematically improves the age estimates (albeit only slightly).  We demonstrate that coronal activity as measured through the fractional X-ray luminosity ($R_X$= L$_{X}$/L$_{bol}$) has nearly the same age- and rotation-inferring capability as chromospheric activity measured through $R^\\prime_{\\rm HK}$. As a first application of our calibrations, we provide new activity-derived age estimates for the nearest 100 solar-type field dwarfs ($d$ $<$ 15 pc). ",
        "introduction": "Age is, arguably, the most difficult basic stellar quantity to estimate for low-mass field dwarfs \\citep[see e.g.][]{Mamajek07}. Yet, the temporal evolution of phenomena such as stellar activity, surface abundances, rotation, and circumstellar matter is of current interest and within observational means for nearby stars.  Our particular motivation for improving field star age estimates stems from our interest in circumstellar disk evolution as executed via the {\\it Spitzer Space Telescope} Formation and Evolution of Planetary Systems (FEPS)\\footnote{http://feps.as.arizona.edu} Legacy Science program which is surveying the dust surrounding solar-type stars between $\\sim$3 Myr and $\\sim$3 Gyr \\citep{Meyer04, Kim05, Stauffer05, Hines06, Meyer06, Silverstone06, Hines07, Moro-Martin07, Bouwman08, Meyer08, Hillenbrand08, Carpenter09}.  The most theoretically grounded stellar age estimator is the Hertzsprung-Russell diagram, which predicts ages based on our understanding of nuclear physics, stellar interior structure, and stellar atmospheres.  It can be employed in stellar clusters for which main sequence turn-off and/or turn-on ages are typically available and to field stars of known distance that are in the pre-MS or post-MS phases of stellar evolution.  Field stars, however, are generally main sequence objects and, by definition, lack co-eval accompanying stellar populations that might enable accurate age dating via standard H-R diagram techniques. Thus proxy indicators of age are necessary.  \\subsection{Chromospheric Activity as an Age Indicator}  Historically, a popular age estimator for field stars of roughly solar mass has been the \\rphk\\, index which measures chromospheric emission in the cores of the broad photospheric Ca II H \\& K absorption lines, normalized to the underlying photospheric spectrum.  Chromospheric activity is generated through the stellar magnetic dynamo, the strength of which appears to scale with rotation velocity \\citep{Kraft67, Noyes84, Montesinos01}.  Both chromospheric emission and rotation are observationally constrained to decay with age \\citep{Wilson63, Skumanich72, Soderblom83, Soderblom91}. The angular momentum loss is theoretically understood as due to mass loss in a magnetized wind \\citep{Schatzman62, Weber67, Mestel68}.  The chromospheric activity index \\rphk\\, is calculated from a band-ratio measurement of the Ca H \\& K emission line strength \\citep[the ``S-index'' or, when converted to the Mount Wilson system, \\smw;][]{Vaughan78,Vaughan80,Duncan91} from which the underlying stellar photospheric contribution is then subtracted.  We refer the reader to papers by \\citet[][and references therein]{Noyes84, Baliunas96, Baliunas95, Henry96, Wright04} for in-depth discussion of how to measure \\smw\\ and \\rphk, as well as the history of studies using this index.  Our simple goal for this study is to provide an \\rphk\\, vs. age relation applicable to sets of \\rphk\\, and \\bvo\\ data (the latter derived from a spectral type or from a color) for solar-type and near-solar metallicity dwarfs.  The activity-age data pair of highest quality is that for the Sun, and our adopted values are listed in Table \\ref{tab:sun}. The solar age is presumed coincident with that of the oldest portions of meteorites \\citep[the Ca-Al-rich inclusions; 4.570 Gyr;][]{Baker05}. However, the Sun and presumably most other stars exhibit activity cycles (with period 11 years in the case of the Sun) as well as longer term variations (e.g. the so-called Maunder minimum in the case of the Sun).  Over the period 1966-1993, covering mostly solar cycles 20, 21, and 22, \\citet{Baliunas95} estimated the solar Mt. Wilson S-index to be $\\overline{S_{\\odot}}$ = 0.179.  Over the period 1994-2006, mostly solar cycle 23, \\citet{Hall07} measured $\\overline{S_{\\odot}}$ = 0.170.  Using a mean solar $S$-value which is approximately weighted by the span of measurements ($\\overline{S_{\\odot}}$ = 0.176; for $\\sim$1966-2006) and a mean solar color of \\bv\\, = 0.65 \\citep{Cox00}, and using the equations from \\citet{Noyes84}, we estimate the mean solar activity to be \\logrphk\\, = --4.91.  We also give in Table \\ref{tab:sun} the 68\\% and 95\\% range of the observed solar \\logrphk\\, due to variability.  \\begin{deluxetable}{lcccl} \\setlength{\\tabcolsep}{0.03in} \\tablewidth{0pt} \\tablecaption{Adopted Solar Data \\label{tab:sun}} \\tablehead{ {(1)}        &{(2)}   & {(3)}   & {(4)}  \\\\ {Parameter}  &{Value} & {Units} & {Ref.}} \\startdata \\bvo\\,                                    & 0.65   & mag & 1 \\\\ Age                                       & 4.570  & Gyr & 2 \\\\ $\\overline{{\\rm S}_{\\odot}}$              & 0.176  & ... & 3 \\\\ $\\overline{\\log {\\rm R}^\\prime_{\\rm HK}}$ & -4.906 & dex & 4 \\\\ \\logrphk\\, 68\\%\\, range      & -4.942 to -4.865 & dex & 5 \\\\ \\logrphk\\, 95\\%\\, range      & -4.955 to -4.832 & dex & 5 \\\\ \\loglx                       & 27.35            & erg\\,s$^{-1}$ & 6\\\\ \\logrx (=\\loglxlbol)         & -6.24            & dex & 6 \\\\ \\enddata \\tablecomments{ References: (1) \\citet{Cox00}, (2) minimum age from \\citet{Baker05}, (3) time-weighted average of \\citet{Baliunas96} and \\citet{Hall07} for 1966-2006, (4) calculated using \\bvo\\, and mean S$_{\\odot}$ via \\citet{Noyes84}, (5) calculating using solar 1$\\AA$ K-index data from \\citet{Livingston07}, using relations from \\citet{Radick98} and \\citet{Noyes84}, and adopting the solar \\bvo\\, color listed, (6) soft X-ray (0.1-2.4 keV) luminosity and fraction luminosity estimated from \\citet{Judge03}, with 50\\% uncertainty. An uncertainty in the solar \\bvo\\, of $\\pm$0.01 mag produces a systematic uncertainty of the \\logrphk\\, values by $\\mp$0.004 dex. Note that the absolute calibration of the \\logrphk\\, values (as a physical metric of chromospheric line losses) are probably only accurate to $\\sim$10\\% \\citep{Hartmann84,Noyes84}.} \\end{deluxetable}  \\subsection{Shortcomings of Previous Activity-Age Calibrations}  Using the Sun as one anchor point, we can look to open clusters with ages derived from other methods (e.g. the H-R diagram) in order to populate an activity-age calibration.  There are four such \\rphk\\, vs. age relations in the literature which have been used in age-dating field stars: two from \\citet{Soderblom91}, and one each from \\citet{Donahue93}, and \\citet{Lachaume99}.  The activity-age relations from \\citet{Soderblom91} include a linear fit to age vs. activity for members of clusters and binaries.  The second relation, often overlooked, assumes a constant star-formation history and takes into account kinematic disk heating. D. Soderblom (priv. comm.) has kindly provided an analytic version of this alternative activity-age relation.  That there are deficiencies with these existing calibrations can be easily demonstrated.  For the \\citet{Lachaume99} calibration, the solar \\rphk\\, value adopted here (-4.91) would imply a solar age of 7.2 Gyr, which is clearly in error. The other two calibrations used the Sun as one of their anchor points, but with slightly different \\rphk\\, values (for the calibrations of \\citet{Soderblom91} and \\citet{Donahue93}, one derives ages of 4.1 and 4.0 Gyr, respectively). \\citet{Soderblom91} do not advocate extrapolating either of their activity-age relations to the young/active regime (\\logrphk\\, $>$ -4.4), however \\citet{Donahue93} explicitly fit their activity-age relation to age $\\sim$10 Myr and \\logrphk\\, $\\simeq$ -4.2 (anchoring their fit to data for NGC 2264).  Given the observed activity levels in the $\\sim$5-15 Myr Sco-Cen OB association (\\logrphk\\, $\\simeq$ --4.05; \\S2), neither the fit from \\citet{Donahue93} or extrapolating the two fits from \\citet{Soderblom91} estimates an age similar to the isochronal value. Indeed, the commonly used fit of \\citet{Donahue93} would estimate an age of {\\it one minute} for a star with \\logrphk\\, $\\simeq$ --4.05.  Given the paucity of young stars in the previous calibrations, we should not be too surprised at the lack of agreement with other age-dating methods at the high-activity end of the relation.  \\subsection{Potential for Improved Activity-Age Calibrations}  Clearly, an improved activity-age calibration is needed. Further, we would like to understand and quantify the limitations of any such relationship and hence its practical application.  We focus this paper primarily on {\\it refining the age-activity relation for solar-type dwarfs}.  By ``solar-type'', we mean $\\sim$F7-K2 or 0.5 $<$ \\bvo\\ $<$ 0.9 mag, which is approximately the color range over which the \\citet{Noyes84} relation for the photospheric contribution to the S-index is applicable, as well as the color range blanketed by recent activity surveys. The F3V-F6V temperature region (0.42 $<$ \\bvo\\, $<$ 0.5) appears to mark the transition where the rotation-activity correlation breaks down, chromospheric activity diminishes, stellar convective envelopes thin, and magnetic breaking becomes inefficient \\citep{Kraft67,Wolff85,GarciaLopez93}. By ``dwarfs'', we mean MS and pre-MS stars, and explicitly exclude evolved stars more than one magnitude above the MS.  There are three developments that make our investigation timely:  (1) Recently measured \\rphk\\, values for stars that belong to age-dated young stellar aggregates (e.g. Sco-Cen, $\\beta$ Pic, etc.). These additions to the literature both broaden and strengthen modern activity-age derivations relative to the data landscape of 1-2 decades ago.  (2) The ages of well-studied nearby open clusters (e.g. $\\alpha$~Per, Pleiades) have been updated during the past decade. The most noticeable difference relative to traditionally accepted age values is the systematic shift towards older ages driven by results using the Li-depletion boundary age estimation method \\citep[e.g.][]{Stauffer98,Barrado04}.  (3) Interest in circumstellar disk and planetary system evolution has increased dramatically over the past five years.  The availability of relevant infrared data, e.g. from {\\it Spitzer} observations, begs for a robust stellar age estimator in order to probe the collisional and radiative evolution of debris disks, and the connection of such phenomena to exo-solar planetary system dynamics. Similarly, exoplanet discoveries over the past decade have motivated interest in the ages of the parent field stars for comparison to the Sun and solar system. In this paper we derive using samples drawn from cluster and moving group populations (\\S2) a new \\rphk\\ activity vs. age relation (\\S3). In \\S4, we tie both chromospheric activity index (\\rphk) and coronal activity index (\\rx\\, = \\loglxlbol) data to stellar rotation rates via the Rossby number (i.e. secure an activity-rotation relation), and attempt to derive independently an activity-age relation based on the ``gyrochronology'' rotation evolution formalism of \\citet{Barnes07}, though with newly derived coefficients.  In an Appendix, we quantify the correlation between fractional X-ray luminosity and Ca H\\&K activity for solar-type stars, and demonstrate that \\rx, like \\rphk, can be used to derive quantitative age estimates. ",
        "conclusions": "The primary goal of this study was derive a well-calibrated conversion between activity and age for stars younger than the Sun. To achieve this, we compiled from the literature \\rphk, \\rx, and rotation period data for members of stellar associations and clusters; in particular, we have populated for the first time the young end of the chromospheric activity-age relation.  We also used updated/modern ages for many young associations and clusters.  We then fit the following relations critical to assessing stellar ages of solar-type dwarfs: a chromospheric activity-age relation, a chromospheric activity-rotation relation, a coronal activity-rotation relation, and a rotation-age ``gyrochronology'' relation. Our main results drawn from study of the rotation and activity observed among binary stars and star cluster members with 0.5 $<$ \\bv\\, $<$ 0.9 can be summarized as:  $\\bullet$ We provide an improved \\logrphk\\, vs. age relation for solar-type stars which constrains especially the young, high-activity end relative to the relations of \\citet{Soderblom91, Donahue93, Lachaume99}. The activity-age relation for solar-color stars appears to be absolutely calibrated to the modern cluster age scale to $\\sim$$\\pm$0.07 dex in log($\\tau$/yr) for stars older than the Pleiades, and perhaps to only $\\sim$$\\pm$0.23 dex accuracy in log($\\tau$/yr) for stars younger than the Pleiades.  For young stars recently arriving on the MS (e.g. the Pleiades), \\logrphk\\, is not very useful as a quantitative age estimator as the inferred r.m.s. spread in ages derived from chromospheric activity is an order of magnitude. For older samples ($>$0.5 Gyr) and typical \\logrphk\\, measurements, it appears that our calibration can estimate the ages of solar-type dwarfs to $\\sim$$\\pm$0.25 dex ($\\sim$60\\%; 1$\\sigma$) accuracy, accounting for systematic errors in the calibration, random errors due to astrophysical scatter, variability of \\logrphk, and measurement errors. This activity-age relation, however, does not account for color-dependent evolution effects which appear to be present.  $\\bullet$ We corroborate previous studies which find a tight relation between chromospheric activity and rotation for stars with -5.0 $<$ \\logrphk\\, $<$ -4.35, as well as coronal X-ray activity and rotation for stars with -7 $<$ \\logrx\\, $<$ -4 (both via the Rossby number). In their respective saturated regimes (\\logrphk\\, $>$ -4.35, \\logrx\\, $>$ -4), the correlation between chromospheric and coronal activity is poor. For stars with long-term \\logrphk\\, averages and well-determined periods, we find that rotation period can predict mean \\logrphk\\, to $\\pm$0.05 dex (1$\\sigma$) accuracy. For stars with multi-decadal average \\logrphk\\, measurements (e.g. Mt. Wilson HK sample), \\logrphk\\, can be used to predict Rossby number (period divided by convective turnover time) to $\\pm$0.1 (1$\\sigma$) accuracy.  For shorter baseline \\logrphk\\, measurements this uncertainty in Rossby number is larger, with the limit of a single \\logrphk\\, measurement probably capable of predicting the Rossby number to $\\sim$0.2-0.3 1$\\sigma$ accuracy.  Similarly, fractional X-ray luminosity \\rx\\, for non-saturated X-ray emitters can be used to infer Rossby number to $\\sim$0.25 1$\\sigma$ accuracy.  $\\bullet$ We provide an improved gyrochronology relation (period as a function of color and age), which fits the young cluster data better than the coefficients provided by \\citet{Barnes07}.  For so-called I-sequence rotators, the new fit is statistically accurate to $\\pm$1.2 days in rotation between the age of the Pleiades and Sun. Our revised gyro relation predicts self-consistent ages with statistical accuracy of $\\pm$0.06 dex (14\\%; 1$\\sigma$) for solar-type stars with well-determined periods.  $\\bullet$ Combining our activity-rotation relation (via the Rossby number; \\S\\ref{Rossby}) and our improved gyrochronology relations (rotation-color-age; \\S\\ref{Gyro}), we predict the evolution of activity as a function of color for solar-type dwarf stars.  Our activity-rotation-age calibration appears to yield slightly better ages than using an activity-age relation alone.  Statistical analysis of binary samples suggest that the activity-rotation-age technique can estimates ages of $\\sim$$\\pm$0.1 dex accuracy, whereas analysis of the cluster samples suggests an accuracy of more like $\\sim$$\\pm$0.2 dex."
    },
    "0807/0807.0680_arXiv.txt": {
        "abstract": "Tides raised on a planet by its host star's gravity can reduce a planet's orbital semi-major axis and eccentricity. This effect is only relevant for planets orbiting very close to their host stars.  The habitable zones of low-mass stars are also close-in and tides can alter the orbits of planets in these locations.  We calculate the tidal evolution of hypothetical terrestrial planets around low-mass stars and show that tides can evolve planets past the inner edge of the habitable zone, sometimes in less than 1 billion years. This migration requires large eccentricities ($>0.5$) and low-mass stars ($\\lsim 0.35 M_\\odot$). Such migration may have important implications for the evolution of the atmosphere, internal heating and the Gaia hypothesis. Similarly, a planet detected interior to the habitable zone could have been habitable in the past. We consider the past habitability of the recently-discovered, $\\sim 5$ $M_\\oplus$ planet, Gliese 581 c.  We find that it could have been habitable for reasonable choices of orbital and physical properties as recently as 2 Gyr ago. However, when we include constraints derived from the additional companions, we see that most parameter choices that predict past habitability require the two inner planets of the system to have crossed their mutual 3:1 mean motion resonance.  As this crossing would likely have resulted in resonance capture, which is not observed, we conclude that Gl 581 c was probably never habitable. ",
        "introduction": "After the discovery of extra-solar planets, the detection of a habitable, terrestrial (rocky) planet became a real possibility. Of particular interest are planets in orbit about low-mass stars ($\\lsim 0.5 M_\\odot$). The lower stellar mass makes detection of lower mass planets easier and leads to habitable orbits with smaller semi-major axes (which also increase the likelihood of detection by current technologies). The recent detection of putative 5 and 8 $M_\\oplus$ planets around an M star (stellar mass $M_* = 0.31 M_\\odot$) (Udry \\etal 2007) near the circumstellar ``habitable zone'' (HZ) (Kasting \\etal 1993; Franck \\etal 2000; von Bloh \\etal 2007; Selsis \\etal 2007; hereafter S07) suggests that Doppler surveys (\\ie Butler \\etal 2006; Naef \\etal 2007) are becoming sensitive enough to detect terrestrial planets in the HZ of low-mass stars. Not surprisingly, M stars now appear to be attractive targets for detecting habitable planets (Tarter \\etal 2007).  The \\textit{habitable zone} is defined, following Kasting \\etal (1993), as the range of distances from a star at which liquid water could be stable at the surface of a planet. The definition usually implicitly assumes that the star is the main source of heat. Thus, it excludes locations where other sources of heat might be available, such as moons, like Jupiter's Europa, where tidal heating is a major factor maintaining near-surface liquid water. In that sense, the phrase \\textit{habitable zone}, as usually used in astrophysics, is a misnomer. Nevertheless, here we use the conventional definition, while explicitly emphasizing that life could exist outside this zone.  The detection of Gl 581 c and d (Udry \\etal 2007) has demonstrated that planets can form near the HZ of low-mass stars. For such stars the habitable zone is so close to the star that tidal effects may be important (Mardling \\& Lin 2004). Tides raised between the star and planet introduce torques on the planet resulting in orbital evolution, usually decay of semi-major axis, $a$, and eccentricity, $e$. Here we quantify the effects of tidal forces in such zones. We find that Earth-sized planets with large initial eccentricities could migrate out of the HZ of low-mass stars in less than 1 billion years.  The proper definition of an HZ is not obvious. As described in S07, the boundaries of the HZ are a complex function of both stellar and planetary properties. For any individual planet, the definitions we choose may not be appropriate. Nonetheless, we will show a general trend that must be considered when assessing a planet's habitability, both in the past and in the future.  We consider the evolution of terrestrial planets beginning from a location that is almost assuredly habitable to a region that is probably not habitable. We call the time to complete this traversal the ``habitable lifetime'', $\\tau_{HL}$. For relatively massive stars ($\\gsim 0.35 M_\\odot$), the HZ is far enough away from the star that tides cannot cause significant tidal evolution\\footnote{Note we only consider tidal effects after the star has reached the main sequence}.  However for planets with relatively large eccentricities ($\\gsim 0.5$) orbiting low-mass stars ($\\lsim 0.35 M_\\odot$) tides can pull a planet out of the HZ in less than 1 Gyr. The large number of low-mass stars and significant fraction of known exoplanets with large $e$ suggest that tides may limit the total number of habitable planets in the galaxy.  We assume that potentially habitable planets can form with a range of masses and orbits in the HZ of low-mass stars (see Chambers [2004]; Raymond, Quinn \\& Lunine [2007] or Lissauer [2007] for a discussion of the formation of such planets). Of course, the formation of a planet in the HZ does not mean a planet is, will be, or has been habitable. Many other factors are required such as favorable composition (Lissauer 2007), a thick atmosphere, substantial mass for sustained plate tectonics ($m \\gsim 0.3 M_\\oplus$) (Williams, Kasting \\& Wade 1997; Raymond, Scalo \\& Meadows 2007), which on Earth drives the carbonate-silicate cycle (Walker, Hays \\& Kasting 1981), and substantial water content (\\eg Raymond, Quinn \\& Lunine 2004), both for life and to facilitate plate tectonics (Regenauer-Lieb, Yuen \\& Barlund 2001). Nonetheless we will call planets ``habitable'' if their orbital properties are such that they permit water to be stable on the surface.  The recently announced Gl 581 system contains a planet, c, that is probably just interior to the HZ (S07). In Table 1 we list the properties of the planets in this system; $m_p$ is the planet mass, $P$ is the orbital period, $a$ is the semi-major axis, $e$ is the eccentricity, $\\varpi$ is the longitude of periastron, and $T_{peri}$ is the time of periastron passage in Julian Days. Tidal theory predicts planet c orbited with larger values of semi-major axis and eccentricity in the past. Therefore it may have been habitable in the past, but tides subsequently moved the planet into an uninhabitable orbit. We find that if planet c was the only planet in the system, plausible physical properties predict it was habitable. However, when we consider constraints from the mutual interactions of the additional planets, we find that planet c has likely never been habitable.  Tides strong enough to induce significant orbital evolution can also heat the interior of a planet through tidal working. The same type of heating causes Io's intense vulcanism (Peale, Cassen \\& Reynolds 1979). Therefore we might expect an increased heating rate on the planets we consider here as well. The orbital evolution we consider could also have implications for the evolution of a habitable planet's atmosphere. In some cases, the orbit-averaged flux can increase by a factor of 2 as the orbit shrinks (assuming constant stellar luminosity).  In $\\S$ 2 we review tidal theory, present our definition of the HZ, and describe how we determine the habitable lifetime. In $\\S$ 3 we present our results for hypothetical terrestrial planets. In $\\S$ 4 we describe how our model applies to the Gl 581 c planet. Finally, in $\\S$ 5 we draw our conclusions and discuss directions for future research. ",
        "conclusions": ""
    },
    "0807/0807.2528_arXiv.txt": {
        "abstract": "s{ We explore the possibility that the present stage of accelerated expansion of the universe is due to the presence of a cosmic vector field. We show that vector theories allow for the generation of an accelerated phase without the introduction of potential terms or unnatural scales in the Lagrangian. We propose a particular model with the same number of parameters as $\\Lambda$CDM and excellent fits to SNIa data. The model is scaling during radiation era, with natural initial conditions, thus avoiding the cosmic coincidence problem. Upcoming observations will be able to clearly discriminate it from standard $\\Lambda$CDM cosmology.}  The fact that today the dark energy density is comparable to the matter energy density poses one of the most important problems in order to find viable models of dark energy. Indeed, to achieve this, most of the models, not only the cosmological constant, but also those based on scalar fields such as quintessence or k-essence, and modified gravity theories such as $f(R)$, DGP, etc, require the  introduction of unnatural scales either in their Lagrangians or in their initial conditions. This is the so called {\\it cosmic coincidence problem} (see \\cite{review} and references therein).  Therefore, we would like to find a model without  dimensional scales (apart from Newton's constant $G$), with the same number of parameters as $\\Lambda$CDM, with natural initial conditions and with good fits to observations. In addition, the model should be stable under small perturbations. We will show that vector models can do the job \\cite{BM}.   Let us start by writing the action of our vector-tensor theory of gravity containing only two fields and two derivatives and without potential terms (see  \\cite{potential} for previous works on vector models for dark energy with potential terms): \\begin{eqnarray} S=\\int d^4x \\sqrt{-g}\\left(-\\frac{R}{16\\pi G} -\\frac{1}{2}\\nabla_\\mu A_\\nu\\nabla^\\mu A^\\nu +\\frac{1}{2} R_{\\mu\\nu}A^\\mu A^\\nu\\right) \\label{action} \\end{eqnarray} Notice that the theory contains no free parameters, the only dimensional scale being the Newton's constant. The numerical factor in front of the vector kinetic terms can be fixed by the field normalization. Also notice that $R_{\\mu\\nu}A^\\mu A^\\nu$ can be written as a combination of derivative terms as $\\nabla_\\mu A^\\mu\\nabla_\\nu A^\\nu-\\nabla_\\mu A^\\nu\\nabla_\\nu A^\\mu$.   The classical equations of motion derived from the action in (\\ref{action}) are the Einstein's and vector field equations: \\begin{eqnarray} R_{\\mu\\nu}-\\frac{1}{2}R g_{\\mu\\nu}&=&8\\pi G (T_{\\mu\\nu}+T_{\\mu\\nu}^A) \\label{eqE}\\\\ \\Box A_\\mu + R_{\\mu\\nu}A^\\nu&=&0 \\label{eqA} \\end{eqnarray} where $T_{\\mu\\nu}$ is the conserved energy-momentum tensor for matter and radiation and $T_{\\mu\\nu}^A$ is the energy-momentum tensor coming from the vector field. For the simplest isotropic and homogeneous flat cosmologies,  we assume that the spatial components of the vector field vanish, so that $A_\\mu=(A_0(t),0,0,0)$ and that the space-time geometry will be given by: \\begin{equation} ds^2=dt^2-a^2(t)\\delta_{ij}dx^idx^j, \\label{metric} \\end{equation} For this metric (\\ref{eqA}) reads: \\begin{eqnarray} \\ddot{A}_0+3H\\dot{A}_0-3\\left[2H^2+\\dot{H}\\right]A_0=0\\label{fieldeq0} \\end{eqnarray} Assuming that the universe has gone through  radiation and matter phases in which the contribution from dark energy was negligible, we can easily solve this equation in those periods.  In that case, the above equation has a growing and a decaying solution: \\begin{eqnarray} A_0(t)=A_0^+t^{\\alpha_+}+A_0^-t^{\\alpha_-}\\label{fieldsol} \\end{eqnarray} with $A_{0}^\\pm$ constants of integration and $\\alpha_{\\pm}=-(1\\pm 1)/4$ in the radiation era, and $\\alpha_{\\pm}=(-3\\pm\\sqrt{33})/6$ in the matter era. On the other hand, the $(00)$ component of Einstein's equations reads: \\begin{eqnarray} H^2=\\frac{8\\pi G}{3} \\left[\\sum_{\\alpha=M,R} \\rho_\\alpha+\\rho_{A}\\right] \\label{Friedmann}\\end{eqnarray} where the vector energy density is given by: \\begin{eqnarray} \\rho_{A}=\\frac{3}{2}H^2A_0^2+3HA_0\\dot A_0-\\frac{1}{2}\\dot A_0^2 \\end{eqnarray} \\begin{figure} \\vspace{0.3cm} \\begin{center}{\\epsfxsize=7.5 cm\\epsfbox{densitie.eps}}\\hspace{1.0cm} {\\epsfxsize=7 cm\\epsfbox{eosg.eps}} \\caption{\\footnotesize (Left) Evolution of energy densities for the best fit model. Dashed (red) for radiation, dotted (green) for matter and solid (blue) for vector dark energy. We show also for comparison the cosmological constant density in dashed-dotted line. (Right) Evolution of dark energy equation of state for the best fit model. The lower panel shows the 1$\\sigma$ confidence interval.} \\end{center} \\end{figure} Using the growing mode solution from (\\ref{fieldsol}), we obtain $\\rho_{A}= \\rho_{A0} a^\\kappa$ with $\\kappa=-4$ in the radiation era and $\\kappa=(\\sqrt{33}-9)/2 \\simeq -1.63$ in the matter era. Thus, the energy density of the vector field starts scaling as radiation at early times, so that $\\rho_A/\\rho_R=$ const. However, when the universe enters its matter era, $\\rho_A$ starts growing relative to $\\rho_M$ eventually overcoming it at some point, in which the dark energy vector field would become the dominant component (see Fig. 1).  Notice that since $A_0$ is essentially constant during radiation era, solutions do not depend on the precise initial time at which we specify it.  Thus, once the present value of the Hubble parameter $H_0$ and the constant $A_0$ during radiation (which fixes the total matter density $\\Omega_M$) are specified, the model is completely determined, i.e. this model contains the same number of parameters as $\\Lambda$CDM, which is the minimum number of parameters of a cosmological model with dark energy. As seen from Fig.1  the evolution of the universe ends at a finite time $t_{end}$ with a singularity in which $a\\rightarrow a_{end}$ with $a_{end}$ finite, $A_0(t_{end})=M_P/(4\\sqrt{\\pi})$, $\\rho_{DE}\\rightarrow \\infty$ and $p_{DE}\\rightarrow -\\infty$.  We can also calculate the effective equation of state for dark energy as: \\begin{eqnarray} w_{DE}=\\frac{p_A}{\\rho_A}=\\frac{-3\\left(\\frac{5}{2}H^2+\\frac{4}{3}\\dot{H}\\right)A_0^2+HA_0\\dot{A}_0 -\\frac{3}{2}\\dot{A}_0^2}{\\frac{3}{2}H^2A_0^2+3HA_0\\dot A_0-\\frac{1}{2}\\dot A_0^2} \\end{eqnarray} Again, using the approximate solutions in (\\ref{fieldsol}), we obtain: $w_{DE}=1/3$ in the radiation era and $w_{DE}\\simeq -0.457$ in the matter era. As shown in Fig. 1, the equation of state can cross the so called phantom divide, so that we can have $w_{DE}(z=0)<-1$.  In order to confront the  predictions of the model with observations of high-redshift supernovae type Ia, we have carried out  a $\\chi^2$ statistical analysis for two  supernovae datasets, namely, the Gold set \\cite{Gold}, containing 157 points with $z < 1.7$, and the more recent SNLS  data set \\cite{SNLS}, comprising 115 supernovae but with lower redshifts ($z < 1$). In Table 1 we show the results for the best fit together with its corresponding $1\\sigma$  intervals for the two data sets. We also show for comparison the results for a standard $\\Lambda$CDM model. We see that the vector model (VCDM) fits the data considerably better than  $\\Lambda$CDM (in more than $2\\sigma$)  in the Gold set, whereas the situation is reversed in the SNLS set. This is just a reflection of the well-known $2\\sigma$ tension \\cite{tension} between the two data sets.  Compared with $\\Lambda$CDM, we see that VCDM favors a younger universe (in $H_0^{-1}$ units) with larger matter density. In addition, the deceleration-aceleration transition takes place at a lower redshift in the VCDM case. The present value of the equation of state  with $w_0=-3.53^{+0.46}_{-0.57}$ which clearly excludes the cosmological constant value $-1$. Future surveys \\cite{Trotta} are expected to be able to measure $w_0$ at the few percent level and therefore could discriminate between the two models. \\begin{table} \\begin{center} \\footnotesize{\\begin{tabular}{|c|c|c|c|c|} \\hline & & & &\\\\ & VCDM  & $\\Lambda$CDM  & VCDM & $\\Lambda$CDM \\\\ & \\footnotesize{Gold}  &\\footnotesize{Gold} & \\footnotesize{SNLS}   & \\footnotesize{SNLS} \\\\ & & & &\\\\ \\hline & & & & \\\\$\\Omega_M$ &$0.388^{+0.023}_{-0.024}$ & $0.309^{+0.039}_{-0.037}$ & $0.388^{+0.022}_{-0.020}$ & $0.263^{+0.038}_{-0.036}$ \\\\& & && \\\\ \\hline & & & &\\\\ $w_0$ &$-3.53^{+0.46}_{-0.57}$& $-1$ & $-3.53^{+0.44}_{-0.48}$ & $-1$ \\\\& & & & \\\\ \\hline  & & & &\\\\$A_0$ & $3.71^{+0.022}_{-0.026}$ &--- & $3.71^{+0.020}_{-0.024}$&--- \\\\\\footnotesize{$(10^{-4}\\;M_P)$}& & & & \\\\ \\hline  & & & &\\\\$z_T$ &$0.265^{+0.011}_{-0.012}$ &$0.648^{+0.101}_{-0.095}$ & $0.265^{+0.010}_{-0.012}$&$0.776^{+0.120}_{-0.108}$ \\\\& & & & \\\\ \\hline  & & & &\\\\$t_0$   &$0.926^{+0.026}_{-0.023}$ & $0.956^{+0.035}_{-0.032}$& $0.926^{+0.022}_{-0.022}$ &$1.000^{+0.041}_{-0.037}$ \\\\\\footnotesize{$(H_0^{-1})$}& &  & & \\\\ \\hline & & & &\\\\$t_{end}$ &$0.976^{+0.018}_{-0.014}$ & --- &$0.976^{+0.015}_{-0.013}$ & --- \\\\\\footnotesize{$(H_0^{-1})$}& & & &\\\\ \\hline & & & &\\\\ $\\chi^2_{min}$ & 172.9& 177.1 &115.8 & 111.0 \\\\ & & & &\\\\ \\hline \\end{tabular}} \\caption{\\footnotesize{Best fit parameters with $1\\sigma$ intervals for the vector model (VCDM) and the cosmological constant model ($\\Lambda$CDM) for the Gold (157 SNe) and SNLS (115 SNe) data sets. $w_0$ denotes the present equation of state of dark energy. $A_0$ is the constant value of the vector field component during radiation. $z_T$ is the deceleration-aceleration transition redshift. $t_0$ is the age of the universe in units of the present Hubble time. $t_{end}$ is the duration of the universe in the same units.}} \\end{center} \\end{table} We have also compared with other parametrizations for the dark energy equation of state \\cite{models}. Since our one-parameter fit has a reduced chi-squared: $\\chi^2/d.o.f=1.108$,  VCDM provides the best fit to date for the Gold data set.  We see that unlike the cosmological constant case, throughout radiation era $\\rho_{DE}/\\rho_R\\sim 10^{-6}$ in our case. Moreover the scale of the vector field $A_0=3.71 \\times 10^{-4}$ $M_P$ in that era is relatively close to the Planck scale and could arise naturally in the early universe without the need of introducing extremely small parameters.   In order to study the model stability we have considered the evolution of metric and vector field perturbations. Thus, we obtain the dispersion relation and  the propagation speed of  scalar, vector and tensor modes. For all of them we obtain $v=(1-16\\pi G A_0^2)^{-1/2}$ which is real throughout the universe evolution, since the value $A_0^2=(16\\pi G)^{-1}$ exactly corresponds  to that at the  final singularity. Therefore the model does not exhibit exponential instabilities. As shown in \\cite{Mukhanov}, the fact that the propagation speed is faster than $c$ does not necessarily implies inconsistencies with causality. We have also considered the evolution of scalar perturbations in the vector field generated by scalar metric perturbations during matter and radiation eras, and found that the energy density contrast $\\delta \\rho_A/\\rho_A$ is constant on super-Hubble scales, whereas it oscillates with growing amplitude as $a^2$ in the radiation  era and as $\\sim a^{0.3}$ in the matter era for sub-Hubble scales. Therefore  again, we do not find exponentially growing modes.   If we are interested in extending the applicability range of the model down to solar system scales then we should study the corresponding post-Newtonian parameters (PPN). We can see that for the model in (\\ref{action}), the static PPN parameters agree with those of General Relativity \\cite{Will}, i.e. $\\gamma=\\beta=1$. For the parameters associated to preferred frame effects we get: $\\alpha_1=0$ and $\\alpha_2=8\\pi  A_\\odot^2/M_P^2$ where $A_\\odot^2$ is the norm of the vector field at the solar system scale. Current limits $\\alpha_2 \\;\\lsim 10^{-4}$ (or $\\alpha_2\\;\\lsim 10^{-7}$ for static vector fields during solar system formation) then impose a bound $A_\\odot^2\\lsim \\, 10^{-5}(10^{-8}) \\,M_P^2$. In order to determine whether such bounds conflict with the model predictions or not,  we should know the predicted value of the field at solar system scales, which in principle does not need to agree with the cosmological value. Indeed, $A_\\odot^2$  will be determined by the mechanism that generated this field in the early universe characterized by its primordial spectrum of perturbations, and the subsequent evolution in the formation of the galaxy and solar system. Another potential difficulty arising generically in vector-tensor models is the presence of negative energy modes for perturbations on sub-Hubble scales. They are known to lead to instabilities at the quantum level,  but not necessarily at the classical level as we have shown previously. In any case, the model proposed is not intended as a quantum theory of the gravitational interaction, which would be beyond the scope of this work.  In conclusion,  vector theories offer an accurate phenomenological description of dark energy in which fine tuning problems could be easily avoided.  {\\em Acknowledgments:} This work has been  supported by DGICYT (Spain) project numbers FPA 2004-02602 and FPA 2005-02327, UCM-Santander PR34/07-15875 and by CAM/UCM 910309. J.B. aknowledges support from MEC grant BES-2006-12059. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0807/0807.0231_arXiv.txt": {
        "abstract": "We present a Gaussianity analysis of the WMAP 5-year Cosmic Microwave Background (CMB) temperature anisotropy data maps. We use several third order estimators based on the spherical Mexican hat wavelet.  We impose constraints on the local non-linear coupling parameter $f_{nl}$ using well motivated non-Gaussian simulations. We analyse the WMAP maps at resolution of 6.9 arcmin for the Q, V, and W frequency bands. We use the $KQ$75 mask recommended by the WMAP team which masks out 28\\% of the sky. The wavelet coefficients are evaluated at 10 different scales from 6.9 to 150 arcmin. With these coefficients we compute the third order estimators which are used to perform a $\\chi^2$ analysis. The $\\chi^2$ statistic is used to test the Gaussianity of the WMAP data as well as to constrain the $f_{nl}$ parameter. Our results indicate that the WMAP data are compatible with the Gaussian simulations, and the $f_{nl}$ parameter is constrained to $-8 < f_{nl} < +111$ at 95\\% CL for the combined V+W map. This value has been corrected for the presence of undetected point sources, which add a positive contribution of $\\Delta f_{nl}=3 \\pm 5$ in the V+W map. Our results are very similar to those obtained by \\citet{komatsu2008} using the bispectrum. ",
        "introduction": "A valuable source of information of the early universe is the cosmic microwave background (CMB) radiation. The temperature fluctuations of this radiation can be used to put tight constraints on the cosmological parameters. The Big Bang theory and the inflationary models are our best theories to describe the universe. In particular, the standard, single field, slow roll inflation \\citep{guth,albrecht,linde1982,linde1983} is one of the most accepted models because of the accuracy of its predictions. In particular, it forecasts that the primordial density fluctuations are compatible with a nearly Gaussian random field. However, different primordial processes (as for example topological defects, etc.) can introduce non-Gaussian features at certain levels that may be detected \\citep{cruzc}. Also, several non-standard models of inflation predict a detectable level of non-Gaussianity in the primordial gravitational potential \\citep[see, e.g.][and refs. therein]{bartolo}. Other non-Gaussian deviations can be explained by the presence of foreground contamination or systematic errors. In any case, the search of non-Gaussian deviations in the CMB has become a question of considerable interest, as it can be used to discriminate different possible scenarios of the early universe and also to study the secondary sources of non-Gaussianity.  Many tests of Gaussianity have been performed on different CMB data sets. For a recent review of this topic see for example \\citet{martinez2008}. Some of the tests have found that the data are compatible with Gaussianity but there are some important detections of non-Gaussian deviations \\citep[see for example][]{eriksen2004,eriksen2005,copi2004,copi2006,vielva2004,cruza,wiaux2006,vielva2007,monteserin2008}, all of them using the {\\it Wilkinson Microwave Anisotropy Probe}\\footnote{http://map.gsfc.nasa.gov/} (WMAP) data.  High precision experiments as the WMAP are sensitive to deviations due to second-order effects in perturbation theory, usually parametrised through the local non-linear coupling parameter $f_{nl}$. Several methods have been used to constrain this parameter using the data of different experiments. We can mention the angular bispectrum on WMAP data \\citep{komatsu2003,creminelli2006,spergel,komatsu2008} and COBE \\citep{komatsu2002}; the Minkowski functionals on WMAP \\citep{komatsu2003,spergel,gott2007,hikage2008,komatsu2008}, BOOMERanG \\citep{troia} and Archeops \\citep{curto2007,curto2008}; different kind of wavelets on WMAP \\citep{mukherjee2004,cabella2005} and COBE \\citep{cayon2003} among others. The previous works suggest that $f_{nl}$ is compatible with zero at least at 95\\% confidence level. However other recent works \\citep{smoot2007,wandelt2008} suggest a positive detection of $f_{nl}$ at a confidence level greater than 95\\% using WMAP data.  In this paper we perform a wavelet-based analysis of the 5-year WMAP data in order to constrain the $f_{nl}$ parameter. We use the high resolution WMAP data maps (6.9 arcmin) and realistic non-Gaussian simulations performed following the algorithms developed by \\citet{liguori2003,liguori2007}.  Our article is organised as follows. Section \\ref{methodology} presents the statistical method, the estimators that we use to test Gaussianity and constrain $f_{nl}$, as well as the data maps and simulations to be analysed. In Section \\ref{results} we summarise the results of this work. The conclusions are drawn in Section \\ref{conclusions}. ",
        "conclusions": "\\label{conclusions} We have analysed the 5-year WMAP data using Gaussian and realistic non-Gaussian simulations through a wavelet-based test. We have considered different combined maps for the analysis: Q+V+W, V+W, Q, V, and W. We have considered two kinds of extended masks to analyse the V+W map (defined in Subsect. \\ref{datasimu}). The third order moments (Eq. \\ref{witowip22}) of the wavelet coefficients of these maps are compatible with Gaussian simulations (see Fig \\ref{estimators_qvw_vw}).  We have performed a $\\chi^2$ analysis and found that the data are indeed compatible with Gaussian simulations (see Fig. \\ref{chi_qvw_vw}). We performed another $\\chi^2$ analysis to constrain the non-linear coupling parameter $f_{nl}$ by using non-Gaussian simulations with $f_{nl}$. The best-fit $f_{nl}$ values of the five analysed maps are compatible with the ones obtained by \\citet{komatsu2008} using the bispectrum, showing similar confidence intervals and also a similar trend with the frequency.  Finally we have estimated the contribution to $f_{nl}$ from unresolved point sources for the V+W map using a simple model that has point sources with constant intensity ($F_{src}=$ 0.5 Jy) and a realistic model given by \\citet{zotti}. We have found that they add a positive contribution of $\\Delta f_{nl} = 11 \\pm 4$ for the simple model and $ \\Delta f_{nl} = 17 \\pm 5$ for the realistic model. These values are larger than the one obtained by \\citet{komatsu2008} for the bispectrum, which can be explained by the enhancement of the point sources produced by wavelets. Using the less restrictive extended masks, the point source distribution add a positive contribution of $\\Delta f_{nl} = 3 \\pm 4$ for the simple model and $ \\Delta f_{nl} = 3 \\pm 5$ for the realistic model. The smaller values are explained because the less restrictive extended masks add some pixels that affect the efficiency of the wavelet to detect point sources.  Taking into account the point source correction, our best estimate of $f_{nl}$ is $-8 < f_{nl} < +111$ at 95\\% CL.  It is important to emphasise the agreement found between the two estimators (bispectrum and wavelets), since they are formed by very different combinations of the data and are affected by systematic effects (like the mask, noise, beam response and foreground residuals) in very different ways."
    },
    "0807/0807.2975_arXiv.txt": {
        "abstract": "We use a semi-analytic model of Ly$\\alpha$ emitters (LAEs) to constrain the reionization history. By considering two physically motivated scenarios in which reionization ends either early (ERM, $z_i\\approx 7$) or late (LRM, $z_i \\approx 6$), we fix the global value of the IGM neutral fraction (e.g. $\\chi_{HI}=3\\times 10^{-4}$, $0.15$ at $z=6.56$ for the ERM and LRM, respectively) leaving only the star formation efficiency and the effective escape fraction of \\Lya photons as free parameters. The ERM fits the observed LAE luminosity function (LF) at $z=5.7$ and $6.56$ requiring no redshift evolution or mass dependence of the star formation efficiency, and LAE star formation rates (SFR) of $3 < \\dot M_\\star/{\\rm M}_\\odot {\\rm yr}^{-1} < 103$, contributing $\\approx 8$\\% of the cosmic SFR density at $z=5.7$. The LRM requires a physically uncomfortable drop of $\\approx 4.5$ times in the SFR of the emitters from $z=6.5$ to $5.7$. Thus, the data seem to imply that the Universe was already highly ionized at $z=6.56$. The mass-dependent \\Lya transmissivity is $0.36 \\simlt T_\\alpha \\simlt 0.51$ (ERM) and $T_\\alpha \\simlt 0.26$ (LRM) at $z=6.56$. The LF data at $z=4.5$ imply an extra \\Lya line damping factor of $\\approx 0.25$ possibly due to dust; the presence of a (clumpy) dust component with $E(B-V) \\simlt 0.28$ is also required to reproduce the observed large \\Lya equivalent widths at the same redshift. Additional useful information can be extracted from the line profile (weighted) skewness, found to be $S_W=10-17$~\\AA~for the two reionization models, which shows an interesting $L_\\alpha-\\chi_{HI}$ anti-correlation, holding under the model assumptions. The shortcomings of the model and strategies to overcome them are discussed. ",
        "introduction": "The Epoch of Reionization (EoR) marks the second major change in the ionization state of the universe after recombination and is directly linked to structure formation.  Reionization begins when the first structures form within dark matter halos and emit neutral hydrogen ionizing photons. In addition to changing the ionization state, these first structures also affect subsequent structure formation due to various radiative, mechanical and chemical feedback effects. Thus, to probe reionization, one needs an excellent understanding of initial density perturbations and their growth, as well as simulations that can trace the evolution of structure formation.  One of the major challenges of reionization models is to be able to simultaneously account for the considerable, and often apparently conflicting, amount of data accumulated by experiments exploiting QSO absorption line spectra (Fan et al. 2006), cosmic microwave background radiation (Page et al. 2007, Spergel et al. 2007) and high redshift galaxy surveys (Bouwens et al. 2006, Stark et al. 2007).  The emerging picture (Choudhury \\& Ferrara 2007) is one in which hydrogen reionization is an extended process starting at $z\\approx 15$  and being 90\\% complete by $z=8$. Reionization is initially driven by metal-free stars in low mass ($M < 10^8 M_\\odot$) halos; the conditions for the formation of these objects are soon erased by the combined action of chemical and radiative feedbacks at $z < 10$.  Given the many assumptions necessarily made by reionization models, the above scenario needs constant confrontation with freshly acquired data sets. In this sense, it has been suggested (Malhotra \\& Rhoads 2004, 2005; Santos 2004; Haiman \\& Cen 2005; Mesinger, Haiman \\& Cen 2004; Dijkstra, Wyithe \\& Haiman 2007; Dijkstra, Lidz \\& Wyithe 2007; Mesinger \\& Furlanetto 2007)  that a class of high redshift galaxies, the Lyman Alpha Emitters (LAEs) can be suitably used to put additional constraints on the reionization history:  the Lyman break and the strength, width and asymmetry of the observed Ly$\\alpha$ line make the detection of LAEs unambiguous. The strength of the method is based on the sensitivity of Ly$\\alpha$ photons to even tiny amounts of \\HI in the intergalactic medium (IGM). At redshifts $z\\sim5$, the optical depth to Ly$\\alpha$ photons is very large. Let $e$ be the electron charge, $f$ the oscillator strength ($0.4162$), $\\lambda_\\alpha$ the wavelength of Ly$\\alpha$ in its rest frame (1216 \\AA), $m_e$ the electron mass, $c$ the speed of light, $H(z)$ the Hubble parameter at the required redshift, $n_{HI}$ the global neutral hydrogen density and $n_H$ the global mean hydrogen density at that redshift. Note that $n_{HI} =\\chi_{HI} n _{H}$ where $\\chi_{HI}$ is the fraction of neutral hydrogen at the redshift under consideration. Further, $\\Omega_b$ represents the baryonic density parameter and $\\Omega_m$ is the total (baryonic + dark) matter density parameter of the universe, $\\Omega_m = \\Omega_b+\\Omega_{dm}$. Then, \\begin{displaymath} \\tau_\\alpha = \\frac{\\pi e^2 f \\lambda_\\alpha}{m_e c H(z)} n_H \\frac{n_{HI}} {n_H}, \\end{displaymath} where \\begin{displaymath} \\frac{\\pi e^2 f \\lambda_\\alpha}{m_e c H(z)} n_H = 1.76 \\times 10^5 h^{-1} \\Omega_m^{-1/2} \\frac{\\Omega_b h^2}{0.022} \\bigg(\\frac{1+z}{8}\\bigg)^{3/2}. \\end{displaymath} Hence, even a \\HI fraction of $10^{-4}$ can lead to a significant attenuation of the Ly$\\alpha$ line. The observed (i.e. transmitted) Ly$\\alpha$ luminosity, $L_\\alpha$, can then be used to infer the ionization state of the IGM at redshifts close to those of the emitter and hence to reconstruct, at least piecewise, the cosmic reionization history.  This simple picture is complicated by a number of important physical effects. First of all, Ly$\\alpha$ photons from the stars have to propagate through and escape from the interstellar medium of the LAE. During their travel they are multiply scattered by \\HI atoms (thus being either removed from or added to the line of sight [LOS]) and possibly absorbed by dust grains (Neufeld 1991; Tasitsiomi 2005; Hansen \\& Oh 2006; Finkelstein et al. 2007). These processes modify both the emerging Ly$\\alpha$ luminosity and the shape and equivalent width of the line.  Second, the ionizing radiation from the same stars builds regions of ionized IGM around the emitters, whose size depends on the star formation rate, age, escape of ionizing photons from the galaxy and the stellar Initial Mass Function (IMF; the case of very massive stars has been explored, for example, by Dijkstra \\& Wyithe 2007). As a result, the flux redwards of the Ly$\\alpha$ line can escape, attenuated only by the red damping wing of the Gunn-Peterson absorption (Miralda-Escud\\'e 1998; Madau \\& Rees 2000).  To a first approximation, the spatial scale imposed by the Gunn-Peterson damping wing on the size of the \\HII region corresponds to a redshift separation of $\\Delta z\\approx 0.01$, i.e. about 200 kpc (physical) at $z=10$. The effects of the damping wing fade away if the emitter is powerful enough to create a large enough \\HII region and/or if the universe is already reionized when the emitter turns on.  Alternatively, one would observe the damping wing if there were even a small fraction of neutral hydrogen left inside the sphere and/or if a \\HI cloud is present along the LOS to the source.  All the above effects combine to shape the observed LAE Luminosity Function (LF), which has been now measured (Rhoads et al. 2000; Taniguchi et al. 2005; Shimasaku et al. 2006; Iye et al. 2006; Kashikawa et al. 2006; Murayama et al. 2007; Ota et al. 2007; Dawson et al. 2007) with different degrees of accuracy up to $z\\approx 7$. Such tremendous progress has been made possible by the  increase of survey fields and available samples. The current observational situation can be summarized as follows. All studies seem to converge toward the conclusion that there is very little indication of evolution of the LF moving from $z=3$ to $z=5.7$. Beyond that epoch there seems to be evidence of a decline in the LF, with $L_*$ at $z=6.6$ being about 50\\% of that at $z=5.7$. Such a high luminosity steepening of the LF can be produced by a number of different physical effects. A rapid evolution of the IGM ionization state can be invoked (Kashikawa et al. 2006) if the overlapping phase of reionization ended around $z=6$; however, the net effect of reionization on the observed Ly$\\alpha$ luminosity of the most luminous (and presumably massive) LAEs is unclear. If these objects are expected to live in more dense and hence more neutral environments, they are also more heavily clustered (McQuinn et al. 2007). The two effects might not change appreciably the size of their \\HII regions. Alternatively, the observed evolution could be simply a result of the evolution of the mass function of dark matter halos housing the LAEs (Dijkstra, Wyithe \\& Haiman 2007; Dijkstra, Lidz \\& Wyithe 2007). Finally, extinction due to dust, which is expected to be more prominent in actively star forming galaxies, may act as a sink for Ly$\\alpha$ photons in the most luminous LAEs.  As of now, it is difficult to firmly assess which of these explanations is more robust. Fortunately, other aspects of the data, such as the line shape and equivalent width, might allow one to make progresses. Here we try to assess to what extent the reionization history can affect the shape of the LF and the observed properties of individual LAEs. Our approach is similar in spirit to some of those mentioned above, but it has the strength of being based on reionization models that simultaneously account for all the available data beyond LAEs including Ly$\\alpha$/Ly$\\beta$ Gunn-Peterson opacity, electron scattering optical depth, Lyman Limit Systems, cosmic SFR history and the number density of high-redshift sources.\\footnote{Throughout the paper, we use the best-fit cosmological parameters from the 3-year WMAP data (Spergel et al. 2007), i.e., a flat universe with ($\\Omega_m$, $\\Omega_{\\Lambda}$, $\\Omega_b h^2 $, $h$)=(0.24, 0.76, 0.022, 0.72). The parameters defining the linear dark matter power spectrum are $\\sigma_8=0.82$, $n_s=0.95$, $dn_s/d\\ln k =0$. We use a value of $\\sigma_8$ much higher that quoted from WMAP3 (0.76) as the combination of WMAP3 and SDSS data give $\\sigma_8 \\sim 0.78$ (0.86) for low (high) resolution Ly$\\alpha$ forest data (Viel et al. 2006). Mpc is comoving unless otherwise specified. }. ",
        "conclusions": "\\begin{figure}  \\center{\\includegraphics[scale=0.45]{weighted_skew_2.eps}} \\caption{Weighted skewness of the observed Ly$\\alpha$ line for different models. The asterisks are the data from \\citet{b10}. The dashed (solid) line correspond to the best fit ERM (LRM) at $z=6.56$.} \\label{weighted_skewness} \\end{figure}   \\begin{figure*} \\center{\\includegraphics[scale=1.0]{contour_all.eps}} \\caption{Dependence of $S_W$ (values are color-coded by the bar on the right) on $\\chi_{HI}$ and $L_\\alpha$ at $z=6.56$ for a set of LAEs with SFR in the range predicted by the two reionization models, i.e. $\\dot M_\\star=2.7-197 M_\\odot$yr$^{-1}$.} \\label{countour} \\end{figure*}  Additional constraints on the model can come from the information embedded in the observed line profiles, as for example the line profile asymmetry. This can be suitably quantified by the {\\it weighted skewness} parameter, $S_W$, introduced by \\citet{b10}, which we calculate for the best fit parameter values for the ERM mentioned above. We adopt the following definition for such a quantity: \\begin{equation} S_W = S \\Delta \\lambda = S (\\lambda_{10,r}-\\lambda_{10,b}), \\label{weighted_skew} \\end{equation} where $\\lambda_{10,r}$ ($\\lambda_{10,b}$) is the wavelength redward (blueward) of the Ly$\\alpha$ line where the flux falls to 10\\% of the peak value. In addition, we have that \\begin{eqnarray} I &=& \\sum_{i=1}^n f_i, \\\\ \\bar x & = & \\frac{\\sum_{i=1}^n x_if_i}{I}, \\\\ \\sigma^2 & = & \\frac{\\sum_{i=1}^n(x_i-\\bar x)^2f_i}{I}, \\\\ S & = & \\frac{\\sum_{i=1}^n(x_i-\\bar x)^3f_i}{I\\sigma^3}, \\label{skew} \\end{eqnarray} where $f_i$ is the line flux in the wavelength pixel $i$ whose coordinate is $x_i$, and the summations are performed over the pixels covered by the \\Lya line. On general grounds one would expect that the observed \\Lya line shape would be more symmetric (i.e. low $S_W$) in reionization models characterized by a lower value of $\\chi_{HI}$. However, given the above definition, just the opposite is true. In fact, for any reasonable value of the relevant parameters (see Fig.1) the blue part of the line is heavily absorbed, thus yielding a high value of $S_W$; as $\\chi_{HI}$ is increased, also the long-wavelength part of the line is affected by the red damping wing, making the line more symmetric around the peak.   The predicted trend of $S_W$ with the observed \\Lya luminosity at $z=6.56$ is reported in Fig.\\ref{weighted_skewness}, for the parameters of the ERM and LRM that best fit the LF data (discussed in the previous Section). For both models, the weighted skewness of the line increases for more luminous objects; however, such dependence is steeper for the ERM than for the LRM. In general, though, the two reionization scenarios predict $S_W$ values in the range 10-17. The data from Kashikawa et al. (2006) spans the somewhat larger range 3-17, with many of the data points lying around $S_W=5$. Given the paucity of the observed points and the large errors associated to them, it is probably premature to draw any strong claim from these results. However, given the constant increase in the amount and quality of LAE data, it is quite possible that the line skewness could represent a very interesting tool to constrain reionization models in the near future. It has to be noted that the data show a large scatter of $S_W$ at a given value of $L_\\alpha$, perhaps indicating that local conditions, including gas infall/outflow, density inhomogeneities and interaction of the Ly$\\alpha$ line with the interstellar medium of the galaxy, might play a dominant role. These can only be investigated in a statistically meaningful manner via high-resolution numerical simulations to which we will defer a forthcoming study.   From the theoretical point of view it is instructive to summarize the response of the skewness to different physical conditions. As we have seen from Fig.\\ref{weighted_skewness}, $S_W$ increases with $L_\\alpha$ (or, equivalently with SFR); this is true for any fixed value of $\\chi_{HI}$.  This is because as the SFR increases, more of the Ly$\\alpha$ line escapes forcing $S_W$ to increase as a result of the larger value of $\\Delta \\lambda$. Further, the long -wavelength part of the observed Ly$\\alpha$ line begins to flatten with increasing $\\chi_{HI}$ due to attenuation by the red damping wing. Hence, $\\Delta \\lambda$ varies slower with $L_\\alpha$ (SFR) for high $\\chi_{HI}$ (LRM) as compared to lower values (ERM); this makes the slope of $S_W$ steeper for the ERM.   A more general view of the dependence of $S_W$ on $L_\\alpha$ (hence on SFR) and $\\chi_{HI}$ is shown in Fig. \\ref{countour}. The plot has been obtained by dividing the observed Ly$\\alpha$ luminosity into bins and averaging the weighted skewness over the number of LAEs in each bin at a given value of $\\chi_{HI}$. The regions with weighted skewness values equal to zero represent a lack of LAEs in that bin.  The most intriguing feature of Fig.\\ref{countour} is a clear anti-correlation between $L_\\alpha$ and $\\chi_{HI}$. Given the range of SFR considered ($\\dot M_\\star=2.7-197 M_\\odot$yr$^{-1}$), LAEs populate progressively fainter \\Lya luminosity bins as the IGM becomes more neutral. Notice that relatively luminous objects ($L_\\alpha \\approx 10^{42.5 }$~erg~s$^{-1}$) would not be detected if $\\chi_{HI}\\simgt 0.25$. Within the range in which these objects are visible, the most luminous objects always show the largest $S_W$ at fixed $\\chi_{HI}$; however, such maximum value is also seen to increase with decreasing $\\chi_{HI}$.  The model does not include important effects such as inflows/outflows and interaction of the Ly$\\alpha$ photons with the ISM, which will definitely leave an imprint on the $S_W$ and hence, weaken the $L_\\alpha$-$\\chi_{HI}$ anti-correlation. However, a composite Ly$\\alpha$ line, built from the observations of a sufficiently large number of LAEs might show such an anti-correlation. In future work, we will endeavor to include these effects, hence obtaining an estimate of how large a sample might be sufficient for this purpose."
    },
    "0807/0807.2372_arXiv.txt": {
        "abstract": "{We have re-analyzed the X-ray spectra of the \\emph{gravitational lensed} high-redshift BAL QSO \\apm, observed with the \\xmm \\sp and \\chan \\sp observatories. Previous studies (Hasinger et al. 2002; Chartas et al. 2002) detected unusual, highly-ionized iron absorption features, but differed in their interpretation of these features, regarding the kinematical and ionization structure.}{We seek one physical model that can be successfully applied to both observations.}{For the first time we have performed detailed photoionization modeling on the X-ray spectrum of \\apm.}{The absorbing gas in \\apm \\sp can be represented by a two-absorbers model with column densities $N_H(1)\\approx 7\\times 10^{22}$ \\cmn, $N_H(2)\\approx 6\\times 10^{22}$ \\cmn, and ionization parameters $\\log\\xi(1)\\approx 1.5$ and $\\log\\xi(2)\\approx 3$, with one of them (the high-ionization component) outflowing at $v\\approx 0.18(\\pm 0.01)c$, carrying large amount of gas out of the system. We find that the \\chan \\sp spectrum of \\apm \\sp requires the same Fe/O ratio overabundance (previously) indicated by the \\xmm \\sp observation, showing that both absorber components underwent similar chemical evolution and/or have similar origin.}{} ",
        "introduction": "Broad absorption line (BAL) quasi-stellar objects (QSOs) are objects displaying in their spectra broad (FWHM $\\approx$ 10\\,000 \\kms) absorption lines in the rest-frame ultraviolet (UV), originated in outflows of matter from the central engine of QSOs \\citep{foltz1990a,weymann1991a}. The outflow velocity may reach up to $0.2c$ \\citep[e.g.,][]{foltz1983a}. Determining the relationship between the material absorbing the X-rays and the one absorbing the UV radiation is key to our understanding of the geometry and the physical state of the medium surrounding the vicinity of supermassive black holes \\citep[e.g.,][]{mathur1995a,murray1995a,hamann1998a,proga2000a}.  Before \\chan/\\xmm \\sp missions, detections of BAL quasars in X-ray were rare. Usually, these object are X-ray weak \\citep[e.g.,][]{gallagher2006a}, sometimes interpreted as strong excess absorption. \\chan \\sp and \\xmm \\sp observations of BAL QSOs, have provided new constraints in the amount of absorption toward selected objects \\citep[e.g.,][]{sabra2001a,oshima2001a,gallagher2002a,gallagher2006a}, indicating large column densities $N_H \\gtrsim 10^{23}$ \\cmn.  The BAL QSO \\apm \\space at redshift $z=3.91$ \\citep{irwin1998a} is one of the most luminous objects in the universe, further magnified by gravitational lensing by a factor of $\\sim 50-100$ \\citep[e.g.,][]{ledoux1998a}. It was detected with the Submilimiter Common-User Bolometric Array, implying an apparent far-infrared luminosity greater than 5$\\times 10^{15} L_{\\sun}$ \\citep{lewis1998a}. The optical spectrum, obtained with the High Resolution Echelle Spectrometer at the Keck telescope \\citep{ellison1999a}, along with a detailed study of the physical conditions in the BAL flow of the QSO by \\cite{srianand2000a}, allowed them to conclude that the corresponding gas stream, outflowing with velocities of up to 12\\,000 \\kms, is heavily structured and highly ionized.  The quasar \\apm \\space was observed twice with \\xmm \\sp \\citep[][ hereafter H02]{hasinger2002a}. In both observations the quasar is observed clearly out to 12 keV, which corresponds to almost 60 keV in the rest frame. The most apparent feature in the \\xmm \\sp spectrum is an absorption-like feature around 1.55 keV (which they interpret as an absorption edge), corresponding to $\\sim 7.7$ keV in the rest frame of \\apm. The high-inferred iron abundance at the high redshift, corresponding to a young age of the universe, is of great interest in the context of chemical enrichment models, and provides constraints on the early star formation history of the universe and on its cosmological parameters \\citep[e.g.,][]{hamann1993a,hasinger2002a,komossa2003a}.  The quasar \\apm \\space was also observed with \\chan \\sp \\citep[][ hereafter C02]{chartas2002a}. The \\chan \\sp spectrum shows a similar absorption feature as the \\xmm \\sp observation, but the feature led to a different interpretation. In particular, C02 modeled the spectrum with two absorption lines at 8.1 keV and 9.8 keV in the rest frame of the quasar, interpreted as \\ion{Fe}{xxv} K lines. If the \\chan \\sp interpretation of the data is right, the {\\it bulk} velocity of the X-ray BALs is $\\sim 0.2c-0.4c$. The presence of similar outflow velocities has been claimed in a few other AGN X-ray spectra \\citep[e.g.,][]{pounds2003a,chartas2003a,pounds2006a}, but alternative interpretations of the same spectra have been proposed, which do not require these relativistic outflow velocities \\citep[e.g.,][]{kaspi2006a}.  New models \\citep[e.g.,][]{elvis2000a,proga2000a,proga2007a} predict that a large fraction of the accreted matter into the region of the compact object, is expelled out again in the form of high-velocity outflows. Broad UV absorption features, with velocities $\\sim$ 0.05$c-0.1c$, are associated with this material through acceleration mechanisms like acceleration by gas pressure \\citep[e.g.,][]{weymann1982a,begelman1991a} due to dust \\citep[e.g.,][]{voit1993a,yun1995a} and acceleration due to radiative pressure by spectral lines \\citep[e.g.,][]{drew1984a,shlosman1985a,arav1994a,dekool1995a,murray1995a,proga2000a} observed in about 10\\% of luminous high-$z$ quasars \\citep{laor2002a}, implying that these outflows are an important component of the general picture of AGNs. Furthermore, these models also predict that in order for the UV material to reach such high velocities a \\emph{shield} made by a high column density of gas ($N_H\\sim 10^{22}-10^{24}$ \\cmn) must absorb radiation in the X-ray band, in this manner preventing the destruction of the UV material. The new generation of X-ray observatories, \\xmm \\sp and \\chan \\sp give us an unique opportunity to study in detail the X-ray component of the BAL QSOs, contributing in our knowledge about the dynamical and physical evolution of the center of these systems, and in the case of high-redshited BAL QSOs, the gas enrichment history in the early universe \\citep{hamann2004a}.  However, none of the previous studies include a self-consistent photoionization modeling of the X-ray spectrum of \\apm. Any constraint on the ionization level of the absorbing gas and its connection with the kinematical properties of the BAL outflow is of major interest to elucidate some of the greatest discrepancies between these two proposed models. Furthermore, a model that can separate the differences between observations (\\xmm \\sp {\\it vs} \\chan), or unify both in a single frame, is highly desirable. We report a spectral analysis (made of these two observations separated by $\\sim 2$ weeks in the rest frame) of the high-redshifted BAL QSO \\apm, and present one model that might reconcile both observations in the same physical context, shedding new light on the ionization degree, kinematics and evolution of this system, in connection with its cosmological consequences. We use the following cosmological parameters: H$_0=70$ \\kms \\space Mpc$^{-1}$, $\\Omega_M=0.3$ and $\\Omega_{\\lambda}=0.7$. ",
        "conclusions": "}  We have found several interesting aspects in the X-ray spectrum of the QSO \\apm. The physical scenarios raised after the interpretation of \\cite{hasinger2002a} as a Fe edge, and of \\cite{chartas2002a} as relativistic ($\\sim 0.2c-0.4c$) Fe~{\\sc xxv} lines, have helped to scrutinize more closely the X-ray spectrum of this object. If we attempt to explain the spectrum by {\\it only} investigating the full band $\\sim 0.2-10$ k~e~V (observed frame), we find it hard to justify which model represents the data better in a statistical sense. After including a local criterion to evaluate the fits in     \t % the band $1-2$ k~e~V, we were able to converge to a model that   % accounts for both low- and high- energy spectral bands.      \t % Indeed, by modeling the spectrum with a photoionization model, which gives a better (over Gaussian lines alone for instance) representation of edges and absorption lines in the whole band $\\sim 0.2-10$ keV, we are able to draw more physical information.  We discuss the implications of our results from Sect. \\ref{photo} in two parts: (i) Assuming that the X-ray spectrum of \\apm \\sp can be modeled by a {\\tt single-absorber} photoionization model. (ii) Assuming that the best representation of the spectrum is a {\\tt two-absorbers} model.  (i) {\\it Changes in the physical and kinematical state of the absorber:} We see that the \\xmm \\sp data is represented by a highly-photoionized gas with column density $\\sim 10^{23}$ \\cmn, and $\\log(\\xi)\\sim 1.5$ at 0 \\kms \\sp in the frame of the quasar. Nevertheless, by allowing iron to be higher than {\\it solar} (model B), we obtain a iron abundance best-fit value of ${\\rm Fe}=5.04^{+0.44}_{-0.23}$, and the ionized gas column no longer needs to be as high as $10^{23}$ \\cmn, but $\\sim 5\\times 10^{22}$ \\cmn \\sp instead, statistically improving the fit at 99\\% of significance over model A (with the same number of parameters than the absorbed power-law modified by an edge), similar to the H02 interpretation.  But $\\sim$ 2 weeks before (in the frame of the quasar), the \\chan \\sp spectrum presented evidence of a material outflowing $\\sim 0.21c$, not only showing an absorption-line-like feature at $\\sim 8.1$ keV (rest-frame), but also pointing to absorption at lower energies $\\sim 0.8-2$ keV (rest-frame), coming from species of iron from Fe~{\\sc xviii} to Fe~{\\sc xxiv}, and H- and He-like ions of Ne, Mg, Si, S, and O. In this context, the \\chan \\sp data appears to be pointing to (through model D): (a) a change in the kinematical state of the absorbing material, or more plausible, change in the projection of the velocity field on our line of sight (which would yield to the same effect), from $\\sim 0.21c$ to $\\sim $ 0 \\kms \\sp (since the \\chan \\sp observation was made first than the \\xmm \\sp observation); and (b) a slight but noticeable change in the ionization state of the gas, from $\\log(\\xi) \\sim 2$ to $\\log(\\xi) \\sim 1.5$ in a timescale of $\\sim 2$ weeks in the frame of the quasar. However, we find it physically hard to explain the deceleration shown by this model, in addition to a change in the ionization state of the gas without a significant change in the intrinsic luminosity of the source, and once we found the {\\tt two-absorbers} model statistically superior, we ruled out (i) and focused on (ii).  (ii) {\\it Both absorbers have been there}: The {\\tt two-absorbers} model consists of: one absorber at     \t\t\t% $v=0$ \\kms, $N_H(1)=7.34 _{-0.15} ^{+0.18}\\times 10^{22}$ \\cmn, \t\t% $\\log\\xi(1)=1.50 _{-0.04} ^{+0.02}$, and a second absorber at    \t\t% $v=0.18c$, $N_H(2)=5.49 _{-0.54} ^{+0.71}\\times 10^{22}$ \\cmn,\t\t\t% $\\log\\xi(2)=3.00 _{-0.01} ^{+0.04}$ \t\t\t\t\t\t% (parameters coming from the simultaneous best-fit).\t\t\t\t% The chemical composition\t\t\t\t\t\t\t% of both is the same, {\\it solar} in all the elements (see model \t\t% composition in Sect. \\ref{photo}) except in the abundance of iron\t\t\t% ${\\rm Fe}=2.85 _{-0.10} ^{+0.10}$.\t\t\t\t\t\t% It is worth mentioning that from the separated fits, the two observations show a change in the power-law of $\\approx 10$\\% and also show changes in the physical parameters. However, they (individually) change little (i.e, no large difference within errors); $\\log\\xi(1)$, from $\\approx$ 1.2 to 1.5 ($\\approx$ 30\\%); $N_H(1)\\times 10^{22}$ \\cmn, from $\\approx$ 5.7 to 6.1 ($\\approx$ 10\\%); $\\log\\xi(2)$, from $\\approx$ 3 to 3.1 ($\\approx$ 3\\%); $N_H(2)\\times 10^{22}$ \\cmn, from $\\approx$ 4.8 to 2.7 ($\\approx$ 70\\%, but see errors). The computation of $\\chi^2$, on the other hand, is very sensitive to the continuum level, and small changes in the power-law are easily detected by the minimization-$\\chi^2$ routine. {\\it We do not think this is a fundamental physical change of state} between observations. At the moment these data do not allow us to discriminate whether the changes seen in the physical parameters are significant (within $\\approx 2\\sigma$) or not. Apparently, apart from the power-law component, both observations can be represented by the same two-absorbers model.  We have verified that the overabundance of Fe can be established\t\t% respect to the abundance of oxygen (since it is the Fe/O ratio, which is\t% cosmologically relevant), by comparing the observed spectra\t\t\t% with the synthetic theoretical spectra at low-energies, if we increase\t\t% the oxygen abundance. The ratio Fe/O must be $\\approx 3$ in order to obtain\t% acceptable fits in the low-energy band.\t\t\t\t\t\t% \\footnote{We verified that the data is sensitive to an increase of oxygen abundance. We take the best-fit two-absorbers model and compute models with the oxygen abundance at 1.5, 3 and 5 $\\times$ \\emph{solar} oxygen abundance. The fit get worse with $\\chi^2({\\rm O=1.5})=304$, $\\chi^2({\\rm O=3})=333$ and $\\chi^2({\\rm O=5})=388$. The conclusions are: 1) the data is sensitive to the Fe/O ratio. 2) This ratio must be $\\approx 3 \\pm 1$, in order to produce acceptable fits (and produce good description of the low-energy band).} This was first noticed by \\cite{hasinger2002a}, and its        \t\t\t% cosmological implications are discussed in \\cite{komossa2003a},\t\t\t% based on chemical evolution studies of \\cite{hamann1993a}.\t\t\t%   This multi-component photoionization approach has been       \t\t\t% successfully applied to other AGNs in different bands (X-ray and UV): \t\t% NGC 3783 \\citep{netzer2003a,krongold2003a,gabel2005a};\t\t\t\t% Mrk 279 \\citep{fields2007a,costantini2007a};\t\t\t\t\t% NGC 4593 \\citep{steenbrugge2003b};\t\t\t\t\t\t% NGC 985 \\citep{krongold2005a};\t\t\t\t\t\t\t% \\mr \\sp \\citep{kaspi2004a};\t\t\t\t\t\t\t% NGC 4051 \\citep{kraemer2006a,armentrout2007a,krongold2007a}, \t\t\t% a survey of a sample of 15 Type I AGNs can be found in \\cite{mckernan2007a}.\t% We will treat, this multi-component photoionization approach, as the best \t% solution for our problem. \t\t\t\t\t\t\t% However, note that a key difference with some of these cited papers is that the solution for \\apm \\sp requires two unrelated absorbers at very different velocities.  A plausible framework in which to place these observational clues is in accretion disk wind models \\citep{murray1995a,elvis2000a,proga2000a,proga2007a}. Recent high-resolution spectroscopy of \\apm, shows UV BALs in a wide range of velocities ($5\\,000 \\lesssim v \\lesssim 12\\,000$ \\kms) and ionization degrees (ions like C~{\\sc vi}, O~{\\sc vi}, N~{\\sc v} and Si~{\\sc iv}), implying that the outflow may be composed of multiple components, with velocity; density and/or ionization parameter gradients \\citep{srianand2000a}. In particular, the model proposed by \\cite{elvis2000a} predicts that the UV BAL is formed in the conical section of a funnel-shaped flow ranging velocities $10\\,000 \\lesssim v \\lesssim 60\\,000$ \\kms, (or up to $\\sim 0.2c$ in our context). Furthermore, the model is very specific as to the location and physical properties of the X-ray absorbing material. The warm highly-ionized medium has column densities of $\\sim 10^{22.5}$ \\cmn, and temperatures of $\\sim 10^{6}$ K, right pressure, and ionization parameter to ensure pressure equilibrium with the BELR clouds. Our XSTAR-based photoionization model estimates that the high-ionization, high-outflow velocity component (HV component), $\\log(\\xi)\\approx 3$ has a temperature $T\\sim$ few times $10^6$ K, and the best-fit column density $N_H(2) \\sim 10^{22.5}$ \\cmn, nicely consistent with the main physical properties of the warm highly-ionized medium. The low-ionization rest-frame velocity component (RV component) with $\\log(\\xi)\\approx 1.5$ has a temperature $T\\sim$ $10^5$ K, in accordance with the temperature of $\\sim$ few times $10^5$ K found by \\cite{kaastra1995a} in NGC 5548, and the degree of ionization is consistent with $\\log(U) \\sim 0.3$ reported by \\cite{mathur1995a} for the same object, where a test to unify UV/X-ray absorbers in Seyfert galaxies was successfully applied.  Recalling the definition of the ionization parameter $\\xi=L_{\\rm ion}/(nR^2)$, where $L_{\\rm ion}$ is ionizing luminosity, $n$ is the gas density, and $R$ is the location of the absorber, we are able to compute the product of two un-observable quantities $(nR^2)$ from the observable $\\xi$ and $L_{\\rm ion}$. We measured an intrinsic X-ray luminosity \\footnote{This is the unabsorbed intrinsic $(1-50)$ keV rest-frame luminosity of the source. From the \\xmm \\sp data $L_{\\rm x}(\\xmm)=1.07(\\pm0.03)\\times 10^{47}k^{-1}$ \\ergs, from the \\chan \\sp data $L_{\\rm x}(\\chan)=1.17(\\pm0.05)\\times 10^{47}k^{-1}$ \\ergs. We are taking the average of both.} of $L_{\\rm x}=1.12(\\pm0.04)\\times 10^{47}k^{-1}$ \\ergs, $k$ being the lensing magnification factor. With this we have: $(nR^2)$[HV]$=1.12(\\pm 0.11)\\times10^{44}k^{-1}$ cm$^{-1}$ and $(nR^2)$[RV]$=3.54(\\pm 0.11)\\times10^{45}k^{-1}$ cm$^{-1}$. This is the first time this quantity is measured for this BAL QSO (since it is the first time $\\xi$ is properly constrained). Despite we still are not able to break the natural degeneracy between $n$ and $R$ (due to lack of variability), it allows us to present order-of-magnitude estimates on the mass-loss rate and energy budgets of the system. Assuming the HV forms part of the high-ionization BAL flows, which reaches velocities up to $\\sim 0.2c$ \\citep[see][ for a recent review of high velocity (HV) outflow in quasars]{rodriguez2007a}, and that the X-ray absorber in \\apm \\sp is part of the conical section of the funnel-shaped flow proposed by \\cite{elvis2000a} (although this estimation is independent of this specific geometry and can be equally applied to a spherical shell of gas for instance, without significatively changing the final conclusions), which forms a shell with covering factor $f_{\\rm cov}$ moving at a velocity $v$, we estimate the mass-loss rate for the system: \\begin{equation} \\dot{M}_{w}= 4\\pi \\mu_H (nR^2) vf_{\\rm cov}, \\end{equation} where $\\mu_H$ is the mean mass per H atom. Using $(nR^2)$[HV]$=1.12\\times10^{44}k^{-1}$ cm$^{-1}$, an outflow velocity $v=0.18c$ and assuming a covering factor of 0.01, we have a mass-loss rate of $\\dot{M}_{w}\\approx 2873k^{-1}$ $M_{\\odot}~{\\rm yr}^{-1}$. The total kinetic power released by the wind in \\apm \\sp is $\\dot{M}_{w}v^2/2=2.6\\times 10^{48}k^{-1}$ \\ergs. The accretion rate is related with the accretion luminosity with: \\begin{equation} \\dot{M}_{\\rm acc} \\approx 17 \\frac{L_{\\rm bol}}{10^{47}~{\\rm ergs~s^{-1}}}\\frac{0.1}{\\eta} M_{\\odot}~{\\rm yr}^{-1}, \\end{equation} where $L_{\\rm bol}$ is the bolometric luminosity, $\\eta$ is the accretion efficiency (assumed to be $\\sim 0.1$), and $c$ is the speed of light. With our measured $L_{\\rm x}$ we are able to estimate $L_{\\rm bol}=\\beta L_{\\rm x}$. We adopt a bolometric correction factor $\\beta=20$, based on the mean AGN SED of \\cite{elvis1994a}. Explicitly, in terms of the lensing magnification factor, $\\dot{M}_{\\rm acc} \\approx 340k^{-1}$~ $M_{\\odot}~{\\rm yr}^{-1}$. This would lead to the conclusion that the central black hole in \\apm \\sp {\\it accretes $\\sim 1$ order of magnitude less matter than the estimated mass-loss rate}. Recently, from high-resolution \\chan \\sp and \\xmm \\sp spectroscopic studies at least three systems show the property of having mass outflow rates $\\sim 1$ order of magnitude larger than the accretion rates (at $\\sim $ 10\\% efficiency): the micro-quasar GRS 1915+105 \\citep{lee2002a}, the Seyfert 1 galaxy MCG-6-30-15 \\citep{turner2004a,sako2003a,lee2002b}, and the Seyfert 2 galaxy IRAS 18325-5926 \\citep{lee2005a}. However, we are aware of two major sources of uncertainties in these estimations. Due to its location, the uncertainty in $f_{\\rm cov}$ is large because the exact geometry in the base of the outflow is not known. On the other hand, the bolometric correction factor could be larger than the one used here, locating $\\dot{M}_{\\rm acc}$ and $\\dot{M}_w$ in the same order of magnitude. Either we detect the outflow in a very small period of activity (which would be very fortuitous) or the outflow plays a major role in the process of sending matter back to the interstellar medium \\citep[e.g.,][]{crenshaw2003a,lee2005a,krongold2007a}. If we assume that the mass of the black hole $M_{\\rm BH}\\approx 2\\times 10^{10}$ $M_{\\odot}$, \\citep[as estimated by][]{shields2006a}, the ratio $L_{\\rm bol}/L_{\\rm edd}\\approx 0.8k^{-1}$ (where $L_{\\rm edd}$ is the Eddington luminosity). This number is in accordance with the general trend of low $L_{\\rm bol}/L_{\\rm edd}$ ratio for BALQSOs reported by \\cite{boroson2002a}. With these estimations, we can derive the launching radius $r_{\\rm launch}$ for a radiatively-driven wind: \\begin{equation} r_{\\rm launch}\\approx \\frac{2GM_{\\rm BH}}{v^2} \\left( \\Gamma_f \\frac{L_{\\rm UV}}{L_{\\rm edd}}-1\\right). \\end{equation}  Here, we have used the relation $v~vs~r$ for a radiatively-driven steady wind and assumed that at very large radii the flow has reached a terminal velocity $v$. Assuming that the wind is driven by a UV luminosity $L_{\\rm UV}\\sim 4\\times 10^{46}$ \\ergs, with a force multiplier $\\Gamma_f\\sim 100$ \\citep[taken from][]{laor2002a}, and using a flowing velocity $v=0.18c$, we find $r_{\\rm launch}\\approx 7\\times 10^{16}$ cm. On the other hand, if we use the maximum observed velocity of the UV absorber (through the C~{\\sc vi} line), for this quasar of $v=0.041c$, we find $r_{\\rm launch}\\approx 1\\times 10^{18}$ cm. One possible explanation for the difference in the two distances, is that the UV absorber is at a lower ionization state, lying in the low ionization BAL region of the funnel-shaped flowing \\citep[for visual help see Fig. 5 of][]{elvis2000a}, and more specifically in the large radii part in which the UV absorber has reached its highest (terminal) velocity. A rigorous study of the hydrodynamical properties of the X-rays/UV absorbers in AGN is beyond the scope of this paper. However, the simplistic approach used in this analysis could be useful as a framework for future works on the study of the central region of this quasar.  To summarize, the absorbing gas in \\apm \\sp can be represented by a {\\tt two-absorbers} model with column densities $N_H(1)\\approx 7\\times 10^{22}$ \\cmn, $N_H(2)\\approx 6\\times 10^{22}$ \\cmn, ionization parameters $\\log\\xi(1)\\approx 1.5$ and $\\log\\xi(2)\\approx 3$, with one of them (the HV component) outflowing at $v\\approx 0.18(\\pm 0.01)c$, carrying large amount of gas out of the system. The feature at $\\sim 8$ keV (rest-frame) is fully predicted and reproduced by our photoionization model, to be a complex of Fe lines coming from high state of ionization, in which the main contributor is the Fe~{\\sc xxv} $\\lambda 1.85$ \\AA, improving the characterization of the kinematics and the quantitative evolution analysis of this high $z$ quasar. We confirm evidence for an overabundance of Fe/O, from the \\xmm \\sp observation of this quasar \\citep[previously inferred and discussed in][ based on calcutations of Hamann \\& Ferland 1993]{hasinger2002a,komossa2003a}. The analysis is made for the first time on the \\chan \\sp observation of \\apm, implying that both absorbers require Fe/O supersolar, placing similar constraints on models as before, and additionally shows that both independent absorbers have a similar chemical history."
    },
    "0807/0807.3697_arXiv.txt": {
        "abstract": "Motivated by the possible existence of other universes, with possible variations in the laws of physics, this paper explores the parameter space of fundamental constants that allows for the existence of stars. To make this problem tractable, we develop a semi-analytical stellar structure model that allows for physical understanding of these stars with unconventional parameters, as well as a means to survey the relevant parameter space.  In this work, the most important quantities that determine stellar properties --- and are allowed to vary --- are the gravitational constant $G$, the fine structure constant $\\alpha$, and a composite parameter $\\conlum$ that determines nuclear reaction rates.  Working within this model, we delineate the portion of parameter space that allows for the existence of stars.  Our main finding is that a sizable fraction of the parameter space (roughly one fourth) provides the values necessary for stellar objects to operate through sustained nuclear fusion.  As a result, the set of parameters necessary to support stars are not particularly rare. In addition, we briefly consider the possibility that unconventional stars (e.g., black holes, dark matter stars) play the role filled by stars in our universe and constrain the allowed parameter space. ",
        "introduction": "The current picture of inflationary cosmology allows for, and even predicts, the existence of an infinite number of space-time regions sometimes called pocket universes \\cite{guth, rees, vilenkin}.  In many scenarios, these separate universes could potentially have different versions of the laws of physics, e.g., different values for the fundamental constants of nature. Motivated by this possibility, this paper considers the question of whether or not these hypothetical universes can support stars, i.e., long-lived hydrostatically supported stellar bodies that generate energy through (generalized) nuclear processes.  Toward this end, this paper develops a simplified stellar model that allows for an exploration of stellar structure with different values of the fundamental parameters that determine stellar properties. We then use this model to delineate the parameter space that allows for the existence of stars.  A great deal of previous work has considered the possibility of different values of the fundamental constants in alternate universes, or, in a related context, why the values of the constants have their observed values in our universe (e.g., \\cite{carr, bartip}).  More recent papers have identified a large number of possible constants that could, in principle, vary from universe to universe.  Different authors generally consider differing numbers of constants, however, with representative cases including 31 parameters \\cite{tegmark} and 20 parameters \\cite{hogan}. These papers generally adopt a global approach (see also \\cite{wilczek, tegold, aguirre}), in that they consider a wide variety of astronomical phenomena in these universes, including galaxy formation, star formation, stellar structure, and biology.  This paper adopts a different approach by focusing on the particular issue of stars and stellar structure in alternate universes; this strategy allows for the question of the existence of stars to be considered in greater depth.  Unlike many previous efforts, this paper constrains only the particular constants of nature that determine the characteristics of stars.  Furthermore, as shown below, stellar structure depends on relatively few constants, some of them composite, rather than on large numbers of more fundamental parameters. More specifically, the most important quantities that directly determine stellar structure are the gravitational constant $G$, the fine structure constant $\\alpha$, and a composite parameter $\\conlum$ that determines nuclear reaction rates. This latter parameter thus depends in a complicated manner on the strong and weak nuclear forces, as well as the particle masses. We thus perform our analysis in terms of this $(\\alpha, G, \\conlum)$ parameter space.  The goal of this work is thus relatively modest.  Given the limited parameter space outlined above, this paper seeks to delineate the portions of it that allow for the existence of stars. In this context, stars are defined to be self-gravitating objects that are stable, long-lived, and actively generate energy through nuclear processes. Within the scope of this paper, however, we construct a more detailed model of stellar structure than those used in previous studies of alternate universes. On the other hand, we want to retain a (mostly) analytic model. Toward this end, we take the physical structure of the stars to be polytropes. This approach allows for stellar models of reasonable accuracy; although it requires the numerical solution of the Lane-Emden equation, the numerically determined quantities can be written in terms of dimensionless parameters of order unity, so that one can obtain analytic expressions that show how the stellar properties depend on the input parameters of the problem. Given this stellar structure model, and the reduced $(\\alpha, G, \\conlum)$ parameter space outlined above, finding the region of parameter space that allows for the existence of stars becomes a well-defined problem.  As is well known, and as we re-derive below, both the minimum stellar mass and the maximum stellar mass have the same dependence on fundamental constants that carry dimensions \\cite{phil}. More specifically, both the minimum and maximum mass can be written in terms of the fundamental stellar mass scale $M_0$ defined according to \\be M_0 = \\alpha_G^{-3/2} m_P = \\left( {\\hbar c \\over G} \\right)^{3/2} m_P^{-2} \\, \\approx \\, 3.7 \\times 10^{33} g \\approx 1.85 M_\\odot \\, , \\label{mscale} \\ee where $\\alpha_G$ is the gravitational fine structure constant, \\be \\alpha_G = {G m_P^2 \\over \\hbar c} \\approx 6 \\times 10^{-39} \\, , \\label{alphag} \\ee where $m_P$ is the mass of the proton.  As expected, the mass scale can be written as a dimensionless quantity ($\\alpha_G^{-3/2}$) times the proton mass; the appropriate value of the exponent (--3/2) in this relation is derived below. The mass scale $M_0$ determines the allowed range of masses in any universe.  In conventional star formation, our Galaxy (and others) produces stars with masses in the approximate range $0.08 \\le M_\\ast/M_\\odot \\le 100$, which corresponds to the range $0.04 \\le M_\\ast / M_0 \\le 50$. One of the key questions of star formation theory is to understand, in detail, how and why galaxies produce a particular spectrum of stellar masses (the stellar initial mass function, or IMF) over this range \\cite{sal}.  Given the relative rarity of high mass stars, the vast majority of the stellar population lies within a factor of $\\sim10$ of the fundamental mass scale $M_0$. For completeness we note that the star formation process does not involve thermonuclear fusion, so that the mass scale of the hydrogen burning limit (at 0.08 $M_\\odot$) does not enter into the process. As a result, many objects with somewhat smaller masses -- brown dwarfs -- are also produced. One of the objectives of this paper is to understand how the range of possible stellar masses changes with differing values of the fundamental constants of nature.  This paper is organized as follows. We construct a polytropic model for stellar structure in \\S 2, and identify the relevant input parameters that determine stellar characteristics. Working within this stellar model, we constrain the values of the stellar input parameters in \\S 3; in particular, we delineate the portion of parameter space that allow for the existence of stars. Even in universes that do not support conventional stars, those generating energy via nuclear fusion, it remain possible for unconventional stars to play the same role.  These objects are briefly considered in \\S 4 and include black holes, dark matter stars, and degenerate baryonic stars that generate energy via dark matter capture and annihilation. Finally, we conclude in \\S 5 with a summary of our results and a discussion of its limitations, including an outline for possible future work. ",
        "conclusions": "In this paper, we have developed a simple stellar structure model (\\S 2) to explore the possibility that stars can exist in universes with different values for the fundamental parameters that determine stellar properties. This paper focuses on the parameter space given by the variables $(G, \\alpha, \\conlum)$, i.e., the gravitational constant, the fine structure constant, and a composite parameter that determines nuclear fusion rates. The main result of this work is a determination of the region of this parameter space for which bona fide stars can exist (\\S 3). Roughly one fourth of this parameter space allows for the existence of ``ordinary'' stars (see Figure \\ref{fig:agspace}). In this sense, we conclude that universes with stars are not especially rare (contrary to previous claims), even if the fundamental constants can vary substantially in other regions of space-time (e.g., other pocket universes in the multiverse).  Another way to view this result is to note that the variables $(G, \\alpha, \\conlum)$ can vary by orders of magnitude from their measured values and still allow for the existence of stars.  For universes where no nuclear reactions are possible, we have shown that unconventional stellar objects can fill the role played by stars in our universe, i.e., the role of generating energy (\\S 4).  For example, if the gravitational constant $G$ and the fine structure constant $\\alpha$ are smaller than their usual values, black holes can provide viable energy sources (Figure \\ref{fig:bhspace}).  In fact, all universes can support the existence of stars, provided that the definition of a star is interpreted broadly. For example, degenerate stellar objects, such as white dwarfs and neutron stars, are supported by degeneracy pressure, which requires only that quantum mechanics is operational. Although such stars do not experience thermonuclear fusion, they often have energy sources, including dark matter capture and annihilation, residual cooling, pycnonuclear reactions, and proton decay. Dark matter particles can also (in principle) form degenerate stellar objects (see \\S 4).  In order to assess the suitability of non-nuclear power sources, one must specify how much power is required, and for how long.  In this work we have used the power that Earth intercepts from the Sun as the minimum benchmark value $\\lmin_0$, and scaled the necessary power according to equation (\\ref{lminscale}) to account for variations in the fine structure constant; similarly, the required amount of time is taken to to $\\sim1$ Gyr, scaled by the atomic time of equation (\\ref{atomic}).  These choices are not definitive, and hence alternative scalings can be explored.  The issue of alternate values for the fundamental constants, as considered herein, is related to the issue of time variations in the constants in our universe. However, current experiments place rather strong limits on smooth time variations, with time scales exceeding the current age of the universe (see the review of \\cite{uzan}). Another possibility is for the constants to have different values at other spatial locations within our universe, although this scenario is also highly constrained \\cite{barrow}.  This paper has focused on stellar structure properties. An important related question (beyond the scope of this work) is whether or not stellar bodies can be readily made in universes with varying values of the constants. Even if the laws of physics allow for stellar objects to exist and actively burn nuclear fuel, there is no guarantee that such bodies will be produced in significant numbers. In our universe, for example, there is a moderate mismatch between the mass range of possible stars and the distribution of masses of stellar bodies produced by the star formation process.  At the present cosmological epoch, star formation produces objects over the entire possible range of stellar masses, with additional bodies produced in the substellar range (brown dwarfs).  The matching is relatively good, in that the fraction of bodies in the brown dwarf range is small, only about 1 out of 5 \\cite{luhman}.  Since the masses of these objects are small, the fraction of the total mass locked up in the smallest bodies is even smaller, less than 5 percent. On the other hand, nearly all of the stars in our universe have small masses. As one benchmark, only about 3 or 4 out of a thousand stars are larger than the $\\sim 8 M_\\odot$ threshold required for stars to experience a supernova explosion, whereas stellar masses can extend up to $\\sim100 M_\\odot$. The high mass end of the possible mass range is thus sparsely populated.  The corresponding match between the range of allowed stellar masses and the mass range of objects produced can be quite different in other universes.  In future work, another issue to be considered is coupling the effects of alternate values of the fundamental constants to the cosmic expansion, big bang nucleosynthesis, and structure formation. Each of these issues should be explored in the same level of detail as stellar structure is studied in this work. With the resulting understanding of these processes, the coupling between them should then be determined.  Finally, we note that this paper has focused on the question of whether or not stars can exist in universe with alternate values of the relevant parameters. An important and more global question is whether or not these universes could also support life of some kind. Of course, such questions are made difficult by our current lack of an {\\it a priori} theory of life. Nonetheless, some basic requirements can be identified (with reasonable certainty): In addition to energy sources (provided by stars), there will be additional constraints to provide the right mix of chemical elements (e.g., carbon in our universe) and a universal solvent (e.g., water). These additional requirements will place additional constraints on the allowed region(s) of parameter space.  \\medskip  Acknowledgment: We thank Greg Laughlin for useful discussions. This work was supported by the Foundational Questions Institute through Grant RFP1-06-1 and by the Michigan Center for Theoretical Physics."
    },
    "0807/0807.1237_arXiv.txt": {
        "abstract": "We present new results from our survey of diffuse \\osix -emitting gas in the interstellar medium with the {\\it Far Ultraviolet Spectroscopic Explorer (FUSE).}  Background observations obtained since 2005 have yielded eleven new \\osix\\ detections of 3\\sig\\ significance, and archival searches have revealed two more.  An additional 15 sight lines yield interesting upper limits.  Combined with previous results, these observations reveal the large-scale structure of the \\osix -bearing gas in the quadrant of the sky centered on the Magellanic Clouds.  The most prominent feature is a layer of low-velocity \\osix\\ emission extending more than 70\\arcdeg\\ from the Galactic plane.  At low latitudes ($|b| < 30\\arcdeg$), the emission comes from narrow, high-density conductive interfaces in the local ISM.  At high latitudes, the emission is from extended, low-density regions in the Galactic halo.  We also detect \\osix\\ emission from the interface region of the Magellanic System, a structure recently identified from \\hone\\ observations.  These are the first detections of emission from high-ionization species in the Magellanic System outside of the Clouds themselves. ",
        "introduction": "The lithium-like \\osix\\ ion is an important tool for studying hot gas in interstellar space, because emission via its 1031.93 and 1037.62 \\AA\\ resonance lines is the dominant cooling mechanism for gas at temperatures of a few times $10^5$ K \\citep{Gnat:Sternberg:07}.  Recent \\osix\\ absorption-line surveys \\citep{Savage:03, Bowen:08} reveal that, within our Galaxy, the \\osix -bearing gas has an exponential scale height of 3.2 kpc at negative latitudes and 4.6 kpc at positive latitudes.  The gas is clumpy: spectra of nearly two dozen stars in the Large and Small Magellanic Clouds reveal significant variations in the \\osix\\ column density of Galactic halo gas over angular scales as small as $3\\farcm5$ \\citep{Howk:02}.  Closer to the plane, \\citet{Bowen:08} find that the average volume density of \\osix\\ towards stars in the Galactic disk, $n = N($\\osix$)/d$, exhibits a dispersion of about 0.26 dex, regardless of the distance to the target stars.  They conclude that the \\osix\\ does not arise in randomly distributed clouds of a uniform size and density, but is instead distributed among regions with a broad range of column densities, with smaller clouds more common and larger clouds less so.  They envision a turbulent, multiphase medium churned by shock-heated gas from multiple supernova explosions.  \\defcitealias{Otte:06}{Paper~I} \\defcitealias{Dixon:06}{Paper~II}  We present the latest results from a survey of diffuse \\osix\\ emission from the interstellar medium (ISM) conducted with the {\\it Far Ultraviolet Spectroscopic Explorer (FUSE)}.\\/  \\citet[hereinafter Paper~I]{Otte:06} and \\citet*[Paper~II]{Dixon:06} discuss observations obtained from launch through 2003 July 25 and 2004 December 31, respectively. In these papers, \\osix\\ emission at 3\\sig\\ significance is detected along 29 lines of sight.  Measured intensities range from 1.8 to 9.1 kLU, with a median of 3.3 kLU.  An additional 35 sight lines provide upper limits of 2.0 kLU or less.  (One line unit, or photon cm$^{-2}$ s$^{-1}$ sr$^{-1}$, corresponds to $1.9 \\times 10^{-11}$ erg cm$^{-2}$ s$^{-1}$ sr$^{-1}$ at 1032 \\AA.)  By combining emission and absorption measurements through the same \\osix -bearing regions, the authors find evidence for two flavors of \\osix -emitting gas: relatively narrow, high-density conductive interfaces in the local ISM and more extended, low-density regions in HVCs.  In this paper, we present data obtained between 2005 January 1 and 2007 July 6, along with two older observations that were missed in the earlier papers.  In \\S \\ref{sec_observations}, we describe the \\fuse\\/ instrument, our criteria for selecting sight lines for study, and our data-reduction and analysis techniques.  In \\S \\ref{sec_discussion}, we derive the physical parameters of the emitting gas and discuss its location along the line of sight.  We present our conclusions in \\S \\ref{sec_conclusions}.  Because the \\fuse\\/ pointing system failed in July of 2007, these observations represent the final installment of our survey. ",
        "conclusions": "We present the final installment of the \\fuse\\/ survey of diffuse \\osix\\ emission in the interstellar medium.  These observations reveal the large-scale structure of the \\osix -bearing gas in the quadrant of the sky centered on the Magellanic Clouds.  Its most prominent feature, a layer of low-velocity \\osix\\ emission extending more than 70\\arcdeg\\ from the Galactic plane, is produced by two distinct populations of emitting clouds.  At latitudes $|b| < 30\\arcdeg$, sight lines suffer significant extinction, suggesting that they probe relatively nearby emitting regions; their emission features are more likely to be narrow or unresolved, suggesting single interfaces less than about 0.15 pc across; and their velocities are generally consistent with Galactic rotation.  We identify these emitting regions with narrow, high-density interfaces in the local ISM.  At latitudes $|b| > 30\\arcdeg$, extinction is low.  Intensities fall with distance from the plane.  Emission features are generally broad, suggesting that they sample turbulent gas or multiple emitting regions spanning a range of velocities.  The \\osix -bearing gas seems to be moving away from the Galactic plane.  For two sight lines, Fairall~9 and NGC~625, we combine \\osix -emission and -absorption measurements to constrain the physical parameters of the emitting regions.  From the resulting path lengths, we estimate that these regions lie at least 800 pc from the Galactic plane.  The presence of nearby sight lines yielding null detections but strong upper limits suggests that the emission is patchy, consistent with the results of \\citet{Howk:02}.  High-velocity \\osix\\ emission is seen along three sight lines that probe the interface region of the Magellanic System.  The \\osix\\ velocities are generally consistent with those of the underlying \\hone.  High-velocity \\osix\\ emission is not seen in sight lines that probe the Magellanic Stream, suggesting that the dynamics of the interface region differ from those of the Stream itself.  These are the first detections of emission from high-ionization species in the Magellanic System outside of the Clouds themselves."
    },
    "0807/0807.3004_arXiv.txt": {
        "abstract": "{During Big Bang Nucleosynthesis (BBN), in the first $\\sim$20 minutes of the evolution of the Universe, the light nuclides, D, \\3he, \\4he, and \\7li were synthesized in astrophysically interesting abundances. The Cosmic Microwave Background Radiation (CMB) observed at present was last scattered some $\\sim$400 thousand years later.  BBN and the CMB (supplemented by more recent, $\\sim$10 Gyr, Large Scale Structure data), provide complementary probes of the early evolution of the Universe and enable constraints on the high temperature/energy physical processes in it.  In this overview the predictions and observations of two physical quantities, the baryon density parameter and the expansion rate parameter, are compared to see if there is agreement between theory and observation at these two widely separated epochs.  After answering this question in the affirmative, the consequences of this concordance for physics beyond the standard models of particle physics and cosmology is discussed.}  \\normalsize\\baselineskip=15pt ",
        "introduction": "Big Bang Nucleosynthesis (BBN) probes the evolution of the Universe during its first few minutes, providing a glimpse into its earliest epochs.  The photons in the presently observed Cosmic Microwave Background Radiation (CMB) were last scattered at``recombination\", when the plasma of ions and electrons became predominantly neutral, some 400 thousand years later.  Comparison of observations relating to BBN (the primordial abundances of the light elements) and the CMB (the temperature anisotropy spectrum) often provide unique tests of the consistency of the standard models of particle physics and cosmology, as well as enable the possibility of constraining alternate models of cosmology and of physics beyond the standard model of particle physics. In this article, based on my invited talk at the NO-VE IV International Workshop on: ``Neutrino Oscillations in Venice\", the complementarity between BBN and the CMB (supplemented by the Large Scale Structure (LSS) data needed to break some degeneracies) is explored, concentrating on the universal baryon asymmetry and the expansion rate of the early, radiation-dominated Universe.  In particular, do observations of the baryon density parameter and of the expansion rate parameter (both to be defined in more detail below) agree at $\\sim$20 minutes and $\\sim$400 thousand years and, if so, what does this tell us about non-standard physics and/or cosmology in the time interval between these two, so widely separated, epochs?  After introducing and defining the cosmological parameters in \\S1.1 and \\S1.2, their influence on the predictions of BBN (\\S2) and the CMB observations (\\S3) is discussed and compared to the standard model predictions and, with the observational data (\\S4).  After establishing the concordance of the standard models of particle physics and cosmology, some general constraints on possible new physics in the interval between BBN and recombination are presented in \\S5.  The results are summarized in \\S6.  In my recent review of BBN\\cite{steigman07} these issues are discussed in more detail (including a more extensive list of references).  My talk at NO-VE IV and this article are drawn from the recently published paper with V. Simha\\cite{vs1}.  \\subsection{The Baryon Density Parameter: $\\eta_{\\rm B}$}  During the much earlier evolution of the Universe than is considered here, a universal matter-antimatter (baryon-antibaryon) asymmetry was established by particle physics processes yet to be uniquely determined. Thereafter, in the evolution of the Universe up to the present, the number of baryons in a comoving volume is preserved, although the number density of baryons decreases as the Universe expands.  During the same early evolution of the Universe, very rapid electromagnetic processes establish and maintain a Bose-Einstein (black body) distribution for the cosmic background photons.  This guarantees that when the photons are in equilibrium or, when they are entirely decoupled, the number of CMB photons in a comoving volume is preserved (except for the new photons added when various particle-antiparticle pairs annihilate or, when unstable particles decay).  In the standard models of cosmology and particle physics the numbers of baryons and of CMB photons in a comoving volume have been unchanged since \\epm anniliation.  The ratio of the numbers of baryons and CMB photons in a comoving volume in the post-\\epm annihilation Universe provides us with a dimensionless, time-invariant parameter -- in the context of the {\\bf standard} models of particle physics (baryon conservation) and cosmology (entropy conservation). \\be \\eta_{\\rm B} \\equiv n_{\\rm B}/n_{\\gamma} \\equiv 10^{-10}\\eta_{10}. \\ee $\\eta_{10}$ is related to the baryon mass density parameter $\\Omega_{\\rm B}$, the ratio of the baryon mass density to the critical mass density, and the reduced, present value of the Hubble parameter $h$ (the Hubble ``constant\": $H_{0} \\equiv 100h$~km~s$^{-1}$~Mpc$^{-1}$) by\\cite{steigman06}, \\be \\eta_{10} = 273.9~\\Omega_{\\rm B}h^2. \\ee $\\eta_{10}$ ``measured\" at BBN ($\\sim$20 minutes) and inferred from the CMB ($\\sim$400 thousand years later) should agree, enabling constraints to be placed on non-standard entropy production and/or baryon number violation.  \\subsection{The Universal Expansion Rate Parameter: $S$}  In the standard model of cosmology, the expansion rate is described by the Hubble parameter ($H$) which, in the early Universe, is related to the total energy density through the Friedman equation by, \\be H^{2} = 8\\pi G\\rho_{\\rm TOT}/3. \\label{hubble} \\ee During its early evolution, the Universe is ``radiation dominated\" ($\\rho_{\\rm TOT} \\rightarrow \\rho_{\\rm R}$); that is, the total energy density is dominated by massless particles (\\eg photons) or relativistic, massive particles ($m~\\la T$).  Prior to \\epm annihilation, for temperatures $T~\\la$~few MeV, the standard-model relativistic particles present are photons, \\epm pairs, and three flavors of left-handed neutrinos along with their right-handed antineutrinos.  Counting degrees of freedom and accounting for the relativistic bosons and fermions, \\be \\rho_{\\rm R} = \\rho_{\\gamma} + \\rho_{e} + 3\\rho_{\\nu} = (1 + 7/8(2 + 3))\\rho_{\\gamma} = 43\\rho_{\\gamma}/8. \\ee  As may be seen from eq.~\\ref{hubble}, any modification to the Friedman equation will change the expansion rate ($H \\rightarrow H' \\equiv SH$). For example, a non-standard, early-Universe expansion rate may be due to the presence of a non-standard energy density of relativistic particles ($\\rho_{\\rm R}' \\neq \\rho_{\\rm R}$), or by a non-standard value of the gravitational constant $G' \\neq G$), or by some form of new physics beyond the standard model which modifies the Friedman equation.  A non-standard particle content may be conveniently parameterized by the equivalent number of ``additional\" neutrinos defined, prior to \\epm annihilation, by $\\Delta$N$_{\\nu} \\equiv$ N$_{\\nu} - 3$, so that $\\rho_{\\rm R}' \\equiv \\rho_{\\rm R} + \\Delta$N$_{\\nu}\\rho_{\\nu}$. \\Deln and the expansion rate parameter $S$ are related by\\cite{steigman07}, \\be S^{2} \\equiv \\left({H' \\over H}\\right)^{2} = {\\rho_{\\rm R}' \\over \\rho_{\\rm R}} = {G' \\over G} \\equiv 1 + {7\\Delta{\\rm N}_{\\nu} \\over 43} = 1 + 0.163\\Delta{\\rm N}_{\\nu}. \\ee However, {\\bf any} change in $H$ ($S \\neq 1$) which is traceable to a term evolving like the radiation density (as the inverse fourth power of the scale factor) can equally well be parameterized by $\\Delta$N$_{\\nu}$.  In this sense, \\Deln $\\neq 0$ should be thought of as a proxy for any non-standard, early-Universe expansion rate ($S \\neq 1$) and need have {\\bf nothing at all} to do with ``extra\" (or fewer!) neutrinos.  After \\epm annihilation the only relativistic particles present during the radiation dominated epoch are the photons (which redshift to become the presently observed CMB) and the neutrinos, which decoupled prior to \\epm annihilation.  In the approximation that the neutrinos were {\\bf fully} decoupled at \\epm annihilation, the post-annihilation photons are hotter than the neutrinos by a factor of $T_{\\gamma}/T_{\\nu} = (11/4)^{1/3}$ and\\cite{steigman07} \\be S^{2} = \\rho'_{\\rm R}/\\rho_{\\rm R} \\rightarrow 1 + 0.135 \\Delta {\\rm N}_{\\nu}, \\ee where, in the post-BBN, pre-recombination Universe, $\\Delta$N$_{\\nu} \\equiv$~N$_{\\nu} - 3$.  Since the standard-model neutrinos were {\\bf not} fully decoupled at \\epm annihilation, they share some of the energy/entropy released during \\epm annihilation\\cite{dicus}.  As a result, they are warmer than is predicted by the fully decoupled approximation, increasing the ratio of the post-\\epm annihilation radiation density to the photon energy density.  In the standard model, this additional contribution to the total energy density can be accounted for by replacing \\Nnu = 3 with \\Nnu = 3.046\\cite{mangano}.  Any post-BBN deviations from the standard model that can be treated as equivalent to contributions from {\\bf fully decoupled} neutrinos can thus be accounted for by\\cite{steigman07}, \\be S^{2} = \\rho'_{\\rm R}/\\rho_{\\rm R} \\rightarrow 1 + 0.134\\Delta {\\rm N}_{\\nu}, \\ee where in the post-\\epm annihilation Universe relevant for comparison with the CMB and LSS, $\\Delta$N$_{\\nu} \\equiv$~N$'_{\\nu} - 3.046$. Note that with these definitions, the standard model corresponds to \\Deln = 0 in {\\bf both} the pre- and post-\\epm annihilation Universe. ",
        "conclusions": "In the standard models of cosmology and of particle physics the particle content is fixed (\\eg \\Nnu = 3) and, baryon number has been conserved since the very earliest epochs.  By comparing BBN with the CMB and LSS, baryon conservation can be tested between $\\sim$20 minutes (BBN) and $\\sim$400 thousand years.  So, too, can deviations from the standard-model particle content (\\Nnu $\\neq 3$?), as well as deviations from the standard-model predicted early Universe expansion rate.  In this talk and review, $\\eta_{\\rm B}$ evaluated from BBN and from the CMB have been compared and shown to be in excellent agreement, consistent with the standard model expectation.  BBN and the CMB have also been employed to compare \\Nnu (or, $S$) at these two, widely separated epochs and, to compare them to the standard-model expectation that \\Nnu = 3 ($S=1$).  While the central values of \\Nnu evaluated at BBN and from the CMB differ from each other and from the standard-model value, they are, in fact, in agreement with each other, and with the standard-model value within the uncertainties.  This concordance of the standard models of particle physics and cosmology permit constraints on ``new\" physics at the times (and energies) between BBN and the CMB (and, at present).  Prior to WMAP and the other ground- and balloon-based CMB experiments, BBN and deuterium provided the best cosmological baryometer, while \\4he provided the only early-Universe chronometer.  The WMAP 5-year data, combined with other CMB and LSS data, now lead to a determination of the baryon density at the $\\sim$2-3\\% level, a factor of $\\sim$2 better than that from BBN.  However, although the CMB/LSS constraint on \\Nnu has improved significantly and is consistent with that from BBN, it still remains weaker than the corresponding BBN constraint.  For some time now BBN has established at high confidence that \\Nnu $> 1$ at BBN, when the Universe was $\\sim$20 minutes old.  The more recent CMB data, combined with other SNIa and LSS datasets, now confirm that \\Nnu $> 1$ (or, \\Nnu $> 2$~\\cite{dunkley08}) when the Universe was $\\ga 400$~kyr old.  It may be of interest to some readers that, since the NO-VE IV Workshop, V. Simha and I have presented the results of a similar analysis, updating the constraints from BBN and the CMB on the lepton asymmetry parameter\\cite{vs2}."
    },
    "0807/0807.2414_arXiv.txt": {
        "abstract": "Single field inflationary models with a non-minimal kinetic term (also called k-inflationary models) can be characterised by the so-called sound flow functions, which complete the usual Hubble flow hierarchy. These parameters appear in the primordial power spectra of cosmological perturbations at leading order and, therefore, affect the resulting Cosmic Microwave Background (CMB) anisotropies. Using the fifth year Wilkinson Microwave Anisotropy Probe (WMAP5) data, we derive the marginalised posterior probability distributions for both the sound and Hubble flow parameters. In contrast to the standard situation, these parameters remain separately unbounded, and notably there is no longer any upper limit on $\\epsilon_1$, the first Hubble flow function. Only special combinations of these parameters, corresponding to the spectral index and tensor-to-scalar ratio observables, are actually constrained by the data. The energy scale of k-inflation is nevertheless limited from above to $\\Hinf \\le 6 \\times 10^{-6} \\mpl$ at two-sigma level. Moreover, for the sub-class of Dirac--Born--Infeld models, by considering the non-gaussianity bounds on the sound speed, we find a weak limit $\\epsilon_1 < 0.08$ at $95\\%$ confidence level. ",
        "introduction": "\\label{sec:intro}  Accelerated expansion of the early Universe in scalar field models is usually achieved by domination of the field's potential over its kinetic energy. However, it has been shown that inflation can also be supported through modifications of the kinetic term, these models being known as ``k-inflation'' scenarios~\\cite{ArmendarizPicon:1999rj}. Assuming the gravity sector to be described by General Relativity [the signature of the metric tensor being $(-,+,+,+)$], the action describing such a ''k-inflaton'' $\\varphi(x^\\mu)$ reads \\begin{equation} S=\\dfrac{1}{2\\kappa } \\int \\ud^4x \\sqrt{-g} \\left[ R + 2\\kappa P\\left(X,\\varphi\\right) \\right], \\end{equation} where $\\kappa \\equiv 8\\pi/\\mpl^2$ and where $X$ is defined by the following expression \\begin{equation} X \\equiv -\\dfrac{1}{2} g^{\\mu \\nu} \\partial_\\mu \\varphi \\partial_\\nu \\varphi . \\end{equation} The quantity $P(X,\\varphi)$ is any acceptable functional (see below) for the scalar field Lagrangian. For instance, the canonical scalar field model is recovered for $P(X,\\varphi) = X - V(\\varphi)$, where $V(\\varphi)$ is the scalar field potential.  \\par  The theory of k-inflation is interesting not only because it enlarges the class of inflationary models that one can consider but also because this is precisely this type of scenarios that emerges from string theory. This is the case of the so-called Dirac--Born--Infeld (DBI) scenarios of brane inflation, which have recently received particular attention as effective field theory realizations of inflation within string theory~\\cite{Dvali:1998pa, Kachru:2003sx, HenryTye:2006uv, Hertzberg:2007ke, Bean:2007eh, Peiris:2007gz, Bean:2007hc, Lorenz:2007ze}. The DBI inflaton typically represents the distance (\\ie, an open string mode) between branes in a higher-dimensional bulk, and the Lagrangian functional is given by the determinant of the induced metric along the moving brane, \\begin{equation} \\label{eq:Pdbi} \\Pdbi(X,\\varphi) = - T(\\varphi) \\sqrt{1 -\\dfrac{2 X}{T(\\varphi)}} + T(\\varphi)  - V(\\varphi). \\end{equation} Here, $T(\\varphi)$ is a warping function encoding the bulk geometry, and $V(\\varphi)$ is the interaction potential between the branes, assembled from contributions of various origin.  \\par  In general, a model of k-inflation is sensible if the function $P(X,\\varphi)$ satisfies the two following conditions. Firstly, $\\deriv{P}{X}>0$, coming from the requirement that the Hamiltonian is bounded from below and, secondly, $\\deriv{P}{X}+2X \\dderiv{P}{X}{X}>0$, which is necessary in order for the Cauchy problem to be well-posed, see for instance Ref.~\\cite{Bruneton:2007si}. Let us notice that the DBI models satisfy these requirements. Moreover, these two conditions play an important role at the perturbative level. Indeed, all kinetically modified inflationary models share the common feature that the ``speed of sound'' \\begin{equation} \\label{eq:csdef} \\cs^2 \\equiv \\dfrac{\\deriv{P}{X}}{\\deriv{P}{X} + 2 X \\dderiv{P}{X}{X}}\\,, \\end{equation}is generically different from the speed of light and may be space-time dependent. Clearly, the two above-mentioned inequalities guarantee that $\\cs^2>0$. This is relevant for the treatment of cosmological perturbations. Indeed, for a flat Friedmann-Lema\\^{\\i}tre--Robertson--Walker (FLRW) universe, it was shown in Ref.~\\cite{Garriga:1999vw}, that the Mukhanov--Sasaki mode function $\\qmode (k,\\eta)= \\zeta \\sqrt{2\\kappa \\epsilon_1}/\\cs$ ($\\zeta$ being the comoving curvature perturbation) satisfies the modified equation of motion \\begin{equation} \\label{eq:qmode} \\qmodeS'' + \\left[\\cs^2 k^2 - \\dfrac{(a \\sqrt{\\epsilon_1}/\\cs)''}{a \\sqrt{\\epsilon_1}/\\cs} \\right] \\qmodeS = 0. \\end{equation} Derivatives are with respect to conformal time $\\eta$, $a(\\eta)$ is the FLRW scale factor and $\\epsilon_1 = -\\ud \\ln H/\\ud \\ln a$ is the first Hubble flow function ($H$ being the Hubble parameter). In addition to the induced modification of the perturbations' effective potential, $\\cs$ also appears as their propagation speed. Let us notice that for DBI inflation, Eqs.~(\\ref{eq:Pdbi}) and~(\\ref{eq:csdef}) give \\begin{equation} \\label{eq:csdbi} \\cs = \\sqrt{1 - \\dfrac{{\\varphi'}^2}{a^2 T(\\varphi)}}\\, , \\end{equation} and therefore ensure that $\\cs\\le1$. In the following, we will only consider the k-inflation models for which $\\cs \\le 1$ (see, however, Refs.~\\cite{Bruneton:2006gf, Bonvin:2006vc, Ellis:2007ic,Babichev:2007dw, Bonvin:2007mw, Kang:2007vs}).  \\par  In this paper, we are interested in constraining the shape of the scalar and tensor primordial power spectra for all the kinetically modified models admitting an observable ``slowly-evolving'' phase. To this end, we assume the existence of a background solution supporting a long enough inflationary era to solve the FLRW problems.  \\par  The fact that $\\cs$ is time-dependent and appears in front of the $k^2$ term in Eq.~(\\ref{eq:qmode}) has important consequences and brings new complications into the problem. In particular, one can no longer use the standard techniques consisting in solving the mode equation in terms of Bessel functions. In the case of k-inflation, one is therefore forced to rely on a new method of approximation. This is done in Ref.~\\cite{Kinney:2007ag} and in Ref.~\\cite{Lorenz:2008xx}, this last article making use of the Wentzel--Kramers--Brillouin (WKB)/uniform approximation~\\cite{Habib:2002yi, Habib:2004kc,Martin:2002vn}. In this manner, one can solve Eq.~(\\ref{eq:qmode}) perturbatively in terms of a double hierarchy of functions~\\cite{Alishahiha:2004eh, Chen:2006nt, Kinney:2007ag, Peiris:2007gz, Lorenz:2008xx} defined by \\begin{eqnarray} \\label{eq:hierarchies} \\begin{aligned} \\epsilon_{i+1} & = \\dfrac{\\ud \\ln |\\epsilon_i|}{\\ud \\ln a}\\,, \\quad \\epsilon_0 = \\dfrac{H_\\uini}{H}\\,, \\\\ \\delta_{i+1} & = \\dfrac{\\ud \\ln |\\delta_i|}{\\ud \\ln a}\\,, \\quad \\delta_0 = \\dfrac{{\\cs}_\\uini}{\\cs}\\,. \\end{aligned} \\end{eqnarray} The $\\epsilon_i$ are the usual Hubble flow functions introduced in Refs.~\\cite{Schwarz:2001vv,Leach:2002ar,Schwarz:2004tz}, while the $\\delta_i$ are their equivalent based on the sound speed and therefore encode its rate of change. At first order in these parameters, the scalar and tensor primordial power spectra\\footnote{For single field models, $\\zeta$ is a conserved quantity on ``super-sonic'' length scales. As in the standard case~\\cite{Martin:1997zd}, this allows us to propagate the primordial spectrum from the end of inflation to the time of decoupling.} can be expressed as~\\cite{Lorenz:2008xx} \\begin{equation} \\label{eq:powerspectra} \\begin{aligned} \\calP_\\zeta & = \\dfrac{\\Hinf^2}{\\pi \\mpl^2 \\epsilon_1 \\cs} \\left[1 - 2(D+1) \\epsilon_1 - D \\epsilon_2 + (D+2) \\delta_1 \\phantom{\\dfrac{k}{\\kpiv}}\\right. \\\\ & - \\left. \\left(2 \\epsilon_1 + \\epsilon_2 - \\delta_1 \\right) \\ln \\dfrac{k}{\\kpiv} \\right], \\\\ \\calP_h & = \\dfrac{16 \\Hinf^2}{\\pi \\mpl^2} \\left[1 - 2(D+1-\\ln \\cs) \\epsilon_1 -2 \\epsilon_1 \\ln \\dfrac{k}{\\kpiv} \\right]. \\end{aligned} \\end{equation} The wavenumber $\\kpiv$ is a pivot scale around which these power spectra are expanded, see Ref.~\\cite{Lorenz:2008xx}. All the functions appearing in this equation have to be evaluated at the conformal time $\\etapiv$ at which this pivot mode crossed the sound horizon, namely the solution of \\begin{equation} \\kpiv \\etapiv  = - \\dfrac{1}{\\cs(\\etapiv)}\\,. \\end{equation} Notice that although the tensor perturbations are not affected by the non-minimal kinetic terms in the scalar sector, $\\cs$ nevertheless appears in $\\calP_h$ due to the evaluation of the flow functions at the sound horizon crossing time instead of Hubble crossing. In Eq.~(\\ref{eq:powerspectra}), $D = 1/3 - \\ln 3\\simeq -0.76$ is a constant coming from the uniform approximation, and an overall constant factor of $18 \\exp(-3)$ has been absorbed into $\\Hinf^2$~\\cite{Martin:2002vn}. From Eq.~(\\ref{eq:powerspectra}) we can read off the spectral indices to first order, \\begin{equation} \\label{eq:indices} \\begin{aligned} \\nS-1&=-2\\epsilon_{1}-\\epsilon_{2}+\\delta_{1}\\, ,\\quad \\nT=-2\\epsilon_{1}, \\end{aligned} \\end{equation} and the k-inflation tensor to scalar ratio $r=16 \\epsilon_1 \\cs$.  \\par  As expected, the power spectra in Eq.~(\\ref{eq:powerspectra}) are not the same than those obtained for a canonically normalised scalar field due to the presence of the two sound flow parameters $\\cs$ and $\\delta_1$. Notice that in the limit $\\cs = 1 - \\order{\\epsilon_i}$, one has $\\delta_1 = \\order{\\epsilon_i^2}$ and the above spectra only differ from their usual analogue at second order in the Hubble flow functions~\\cite{Steer:2003yu}.  \\par  In the following, we compare the induced Cosmic Microwave Background (CMB) anisotropies seeded by these power spectra to the WMAP5 data~\\cite{Gold:2008kp, Hill:2008hx, Hinshaw:2008kr, Nolta:2008ih, Dunkley:2008ie, Komatsu:2008hk}. Using Markov--Chains--Monte--Carlo (MCMC) methods, we extract the posterior marginalised probability distributions on the sound and Hubble flow parameters and discuss the physical consequences. Moreover, as shown in Refs.~\\cite{Seery:2005wm, Chen:2006nt, Bean:2008ga, Gauthier:2008mq}, kinetically modified inflation models may induce a large amount of non-gaussianities in the statistics of the CMB anisotropies, usually measured in terms of the $\\fNL$ parameter~\\cite{Komatsu:2001rj}. The expected value for $\\fNL$ strongly depends on the model at hand: as shown in Ref.~\\cite{Chen:2006nt}, $\\fNL$ for equilateral configurations demands the knowledge of the third $X$-derivative of $P(X,\\varphi)$ in addition to the sound and Hubble flow functions. For these reasons, we have chosen to split our analysis into two parts. In the first section, we do not consider the non-gaussianity constraints, and our results are only based on the observed two-point functions of the CMB temperature and polarisation anisotropies. In the second part, we implement the WMAP5 bounds on $\\fNL$ as prior knowledge on the sound speed values in the restricted class of DBI models: for those, we indeed have the universal expression~\\cite{Chen:2006nt,Gao:2008dt, Langlois:2008qf} \\begin{equation} \\label{eq:fnl} \\fNL=\\dfrac{35}{108}\\left(1-\\dfrac{1}{\\cs^2}\\right ) + \\order{\\epsilon_i,\\delta_i}. \\end{equation} {}From now on, $\\fNL$ is implicitly assumed to refer to equilateral configurations, the ``local'' estimator being $\\fNL^{\\mathrm{local}} \\sim \\order{\\epsilon_i,\\delta_i}$ and well below the current detection thresholds, see Ref.~\\cite{Chen:2006nt, Komatsu:2008hk, Bean:2008ga}. ",
        "conclusions": "We have used the WMAP fifth year data to constrain all single field k-inflationary models admitting a phase of ``slow-rolling'' evolution during the generation of the cosmological perturbations.  \\par  The probability distributions of the Hubble and sound flow parameters, $\\epsilon_i$ and $\\delta_i$, evaluated when the scalar perturbations crossed the sound horizon, are significantly different than the ones associated with canonically normalised single field models. The presence of two additional parameters in the primordial power spectra, $\\cs$ and $\\delta_1$, leads to new degeneracies in the parameter space: there is no longer any constraint on $\\epsilon_1$ and $\\epsilon_2$ separately, neither on $\\cs$ or $\\delta_1$ alone. The energy scale of inflation is nevertheless upper limited: $\\Hinf/\\mpl \\le 6\\times 10^{-6}$ at $95\\%$ of confidence. In fact, only two particular combinations of the primordial parameters end up being constrained by the CMB power spectra: the quantity $\\epsilon_1\\cs$ in Eq.~(\\ref{eq:cseps1twosig}), which encodes the tensor-to-scalar ratio in the CMB spectra, and $2(\\epsilon_1-\\delta_1) + (\\epsilon_2 + \\delta_1)$, which fixes the scalar spectral index in Eq.~(\\ref{eq:nstwosig}). Under our prior assumptions, Eq.~(\\ref{eq:nstwosig}) shows that the data favour a ``red-tilted'' scalar power spectrum, at two-sigma level, with tensor modes and varying sound speed.  \\par  Implementing the current WMAP5 bound on non-gaussianity through the $\\fNL$ parameter (for equilateral configurations) breaks some of the above-mentioned degeneracies. For the DBI inflationary models, more than restricting the acceptable values of the sound speed $\\cs$, we find that the non-gaussianity bound induces an upper limit on the first Hubble flow parameter: $\\log(\\epsilon_1) < -1.1$ at two-sigma confidence level.  \\par  The present results could be improved by various means. A first step may be to add more data sets, but according to the previous discussion, this will certainly improve only the limits on the spectral index and tensor-to-scalar ratio combinations without breaking the degeneracies between the $\\epsilon_i$ and $\\delta_i$. Therefore, it is clearly more promising to improve the tests that bring new information on non-gaussianities. We have implemented them for DBI models only through a prior knowledge over $\\cs$. Clearly, our approach could be improved by directly evaluating the three-point functions and comparing it to the observed one, provided one can define an tractable likelihood for MCMC exploration. Such an approach is currently technically difficult. Evaluating the three-point functions from CMB data is not trivial (see Appendix A in Ref.~\\cite{Komatsu:2008hk}); and, from a theoretical point of view, new parameters have to be introduced to quantify the non-gaussian signals in a generic way for k-inflation. Therefore, one faces the risk to add even more degeneracies. On the other hand, with the incoming flow of data on cosmological observables, this represents an interesting challenge for the forthcoming years."
    },
    "0807/0807.0411_arXiv.txt": {
        "abstract": "{ We present \\emph{Chandra X-ray Observatory} ACIS-S3 X-ray imaging observations and VLT/FORS2 and \\emph{Hubble Space Telescope} optical observations of two low-density Galactic globular clusters; NGC\\,6366 and M55.  We detect 16 X-ray sources with 0.5--6.0\\,keV luminosities above $L_\\mathrm{X}=4\\times10^{30}$\\,erg\\,s$^{-1}$ within the half-mass radius of M55, of which 8 or 9 are expected to be background sources, and 5 within the half-mass radius of NGC\\,6366, of which 4 are expected to be background sources. Optical counterparts are identified for several X-ray sources in both clusters and from these we conclude that 3 of the X-ray sources in M55 and 2 or 3 of the X-ray sources in NGC\\,6366 are probably related to the cluster. Combining these results with those for other clusters, we find the best fit for a predicted number of X-ray sources in a globular cluster $\\mu_\\mathrm{c}=1.2\\Gamma+1.1M_\\mathrm{h}$, where $\\Gamma$ is the collision number and $M_\\mathrm{h}$ is (half of) the cluster mass, both normalized to the values for the globular cluster M4. Some sources tentatively classified as magnetically active binaries are more luminous in X-rays than the upper limit of $L_\\mathrm{X}\\simeq0.001L_\\mathrm{bol}$ of such binaries in the solar neighbourhood. Comparison with \\textit{XMM} and \\textit{ROSAT} observations lead us to conclude that the brightest X-ray source in M55, a dwarf nova, becomes fainter in X-rays during the optical outburst, in accordance with other dwarf novae. The brightest X-ray source in NGC\\,6366 is a point source surrounded by a slightly offset extended source. The absence of galaxies and H$\\alpha$ emission in our optical observations argues against a cluster of galaxies and against a planetary nebula, and we suggest that the source may be an old nova. ",
        "introduction": "All stars emit X-rays, but some emit more than others.  Thanks to the \\emph{Chandra} X-ray Observatory the study of X-ray sources in globular clusters now includes sources down to luminosities of $L_\\mathrm{X}\\sim10^{29-30}$\\,erg\\,s$^{-1}$ in the 0.5--2.5\\,keV band.  At these low luminosities, most sources are low-mass main-sequence stars that rotate rapidly, which in old stellar clusters is the case only when they have been spun up by tidal forces or by accretion in close binaries.  At the high luminosity end, $L_\\mathrm{X}\\gtap10^{32}$\\,erg\\,s$^{-1}$, most sources are low-mass X-ray binaries in which a neutron star accretes matter from a companion.  At intermediate luminosities, most sources are cataclysmic variables (CVs).  The clusters best studied at low X-ray luminosities are 47\\,Tuc (\\citealt{hge+05} and references therein) and M4 \\citep{bph+04}.  Whereas the binaries of main-sequence stars are primordial, i.e.\\ formed as binaries when the component stars formed, the low-mass X-ray binaries in globular clusters are thought to be formed in close stellar encounters that bring a previously single neutron star into a binary (see the review by \\citealt{vl06} and references therein).  For cataclysmic variables both formation mechanisms, evolution from a primordial binary and capture of a previously single white dwarf in a close encounter, are viable, depending on the circumstances. The progenitor binary of a cataclysmic variable must be wide enough to allow the more massive star to evolve into a fairly big giant, before it reaches its Roche lobe. Such wide binaries are destroyed (`ionized') in cores of globular clusters with densities $\\rho_0\\gtap10^3\\,M_\\odot$\\,pc$^{-3}$ \\citep{dav97}.  One would naively expect that all cataclysmic variables in these cores were formed in stellar encounters.  However, it has been pointed out that cataclysmic variables that evolved in the low-density outskirts of globular clusters can sink towards the core at late times.  Thus, the population of cataclysmic variables in dense cores can be a mixture of locally produced products of stellar encounters and recently arrived products of the evolution of primordial binaries (\\citealt{ihr+06}; see also \\citealt{has07}).  The number of stellar encounters in a globular cluster scales roughly with the collision number $\\Gamma\\equiv{\\rho_0}^2{r_\\mathrm{c}}^3/v$, where $\\rho_0$ is the central density of the cluster, $r_\\mathrm{c}$ the core radius, and $v$ the central velocity dispersion.  Through the virial theorem $v\\propto r_\\mathrm{c}\\sqrt{\\rho_0}$ and thus $\\Gamma\\propto {\\rho_0}^{1.5}{r_\\mathrm{c}}^2$ \\citep{vh87,ver03}. The mass $M$ of a cluster can be estimated from the total luminosity of a cluster, with use of a mass-to-light ratio appropriate for a cluster star population, and if necessary a bolometric correction.  The number of X-ray sources in a globular cluster then may be expected to depend on the collision number (dynamical origin) and on the cluster mass (primordial origin); the form of the dependence depends on the type of binary, and therefore on the luminosity range that is studied.  The brightest low-mass X-ray binaries, with $L_\\mathrm{X}>10^{35}$\\,erg\\,s$^{-1}$, are sufficiently rare that a cluster contains either none or one; in one case (M15) two. The probability that a cluster contains such a source scales with $\\Gamma$ \\citep{vh87}.  The fainter low-mass X-ray binaries, with $10^{32}\\ltap L_\\mathrm{X}\\ltap10^{35}$\\,erg\\,s$^{-1}$, are more numerous; the number of them in a cluster also scales with $\\Gamma$ \\citep{hgl+03,gbw03b,pla+03}.  These results appear to confirm the theoretical idea that neutron stars enter binaries in globular clusters via stellar encounters. Studies of globular clusters within other galaxies complicate the issue, because they indicate that the probability that a cluster contains a bright source depends on its metallicity, being higher for metal-rich clusters (e.g.\\ \\citealt{kmz02}), and furthermore scales with ${\\rho_0}^{1.5}$, a shallower dependence on density than implied by a proportionality with $\\Gamma$ \\citep{jfb+04}.  The numbers of less luminous sources, mostly cataclysmic variables -- hitherto only observed in globular clusters of our Milky Way -- appear to have a more shallow dependence on central density $N\\propto{\\rho_0}^\\gamma$ with $\\gamma=0.6-0.7$ \\citep{hgl+03,pla+03}. \\citet{hwc+06} find different values of $\\Gamma$ for bright and faint cataclysmic variables, with the boundary near $10^{32}$\\,erg\\,s$^{-1}$.  Finally, \\citet{kbp+06} show that the number of sources (2-5 only!)  in the low-density globular cluster NGC\\,288, a mixture of cataclysmic variables and binaries of main-sequence stars, scales with the mass of this cluster. \\citet{ph06} successfully describe the number $N$ of X-ray sources in globular clusters with an equation of the form $N=a M + b \\Gamma$.  There are many uncertainties in the theoretical description of the formation of X-ray sources in globular clusters. We mention only the retained fraction of neutron stars; and the details of the spiral-in process which changes the wide orbit of the progenitor binary into the close orbit of a cataclysmic variable.  To help untangle these uncertainties empirically we have observed a variety of globular clusters both in X-rays and in the optical.  We have selected these clusters to span a wide range in central density, core radius, and mass. In this paper we discuss our observations of two clusters with relatively large core radii and low central densities: M55 and NGC\\,6366. Some parameters of both clusters are given in Tables\\,\\ref{t1} and \\ref{t6}.  \\begin{table} \\begin{minipage}[t]{\\columnwidth} \\centering \\caption{Parameters of M55 and NGC\\,6366 used in this paper; core radius $r_\\mathrm{c}$, half-mass radius $r_\\mathrm{h}$, distance $d$ and reddening $E_{B-V}=A_V/3.1$ (taken from \\citealt{har96}, version Feb 2003). The interstellar column $N_\\mathrm{H}$ is computed from $N_\\mathrm{H}=1.79\\times10^{21}\\,\\mathrm{cm}^{-2}\\,A_V$, see \\citet{ps95}. \\label{t1}} \\begin{tabular}{lccccc} \\hline \\hline Cluster & $r_\\mathrm{c}$ (\\arcmin) & $r_\\mathrm{h}$ (\\arcmin) & $d$ (kpc) & $E_{B-V}$ & $N_\\mathrm{H}$ (cm$^{-2}$) \\\\ \\hline M55 & 2.83 & 2.89 & 5.3 & 0.08 & $4.44\\times10^{20}$ \\\\ NGC\\,6366 & 1.83 & 2.63 & 3.6 & 0.71 & $3.94\\times10^{21}$ \\\\ \\hline \\end{tabular} \\end{minipage} \\end{table}  Both clusters were observed with the \\emph{ROSAT} PSPC, which detected one source in the core of each cluster \\citep{jvh96}. The position of the source in M55 was more accurately determined with a \\emph{ROSAT} HRI observation \\citet{ver01}, and this source was later optically identified with a dwarf nova \\citep{kpt+05}. \\emph{XMM} detected five sources in the core of M55, one of them coincident with the \\emph{ROSAT} source \\citep{wwb06}. ",
        "conclusions": "\\subsection{Comparison with earlier observations} Both M55 and NGC\\,6366 have been observed with \\emph{ROSAT}, and \\emph{XMM} also observed M55.  \\subsubsection{M55}\\label{s511} In Fig.\\,\\ref{f8} we compare the countrates of sources in M55 observed with \\emph{Chandra} (from Table\\,\\ref{t2}) with those observed with \\emph{XMM} EPIC (from Table\\,2 of \\citealt{wwb06}). All \\emph{XMM} sources within the field-of-view of \\emph{Chandra} have been detected; the \\emph{XMM} sources are the sources with the highest countrates in Chandra, plus CX29 (XMM45). Comparison of the \\emph{Chandra} and \\emph{XMM} source positions shows that the error estimates by Webb et al.\\ for the \\emph{XMM} source positions are overestimated by about a factor 2.  As expected, the \\emph{XMM} countrate for most sources is $\\sim4$ times higher than the \\emph{Chandra} countrate: for example, a 2\\,keV bremsstrahlung spectrum absorbed by the column towards M55, with 0.001 cts\\,s$^{-1}$ in \\emph{Chandra} ACIS-S will produce about 0.0026\\,cts\\,s$^{-1}$ on the PN and 0.00075\\,cts\\,s$^{-1}$ on each of the two MOS detectors, for a total of 0.0041 cts\\,s$^{-1}$ in \\emph{XMM}\\footnote{We estimate this with the PIMMS tool at {\\tt http://heasarc.gsfc.nasa.gov/Tools/w3pimms.html}; we have multiplied the countrates within a $15\\arcsec$ extraction radius as given by PIMMS with 1.47 to estimate the total countrates}.  \\begin{figure} \\resizebox{0.95\\hsize}{!}{\\includegraphics[angle=270]{9350f8.pdf}} \\caption{Detected countrate in the 0.5--6.0\\,keV energy band ($X_\\mathrm{soft}+X_\\mathrm{hard}$, Table\\,\\ref{t2}) of the \\emph{Chandra} ACIS-S observation as a function of the \\emph{XMM} EPIC (PN and both MOS detectors) countrate in the 0.5--10\\,keV energy band (Table\\,2 of \\citealt{wwb06}). The solid and dashed lines show equal countrates and 4 times higher/lower countrates, respectively. For most sources, the \\emph{XMM} countrate is about 4 times higher than the \\emph{Chandra} countrate, as expected for a constant source.} \\label{f8} \\end{figure}  Remarkably, all three sources with the lowest \\emph{XMM} count rates were brighter during the \\emph{Chandra} observation. One of these is the cataclysmic variable CV1, our source CX1. \\citet{wwb06} state that this source was three times fainter, in the 0.5--10\\,keV band, during the \\emph{XMM} observation than during earlier \\emph{ROSAT} observations.  Since the \\emph{ROSAT} measurements are limited to the 0.2--2.5\\,keV band, we prefer to make the comparison of the luminosities in this energy range.  The hardness ratio of CX1 as measured with \\emph{Chandra} (see Fig.\\,\\ref{f3}) is that of a bremsstrahlung spectrum with temperature $kT_\\mathrm{b}\\simeq2$\\,keV. Assuming this spectrum, the distance and absorption of M55 as above, and the countrates given in Table\\,3 of \\citet{ver01}; cf.\\ \\citet{jvh96} and in Table\\,2 of \\citet{wwb06}, we obtain 0.5--2.5\\,keV luminosities, in units of $10^{31}$\\,erg\\,s$^{-1}$, of $6.0\\pm1.5$ for both the \\emph{ROSAT} PSPC observation (March 1991) and the \\emph{ROSAT} HRI observation (September 1997), and of $0.44\\pm0.09$ for the \\emph{XMM} observation (October 2001); our \\emph{Chandra} observation (May 2004, Table\\,1) gives $4.1\\pm0.4$.  Thus, the X-ray flux was an order of magnitude lower during the \\emph{XMM} observation than during the \\emph{ROSAT} and \\emph{Chandra} observations. \\citet{wwb06} assume that the X-ray flux is higher during an optical/ultraviolet dwarf nova outburst, and conclude that CV1 was in outburst during both \\emph{ROSAT} observations.  However, most dwarf novae have {\\em lower} X-ray fluxes during outburst than in quiescence (\\citealt{wvb+96,vwm99} and the discussion of Fig.\\,2 in \\citealt{vbrp97}). We conclude that CV1 was \\emph{in outburst} during the \\emph{XMM} observations, and in quiescence during the \\emph{ROSAT} and \\emph{Chandra} observations, in good agreement with the observations by \\citet{kpt+05} that CV1 spends most of the time in quiescence.  We detect the optical counterpart to CV1 at $R=20.7$ and $B-R=1.34$, where it lies only about 0.3\\,mag blue-wards of the cluster main-sequence. With $R-\\mathrm{H}\\alpha=0.58$, it exhibits a large excess of H$\\alpha$ emission, reminiscent of cataclysmic variables. At $B=22.1$, our photometry of CV1 puts it close to the quiescent level of about $B=22.5$ seen in the light-curve by \\citet{kpt+05}. Furthermore, the proximity in color of the source to the cluster main-sequence was also seen by \\citet{kpt+05} during quiescence, where the source actually coincides with the main-sequence in $B-V$ and $V-I$ (see Fig.\\,4 of Kaluzny et al.). Hence, we conclude that CX1/CV1 was also in quiescence during our FORS2 observations.  Fig.\\,\\ref{f8} shows that CX20/XMM14 had \\emph{Chandra} and \\emph{XMM} countrates as expected for a constant luminosity between the two observations.  The source was not detected with \\emph{ROSAT}; in re-analysing the \\emph{ROSAT} PSPC observation, using the method outlined in \\citet{ver01}, we obtain a $2\\sigma$ upper limit to the countrate of 0.0006 cts\\,s$^{-1}$ (channels 50--240), which corresponds roughly to a flux at 0.5--2.5\\,keV of $8.9\\times10^{-15}$\\,erg\\,cm$^{-2}$\\,s$^{-1}$. Analysis of the HRI data gives an upper limit of 10 source counts, or $8\\times10^{-15}$\\,erg\\,cm$^{-2}$\\,s$^{-1}$.  Within the error, these limits are compatible with the flux found with \\emph{Chandra} (Table\\,\\ref{t2}) and with \\emph{XMM}. \\citet{wwb06} argue that this source has been variable by at least a factor 22 between the \\emph{ROSAT} and \\emph{XMM} observations. We find no evidence for such variation.  \\citet{kpt+05} discuss an interesting blue variable star, which they name M55-B1. From its variability between observing seasons, in the absence of short-term ($\\sim$ hr) variability, they argue that it is a background quasar rather than a cataclysmic variable. This source is not detected on the \\emph{Chandra} ACIS-S3 chip, but on an adjacent chip.  It is clearly detected in X-rays, with an unabsorbed flux in the 0.5--8.0\\,keV band of $5.1\\times10^{-15}$\\,erg\\,cm$^{-2}$\\,s$^{-1}$ for absorption fixed at that of M55.  A free power law fit gives a higher absorption ($N_\\mathrm{H}=7.4\\times10^{21}$\\,cm$^{-2}$) and higher unabsorbed flux of $8.8\\times10^{-15}$ erg\\,s$^{-1}$.  \\subsubsection{NGC\\,6366} Only the brightest \\textit{Chandra} source was detected with \\textit{ROSAT}, at a 0.5--2.5\\,keV luminosity \\citep{ver01} compatible with that during the \\textit{Chandra} observation. \\textit{ROSAT} observed CX1a and CX1b as one source.  The observed spectrum of the point source is relatively soft (see Fig.\\,\\ref{f2}, and Sect.\\,\\ref{s2.3}), which with the high absorption towards NGC\\,6366 implies a very soft intrinsic spectrum. Whereas many cataclysmic variables have hard spectra, a very soft spectrum is still compatible with identification as an old nova.   \\subsection{Ratio of optical to X-ray luminosity}\\label{s5.2} Our classification of the X-ray sources for which we have an optical identification is based in part on the X-ray to optical luminosity ratio. \\citet{bph+04} showed that a line of constant \\textit{Chandra} X-ray to optical luminosity roughly separates the cataclysmic variables from the active binaries in 47\\,Tuc, NGC\\,6752 and M4 (data from \\citealt{ghem01,eghg03a,plh+02,bph+04}). This separatrix corresponds to the dashed line in Fig.\\,\\ref{f7}, and is given by \\begin{equation} \\log L_\\mathrm{X} \\mbox{(erg\\,s$^{-1}$) [0.5--2.5\\,keV]}=34.0 -0.4M_V \\label{sepbas} \\end{equation}  Remarkably, and perhaps disconcertingly, this line is rather higher than the upper bound to the X-ray luminosities of stars and RS~CVn binaries near the Sun. An upper bound is given in \\citet{vpb07} as \\begin{equation} \\log L_\\mathrm{X} \\mbox{(erg\\,s$^{-1}$) [0.5--2.5\\,keV]}=32.3 -0.27M_V \\label{sepver} \\end{equation} This line is an absolute upper bound in that it lies above all \\textit{ROSAT} measurements for nearby stars (selected from \\citealt{hssv99}) and for RS\\,CVn systems \\citep{dlfs93}, where we multiply the 0.1--2.4\\,keV luminosities given by H\\\"unsch et al.\\ and Dempsey et al.\\ with 0.4 to obtain the 0.5--2.5\\,keV luminosities.  It is shown as a solid line in Fig.\\,\\ref{f7}. It has been remarked already on the basis of \\textit{Einstein} 0.2--4\\,keV measurements that the X-ray luminosity of main-sequence stars increases with increasing stellar rotation velocity, until it saturates at  \\begin{equation} L_\\mathrm{X}\\simeq0.001 L_\\mathrm{bol} \\label{satur} \\end{equation} where $L_\\mathrm{bol}$ is the bolometric luminosity of the star (e.g.\\ \\citealt{vw87}, esp.\\ their Fig.\\,6).  We show in Fig.\\,\\ref{f7} the line of $0.001L_\\mathrm{bol}$ for the 11.2\\,Gyr isochrones calculated by \\citet{gbbc00} for metal-poor stars ($z=0.001$), appropriate for globular clusters, and for stars of solar abundance ($z=0.019$), appropriate for stars in the solar neighbourhood.  These lines differ little, and give a more conservative upper bound to the maximum 0.5--2.5\\,keV luminosity of stars near the Sun.  The region between the separatrix Eq.\\,\\ref{sepbas} on the one hand and the upper bounds given by Eq.\\,\\ref{sepver} or Eq.\\,\\ref{satur} contains a dozen optical counterparts of X-ray sources in globular clusters that have been classified as magnetically active binaries. If this classification is correct, we must accept that such binaries can have rather higher X-ray luminosities in globular clusters than in the solar neighbourhood. Since the bolometric luminosity as a function of $M_V$ is similar for low and solar metallicities, this would imply that the saturation of X-rays occurs at higher $L_\\mathrm{X}/L_\\mathrm{bol}$ in globular clusters. Alternatively, we may wish to reconsider the classification of these sources. If we want to classify them as cataclysmic variables, we must accept that globular clusters host cataclysmic variables that are not blue with respect to the main sequence, even in the $B$-band filter. These systems clearly warrant further study. We note that CV1 in M55 does lie on the main-sequence in $B-V$ in quiescence (\\citealt{kpt+05}, see Sect.\\,\\ref{s511}).  \\begin{table} \\caption{Values for central density $\\rho_0$, core radius $r_\\mathrm{c}$, distance $d$, and absolute visual magnitude $M_V$, taken from Harris (1996, version Feb 2003), and for M4 from \\citet{rfi+97}. Collision number $\\Gamma$ and mass within the half-mass radius $M_\\mathrm{h}$ are scaled on the values for M4. \\label{t6}} \\begin{tabular}{lrrlrrr} \\hline \\hline cluster & log $\\rho_0$ & $r_\\mathrm{c}$ & $d$ & $M_V$ & $\\Gamma$ & $M_\\mathrm{h}$ \\\\ &  ($L_\\odot$~pc$^{-3}$)  & ($\\arcsec$)  & (kpc) & & \\\\ \\hline M4 &  4.01 &   49.8 &   1.73 & $-$6.9 & $\\equiv$1.00 & $\\equiv$1.00 \\\\ NGC\\,6397 &  5.68 &    3.0 &   2.3 & $-$6.6 &   2.05 &    0.78 \\\\ 47\\,Tuc  &  4.81 &   24.0 &   4.5 & $-$9.4 &  24.91 &   10.19 \\\\ NGC\\,288  &  1.80 &   85.2 &   8.8 & $-$6.7 &   0.04 &    0.86 \\\\ M55 &  2.15 &  169.8 &   5.3 & $-$7.6 &   0.18 &    1.82 \\\\ NGC\\,6366 &  2.42 &  109.8 &   3.6 & $-$5.8 &   0.09 &    0.35 \\\\ \\hline \\end{tabular} \\end{table}  \\subsection{Origin of the X-ray sources} To compare the numbers of X-ray sources in globular clusters, \\citet{pla+03} chose a luminosity limit of $L_\\mathrm{X}>4\\times10^{30}$\\,erg\\,s$^{-1}$ [0.5--6.0\\,keV], which is detectable in almost all the clusters they investigated.  We have detected 16 sources above this limit within the half-mass radius of M55, with an expected number of 8-9 unrelated sources; and 5 within the half-mass radius of NGC\\,6366, with an expected number of 4 unrelated sources. As already argued above, from these numbers alone we cannot exclude that all detected sources are unrelated.  The situation changes if we take into account our optical identifications (Section\\,4), which indicate that at least three sources (CX1, CX7, CX10) are highly probable members of M55, and that two sources (CX4, CX5) are highly probable members of NGC\\,6366. We have also argued that the point source plus extended source is an old nova in NGC\\,6366.  In Table\\,\\ref{t6} we compare the collision numbers $\\Gamma\\equiv{\\rho_0}^{1.5}{r_\\mathrm{c}}^2$ and the masses within the half-mass radii $M_\\mathrm{h}$ (as measured with the absolute magnitude assuming a fixed mass-to-light ratio) for several well-observed clusters; where we scale both $\\Gamma$ and $M_\\mathrm{h}$ on the globular cluster M4.  \\citet{pla+03} report $41\\pm2$ sources above the luminosity limit in the core of 47\\,Tuc.  If the number of X-ray sources would scale with the collision number, we would predict 0.3 sources in M55, and 0.15 in NGC\\,6366. The number of detected very probable member X-ray sources is significantly higher in both clusters. If the number of X-ray sources would scale with the mass, we would predict about 7 sources in M55, and 3 to 4 in NGC\\,6366. The number of detected member X-ray sources is remarkably close to this.  To study the origin of X-ray sources in globular clusters in further detail, \\citet{ph06} investigate the hypothesis that the number of sources scales with the collision number and mass as \\begin{equation} N = a\\Gamma + b M\\label{e:numbers} \\end{equation} and show that the dependence on mass is significantly present in the observations. To overcome the small-number statistics of clusters with small source numbers, they combine the numbers of such clusters before fitting them to Eq.\\,\\ref{e:numbers}.  In a recent paper, \\citet{vpb07} suggest a more general way of fitting the observed numbers to Eq.\\,\\ref{e:numbers}.  Briefly, they determine for each cluster the most probable combination of numbers $N_\\mathrm{c}$ of member and $N_\\mathrm{b}$ of unrelated sources that gives the observed number of sources within the half-mass radius $N_\\mathrm{h}$, based on the expected number of cluster sources $\\mu_\\mathrm{c}$ according to the model of Eq.\\,\\ref{e:numbers}, and on the expected number of unrelated sources $\\mu_\\mathrm{b}$. In doing so they assume that the probability of finding $N$ sources when $\\mu$ are expected is given by the Poisson function \\begin{equation} P(N,\\mu) = {\\mu^N\\over N!}e^{-\\mu} \\end{equation} The fitting procedure consists of varying $a$ and $b$ to maximize \\begin{equation} P = \\prod_j [P(N_\\mathrm{c},\\mu_\\mathrm{c})P(N_\\mathrm{b},\\mu_\\mathrm{b})]_j\\label{totprob} \\end{equation} where $j$ indexes the clusters.  The values for $N_\\mathrm{h}$ and $\\mu_\\mathrm{c}$ that we use are taken from \\citet{pla+03}. We ignore the clusters for which only a lower limit to $N_\\mathrm{h}$ is known. We add three clusters: NGC\\,288, M55 and NGC\\,6366; the values we use are 11/8.0, 16/8.3 and 5/4.0 for $N_\\mathrm{h}/\\mu_\\mathrm{c}$ for these three clusters, respectively.  The best solution is found for $a=1.3\\Gamma_\\mathrm{M4}$ and $b=0.6M_\\mathrm{M4}$. The contours of the acceptable solutions are given in Fig.\\,\\ref{fig:probs}, and it is seen that $b=0$ is acceptable at about the $2\\sigma$ level. Thus, from the observed numbers alone, we have no firm evidence for a dependence on mass.  We repeated the fitting procedure, but excluding solutions for which $N_\\mathrm{c}$ is less than the minimum number of highly probable cluster members.  For these fits we assume that 2 sources are certain members of NGC\\,288, 3 of M55 and 2 of NGC\\,6366. As expected, these constraints force the solutions to a stronger dependence on mass: the best solution is found for $a=1.2\\Gamma_\\mathrm{M4}$ and $b=1.1M_\\mathrm{M4}$. As seen in Fig.\\,\\ref{fig:probs}, the dependence on mass is now significant.  The best values of $a$ and $b$ are relatively stable against small variations in the numbers $\\mu_\\mathrm{c}$ and $N_\\mathrm{h}$ for individual clusters, as comparison with the results of \\citet{vpb07} confirm. The probability of the best solution does vary with small variations in $\\mu_\\mathrm{c}$ and $N_\\mathrm{h}$ for individual clusters.  For example, increasing the minimum of cluster member X-ray sources in NGC\\,6366 from 2 to 3 does not change the best values for $a$ and $b$, but decreases the probability of the best solution for NGC\\,6366, for which the best model predicts $\\mu_\\mathrm{c}=0.4$, and thus also slightly decreases the overall probability.  The cluster which is responsible for the lowest probability is NGC\\,6397, which houses 12 sources, of which our best model assigns $N_\\mathrm{c}=11$ to be cluster members, even though it expects only $\\mu_\\mathrm{c}=3.6$. The remarkably large number of sources in this cluster has been remarked upon before, and various explanations have been offered. In our model, however, whereas still the worst fit cluster, it is not altogether unacceptable. The main difference between this result and the strong deviation for NGC\\,6397 as found by e.g. \\citet{pla+03} appears to be due to a different definition of $\\Gamma$. \\citet{pla+03} integrate a King model to estimate $\\Gamma$, whereas we compute it as $\\Gamma={\\rho_0}^{1.5}{r_\\mathrm{c}}^2$.  Our probabilities do not take into account the uncertainties in $\\Gamma$ and $M_\\mathrm{h}$, nor the observed number $N_\\mathrm{h}$ which through its dependence of the flux limit is affected by the assumed distance to the cluster. Clearly these uncertainties warrant further investigation.  An important result illustrated by Fig.\\,\\ref{fig:probs} is that the masses of the clusters may {\\em not} be taken as proxies for their collision numbers. Even though there is some correlation between mass and collision number, the spread in collision number at given mass, and in mass at given collision number, is very large. Indeed our fits show that the number of X-ray sources within a cluster is very badly predicted by the mass alone, as illustrated by the highly significant offset of the best solution from the $a=0$ line.  \\begin{figure} \\resizebox{0.95\\hsize}{!}{\\includegraphics[angle=270]{9350f9.pdf}} \\caption{Probability contours around the best solutions for a fit of the observed number of sources within the half-mass radii of 12 globular clusters to Eq.\\,\\ref{e:numbers} where we allow zero cluster members in some clusters (dashed contours), and where we impose lower limits to the number of cluster members of 2 for NGC\\,288, 3 for M55 and 2 for NGC\\,6366 (solid contours). We assume that $-2(P-P_\\mathrm{bf})$, with the probability $P$ given by Eq.\\,\\ref{totprob} and $P_\\mathrm{bf}$ indicating the best fit probability, follows a $\\chi^2$ distribution, and show $\\Delta\\chi^2=1$, 4 and 9.  The imposition of minimum numbers of cluster member X-ray sources strongly increases the significance of the mass dependence in Eq.\\,\\ref{e:numbers}.}\\label{fig:probs} \\end{figure}"
    },
    "0807/0807.2622_arXiv.txt": {
        "abstract": "{In the borders of the dead zones of protoplanetary disks, the inflow of gas produces a local density maximum that triggers the Rossby wave instability. The vortices that form are efficient in trapping solids.} {We aim to assess the possibility of gravitational collapse of the solids within the Rossby vortices.} {We perform global simulations of the dynamics of gas and solids in a low mass non-magnetized self-gravitating thin protoplanetary disk with the Pencil Code. We use multiple particle species of radius 1, 10, 30, and 100\\,cm. The dead zone is modelled as a region of low viscosity.} {The Rossby vortices excited in the edges of the dead zone are efficient particle traps. Within 5 orbits after their appearance, the solids achieve critical density and undergo gravitational collapse into Mars sized objects. The velocity dispersions are of the order of 10 \\mps~  for newly formed embryos, later lowering to less than 1\\,\\mps~ by drag force cooling. After 200 orbits, 38 gravitationally bound embryos were formed inside the vortices, half of them being more massive than Mars.}% {}% ",
        "introduction": "The formation of planets is one of the major unsolved problems in modern astrophysics. In the standard core accretion scenario, sub-$\\mu$m grains assemble into progressively larger bodies through electrostatic interactions (Natta et al. 2007), eventually growing into centimeter and meter sized boulders. Growth beyond this size, however, is halted since these boulders have very poor sticking properties and are easily destroyed by collisions at the velocities assumed to be prevalent in circumstellar disks (Benz 2000). Furthermore, centimeter and meter sized solids are loosely decoupled from the gas, but remain sufficiently small to be affected by significant gas drag. The resulting headwind from the sub-Keplerian gas reduces their angular momentum and forces them into spiral trajectories onto the star in timescales as short as a few thousand years (Weidenschilling 1977a).  A mechanism for overcoming these barriers was presented by Kretke \\& Lin (2007). In the presence of sufficient ionization, the gaseous disk couples with the ambient weak magnetic field, which triggers the growth of the magneto-rotational instability (MRI; Balbus \\& Hawley 1991). In its saturated state, a vigorous turbulence drives accretion onto the star by means of magnetic and kinetic stresses. However, in the water condensation front (snowline) the abundant presence of snowflakes effectively removes free electrons from the gas, lowering the degree of ionization. The turbulence is weakened locally and the accretion flow is stalled. As the radial inflow proceeds from the outer disk, gas accumulates at the snowline. Since embedded solid bodies move towards gas pressure maxima (Haghighipour \\& Boss 2003), the snowline environment proposed by Kretke \\& Lin (2007) is potentially an efficient particle trap. This scenario was further explored by Brauer et al. (2008), who demonstrated that as solids concentrate at this local pressure maximum, rapid growth into kilometer sized planetesimals occurs by coagulation.  Kretke et al. (2008) emphasized that an identical mechanism is supposed to occur elsewhere in the disk. Ionization ought to be present in the very inner disk due to the high temperatures, as well as in the outer regions where the gas is sufficiently thin for cosmic rays to penetrate to the disk midplane and provide ionization throughout. In between, however, temperatures are too low and gas columns too thick to allow sufficient ionization either by collisions or by cosmic rays. In the midplane of this region, the gas is neutral and the turbulence is largely suppressed (Gammie 1996). As in the snowline, the accretion flow from the MRI-active regions halts at the borders of this ``dead'' zone, where the gas then accumulates.  These models have been tested only by one-dimensional simulations, and these tests have therefore not benefitted from an interesting development. As shown by Varni\\`ere \\& Tagger (2006), the density pileup at the border of the dead zone triggers the Rossby wave instability (RWI; Li et al. 2001). The azimuthal symmetry of the problem is broken and long-lived anticyclonic vortices are formed as the waves break and coalesce. Such entities are of significant interest because, by rotating clockwise in the global counterclockwise Keplerian flow, they amplify the local shear and induce a net force on solid particles towards their center (Barge \\& Sommeria 1995). As shown by Klahr \\& Bodenheimer (2006), the accumulation of solids under these circumstances is likely to lead to high densities. Inaba \\& Barge (2006) continued the study of Varni\\`ere \\& Tagger (2006) by including solids and confirming that the Rossby vortices excited in the borders of the dead zone act as powerful traps, enhancing the local solids-to-gas ratio by at least an order of magnitude. Unfortunately, they used a fluid approach - which limited the maximum particle size they could consider - and they did not include the self-gravity of the solids, which is crucial to follow the gravitational collapse. Studies with interacting particles in the literature include a MRI-unstable local box (Johansen et al. 2007) capable of producing dwarf planets out of meter sized boulders, a global massive disk unstable to gas self-gravity (Rice et al. 2006) in which concentrations of 0.5 $M_{\\oplus}$ are seen in the spiral arms, and a simulation that produces 10-100 km \"sandpile\" clumps formed out of mm-sized particles (Cuzzi et al. 2008).   In this Letter, we build on the studies of Varni\\`ere \\& Tagger (2006) and Inaba \\& Barge (2006) by including self-gravitating centimeter and meter sized Lagrangian particles to model the solid phase. We show that in the vortices launched by the RWI in the borders of the dead zone, the solids quickly achieve critical densities and undergo gravitational collapse into protoplanetary embryos in the mass range 0.1-0.6\\mearth. ",
        "conclusions": "In this Letter, we have shown that when modelling the self-gravity of the gas and solids in protoplanetary disks, gravitational collapse of the solids into Mars sized protoplanetary embryos occurs rapidly at the borders of the dead zone, where particles concentrate. We have also found that tides from the dense gaseous vortices may hinder the formation process significantly  Studies considering the origin of oligarchs usually begin from a collection of 10-20 Mars sized objects (e.g., Kobuko et al. 2006). This Letter presents the first simulation in which a similar number of Mars-sized embryos are formed from centimeter and meter sized sized building blocks.  It is crucial to the model that particles grow to sufficient size, otherwise the drag force from the gas becomes too strong to allow any concentration. Testi et al.\\ (2003) observed grains of up to cm sizes in the disk surrounding the pre-main-sequence star CQ Tauri, which provides some observational evidence that a sufficient number of particles of the required size may exist in true protoplanetary disks.  We emphasize again that the model used in this Letter is simplistic, and the conditions may differ with a more realistic treatment of the dead zone. Nevertheless, the mechanism presented in this Letter (as proposed originally by Varni\\`ere \\& Tagger 2006) appears robust. It only requires the RWI to be excited in the borders between the active and dead zones, which in turn relies only on the slowdown of the accretion flow at this same border. Future studies should include a coagulation/fragmentation model such as those of Brauer et al. (2008) or Johansen et al. (2008), and focus on the precise state of flow at this transition region in global simulations, to address the question of how the RWI interacts with the MRI and non-ideal MHD in three dimensions."
    },
    "0807/0807.0282_arXiv.txt": {
        "abstract": "It is well-known that the universe was in a plasma state both before decoupling and after reionization. However, the conventional wisdom has been that the plasma effects are largely Debye-shielded and can thus be safely ignored when considering large scale evolutions. Recently we showed that large scale structure formation in the universe may actually be suppressed by the plasma collective effect. Indeed, observational data indicate that the conventional theoretical formula tends to overestimate the matter power spectrum at scales $k>1h{\\rm Mpc}^{-1}$. In this paper, we further develop our theory through a more thorough and general derivation of the Maxwell-Einstein-Boltzmann equation. In addition to baryon density perturbation post reionization, we apply this general formulation to investigate the possible plasma effect on CMB anisotropy. As expected, while the plasma effect does render an observable effect to the former, its impact on the latter is totally negligible. ",
        "introduction": "Recently we showed that large scale structure formation in the universe would be suppressed by the plasma collective effect\\cite{CnL}. Indeed, observational data seem to indicate that the conventional theoretical formula tends to overestimate\\cite{Teg} the matter power spectrum at scales $k>1h{\\rm Mpc}^{-1}$. It is implied that the underlying theory may have to be modified. Moreover, our work\\cite{CnL} is performed in the regime of Newtonian limit and we expect to generalize the plasma collective effect  in the general relativistic regime.  On the other hand, the cosmological perturbation theory(CPT) is not complete in the sense that the electromagnetic interaction is absent. Einstein-Boltzmann equations accounting for evolution of cosmological perturbations have been explored for decades\\cite{Lifshitz,PnY,BnS}. Concerning anisotropies, the Sachs-Wolfe effect\\cite{SW} and Silk damping\\cite{Silk} were predicted just after the discovery of CMB\\cite{3K} and the approach of tightly coupled limit\\cite{DZS} was perfected\\cite{Seljak,HnS1995} in 1990's to account for the CMB spectrum. Neither these studies include EM fields even though  tightly coupled limit approach combines photons and baryons as a cosmic plasma, nor have investigations been performed on the likeliness of plasma effects in CMB study. In parallel, the efforts on matter inhomogeneities including study on the growth and transfer functions\\cite{Blu,Pee,HnS1996} comprise the evolution under only gravity. Later modification\\cite{EnH} taking baryons into account neither addressed EM interactions, though it derived more accurate fitting formula to replace the pure cold dark matter BBKS fitting\\cite{BBKS}. Though the CPT as well as its predictions and description of CMB and MPS do not take the EM interaction into account, we expect plasma collective effects for the universe was in plasma state before decoupling and after reionization.  In the paper, we generalize the Einstein-Boltzmann(EB) equations in the CPT to our Maxwell-Einstein-Boltzmann(MEB) equations, which describe the perturbation evolution under both the gravity and EM interaction. With the inclusion of plasma effects we derive not only CMB ion-acoustic oscillations but also baryon ion-acoustic oscillations.  The paper is organized as follow. In Sec. II, we justify that the plasma conditions apply in the universe during pre-decoupling and post-reionization era. Based on this justification, we generalize the EB equation to  the MEB equation in Sec.III. The CMB and baryon ion-acoustic wave equations are then derived, respectively, in Sec.IV and V. Finally, we discuss the significance of plasma effects on cosmological perturbation. ",
        "conclusions": "The acoustic behavior is an impressive feature in the CMB spectrum and matter power spectrum. Our study shows manifestly the EM fields also contribute to this acoustic oscillation collectively. Though the whole system is neutral in total, the nonzero temperature induced charge density fluctuations form a collective electric field which turns to be an extra portion of pressure exerting on the densithy perturbations in both cases. However, the contribution on the CMB is too small to be observed. Both the order of magnitude estimation and the ion fluid equation analysis tell that the EM contribution is of the same order as of the pressure from charged particles. As the adiabatic condition indicates baryon and photon density contrasts being parallel, the baryon to photon number ratio $n_b/n_\\gamma\\sim10^{-10}$ states that the EM effect on the sound speed is $\\sim10^{-5}$ smaller than ordinary CMB fluctuation.  Our generalization of the Boltzmann equation proves this number by incorporating the Lorentz force into the ordinary geodesic equation for baryons.  Though the plasma collective effect is negligibly small in CMB spectrum, it should has sizable imprints in the matter power spectrum as discussed in our work\\cite{CnL}."
    },
    "0807/0807.2764_arXiv.txt": {
        "abstract": "{Over the past decade the use of the red clump (RC) as distance indicator has increased in importance as this evolutionary phase is well populated and a good local calibration exists.} {The absolute calibration of the RC in the $I$ and $K$ band is investigated again based on the recently published revised Hipparcos parallaxes.} {A numerical model is developed that takes the various selection criteria and the properties of the Hipparcos catalogue into account. The biases involved in applying certain selections are estimated with this model.} {The absolute magnitudes that are found are $M_{\\rm I} = -0.22 \\pm 0.03$ and $M_{\\rm K} = -1.54 \\pm 0.04$ (on the 2MASS system). The $I$-band value is in good agreement with previous determinations, the $K$-band value is fainter than previously quoted, and this seems to be related to a selection bias whereby \\it accurate \\rm $K$-magnitudes are only available for relatively few bright stars. Applying population corrections to the absolute $K$ magnitude of RC stars in clusters supports the fainter magnitude scale.} {} ",
        "introduction": "Paczy\\'nski \\& Stanek (1998) were the first to advocate the use of red clump (RC) stars in the $I$-band as distance indicators. Two advantages are that this phase of stellar evolution is usually well populated, and that many RC stars exist with good parallaxes in the Hipparcos catalogue.  They consider stars with an error in parallax smaller than 10\\%, located in the box 0.8 $< (V-I) <$ 1.25 and $-1.4 < M_{\\rm I} <$ 1.2 and with Hipparcos flag H42 equal to one of ``A,C,E,F,G'', which means that the star has one or more $I$-band measurements in the Cousins, Johnson, or Kron-Eggen system.  They find no significant trend of $M_{\\rm I}$ on colour and derived $M_{\\rm I} = -0.185 \\pm 0.016$ for the entire sample. They then considered the 228 stars with $d<$70 pc and average distance of 50 pc to find $M_{\\rm I} = -0.192 \\pm 0.023$, and $M_{\\rm I}= -0.094 \\pm 0.027$ for the rest of the sample at an average distance of 106 pc. The difference is ascribed to reddening, and when extrapolating to 0 pc, they arrive at their final result of find $M_{\\rm I} = -0.279$ and assign an error bar of 0.088.  In deriving these number, they also applied a correction for ``a distance bias''.   Stanek \\& Garnavich (1998) considered the same sample, and their final value of $M_{\\rm I} = -0.23 \\pm 0.03$ is based on the sub sample of stars within 70 pc and they argue that no correction for reddening is needed, and they do not apply a correction for selecting the sample on parallax.   Udalski (2000) considered the effect of metallicity. He took the same cut in parallax error and Hipparcos H42 flags, but also included stars with flag ``H'' to increase the sample and considered stars with a metallicity determination from high-resolution spectroscopy by McWillam (1990). Based on this sample of 284 stars, he finds a weak dependence on metallicity: $M_{\\rm I} = (0.13 \\pm 0.07) ($[Fe/H]$ +0.25) + (-0.26 \\pm 0.02)$.   Alves (2000) was the first to consider the $K$-band for RC stars as a standard candle. He considered the 284 stars from Udalski (2000) and finds that 238 of them have a $K$-magnitude in the Two Micron Sky Survey (TMSS, Neugebauer \\& Leighton 1968; stars are designated by ``IRC'' numbers).  Assuming no reddening and considering stars with $-2.5 < M_{\\rm K} < -0.8$ and all $(V-K)$ colours,  he finds $M_{\\rm K} = -1.61 \\pm 0.03$ and a weak dependence on metallicity $M_{\\rm K} = (0.57 \\pm 0.36) $[Fe/H]$ + (-1.64 \\pm 0.07)$.  The RC has gained importance as a distance indicator since the Paczy\\'nski \\& Stanek (1998) paper as shown for example by distance estimates based on the RC to Fornax (Rizzi et al. 2007), M33 (Kim et al. 2002), or IC 1613 (Dolphin et al. 2001).  With the recent publication of the revised Hipparcos parallax data, it is timely to investigate the absolute magnitude of the RC again.  In Sect.~2 the revised Hipparcos data is introduced with auxiliary data used in the analysis. Section~3 describes the numerical model used for estimating selection biases.  Section~4 describes the results, and Section~5 presents a brief summary and discussion. ",
        "conclusions": "Using revised Hipparcos parallaxes and extensive numerical simulations that take into account de properties of the Hipparcos catalogue in terms of completeness and parallax error, and also taking into account the selection functions that are involved in selecting stars with certain properties ($K$-band, high-quality $I$-band, spectroscopic metallicities), the absolute magnitude of the red clump in $I$ and $K$ is investigated.  In the $I$-band, the value of $M_{\\rm I} = -0.22 \\pm 0.03$ essentially agrees with the value quoted by Stanek \\& Garnavich (1998). There is no dependence of $M_{\\rm I}$ on $(V-I)$ colour, and the dependence on metallicity is marginal: $(0.08 \\pm 0.07)$ ([Fe/H] + 0.15).   In the $K$-band, the influence of selection biases appears stronger and the value derived $M_{\\rm K} = -1.54 \\pm 0.04$ differs significantly from the value originally derived by Alves (2000), which is almost 0.1 mag brighter, when his value of $-1.61$ on the Bessell \\& Brett (1988) system is converted to the 2MASS system (Carpenter 2001).  There is no dependence of $M_{\\rm K}$ on metallicity, and the dependence on colour is weak: $(-0.15 \\pm 0.07)$ ($(V-K)_0 - 2.32$).   The RC stars in clusters have also been used to derive $M_{\\rm K}$, originally by Grocholski \\& Sarajedini (2002). Absolute magnitudes are derived by combining the observed mean magnitude in a box in colour and magnitude with a reddening estimate and main-sequence fitting distances.  Grocholski \\& Sarajedini (2002) used the second incremental data release of 2MASS and 14 clusters to find an absolute $K$-magnitude that is in agreement with the value in Alves (2000). Using the All-Sky data release and increasing the sample to 24 clusters, Van Helshoecht \\& Groenewegen (2007; hereafter vHG) find a value fainter by 0.05 mag, or $M_{\\rm K} = -1.57 \\pm 0.05$ on the Bessell \\& Brett system.  Both Grocholski \\& Sarajedini (2002) and vHG used averages over the cluster sample to arrive at the quoted means.  To compare these values to the absolute magnitude of the local Hipparcos sample, a ``population correction'' has to be made, to account for the difference in metallicity and age of the RC population.  The calculation of this correction is outlined and tabulated in Girardi \\& Salaris (2001) and Salaris \\& Girardi (2002). In the $K$-band, Salaris \\& Girardi (2002) show that the correction is a strong function of age for ages below 3 Gyr.  The top panel in their Fig.~3 shows that the correction is well behaved for ages above about 4 Gyr, and for each metallicity a linear function was fitted.  From vHG the five clusters older that 4 Gyr were taken and the population correction was determined by linear interpolation in age and [Fe/H] in the results in Salaris \\& Girardi.  The error in $\\Delta M_{\\rm K}$ is based on a 1 Gyr error in age and 0.1 dex in metallicity.  The results are listed in Table~\\ref{Tab-Clusters}, which lists the age and metallicity (see vHG for the references for age and metallicity determination), the $M_{\\rm K}$ value and its error that takes into account the error in reddening and the assumed distance to the cluster (from vHG), the population correction with error, and the corrected value. The weighted mean of the corrected $M_{\\rm K}$ values is $-1.39 \\pm$ 0.06 with a dispersion of 0.2 mag.  Correcting to the 2MASS system (Carpenter 2001) this becomes $-1.43 \\pm$ 0.06.  Although the dispersion is large, the mean value is also fainter than the straight mean for the cluster sample and agrees within 2-sigma with the determination from the Hipparcos sample.  To settle the issue on the importance of the bias and the absolute $K$-magnitude of RC stars would require accurate NIR magnitudes of a 100 to a few hundred (cf. Table~2) bright ($K \\less 5$) RC stars. Given the brightness, this represents a challenge to modern instrumentation because of saturation.  The absolute magnitudes derived here are in better agreement with theory than previously. Using a plausible star formation history for the solar neighbourhood, Salaris \\& Girardi (2002) derived absolute magnitudes of $M_{\\rm I}= -0.17$ and $M_{\\rm K}= -1.584$ (on the 2MASS system), which they compared to the Hipparcos-based results quoted in Alves et al. (2002) of $M_{\\rm I}= -0.26 \\pm 0.03$ and $M_{\\rm K}= -1.644 \\pm 0.03$ (on the 2MASS system), which indicates differences at the 2-3 sigma level.  The new results of $M_{\\rm I}= -0.22 \\pm 0.03$ and $M_{\\rm K}= -1.54 \\pm 0.04$ agree with theory at the 1-2 sigma level.     \\begin{table*} \\caption{The RC in old clusters } \\begin{tabular}{lcccccc} \\hline \\hline Name & $\\log t$ & [Fe/H] &     $M_{\\rm K}$   & $\\Delta M_{\\rm K}$ &  $M_{\\rm K}$ (true) \\\\ \\hline  Be 39    & 9.90 & -0.15 & $-1.56 \\pm 0.11 $ & $-0.058 \\pm 0.03$ & $-1.61 \\pm 0.12$ \\\\ NGC 188  & 9.63 & -0.12 & $-1.36 \\pm 0.11 $ & $+0.057 \\pm 0.03$ & $-1.30 \\pm 0.12$ \\\\ NGC 6791 & 9.64 & +0.40 & $-1.39 \\pm 0.11 $ & $+0.223 \\pm 0.05$ & $-1.17 \\pm 0.12$ \\\\ 47 Tuc   &10.08 & -0.70 & $-1.34 \\pm 0.21 $ & $-0.360 \\pm 0.11$ & $-1.70 \\pm 0.24$ \\\\ NGC 362  &10.08 & -1.15 & $-0.81 \\pm 0.24 $ & $-0.680 \\pm 0.11$ & $-1.49 \\pm 0.26$ \\\\   \\hline \\end{tabular} \\label{Tab-Clusters} \\end{table*}     Finally, some implications for existing distance determinations using the RC are discussed. The derived value of the absolute magnitude in the $I$-band of $M_{\\rm I} = -0.22 \\pm 0.03$ is not very different from previously adopted values in the literature of $-0.23$ or $-0.26$, so the impact on derived distances is not so great.  The distance to the LMC quoted by Pietrzy\\'nski \\& Gieren (2002) of 18.501 $\\pm$ 0.049 (random+systematic error and assuming no population correction) based on $JK$ observations of two fields in the bar would become 18.35 $\\pm$ 0.05 for an assumed population correction of $-0.03$ (Salaris \\& Girardi 2002).  Babusiaux \\& Gilmore (2005) obtained infrared data on some fields in the direction of the Galactic bulge.  Their distance scale is tied to one field for which they obtain $M_{\\rm K} = -1.72$ for an assumed distance of 8 kpc. Taking $M_{\\rm K} = -1.54$ would result in a distance of 7.4 kpc ($\\pm$0.2 kpc judging from the information in their Table~2). Adding the error in $M_{\\rm K}$ results in a total error of 0.3 kpc.  Using similar type of observations Nishiyama et al. (2006) obtained 7.52 $\\pm$ 0.1 kpc for $M_{\\rm K} = -1.59 \\pm 0.03$, which translates to 7.33 $\\pm$ 0.14 kpc with the new DM.  This can be compared to a distance of 7.94 $\\pm$ 0.37 (random) $\\pm$ 0.26 (systematic) kpc from Groenewegen et al. (2007) based on $K$-band monitoring observations of Population-{\\sc ii} cepheids and RR Lyrae stars (distance tied to an assumed LMC distance modulus (DM) of 18.50) and the direct determination of 7.73 $\\pm$ 0.32 kpc based on the orbit of the star called S2 around the central black hole when post-Newtonian physics is taken into account (Eisenhauer et al. 2005, Zucker et al. 2006), although neglect of the space motion of Sgr A$^*$ may result in systematic errors of about 0.1-0.45 kpc (Nikiforov 2008). All these independent observation are consistent and seem to indicate a DM to the GC of slightly less than 8 kpc."
    },
    "0807/0807.4468_arXiv.txt": {
        "abstract": "We provide a novel mechanism that resolves the Big Bang Singularity present in FRW space-times without the need for ghost fields.   Building on the fact that a four-fermion interaction arises in General Relativity  when fermions are covariantly coupled, we show that at early times the decrease in scale factor enhances the correlation between pairs of fermions.  This enhancement leads to a BCS-like condensation of the fermions and opens a gap dynamically driving the Hubble parameter $H$ to zero and results in a non-singular bounce, at least  in some special cases. ",
        "introduction": "As is well known the Standard Big Bang cosmology (SBB) successfully predicts the observed large scale expansion, thermal properties, nucleosynthesis and the cosmic microwave background in our universe.  However, the theorems of Hawking and Penrose prove that a curvature singularity exists in the SBB at the ``birth'' of the universe's space-time. Cosmologists have had long held expectations that at the Planck time quantum gravitational effects might resolve the singularity and provide a quantum bridge that connects a collapsing phase to an expanding one; otherwise known as a bouncing cosmology.  The implementation of the bouncing scenario in quantum theories of gravity are still at the toy model stage, since a complete description of quantum gravity is lacking.   Given this fact, we will argue that a non-singular cosmology is possible from the effects of fermions on space-time when the scale factor starts becoming small.  Most modern approaches towards a unified quantum theory of gravity, such as supergravity~\\cite{sugra}, or LQG~\\cite{loop-review} has found the first order formalism of gravity, where one treats the vielbein and connections as independent variables, to be the natural starting point. It is also well known that while passing from the first to the second order formalism (by integrating out the torsion field) the covariant coupling of fermions to gravity yields a four fermion interaction in these theories~\\cite{sugra,Jacobson,alexander,perez,freidel,das}\\footnote{See however~\\cite{simone} for a counter example where even in the presence of the Immirzi parameter, the fermions don't change the classical equations.}.  Now, the BCS theory of superconductivity states that if there exists a weak attractive four fermion interaction and a fermi surface, fermions will condense into a new ground state with lower energy comprising of spin-0 Cooper Pairs.  In the context of a homogeneous and isotropic  cosmological space-time\\footnote{The interplay between  inhomogeneities in the Fermi-gas and the metric, and especially it's implications for CMB remains an important open question.}, we will show that fermions indeed form  Cooper pairs.   Moreover, the FLRW scale factor (or the ``volume'' of the universe) plays the role as an effective coupling constant such that at early times the correlations between the fermions get stronger.  As a result, we show that as we approach the singularity an energy gap opens up, contributing negative energy which, as we shall explain shortly, is one of the fundamental requirements for obtaining realistic bounces in a spatially flat or open universe. We were also able to explicitly show that at least  when the usual (positive) matter energy density has an equation of state parameter, $\\om<1/3$, the negative gap energy can drive the expansion rate of the universe to zero, leading to a bounce between the contracting and expanding phase.  For this mechanism to work, an {\\it a priori} assumption  that is required is the existence of a finite non-zero density of the Dirac fermions (\\ie we assume a particle/anti-particle asymmetry from the very beginning).  We employ robust methods of effective field theory to demonstrate the non-singular behavior of this cosmology. We note in passing that similar ideas have been previously employed to resolve the black hole singularity and information loss problems~\\cite{blackhole}.  In contrast, previous attempts to resolve Big Bang singularity (BBS) has  mostly relied on  introducing pathological ghost fields~\\cite{narlikar,freese,frampton} which either violate quantum unitarity and/or leads to catastrophic instability~\\cite{jim}. In some other bouncing and cyclic universe scenarios~\\cite{tolman,barrow,ekpyrotic}, one leaves the singularity unresolved in the hope that ``quantum gravity'' effects will eventually smooth out the singular Big Crunch/Bang transition. Although rigorous but involved no-go theorems exist in the literature~\\cite{paris-ec} as to  why it is so challenging to avoid the singularity in General Relativity (GR),   in the context of a flat\\footnote{Although our analysis is not specific to a flat universe and readily applies to open or closed universes, for simplicity and phenomenological reasons (WMAP and other measurements strongly constrain the flatness~\\cite{wmap}) we specialize to the flat case. However, we point out that in a closed universe setting the presence of the spatial curvature itself can allow for a resolution of the big bang singularity~\\cite{closed}, but obviously such a resolution mechanism will not work for a general manifold, such as an open or a flat universe which makes it  somewhat less appealing.} homogeneous  isotropic cosmology it is easy to qualitatively see this. One just has to look at the Hubble equation \\be H^2\\equiv\\LF{\\dot{a}\\over a}\\RF^2={1\\over 3M_p^2}\\sum_i\\rho_i \\ee which governs the evolution of the scale-factor $a(t)$ appearing in the Friedman-Lemaitre-Robertson-Walker (FLRW) metric \\be ds^2=-dt^2+a^2(t)dx_3^2\\ . \\ee Here, $i$ labels the different components of matter present in the universe. Now, in the bouncing universe construction the scale factor goes through a minimum, where $\\dot{a}$, or equivalently the Hubble expansion rate, $H$, must vanish. This means that some of the matter components must have negative energy, a property hard to realize in conventional matter. Ghost fields have negative kinetic energies and  can mediate a bounce and therefore have been used in previous literature. One may wonder whether ordinary scalar fields with negative potential energy can work? Unfortunately, having a negative energy component is not sufficient for a bounce, the negative energy component has to precisely cancel  the  positive matter component at the bounce and then redshift away faster than it, to leave the total energy density positive. In other words one requires violation of not only the weak energy condition ($\\rho< 0$) but also the null energy condition ($\\rho+p< 0$)~\\cite{paris-ec}. As a consequence the energy density has to increase with expansion of the universe! It is easy to check that  ordinary scalar fields cannot achieve this.  The reason why an attractive four fermion coupling can give rise to a consistent bouncing scenario is because, firstly the interaction energy between fermions gives rise to a negative contribution to the energy density (binding energy between Cooper pairs). Secondly,  this energy density depends on the ``Gap'' which in turn, depends rather nontrivially on the scale factor (or volume of the universe) via the chemical potential. It turns out that this rather non-trivial volume dependence can violate the null energy condition temporarily as needed to have a bounce. On hind sight the fact that fermions can violate the energy conditions probably should not  come as a real surprise, even ``classically''  fermionic condensates are known to violate energy conditions and give rise to bounces~\\cite{classical}.  We should point out that in the context of effective field theory there have been some successful efforts in constructing toy models involving higher derivative actions which can evade the problem of ghosts (for instance involving non-local modifications of gravity~\\cite{BMS}, or ghost condensation mechanism~\\cite{new-ekpyrotic}) but preserve some of the useful ghost-like properties  of higher derivative theories to resolve the  BBS. (For a more detailed review of various bouncing scenarios the readers are referred to~\\cite{novello}.) However, the mechanism that we propose is not a toy model, and has a much more generic scope as it essentially only relies on having covariant coupling of fermions to gravity  and the presence of a  Fermi surface.  A third requirement to achieve a bounce intriguingly, turns out to be  the presence of the Immirzi parameter (associated with Holst's generalization of the Einstein-Hilbert action) which also seems to play a key role in the nonperturbative quantization of gravity~\\cite{loop-review}. For instance,  in Loop Quantum Gravity (LQG), the quantization of the area operator, and the existence of bouncing cosmologies~\\cite{loop} depend crucially on the Immirzi parameter. What is also suggestive is that like in LQG, we have to rely on a completely nonperturbative effect, the formation of a BCS gap to evade the Big Bang.    This paper is organized as follows: In section II we present the first order formalism of GR and derive the four fermion interaction.  In section III we develop the cosmological BCS theory and derive the effective potential for the cosmological energy gap.  In section IV we analyze a special case in details to show how and when singularity resolution can occur.  Finally we conclude and provide future outlook in section V. ",
        "conclusions": "In this paper we have presented a novel physical mechanism which self consistently resolves the initial big bang singularity. The crucial ingredient relies on the transient violation of the null energy condition during the bounce.  The interactions responsible for the negative gap energy at the bounce is enhanced as the universe contracts to the singularity.  Typically it is difficult to obtain a bounce without the introduction of dangerous ghost states, but in this mechanism the negative energy arises from the binding energy associated with the bound state of the Cooper pair.   Moreover, the gap only becomes significant at early times and consistently redshifts away at late times, in other words the late time cosmology is protected from the negative energy contribution that generated the bounce. In a special region of the parameter space we were explicitly able to obtain bouncing solutions in the presence of matter provided it's equation of state parameter was less than one-third.  While our results are promising, a few developments are in order.  First, we have not included perturbations of the condensate which is expected to affect the homogeneity and isotropy both at early and late times; how robust is the BCS bounce in the presence of inhomogeneities?   Second, what would be the observational consequences of this mechanism? How would inhomogeneities responsible for large scale structure be seeded in this model?   Since we have a bounce, we should be able to employ the techniques and physical picture in close semblance to the Ekpyrotic scenario.  A plausible structure generating scenario might arise from inhomogeneous excitations of the condensate something we plan to investigate in the future.  A promising route would be to employ the techniques developed in~\\cite{ABB} to propagate perturbations (possibly generated by an ekpyrotic scalar field) across the gap-mediated  bounce as the gap only depends on the overall volume of the universe and is not affected by fluctuations in the metric, a pre-requisite for the success of the mechanism advocated in~\\cite{ABB}."
    },
    "0807/0807.3554_arXiv.txt": {
        "abstract": "{The collapsar model predicts that the progenitors of Gamma-ray Bursts (GRBs) are metal poor in Fe group elements. The existence of low metallicity stellar populations could manifest itself in the characteristics of the GRB site immediate environment in the host galaxy. } {We analyse the strong emission lines from the sub-luminous host galaxy of GRB~980425, which showed the first connection with a supernova explosion (SN~1998bw). The host is of sufficient size to allow detailed resolved spectroscopy of individual \\ion{H}{ii} regions and to search for regions with peculiar properties close to the the GRB site.} {Using integral field spectroscopy with VIMOS we study most of the high surface brightness part of the host including the \\ion{H}{ii} region where the supernova and GRB occurred.} {The star formation rate, reddening, equivalent width and stellar mass in the GRB region is similar to other \\ion{H}{ii} regions in the host. Extreme values arise in the only region that shows emission lines from Wolf-Rayet stars, a region that is located 800 pc in projection from the GRB site. Strong emission line diagnostics of all \\ion{H}{ii} regions imply oxygen abundances between 0.3 and 0.8 solar with the lowest values arising in the WR and GRB regions. Including uncertainties from the metallicity diagnostics, all metallicities are similar to within 3$\\sigma$. We demonstrate that there is a good agreement between the luminosity weighted and mass weighted specific star formation rates (SSFR) in individual young \\ion{H}{ii} regions.  While the global average of the SSFR is similar to high redshift GRB hosts, there are significant variations between individual resolved \\ion{H}{ii} regions.  Comparing the measured emission line ratios of low redshift GRB hosts to theoretical models and observations of field galaxies, we find that GRBs are present in different environment metallicities while the regions of their origin are consistently very young. Similar line ratios of GRB hosts compared with those of the WR region can arise in spatially unresolved galaxies with bright \\ion{H}{ii} regions close to the GRB location. } {} ",
        "introduction": "The association of the Gamma-ray Burst \\object{GRB~980425} with the supernova \\object{SN~1998bw} \\citep{gal98} was initially questioned because of the low redshift of the supernova host galaxy \\object{ESO 184-G82}, $z=0.0085$.  The corresponding isotropic energy output of GRB~980425 was low (E~$\\sim~10^{49}$ erg) compared with other GRBs known at that time (at $z~\\gtrsim~1$ with E$~\\sim~10^{52}$ erg). Since then, more apparently under-luminous, low redshift GRB events have been detected. The definite proof of a supernova connection originated with the discovery of SN Type Ic signatures in the afterglow spectra of a couple of long duration GRBs \\citep{hjorth03,stanek03,malesani04}.  While the energetics of GRB~980425 suggested an underluminous event, there was no indication that its accompanying supernova was peculiar for a GRB-SN \\citep{kaneko07}.    About 500 GRBs have been observed and optical afterglows have been detected for about a third\\footnote{\\tt http://www.mpe.mpg.de/$\\sim$jcg/grbgen.html}.  While bursts detected by the BeppoSAX satellite typically had redshifts $z\\sim1$, those detected by the Swift satellite \\citep{gehrels04} have an average redshift of 2.2 \\citep{jakobsson06}\\footnote{\\tt http://www.astro.ku.dk/$\\sim$pallja/GRBsample.html}.  The small angular size of the high redshift host galaxies (typically less than a few arcseconds) implies that a closer examination of the nature of the hosts requires integrated spectra or broad band images for the faintest hosts. Images of GRB hosts indicate that GRBs are located preferentially in sub-luminous \\citep{lefloch03} actively star-forming, young galaxies \\citep{christensen04b}. The distribution of GRB afterglow locations within their hosts follow in general the bright continuum light distribution of the host galaxies \\citep{bloom02a,fruchter06}.  Models of single star collapsar progenitors of long duration bursts predict that they are more likely to reside in metal poor environments \\citep{woosley93}.  Low metallicities in the surrounding medium of GRBs are in general measured by absorption line studies of GRB afterglows at $z\\gtrsim2$ \\citep{savaglio03,vreeswijk04,chen05,dessauges06,prochaska07}. This technique is used mostly to infer iron group elements along the line of sight towards the GRB. High resolution spectra of GRB afterglows have indicated that the medium probed by the spectra resides at distances larger than several tens of pc from the burst site \\citep{prochaska06b,chen07}, and a detailed study of one burst (GRB~060418) discovered that the environment probed by the absorption lines was at a distance of 1.7~kpc from the burst site \\citep{vreeswijk07}.  The GRB afterglow does therefore not probe the medium immediate surrounding the burst itself, but rather the interstellar medium of the host; this implies that perhaps we are unable to estimate the metallicity of the progenitor directly using this technique.   Since the hosts of long duration bursts are undoubtedly star-forming galaxies, the emission lines from the integrated optical spectra can be used to infer the metallicities provided that the redshift of the GRB is less than approximately one where the strong lines are visible in the optical. This technique is generally used to derive the oxygen abundance. Diagnostics of strong emission lines for the hosts of \\object{GRB~980425}, \\object{GRB~020912}, \\object{GRB~030329}, \\object{GRB~031203}, and \\object{GRB 060218} have indicated sub solar metallicities \\citep{prochaska04,hammer06,gorosabel05,sollerman05,wiersema07}, and systematically lower host galaxy metallicities than found in nearby ordinary Type Ic supernovae sites \\citep{modjaz07}.  It is, however, unclear how measurements derived from the observed integrated GRB host spectra relate to the overall host properties, or to the site in which the GRB exploded.  If the host galaxy is clearly extended, an analysis of the spatially resolved emission lines can uncover the characteristics of the precise GRB region, or alternatively the specific environment conditions required to make a GRB. \\citet{thoene08} investigated the spectra of the host of GRB~060505, and found that the GRB site was the youngest and most metal poor of all regions probed by the slit spectrum.  For the host of GRB~980425/SN~1998bw (\\object{ESO~184-G82}), a dwarf SBc galaxy, the SN occurred in a star-forming region in the spiral arm. The region was identified by high spatial resolution images acquired with HST/STIS \\citep{fynbo00} and confirmed by the observation of a fading source in this region \\citep{sollerman02}. The location of the SN was offset by a projected distance of 0.8~kpc from the strongest star-forming region in the host, which had distinct spectral signatures of Wolf-Rayet stars \\citep{hammer06}. The spectrum of the SN region did not have any detected WR features.  Other GRB hosts at $z\\lesssim0.1$ did not show clear signatures of WR stars in emission either \\citep{margutti07,wiersema07}.  At higher redshifts, the absence of \\ion{C}{iv} $\\lambda$ 1550 absorption lines in winds of WR stars can be explained by the afterglow, which ionises the surrounding medium to a large radius (\\citet{chen07}, but see also \\citet{starling05})  All previous analyses of emission lines from GRB hosts were limited to regions within the slit. Hence, the data could be affected by slit-losses and leading to an under-estimated integrated star formation rate (SFR).  In this paper, we present data for the GRB~980425 host galaxy obtained with integral field spectroscopy.  The data from the VIMOS integral field unit (IFU) presented in Sect.~\\ref{sect:data} allows us to study a significant fraction of the surface of the host with a resolution of 0.27~kpc, and to analyse most of its high surface brightness regions (see Fig.~\\ref{fig:hostim}). We use the data to compile maps of emission lines and their ratios and determine the reddening, SFRs, specific SFRs, metallicities, densities, temperatures, and kinematic and stellar masses, as presented in Sect.~\\ref{sect:results}. We compare the emission line ratios with those of other GRB hosts in Sect.~\\ref{sect:ratios}, and show evidence that, while abundances may vary, the regions are in all cases very young.  In Sect.~\\ref{sect:disc}, we discuss the implications for the unresolved high redshift GRB hosts. We adopt a cosmology with H$_0$=70 km~s$^{-1}$~Mpc$^{-1}$ throughout the paper. ",
        "conclusions": "\\label{sect:disc} \\subsection{Implication for high-z GRB hosts} In this paper, we have studied the properties of resolved regions in the very nearby host galaxy of GRB~980425/SN~1998bw. We found that the \\ion{H}{ii} region where the GRB/SN occurred has the second lowest oxygen abundance, while the lowest value arises in the WR region. Apart from the metallicity, other physical properties in the SN region are representative of the integrated spectrum of the host galaxy.  The variations, especially those of the metallicity and extinction, over the face of the galaxy have important implications for the interpretation of the observations of higher redshift GRB hosts. In the case that $z>1$ GRB hosts do not contain just one (giant) star forming region or star cluster, some properties will not be recoverable when the only information available is the integrated properties of the entire galaxy. In several GRB hosts, the environment appears to be complex with a few or several individual star-forming regions that can only be resolved at the spatial resolution of the HST \\citep[e.g.][]{hjorth02}.  Based on the study of the resolved population of the GRB 980425 host, we identified the following effects that would occur if the host was at higher redshift and unresolved   \\begin{itemize} \\item The integrated spectra from the separate regions would decrease the emission line equivalent widths corresponding to individual regions, because older stellar populations would be covered. This would then substantially decrease the inferred age of unresolved star burst regions.\\\\  \\item The specific star formation rates would be much higher in individual \\ion{H}{ii} regions than the integrated one where less active regions are also covered. \\\\  \\item For the GRB 1998bw host galaxy, the scatter in oxygen abundances between different regions would be relatively small compared to the overall uncertainty in the calibrations. Abundances would be similar to within the 3$\\sigma$ uncertainties. For high redshift GRB hosts, the integrated spectra may produce a fairly good representation of the abundances in the GRB regions.\\\\  \\item We found that the luminosity SSFR introduced in \\citet{christensen04b} for GRB hosts is a good proxy for the mass SSFR, which is more frequently used in the literature. Furthermore, the integrated SSFR of the GRB~980425 host is similar to other GRB hosts.\\\\  \\item The extinction for small isolated regions can be severely underestimated. The overall host has a low reddening of \\ebv=0.3 mag, which is normal for (unresolved) higher redshift GRB hosts. The SN region itself has a higher than average value. Relative to spectral energy distribution fitting of GRB host broad band data, the emission line ratios indicated a higher extinction.\\\\  \\item The emission line ratios for the different regions in the host indicated that most other GRB hosts (at $z<0.4$) have properties that would seem similar to the WR region in the GRB~980425 host. This implies that even though we cannot resolve spatially some GRB hosts, they must be dominated by very young stars, even though we cannot always distinguish any WR features in their spectra.\\\\  \\item The separation of 800 pc between the WR and SN region corresponds at $z=1$ to a single pixel with the HST spatial resolution of 0\\farcs1. Had the host been at $z=1$, we would have been able to resolve only two distinct regions: the centre of the galaxy and the WR region. In a broad band HST image, the GRB location would coincide with the brightest pixel as seen for many cosmological GRBs \\citep{fruchter06}.  \\end{itemize}  \\subsection{GRB host metallicities} If low metallicities are a necessary prerequisite to creating GRBs, this effect should be immediately apparent from the properties of the host galaxies. Theoretical models suggest that GRBs correspond to star forming galaxies with lower than 10\\% solar metallicity \\citep{yoon06}.  Models for winds from WR stars are based on Fe as drivers \\citep{vink05}, and likewise the GRB progenitor models are based on the Fe abundance \\citep{yoon05} where no GRB progenitors are expected from stars with metallicities larger than 0.1--0.3 solar \\citep{woosley06,langer06}.  However, while the models are based on the Fe metallicity, the emission lines from the host galaxies trace the oxygen abundance.  As demonstrated in Fig.~\\ref{fig:n2} and \\ref{fig:o2}, regions with higher metallicities than the model predictions are generally found among GRB hosts, although these are still below the solar value \\citep[see also][]{savaglio08}.  A potential bias comes from the fact that the hosts are generally very faint, so the metallicity determination has to rely on strong emission line diagnostics, which have large uncertainties.  The use of strong emission lines to derive abundances can in some cases infer a higher metallicity estimate than direct measurement from the [\\ion{O}{iii}]~$\\lambda$4363 or the [\\ion{N}{ii}]~$\\lambda$5755 emission lines. In the WR region, the strong line diagnostics infer a 0.3 solar oxygen abundance, while the direct measurement indicates a near solar value. In contrast, the O3N2 diagnostics for the GRB~060218 host indicates a 30\\% solar oxygen abundance, while the direct method gives just 8\\% solar \\citep{wiersema07}.   GRB hosts at $z>2$ have metallicities of about 10\\% solar in general as estimated from absorption lines in the afterglows \\citep{savaglio06}.  However, as shown by \\citet{vreeswijk07} for the case of GRB~060418, the rest frame UV absorption lines do not arise in the immediate GRB progenitor environment, but instead the interstellar medium of the host galaxy along one single line-of-sight.  On the other hand, as we have discussed in this paper, the GRB hosts at low redshift are not necessarily very metal poor. The question is how to relate the low and high redshift sample of GRB hosts to each other when the observing techniques are completely different.  Different biases are included from the assumption that the metallicity of the progenitor can be constrained from absorption lines in the GRB afterglow, and that emission lines from the host galaxy can infer the abundance at the progenitor site. Any of these metallicities may not be representative of the GRB progenitor itself.  To understand how these two methods compare with each other, it is necessary to detect the emission lines and derive abundances from GRB hosts at $z>2$ in a similar way as performed routinely for local galaxies, as shown in this paper.  By comparing the absorption metallicity, determined from the GRB afterglow, with abundances derived from emission lines it can be verified if the metallicity is representative of the host as a whole, and to place the GRB hosts in relation to star-forming field galaxies at high and low redshifts.  The comparison of absorption lines and emission line spectroscopy for a high redshift GRB host galaxy can provide a measure of the Fe/O ratio.  Since the GRB hosts are young and show intense SSFRs, it is likely that the oxygen abundance is higher than that derived from Fe due to the relatively small contribution from SN Type Ia. Hence, the determination of the Fe/O ratio can provide constraints on the star formation history in galaxies fainter than those generally studied in flux-limited surveys."
    },
    "0807/0807.3624_arXiv.txt": {
        "abstract": "{} {We describe a method for identifying XMM-Newton observations that have been affected by Solar Wind Charge Exchange (SWCX) emission and present preliminary results of previously unidentified cases of such emission within the XMM-Newton Science Archive.} {The method is based on detecting temporal variability in the diffuse X-ray background. We judge the variability of a low energy band, taken to represent the key indicators of charge exchange emission. We compare this to the variability of a continuum band, which is expected to be non-varying, even in the case when SWCX enhancement has occurred.} {We discuss previously published results with SWCX contamination that have been tested with the above method. We present a selection of observations that we consider to show previously unpublished SWCX enhancements, and further investigate these observations for correlation with data from the solar wind observatory, ACE. We also consider the geometry and viewing angle of XMM-Newton at the time of the observation to examine the origin of the charge exchange emission, whether it be from interactions with geocoronal neutrals in Earth's magnetosheath or from the heliosphere and heliopause.} {} ",
        "introduction": "This paper is the first of two articles which describe a new method developed to search for instances of Solar Wind Charge Exchange (SWCX) contamination of XMM-Newton observations. In this paper we outline the method and provide an initial sample of results. In paper II we will present the results of applying this method to all suitable observations within the public XMM-Newton Science Archive (XSA).  In this introduction we briefly discuss the SWCX process including examples where this effect occurs in the solar system, and then describe the XMM-Newton observatory and how its orbit affects the likelihood of observing such a process. We describe the Advanced Composition Explorer (ACE) satellite which we use to monitor the solar wind at times of interest.  In Section \\ref{secmethod} of the paper we detail the steps employed to identify observations possibly affected by SWCX emission and in Section \\ref{secres} we present the results from the use of this method for an initial small sample of approximately 170 observations. We conclude by discussing our plans for the application of this method to incorporate data from the entire publicly available XSA from launch until the commencement of this second phase. As this time period will incorporate more than seven years of XMM-Newton operations, we hope to detect any relationship between our results and the solar cycle.  \\subsection{Solar wind charge exchange}\\label{secswcx} SWCX occurs when an ion in the solar wind interacts with a neutral atom and gains an electron in an excited state (Eq. \\ref{eqnswcxgive}). If the ion is in a sufficiently high charge state in the subsequent relaxation of the ion, an X-ray or UV photon is released (Eq. \\ref{eqnswcx}).  \\begin{equation} A{^q}{^+} + B \\to  A^{(q-1)^{\\ast}} + B{^+} \\label{eqnswcxgive} \\end{equation} \\begin{equation} A^{(q-1)^{\\ast}} \\to A^{(q-1)} + h\\nu \\label{eqnswcx} \\end{equation}  A is the solar wind ion, for example oxygen, and B the neutral atom, commonly hydrogen or helium in the heliosphere or just hydrogen in the case of geocoronal neutrals in the Earth's magnetosheath.  A SWCX spectrum is characterised by emission lines corresponding to the ion species present in the solar wind. Although the solar wind is approximately 99\\% protons, electrons and $\\alpha$-particles, the remainder of the wind consists of heavier elements such as C, O, Mg, Si, Fe and Ni. The heavier elements found in the solar wind are often of high charge state due to the high temperatures of the solar corona. The relative abundances of the elements are dependent on the conditions of the Sun when the solar wind leaves the solar corona. In this project we concentrate our search on oxygen emission lines as these are typically the strongest SWCX lines in previously published SWCX-enhanced XMM-Newton observations, e.g. \\citet{snowden2004}.  SWCX occurs at several locations in the solar system including cometary tails \\citep{greenwood2000, cravens1997}, at the heliospheric boundary \\citep{cravens2000, koutroumpa2007}, planetary atmospheres \\citep{Dennerl2002, branduardiraymont2007} and from the Earth's magnetosheath or interplanetary neutrals \\citep{cravens2001, wargelin2004, snowden2004, fujimoto2007}. We concentrate on charge exchange between the solar wind and geocoronals in the Earth's magnetosheath and contributions within the heliosphere in the vicinity of the Earth, as this X-ray emission is the most likely to affect the observations of XMM-Newton in terms of a short, temporally varying background source. Local x-ray intensities vary on much shorter timescales (on the order of hours) than those of the heliospheric-ISM boundary contributions, which vary on much longer timescales and with less variation in intensity \\citet{robertson2003b}. The main variation seen in the heliospheric contribution occurs in conjunction with changes in the line of sight. As the solar wind travels at approximately $\\frac{1}{4}$ AU per day, the heliospheric emission is the result of the average conditions of the solar wind over several months.  \\citet{robertson2003a}, \\citet{robertson2003b} and \\citet{robertson2006} have modelled the expected X-ray emission from charge exchange with geocoronal neutrals, using \\citet{hodges1994} model of the exosphere to account for the densities of the geocoronal neutrals with height for differing solar wind conditions, mapping the area of strongest X-ray emission on the sunward side of the magnetosheath (at an intensity of approximately 8 keV cm$^{-2}$s$^{-1}$ sr$^{-1}$ for nominal solar wind conditions, rising to 200 keV cm$^{-2}$ s$^{-1}$ sr$^{-1}$ for simulated solar storm conditions). The solar wind speed can vary dramtically depending on storm or slow wind conditions. A slow solar wind of approximately 300 km s$^{-1}$ and average solar wind densities results in a nominal solar proton flux of approximately 2$\\times$10$^{8}$ cm$^{-2}$ s$^{-1}$. The exact distances from the Earth to the magnetopause (the boundary layer between the solar wind plasma and the plasma of the Earth's magnetosphere) and bow shock (which together define the magnetosheath region) depend on the strength of the solar wind at any given time \\citep{khan1999} but are on average at approximately 9 and 15 R$_{E}$ respectively.  SWCX was first suggested as an explanation for the temporal variations of a few hours to a few days length in the diffuse emission data of the ROSAT All-Sky Survey \\citep[long-term enhancements,][]{snowden1995} by \\citet{cox1998} and \\citet{cravens2000}. SWCX has been seen by XMM-Newton, Chandra and Suzaku, and has been attributed to geocoronal emission (e.g. \\citet{fujimoto2007}). Multiple observations with the same pointing direction of the Hubble Deep Field North region taken with XMM-Newton have shown temporal variations in count rates and line emission that has been correlated to enhancements in the level of the solar wind flux and the solar storm of May 2001 \\citep{snowden2004}. Also, depending on viewing geometry, it would be possible to look along the front of a Coronal Mass Ejection (CME) density enhancement, which could cause short term and extreme variations in the level of X-ray flux observed \\citep{collier2005}.  \\subsection{The XMM-Newton observatory and viewing geometry}\\label{secxmm} The XMM-Newton observatory \\citep{jansen} has been used extensively in the study of extended and diffuse X-ray sources. For such studies, a good understanding of the X-ray background is required, including the SWCX, which may affect a user's observation.  Due to the supposed short term variations in the SWCX at the Earth's magnetosheath an observation may be partly or wholly affected by enhancements at energies characteristic of the SWCX line emission. The effect is also of interest in its own right as a monitor of the ionic composition of the solar wind and the interaction between the Earth's magnetosheath and the bombardment of the solar wind.  XMM-Newton is not able to study the SWCX at all times due to the configuration of it's highly elongated elliptical orbit. Optimum viewing of the magnetosheath and areas of the expected highest X-ray emission within the magnetosheath \\citep{robertson2003b} is expected to occur at only certain points during the 48 hour orbit at certain times of the year, see Figure \\ref{figxmmorbit}. The situation is further complicated by restrictions on the viewing angle, which is strongly constrained by the fixed solar panels and those imposed to protect the instruments from directly viewing the Sun and the X-ray and optically bright Earth and Moon, and the non-operational period during the telescope's passage through perigee as it passes through the radiation belts. This results in the observatory being able to view the area of maximum SWCX emission, i.e. the nose of the magnetosheath, only at certain times of the year. XMM-Newton is therefore more likely to experience SWCX-emission enhancement during the summer months (northern hemisphere) as this is when the telescope is able to be in a position to look through the nose and bright flanks of the magnetosheath. However, our goal is to study all observations in the XMM-Newton archive assuming they were conducted in an appropriate observational mode. Coupled with the information regarding the solar wind flux we will, with this project, gain understanding regarding the likelihood and level of SWCX-contamination an XMM-Newton observation may experience.  \\begin{figure}[] \\vspace{1cm} \\includegraphics[width=0.45\\textwidth]{new1507_xmm_orbit_poster_compressed.ps} \\caption{A schematic illustratration of the orbit of XMM-Newton in relation to the magnetopause and bow shock (dotted lines), which define the boundaries of the magnetosheath, at different times during the year. The left image shows the highly elliptical orbit of XMM-Newton during mid summer (northern hemisphere), and the right image shows the winter configuration. The Sun has been added to show the direction of the Sun and therefore incoming solar wind.} \\label{figxmmorbit} \\end{figure}  \\subsection{The Advanced Composition Explorer (ACE)}\\label{secace} The Advance Composition Explorer (ACE) was launched in 1997, into orbit at the Earth-Sun Langrangian L1 point about 1.5 million miles from Earth. The instrument package on ACE was designed to study solar, galactic and extragalactic energetic particles (see: \\textit{http://www.srl.caltech.edu/ACE/ace\\_mission.html}).  In this project we use ACE to provide additional supporting evidence that a SWCX-enhancement may have been seen in an XMM-Newton observation. We use hourly average data (\\textit{Level 2}) from the Solar Wind Electron, Proton, and Alpha Monitor (SWEPAM) instrument on board ACE. Correlations of increased emission in the SWCX lines with increases in the solar wind proton flux lend credence that a particular observation may have experienced a SWCX-enhancement. However the exact scenario maybe complicated by the orientation of the interplanetary magnetic field, adjusting the delay time between ACE and Earth \\citep{weimer2003}, so conclusions drawn from correlations between increased X-ray and solar proton flux should be treated with caution. We use data from ACE as a general monitor of the condition of the solar wind flux. However, the delay between ACE and the Earth is approximately one hour. ",
        "conclusions": "We have presented a selection of results from a preliminary sample of approximately 170 XMM-Newton observations that have been tested for oxygen line variability as compared to a continuum band that is not expected to vary. Our primary measure of the likelihood of SWCX-enhancement is found through the lack of correlation between a count rate for the prime SWCX indicators, namely the band representing the OVII and OVIII major emission lines, and the count rate for a band that is SWCX-enhanced emission free. SWCX-enhanced observations showed both a high \\redc\\ and \\cratio\\ value. In the event of soft proton contamination, correlation between the count rates from both bands is still expected. The level of anti-correlation is assessed using the metric of the \\redc\\ from the relationship between a band representing the SWCX indicators and a band representing the emission line free continuum. We have demonstrated that the method has been successful in identifying cases of SWCX from a medium sized sample, providing careful consideration of various other diagnostic tests are taken into account. In addition, our method was also able to identify all but one of the control observations with previously identified SWCX that had been placed in the sample, within the top 35 \\redc-ranked results.  Our strongest case of SWCX-enhancement, with the richest spectrum showing many emission lines, gave the highest level of the \\redc\\ value and presented a high \\cratio. This observation was contaminated by residual soft protons, although the resultant spectrum, between the non and SWCX-enhanced periods, was able to show many emission lines corresponding to highly charged ion species.  Cases which we believe present the best evidence of SWCX-enhancement exhibit enhanced spectral features such as emission lines of OVII and OVIII, along with lower energy lines probably associated with carbon. Two of the strongest cases offer emission lines that may be attributed to NeIX and MgXI. The pattern of observed emission lines depends on the exact composition of the solar wind at 1AU at the time of the observation.  We have shown that several of the cases where SWCX-enhancements are suspected occurred during periods when XMM-Newton was found to be on the sub-solar side of the magnetosheath. The line of sight of the telescope in these cases intersects the area of the magnetosheath of brightest expected solar wind-induced geocoronal X-ray flux. The SWCX-enhancement seen for other cases, when XMM-Newton was found to have a line of sight which intersects the magnetosheath through the weaker areas of X-ray flux, is probably due to a non-geocoronal origin, such as a CME density enhancement passing through the heliosphere.  In Paper II we will present results from the processing of a larger sample of observations to identify additional observations with SWCX-enhancement, extending our analysis to include data from the PN camera. We expect to find approximately 2000 observations of XMM-Newton affected by temporally variable SWCX, based on frequency of detections from the sample described in this paper. However, the number of detections is dependent on solar activity and therefore may vary considerably from the number stated. Using a larger sample, we aim to refine the identification of SWCX-enhancement, using the relationship between \\redc\\ and \\cratio.  We will include an improved algorithm to remove sources in the field, reducing the danger of residual source contamination. We will investigate for any trends within the results that are related to the solar cycle and the position angle and orientation of XMM-Newton during each observation. This will help us to constrain the conditions when XMM-Newton is most likely to experience SWCX, whether that be due to the satellite's position or viewing geometry. We will also explore multiple pointings of the same target to consider variations between observations. We will investigate the observed SWCX line ratios and what this implies for the composition of the solar wind."
    },
    "0807/0807.0210_arXiv.txt": {
        "abstract": "A new multi-epoch H$\\alpha$ imaging study of M101 (NGC~5457) has been carried out as part of a larger campaign to study the rate and stellar population of extragalactic novae. The survey yielded a total of 13 nova detections from 10 epochs of M101 observations spanning a three year period. After correcting for the temporal coverage and survey completeness, a global nova rate of $11.7^{+1.9}_{-1.5}$ yr$^{-1}$ is found. This value corresponds to a luminosity-specific nova rate of $1.23\\pm0.27$ novae per year per $10^{10}~L_{\\odot,K}$ when the $K$ luminosity is derived from the $B-K$ color, or $1.94\\pm0.42$ novae per year per $10^{10}~L_{\\odot,K}$ when the $K$ magnitude from the Two Micron All Sky Survey is used. These values are consistent with previous estimates by Shafter et al. that were based on more limited data. The spatial distribution of the combined nova sample from the present survey and from the earlier Shafter et al. survey shows that the specific frequency of novae closely follows the integrated background light of the galaxy. ",
        "introduction": "Novae occur in semi-detached binary systems in which a white dwarf (WD) primary accretes material from a Roche-lobe filling late-type companion (see Warner 1995 for a review). For sufficiently low mass accretion rates, the accreted matter builds up on the surface of the WD under degenerate conditions. When the Fermi temperature is reached at the base of the accreted layer, a thermonuclear runaway (TNR) ensues, blowing off the outer layers of the star, and resulting in an outburst of 10--20 magnitudes. In the resulting explosion, novae reach peak luminosities of up to $M_{V} \\sim -10$, making them visible in external galaxies up to and beyond the Virgo cluster. The high luminosity, coupled with a well-known relationship between a nova's maximum magnitude and its rate of decline, the MMRD relation (Zwicky 1936, Mclaughlin 1945, Downes \\& Duerbeck 2000), has enabled novae to be useful as an extragalactic distance indicator (e.g. Pritchet \\& van den Bergh 1987, Della Valle \\& Gilmozzi 2002). However, difficulties in calibrating the MMRD relationship, as well as uncertainties about the uniformity of the progenitor populations in different galaxy types have hindered the use of novae as reliable distance indicators (Ferrarese et al. 2003). Recent studies have focused mainly on the value of using novae to probe extragalactic stellar populations.  Galactic nova studies have suggested that there are actually two distinct classes of classical novae, a disk population and a bulge population (e.g., Duerbeck 1990, Della Valle et al. 1992; Della Valle \\& Livio 1998). As shown by Della Valle et al. (1992), ``fast\" novae\\footnote{Fast novae can be loosley defined as those novae which decline at least two magnitudes within 12 days (i.e. $t_2\\leq12$~days).} tend to be associated with a disk population, having an average scale height above the Galactic plane generally less that that of slower novae. At about the same time Williams (1992) pointed out that novae could also be segregated into two classes based upon their spectroscopic characteristics. Novae whose spectra display prominent Fe~II lines (the Fe~II novae) tend to evolve more slowly, have lower expansion velocities, and have a lower level of ionization compared with novae that exhibit strong lines of He and N (the He/N novae). In addition, these He/N novae often display very strong neon lines, which suggests that the seat of the eruption is a relatively massive ONe white dwarf. Such novae do not appear to produce the copious carbon-rich dust that is often formed in nova ejecta arising from the lower mass CO white dwarfs (Gehrz et al. 1998). Additional support for the two-population scenario was provided by Della Valle \\& Livio (1998), who looked into the spatial distribution of a sample of 27 Galactic novae with known spectral class and reliable distance estimates. They noted that Galactic novae that could be classified as He/N were more concentrated to the Galactic plane, and tended to be faster and more luminous compared with their Fe~II counterparts.  Despite the progress made in understanding nova populations with Galactic data, observations of novae in nearby galaxies offer several advantages over Galactic studies. For a given galaxy, the nova sample is essentially equidistant, with constant (and usually minimal) foreground reddening, making it easier to study relative nova luminosities and fade rates. For spatially-resolved spiral systems it is often possible to segregate novae from different spatial positions within a galaxy (e.g. between the bulges and disks of spirals), revealing possible differences between novae from differing stellar populations. Perhaps most importantly, extragalactic nova observations make it possible to study novae from isolated stellar populations directly by comparing their properties in galaxies along the Hubble sequence from ellipticals to irregulars.  A number of studies have looked at the question of whether the luminosity-specific nova rate (LSNR) of a galaxy varies with Hubble type. In these studies, the nova rate of a galaxy is normalized to its infrared luminosity (usually $K$-band), which acts as a proxy for stellar mass. In the first such study, Ciardullo et al. (1990b) compiled a list of galaxies with measured nova rates and concluded that the LSNR was essentially independent of Hubble type. Shortly thereafter, Della Valle et al. (1994) re-analyzed much of the same data (using different estimates for completeness and for the normalization of the nova rate by galaxy luminosity) and came to a different conclusion. Namely, that the LSNRs of the late-type, low-mass galaxies in the local group (SMC, LMC, and M33) were significantly higher than those of the other, earlier-type galaxies, such as the Virgo ellipticals.  In an important paper, Yungelson et al. (1997) published the results of population synthesis models demonstrating that the LSNR in a galaxy should be strongly dependent on the galaxy's star formation history. There are two reasons for this dependence. The first is that the mean mass of the accreting white dwarfs in newly-formed nova binaries declines with increasing time elapsed since the formation of the zero-age main sequence system (Tutukov \\& Yungelson 1995). Since the frequency of nova eruptions increases with increasing white dwarf mass (Ritter et al. 1991; Livio 1992), galaxies experiencing recent ($1-2\\times10^9$~yr) star formation, such as late-type spirals, should be more prolific nova producers compared with galaxies that formed the bulk of their stars in a burst of star formation early in their history (i.e. ellipticals). Secondly, the color of a galaxy is dependent on its star formation history. The higher the recent star formation rate, the ``bluer\" the galaxy (Bruzual \\& Charlot 2003). Thus, for a given nova rate per unit mass, late-type galaxies will have a higher $K$-band luminosity-specific nova rate when compared with early-type galaxies.  In addition to the dependence on Hubble type, Yungelson et al. (1997) pointed out that galaxy mass was a key factor affecting the LSNR. In particular, they pointed to work by Gavazzi \\& Scodeggio (1996) who have argued that giant late-type spiral galaxies ($L_H\\sim10^{11}~L_{\\odot}$) have an exponentially decaying star formation rate ($\\tau\\sim0.5\\times10^9$~yr), while dwarf late-type galaxies ($L_H\\sim10^{8}~L_{\\odot}$) have a more nearly constant star formation rate characterized by $\\tau\\grtsim5\\times10^9$~yr. Thus low-mass spirals and irregular galaxies have had more active star formation in the recent past compared with the more massive spiral systems.  Overall, the models of Yungelson et al. (1997) seem to support the conclusions of Della Valle et al. (1994) who found the highest LSNRs in the low-mass, late-type galaxies such as the LMC, SMC and M33. The one galaxy type missing from the analysis were the massive, late-type spirals, for which no nova rates had yet been measured. In an attempt to remedy this deficiency, Shafter et al. (2000) undertook nova surveys in the massive late-type galaxies M51 and M101, with the giant elliptical M87 included for comparison as a massive early-type galaxy. Ultimately, the Shafter et al. (2000) study and a similar study targeting M33 by Williams \\& Shafter (2004) failed to confirm any significant correlation between LSNR and Hubble type.  In recent years, the results of the Shafter et al. (2000) have been called into question by the work of Shara and collaborators. In an analysis of archival HST images of M87, Shara \\& Zurek (2002) reported a (preliminary) nova rate significantly higher than that found by Shafter et al. in their ground-based study. This result prompted Neill \\& Shara (2004) to contemplate the possibility that many ground-based, multi-epoch nova surveys with sparse temporal coverage may have systematically underestimated nova rates because the correction for the effective survey times is often large and uncertain.  Given the importance of establishing reliable nova rates for galaxies spanning a wide range of Hubble types, and for the massive late-type spiral galaxies in particular, we have undertaken a new survey for novae in M101 using the 90Prime camera on the Bok 2.3-m reflector. As a relatively nearby ($\\mu_0=29.13\\pm0.11$ [Freedman et al. 2001]), nearly face-on, late-type spiral, M101 provides an ideal candidate for studying the nova rate and the distribution of nova progenitor populations. This survey improves on the one conducted by Shafter et al. (2000) both by monitoring the galaxy more frequently, and by improving the corrections for spatial and temporal completeness. In this paper, we report the results of this new survey. ",
        "conclusions": "An H$\\alpha$ survey for novae in M101 was carried out in 10 epochs over a period of $\\sim$3~yr, with a total of 13 novae identified. After correcting for the effective survey time, we arrive at a global nova rate of $11.7^{+1.9}_{-1.5}$~yr$^{-1}$, which corresponds to a luminosity-specific nova rate of $1.23\\pm0.27$ novae per year per $10^{10}~L_{\\odot,K}$ when the $K$ luminosity is derived from B$_{tot}$ and the $B-K$ color, and $1.94\\pm0.42$ novae per year per $10^{10}~L_{\\odot,K}$ when the $K$ magnitude is taken from the 2MASS measurements (Jarrett et al. 2003). Both values are consistent, within the stated uncertainties, with the previously published value of $1.27\\pm0.46$ novae per year per $10^{10}~L_{\\odot,K}$ given in Williams \\& Shafter (2004). As in previous determinations of the nova rates in spiral galaxies, a major source of uncertainty in the derived rates is introduced by our inability to properly account for the effects of extinction internal to M101. In this regard our estimates for the global nova rate and the LSNR in M101 could more conservatively be considered as lower limits to their true values.  Future observations should be focused on addressing three remaining sources of uncertainty in computing the LSNR. As described earlier, more frequent temporal sampling would be useful, not only to minimize uncertainty in computing the effective survey time, but also in establishing a larger sample of template light curves to be used in the Monte Carlo simulations. The availability of light curves from novae arising from galaxies with a wide range of Hubble types will establish whether the properties of the light curves are sensitive to stellar population. In this regard, we expect that synoptic survey telescopes such as the PanStarrs and the LSST will revolutionize the study of novae in extragalactic systems. Finally, there remains the possibility that a significant fraction of the novae, particularly in late-type galaxies such as M101, may be missed because of extinction within the galaxies themselves. Perhaps future observations with wide-field infrared imagers can help alleviate this concern."
    },
    "0807/0807.0026_arXiv.txt": {
        "abstract": "{ Radio telescopes with off-axis feeds, such as the (E)VLA, suffer from ``beam squint'' in which the two orthogonal circular polarizations sampled have different pointing centers on the sky. Its effects are weak near the beam center but become increasingly important towards the edge of the antenna power pattern where gains in the two polarizations at a given sky position are significantly different. This effect has limited VLA measurements of circular polarization (Stokes~V) and introduced dynamic range limiting, wide-field artifacts in images made in Stokes~I. We present an adaptation of the visibility-based deconvolution CLEAN method that can correct this defect ``on the fly'' while imaging, correcting as well the associated self-calibration. We present two examples of this technique using the procedure ``Squint'' within the Obit package which allows wide-field imaging in Stokes~V and reduced artifacts in Stokes~I.  We discuss the residual errors in these examples as well as a scheme for future correction of some of these errors. This technique can be generalized to implement temporally- and spatially-variable corrections, such as pointing and cross-polarization leakage errors. \\keywords {Techniques: image processing, Techniques: interferometric} } ",
        "introduction": " ",
        "conclusions": ""
    },
    "0807/0807.2353_arXiv.txt": {
        "abstract": "{} {It is commonly assumed that the two indirect exoplanet detection methods, the radial velocity method and astrometric method, require observational periods exceeding the orbital period to produce positive results. Here we test this assumption in detail. We also investigate the smallest ratio of observational timeline and orbital period required for positive detections.} {We obtain full information on the orbital parameters by combining radial-velocity and astrometric measurements by means of Bayesian inference, and sample the parameter probability densities of orbital and other model parameters with a Markov Chain Monte Carlo (MCMC) method in simulated observational scenarios to test the detectability of planets with orbital periods longer than the observational timelines.} {We show that, when fitting model parameters simultaneously to measurements from both sources, it is possible to extract much more information from the measurements than when using either source alone. Currently available high-precision measurements of radial velocity (with 1ms$^{-1}$ precision) and astrometric measurements achievable with the SIM space telescope (with a precision of 1$\\mu$as) can be used together to detect a Jupiter analog 30pc away with an observational timeline of only three years, approximately one fourth of the orbital period. Such measurements are sufficient for determining all its orbital parameters, including inclination and the true mass. Also, with accurate radial velocity measurements covering a timeline of 20 years, the true mass could be determined by astrometric observations within a single year. These case studies demonstrate the potential power of the Bayesian inference of multiple data sources in exoplanet observations. As an example, we show that using the currently available radial velocity measurements, the inclination of HD 154345b could be determined with SIM in a year.} {} ",
        "introduction": "Since the discovery of 51 Peg b (Mayor \\& Queloz 1995), high-precision radial velocity (RV) measurements have been used successfully to detect planetary companions of nearby stars (Butler et al. 2006). The finest instruments to date are capable of achieving an RV precision of 1ms$^{-1}$ (Santos et al. 2004; Moorwood and Masanori, 2004). However, the exact nature of the companions remains unknown when only RV measurements are available, as these only yield the product of mass and sine of the orbital inclination, giving the lower limit for the mass. Recently, it has been claimed that an RV precision of 1cms$^{-1}$ could be technically possible in the future (Li et al. 2008). However, the RV variations of stars of approximately 5ms$^{-1}$ to 50ms$^{-1}$ for K to F stars, caused by star spots and irregular convective zones, will likely prevent the detection of signals of the cm scale (Saar \\& Donahue 1997; Saar et al. 1998).  Astrometric (A) measurements of the position of the target star in the sky as a function of time can be used to detect the inclination, and consequently the true mass of the companion. But despite several trials (van de Kamp 1969; Neuhaeuser et al. 2008), the precision of these measurements has not been high enough to verify the planetary nature of RV companions.  With the aid of telescopes and instruments capable of high-precision astrometry (SIM, GAIA, PRIMA, etc.); however, this situation is about to change. With an estimated astrometric precision of future telescopes of 1.0$\\mu$as (Unwin et al. 2008) or 8-10$\\mu$as (Casertano et al. 2008; Derie et al. 2002), it will be possible to determine the masses of the already detected RV companions. It is commonly assumed that these detections require observational periods longer than the orbital period of the target system to be able to detect the periodic signal. Here we test this assumption in detail.  In this article we simulate astrometric and RV measurements to study the possibility of detecting planetary companions of nearby stars with various observational timelines. The goal is to find the minimum timeline required for detecting a planetary companion using high-precision RV and astrometric measurements. These two sources of data are combined by means of Bayesian inference, and the probability densities of orbital and inertial reference frame parameters are sampled using Markov Chain Monte Carlo (MCMC) (Metropolis et al. 1953; Hastings 1970) to find the full global solution to this multidata inverse problem (see Ford 2005; Ford 2006; Maness et al. 2007). Also, the probability densities are sampled to calculate the realistic error bars for parameters and to determine whether a positive detection has been made or not. We also analyse the RV measurements of HD 154345 (Wright et al. 2008) alone and together with simulated astrometric measurements to estimate the minimum observational timeline of astrometry for the detection of the true mass of HD 154345 b. ",
        "conclusions": "The time needed to make a positive detection of an extrasolar planetary conpanion candidate depends essentially on its orbital period. It is commonly assumed that, to be able to detect the signature of such companion, an observational timeline longer than the orbital period is required. Also, since most of the exoplanet candidates have been detected using the RV method, only the lower limit of their mass is available. With the aid of future space telescopes and accurate astrometric measurements, it will be possible to detect the inclination and thus the true mass of planetary candidates.  We have shown that when high-precision RV and accurate astrometric measurements are both available, it is possible to detect the true mass of stellar companions with observational timelines considerably shorter than their orbital periods. Also, when the RV measurements have a long time span, astrometric measurements can reveal the true mass of a stellar companion in less time than one tenth of the orbital period of the system. This ability is also demonstrated using the RV measurements of HD154345 as an example. We find that, having these measurements with $T_{RV} = 10.4$ years in hand, astrometric observations with SIM telescope are sufficient for obtaining the true mass, within a single year.  Bayesian inference plays an important role when extracting information from several sources of measurements. The ability to use RV and astrometric measurements simultaneously makes it possible to employ observational timelines below the orbital ones and still be able to make positive exoplanet detections, thus helping to extract the maximum amount of information from measurements and increasing the time efficiency of observations.  In a forthcoming study, we plan to study the inclusion of additional transit-photometry measurements to further tighten the parameter probability densities in transiting scenarios. Also, the approach used here should be extended to systems with two or more planetary companions."
    },
    "0807/0807.5059_arXiv.txt": {
        "abstract": "This paper is concerned with density estimation of directional data on the sphere. We introduce a procedure based on thresholding on a new type of spherical wavelets called {\\it needlets}. We establish a minimax result and prove its optimality. We are motivated by astrophysical applications, in particular in connection with the analysis of ultra high energy cosmic rays. ",
        "introduction": "We consider the problem of estimating the density $f$ of an independent sample of points $X_1,\\ldots,X_n$ observed on the $d$-dimensional sphere $\\bS^d$ of $\\bR^{d+1}$. Obviously, the most immediate examples of applications appear in the case $d=2$. However, no major differences arise from considering the general case.  There is an abundant literature about this type of problems. In particular, minimax $L^2$ results have been obtained (see \\cite{klemelamini},\\cite{klemelalower}). These procedures are generally obtained using either kernel methods (but in this case the manifold structure of the sphere is not well taken into account), or using orthogonal series methods associated with spherical harmonics (and in this case the 'local performances of the estimator are quite poor, since spherical harmonics are spread all over the sphere).   In our approach we  focus on two important points. We aim at a procedure of estimation which is efficient from a $L^2$ point of view (as it is a tradition in statistics to evaluate the procedure with the mean square error). On the other hand, we would like it to perform satisfactorily also  from a local point of view (in infinity norm, for instance). To have these two requirements together seems to us a warrant to have good results in practice. In effect, it is very difficult to produce a loss function which reflects at the same time the requirement of clearly seeing the bumps of the density, of being able to well estimate different level sets, of testing whether there is a difference between the northern and southern hemispheres and so on.  In addition, we require this procedure to be simple to implement, as well as adaptive to inhomogeneous smoothness.  This type of requirements is generally well handled using thresholding estimates associated to wavelets. The  problem requires a special construction adapted to the sphere, since usual tensorized wavelets will never reflect the manifold structure of the sphere and will necessarily create unwanted artifacts. Recently in (\\cite{pnarco},\\cite{NPW}) a tight frame (i.e. a redundant family) was produced which enjoys enough properties to be successfully  used for density estimation.  The fundamental properties of wavelets are  their concentration in the Fourier domain as well as in the space domain. Here, obviously the 'space' domain is the sphere itself whereas the Fourier domain is now obtained by replacing the 'Fourier' basis by the basis of Spherical Harmonics which plays an analogous role on the sphere.  The construction \\cite{pnarco},\\cite{NPW} produces a family of functions which very much resemble to wavelets, the needlets, and in particular have very good concentration properties.  We use these needlets to construct an estimation procedure, and prove that this procedure attains optimal rates over various spaces of regularity.  Again, the problem of choosing appropriated spaces of regularity on the sphere in a serious question, and we decided to consider the spaces which may be the closest to our natural intuition: those which generalize to the sphere case the classical H\\\"older spaces.  In the first section we present (\\cite{pnarco}) needlets, and describe spaces of regularity on the sphere. In the second one we define our estimation procedure, and describe its properties.  The novelties of this paper lie in the application of thresholding to the needlet coefficients, which gives a very simple and adaptive procedure which works on the sphere. We also focus here on giving the results in $\\bL_\\infty$ norm, and obtain the rates of convergence for many other loss functions  as a consequence of the previous ones.  Our results are motivated by many recent developments in the area of observational astrophysics. As an example, we refer to experiments measuring incoming directions of Ultra High Energy Cosmic Rays, such as the AUGER Observatory (\\verb@http://www.auger.org@). Here, efficient estimation of the density function of these directional data may yield crucial insights into the physical mechanisms generating the observations. More precisely, a uniform density would suggest the High Energy Cosmic Rays are generated by cosmological effects, such as the decay of massive particles generated during the Big Bang; on the other hand,  if these Cosmic Rays are generated by astrophysical phenomena (such as acceleration into Active Galactic Nuclei), then we should observe a density function which is highly non-uniform and tightly correlated with the local distribution of nearby Galaxies. Massive amount of data in this area are expected to be available in the next few years. The Auger observatory will be based on two arrays of detectors; the first one covers an area larger than 3000 Km$^2$ in Pampa Amarilla (Argentina), and has already started to collect observations: some preliminary evidence was provided in \\cite{collaboration-2008-29}, and a non-uniform distribution seems to be favored. The whole celestial sphere will actually be covered only when the construction of the northern hemisphere array, due to be built in eastern Colorado, will be completed, a few years from now. Hence, in the immediate future efficient statistical techniques will be eagerly requested for the analysis of the forthcoming datasets.  A survey of statistical methodologies dealing with directional data on the sphere may be found in \\cite{Mardia}, \\cite{Jupp}, \\cite {Mardia-jupp}. The generalization of estimation using orthogonal series methods to the case of compact Riemannian manifold can be found in \\cite{hendriks}. See related works in \\cite{Kim}, \\cite{Ruymgaart}, \\cite{hendriks-ruymgaart}, \\cite{Pelletier}, \\cite{jupp08}. Kernel methods on the sphere have been investigated in \\cite{hallsphere}. Minimax rates for the equivalent of Sobolev spaces on the sphere associated can be found in \\cite{klemelamini}, \\cite{klemela}, \\cite{klemelalower}.  The plan of the paper is as follows. In \\S2 and \\S3 we review some background material on needlets and Besov spaces. \\S4 introduces our thresholding estimator, whose minimax performances are stated in \\S5. \\S6 shows the performance of the estimators on some simulated data. \\S7--\\S9 contain the proofs. ",
        "conclusions": ""
    },
    "0807/0807.3952_arXiv.txt": {
        "abstract": "$21$~cm tomography is emerging as a promising probe of the cosmological dark ages and the epoch of reionization, as well as a tool for observational cosmology in general.  However, serious sources of foreground contamination must be subtracted for experimental efforts to be viable. In this paper, we focus on the removal of unresolved extragalactic point sources with smooth spectra, and evaluate how the residual foreground contamination after cleaning depends on instrumental and algorithmic parameters. A crucial but often ignored complication is that the synthesized beam of an interferometer array shrinks towards higher frequency, causing complicated frequency structure in each sky pixel as ``frizz'' far from the beam centre contracts across unresolved radio sources. We find that current-generation experiments should nonetheless be able to clean out this points source contamination adequately, and quantify the instrumental and algorithmic design specifications required to meet this foreground challenge. ",
        "introduction": "The use of the $21$ cm line of neutral hydrogen as a probe of cosmological structure has generated much excitement over the last few years. $21$ cm tomography promises to place extremely tight constraints on cosmological parameters \\citep{Matt3,Santos2,Wyithe, juddjackiemiguel1,Yi}, as well as to act as a direct probe of the cosmological dark ages and the Epoch of Reionization (EOR) \\citep{Rees, Tozzi2, Tozzi, Iliev,   furlanetto1, Loeb1,furlanetto2, Barkana1,Mack}.  As a result, many $21$~cm experiments in the form of large radio arrays have recently been funded and are currently being built, e.g., LOFAR\\footnote{http://www.lofar.org}, 21CMA\\footnote{http://21cma.bao.ac.cn/}, MWA\\footnote{http://www.haystack.mit.edu/ast/arrays/mwa/} and PAPER\\footnote{http://astro.berkeley.edu/~dbacker/eor/}.  \\begin{figure} \\centering \\includegraphics[width=0.5\\textwidth]{awesomeplot.pdf} \\caption{Sample synthesized beam profile for a typical 21-cm tomography experiment with 500 antenna tiles.  The tiles are distributed within a diameter $D_{\\rm array} \\sim 1500\\,\\textrm{m}$ according to the density function $\\rho (r) \\sim r^{-2}$.  From the profile, it is clear that a realistic evaluation of foreground subtraction techniques should take into account the fact that beam widths vary as $\\lambda /D$.  This is particularly important when subtracting off unresolved point sources, because a point source can fall on an off-beam ``spike\" that has an effectively unpredictable dependence on frequency.} \\label{awesome} \\end{figure}  The observational challenges, however, are daunting.  At radio frequencies, one must deal with a long list of contaminants that have the ability to swamp the cosmological $21$~cm signal in collected data.  Terrestrially, for instance, one must contend with radio frequency interference from radio communications.  At these wavelengths, the ionosphere also interacts significantly with incoming photons, which may pose a challenge for instrumental calibration.  Perhaps most worrisome of all are astronomical foreground sources \\citep{Oh, santos, Miguel, BowmanThesis,angelica,LOFAR}.  These act as contaminants to the signal and include unresolved extragalactic point sources, resolved bright point sources, and Galactic synchrotron radiation.  Taken together, these sources contribute a brightness temperature on the order of hundreds of Kelvins, which easily overwhelms the cosmological $21$~cm signals that are expected to be on the order of mK.  The ability to subtract off foregrounds is therefore crucial.  Foreground subtraction in 21-cm tomography is particularly challenging because current-generation instruments are not optimized for imaging.  Instead, most experiments are designed to provide cosmological information through \\emph{statistical} measurements such as the power spectrum.  Images of the 21-cm sky are expected to be of a rather poor quality in terms of both signal-to-noise and synthesized beam, and image-processing techniques devised to optimally mitigate this problem may be too computationally intensive to be useful.  In this paper, we take a conservative approach and investigate how well foreground cleaning can be done with a computationally simple method using only the the so-called \\emph{dirty maps}  -- distorted maps of the sky as seen by an instrument at many different frequencies. Specifically, we clean out the foregrounds in any given sky direction by fitting a model of the contamination to the corresponding sky pixels in all the dirty maps. This approach of cleaning one sky pixel at a time along the line-of-sight allows one to take advantage of the high spectral resolution offered by most 21~cm tomography experiments. The key idea is that explored in   \\citet{xiaomin}: the point source emission (mainly synchrotron radiation is expected to vary smoothly with frequency while the redshifted 21~cm signal of interest varies rapidly, because a small change in frequency (and hence redshift) corresponds to many Mpc in the radial direction and hence a major difference in local density fluctuations etc. High spectral resolution should thus allow one to separate out a rapidly oscillating cosmological signal from the spectrally smooth foreground using a low-order polynomial fit.  While previous studies \\citep{xiaomin,LOFAR} have already highlighted the potential of this approach, such studies are incomplete because they do not take frequency dependent beam effects into account. At a given frequency, the dirty map is the true sky multiplied by a broad (say $10^\\circ$) function known in radio astronomy as the {\\it primary beam} and then convolved with a function referred to as the {\\it synthesized beam}. The primary beam width is of order $\\theta\\sim\\lambda/D_{\\rm antenna}$, where $D_{\\rm antenna}$ is the physical size of the individual array elements, whereas the synthesized beam is the Fourier transform of the array layout. Figure~\\ref{awesome} shows a synthesized beam example illustrating several key features: \\begin{itemize} \\item There is a narrow central peak whose width is of order $\\theta\\sim\\lambda/D_{\\rm array}$, where $D_{\\rm array}$ is the size of the whole antenna array. This peak width is usually referred to as the angular resolution of the experiment. \\item There is positive and negative ``frizz'' extending far beyond the central peak, corresponding to incomplete coverage of the Fourier plane, causing random-looking low amplitude oscillations on the angular resolution scale. \\item The synthesized profile is the same at all frequencies except for an overall scaling with the wavelength $\\lambda$. \\end{itemize} As one shifts to higher frequencies, the synthesized beam therefore shrinks like $\\nu^{-1}$, causing the contributions from individual point sources to oscillate rapidly as positive and negative parts of the frizz moves across them. This produces an effective foreground signal looking not like a smooth synchrotron spectrum, but more like the rapidly oscillating cosmological signal.  Early work on this problem has provided cautionary but encouraging results \\citep{BowmanThesis}. Our goal in this paper is to quantify in detail how well the simple pixel-by-pixel foreground subtraction strategy can deal with this problem, using simulations including realistic radio arrays and beams.    \\citet{Judd08} have independently studied this problem with a complementary approach, and we compare our results with theirs below. Whereas they perform a detailed case study of how well the MWA experiment can meet this challenge, including all foregrounds, we instead tackle the broader question of how residual point source contamination depends on experimental specifications, and what the experimental design implications are.  We ignore non-point-source contributions to the foregrounds, with the expectation that such components (such as Galactic synchrotron radiation) are rather benign relative to the unresolved sources.  Being spatially smooth, the off-beam contributions of such foregrounds tend to average out, unlike the point sources, even though their expected amplitude is larger.   \\citet{Judd08} confirm this expectation.  In this paper, we deal only with the removal of \\emph{unresolved} point sources, assuming that resolved and detected sources can be eliminated by standard radio astronomy techniques as well as discarding the most contaminated pixels. We thus envision the cleaning of foregrounds as a two-step process: first masking or deconvolving out bright resolvable point sources that exceed some flux cut, then proceeding with our proposed algorithm. We focus solely on the viability of the second step, but explore how the results depend on the flux cut from the first step.  The rest of this paper is organized as follows.  In Section \\ref{method}, we outline the simulation methodology.  We go through the simulation of foregrounds, the simulation of radio interferometers, and the proposed foreground subtraction strategy itself.  The overall results of our analysis are presented in Section \\ref{genresults}, where we consider three possible scenarios for foreground subtraction, ranging from the most pessimistic to the most optimistic. In Section \\ref{howitdepends}, we quantify which design and algorithm choices are most important, and examine how the interplay between various instrumental and algorithmic parameters gives rise to these results. Table \\ref{table1} lists the parameters that we explore as well as the ranges over which we vary them. In Section \\ref{conc}, we summarize our results and discuss future prospects.  \\begin{table*} \\label{table1} \\caption{Range of our parameter space exploration for foreground cleaning. Parameters pertaining both to experimental specifications and to analysis method impact the success of the cleaning.} \\begin{tabular}{p{2.2cm} p{2.5cm}|p{3.0cm}|p{3.0cm}|p{3.0cm}} \\hline \\multicolumn{2}{c|}{\\textbf{Assumptions}}  &  \\textbf{Low-Performance Extreme} & \\textbf{Fiducial Model} & \\textbf{High-Performance Extreme} \\\\ \\hline \\textbf{Experimental} & Tile Arrangement & $\\rho(r) \\sim r^{-2}$  & $\\rho(r) \\sim r^{-2}$  & Monolithic with tiles separated by $40\\,\\textrm{m}$\\\\ \\cline{2-5} & Rotation synthesis & None  & 6 hours, continuous & 6 hours, continuous \\\\ \\cline{2-5} & Noise level & $\\sigma_T \\sim 1\\,\\textrm{mK}$  & Noiseless \\footnotemark& Noiseless \\\\ \\hline \\textbf{Analysis} & Primary beam width adjustments & None  & None  &  Adjusted to be frequency-independent\\\\ \\cline{2-5} & Bright point source flux cut $S_{\\rm cut}$&100 mJy  &10 mJy &0.1 mJy\\\\ \\cline{2-5} & synthesized beam width adjustments & None & None  &Resolutions equalised by extra smoothing\\\\ \\cline{2-5} & $u$-$v$ plane weighting & None (natural)  & Uniform  &  Uniform \\\\ \\cline{2-5} & Order of polynomial fit & Constant  & Quadratic  &  Quintic \\\\ \\cline{2-5} & Range of polynomial fit & $80\\,\\textrm{MHz}$  & $2.4\\,\\textrm{MHz}$  &  $2.4\\,\\textrm{MHz}$ \\\\ \\hline \\end{tabular} \\end{table*} ",
        "conclusions": "\\label{conc}  We have studied the problem of cleaning out point source foreground contamination from 21 cm tomography data, and investigated how the level of residual contamination after the cleaning process depends on various experimental and algorithmic parameters. Our results show that while successful foreground removal is far from guaranteed, the signs are encouraging. For instance, the fact that the post-cleaning residual foregrounds lie below the simulated cosmological signal in the fiducial model of Section \\ref{genresults} is promising, since the scenario in question is considered to be representative of current-generation experiments, and its cleaning algorithms are both simple and computationally cheap. These conclusions agree with those of the independent and concurrent analysis of \\citet{Judd08}, which is complementary by focusing on the specifics of the MWA experiment. As a further consistency check on our calculations and software, we and the authors of that paper agreed on a test example that we both simulated, obtaining consistent results.  We also identified a number of aspects of the problem that can be understood analytically. Noise and cosmic signal decouple from the point source problem as long as the cleaning method is linear. The power spectrum of unresolved point sources in the observed sky simply gets suppressed by a fixed $k$-dependent factor, so any shifts in this input power spectrum due to revised source count estimates or altered flux cuts for resolved source removal simply scale the output power spectrum by the same amount. Finally, simple Fourier space considerations explain why  line-of-sight cleaning works as well as it does.  Based on our analysis, we can make several specific recommendations regarding the foreground subtraction of unresolved point sources, pertaining to both instrumentation and data reduction. \\begin{itemize} \\item Instrumental recommendations: \\begin{itemize} \\item \\textbf{Tile arrangement} should ensure good $u$-$v$ after rotation synthesis; we found both a monolithic arrangement and an $r^{-2}$ arrangement to work well with respect to foreground subtraction. \\item \\textbf{Rotation synthesis} is important, but once an array has achieved $3$ to $4$ hours of integration, the advantage gained from further $u$-$v$ coverage is minimal. \\item \\textbf{Time between snapshots} taken of the sky should be made as short as possible.  Ideally, one should have continuous $u$-$v$ coverage, as is planned for experiments such as the MWA. \\end{itemize} \\item Data reduction/algorithmic recommendations: \\begin{itemize} \\item \\textbf{Uniform weighting} of the $u$-$v$ plane outperforms ``natural weighting''. \\item \\textbf{Adjusting for changes in the primary beam with frequency} through an extra convolution kernel in $u$-$v$ space results in a modest improvement in foreground subtraction.  However, this conclusion may change if the primary beam has sidelobes. \\item \\textbf{Adjusting for changes in the synthesized beam with frequency} appears not to be worthwhile. \\item \\textbf{Quadratic fits} across a few MHz are adequate for subtracting off spectrally smooth foreground components from the data as long as bright point sources can be resolved and removed down to the $10\\,\\textrm{mJy}$ level --- otherwise higher orders should be considered. \\item \\textbf{The frequency range} over which the polynomial fitting is performed should be narrow enough to allow one to see cosmic signal, but broad enough to prevent over-fitting which excessively removes cosmic signal. For quadratic fits, a range of a few MHz appears appropriate. \\item \\textbf{The ability to remove bright point sources} prior to the subtraction of unresolved point sources is crucial, with the corresponding flux cut $S_{cut}$ being the most important parameter of all in our analysis. If one can safely fit a cubic to the spectra without loosing too much cosmic signal, then one needs to be able to remove resolved points sources of flux down to $100\\,\\textrm{mJy}$. \\end{itemize} \\end{itemize}  Our results were deliberately conservative, obtained with an extremely simple line-of-sight cleaning procedure, and it is likely that one can do better. An interesting question for future work is determining the optimal procedure. For example, other classes of fitting functions may be worth considering, and it may be better to replace multiple polynomial fits in disjoint bands by a single fit to a function with more parameters, say a spline. Such work should quantify the extent to which one is removing power from the cosmological signal, and optimize the tradeoff between more foreground removal and less signal removal. Such an optimization should ideally be done using a more complete sky model, including diffuse Galactic synchrotron radiation.  For now, however, what is reassuring for the field of $21$-cm tomography is the fact that all of the recommendations listed above can be followed without any substantial revisions to current experimental designs. The papers mentioned in the introduction have already shown that neither noise nor diffuse Galactic foreground emission appear to pose insurmountable obstacles to 21~cm cosmology. Our results on point source cleaning suggest that arguably the most worrisome remaining foreground problem is also surmountable, allowing 21~cm cosmology to live up to its great potential."
    },
    "0807/0807.1020_arXiv.txt": {
        "abstract": "We use scalar-field Lagrangians with a non-canonical kinetic term to obtain unified dark matter models where both the dark matter and the dark energy, the latter mimicking a cosmological constant, are described by the scalar field itself. In this framework, we propose a technique to reconstruct models where the effective speed of sound is small enough that the scalar field can cluster. These models avoid the strong time evolution of the gravitational potential and the large Integrated Sachs-Wolfe effect which have been a serious drawback of previously considered models. Moreover, these unified dark matter scalar field models can be easily generalized to behave as dark matter plus a dark energy component behaving like any type of quintessence fluid. ",
        "introduction": "If we assume that General Relativity correctly describes the phenomenology of our universe, astronomical observations provide compelling evidence that (1) the dynamics of cosmic structures is dominated by dark matter (DM), a cold collisionless component mostly made of hypothetical elementary particles, and (2) the expansion of the universe is currently accelerating because of the presence of a positive cosmological constant or a more general Dark Energy (DE) component. The DM particles have not yet been detected and there is no theoretical justification for the tiny cosmological constant (or more general DE component) implied by observations (see, e.g. Refs.~\\cite{Spergel:2006hy,Dunkley:2008ie}). Therefore, over the last decade, the search for extended or alternative theories of gravity has flourished.  In this paper we focus on unified models of DM and DE (UDM), in which a single scalar field provides an alternative interpretation to the nature of the dark components of the universe. Compared with the standard DM + DE models (e.g. $\\Lambda$CDM), these models have the advantage that one can describe the dynamics of the universe with a single dark fluid which triggers both the accelerated expansion at late times and the large-scale structure formation at earlier times. Moreover, for these models, we can use Lagrangians with a non-canonical kinetic term, namely a term which is an arbitrary function of the square of the time derivative of the scalar field, in the homogeneous and isotropic background. These models are known as ``k-essence models'' \\cite{Chiba:1999ka,AMS1,AMS2,Rendall:2005fv,Fang:2006yh,Babichev:2007dw} (see also \\cite{Li:2006bx,Kang:2007vs,Babichev:2006cy,Babichev:2007tn,Bazeia:2007df,Calcagni:2006ge}) and have been inspired by earlier studies of k-inflation \\cite{kinf} \\cite{Garriga:1999vw} (a complete list of dark energy models can be found in the review \\cite{Copeland:2006wr}).  Most UDM models studied so far in the literature require non-trivial fine tunings. Moreover, the viability of UDM models strongly depends on the value of the effective speed of sound $c_s$ \\cite{Huss,Garriga:1999vw,Mukhanov:2005sc}, which has to be small enough to allow structure formation \\cite{Sandvik:2002jz,Giannakis-Hu} and to reproduce the observed pattern of Cosmic Microwave Background (CMB) temperature anisotropies \\cite{Huss,Bertacca:2007ux,Sandvik:2002jz,Finelli,Bertacca:2007cv}. The prospects for a unified description of DM/DE (and inflation) through a single scalar field has been addressed also in Ref.~\\cite{LiddleLopez}. \\\\  Several adiabatic or, equivalently, purely kinetic models have been investigated in the literature: for example, the generalized Chaplygin gas \\cite{Kamenshchik:2001cp,Bilic02,Bento:2002ps} (see also Refs.~\\cite{GCG,chap_con1,chap_con2,chap_con3,chap_con4,Giannantonio:2006ij}), the Modified Chaplygin gas \\cite{MCG}, the Scherrer \\cite{Scherrer:2004au} and generalized Scherrer solutions \\cite{Bertacca:2007ux}, the single dark perfect fluid with a simple 2-parameter barotropic equation of state \\cite{Bruni}, or the homogeneous scalar field deduced from the galactic halo space-time \\cite{DiezTejedor:2006qh} (see also Ref.\\cite{Bertacca:2007fc}). \\\\  Moreover, one can build up scalar field models where the constraint that the Lagrangian is constant along the classical trajectories, namely the solutions of the equations of motion, allows to describe a UDM fluid whose average behaviour is that of dark matter plus a cosmological constant \\cite{Bertacca:2007ux} (see also Ref.~\\cite{stability2,stability1,prc}, for a different approach). Alternative approaches to the unification of DM and DE have been proposed in Ref.~\\cite{takayana}, in the framework of supersymmetry, in Ref.\\cite{LukesGerakopoulos:2008rr} in connection with chaotic scalar field solutions in Friedmann-Robertson-Walker cosmologies and in Ref.~\\cite{bono}, in connection with the solution to the strong CP problem. One could also easily reinterpret UDM models based on a scalar field Lagrangian in terms of generally non-adiabatic fluids \\cite{DiezTejedor:2005fz,Brown:1992kc} (see also \\cite{Bilic:2008zk}).  Here we choose to investigate the class of scalar-field Lagrangians with a non-canonical kinetic term to obtain UDM models. In Ref.~\\cite{Bertacca:2007ux}, the authors require that the Lagrangian of the scalar field is constant along the classical trajectories. Specifically, by requiring that $\\mathcal{L}=-\\Lambda$ on cosmological scales, the background they obtain is identical to the background of the $\\Lambda$CDM model. In this case the limited number of degrees of freedom does not leave any room for choosing the evolution of the effective speed of sound $c_s^2$ in agreement with observations \\cite{Bertacca:2007ux}.  Moreover, one of the main issues of these UDM models is to see whether their single dark fluid is able to cluster and produce the cosmic structures we observe in the universe today. In fact, the effective speed of sound can be significantly different from zero at late times; the corresponding Jeans length (or sound horizon), below which the dark fluid can not cluster, can be so large that the gravitational potential first strongly oscillates and then decays \\cite{Bertacca:2007cv}, thus preventing structure formation. Previous work attempted to solve this problem by a severe fine-tuning of the parameters appearing in the Lagrangian (see for example \\cite{chap_con1,chap_con2,chap_con3,chap_con4,Scherrer:2004au,Giannakis-Hu}).  In Section \\ref{basics}, we layout the basic equations; in Sections \\ref{UDM}, \\ref{BIL} and \\ref{UDMg(X)-generic}, we suggest a reconstruction technique to find models where the effective speed of sound is small enough that the scalar field can cluster. Specifically, in Section \\ref{BIL} we consider a model with kinetic term of Born-Infeld type \\cite{prc,abramo1,abramo2,Alishahiha:2004eh,stability1,stability2} that does not allow a strong time evolution of the gravitational potential and the large Integrated Sachs-Wolfe (ISW) effect which have been a serious drawback of previous models. In Section \\ref{Class-eq}, we consider a more general class of UDM Lagrangians, with a non-canonical kinetic term, whose equations of motion are dynamically equivalent to those of the previous models. Finally, in Section \\ref{GUDM} we study a possible way to generalize UDM models so that they can mimic dark matter and dark energy in the form of a general quintessence fluid. ",
        "conclusions": "A general severe problem of many UDM models considered so far is that their large effective speed of sound causes a strong time evolution of the gravitational potential and generates an ISW effect much larger than current observational limits. In this paper we have outlined a technique to reconstruct UDM models such that the effective speed of sound is small enough that these problems are removed and the scalar field can cluster.  We have also considered a more general class of UDM Lagrangians with a non-canonical kinetic term. Specifically, we have studied some invariance properties of general Lagrangians of the form $\\mathcal{L} = \\mathcal{L}(h(\\varphi)X,\\varphi)$ which allows to define different models whose equations of motion are dynamically equivalent.  Finally, we have studied a possible way to generalize UDM models that can mimic a fluid of dark matter and quintessence-like dark energy.  The Lagrangians that we obtained appear rather contrived. Indeed, these models should be understood as examples which show that the mechanism itself can work. These models can however help to search for physically motivated models with the desired properties. Moreover many previous work attempted to solve this problem by a severe fine-tuning of the parameters appearing in the Lagrangian. This drawback does not belong to our models.  In future work, we will consider models with Lagrangians $\\mathcal{L}=\\mathcal{L}(X,\\varphi)$ to estimate astrophysical observables, like the cross-correlation of CMB anisotropies and large-scale structure or the weak lensing shear signal power-spectrum."
    },
    "0807/0807.3739_arXiv.txt": {
        "abstract": "{} {The CoRoT satellite has announced its fourth transiting planet (Aigrain et~al. 2008) with space photometry. We describe and analyse complementary observations of this system performed to establish the planetary nature of the transiting body and to estimate the fundamental parameters of the planet and its parent star.} {We have analysed high precision radial-velocity data, ground-based photometry, and high signal-to-noise ratio spectroscopy.} {The parent star CoRoT-Exo-4 (2MASS 06484671-0040219) is a late F-type star of mass of 1.16 \\Msun\\ and radius of 1.17 \\Rsun.  The planet has a circular orbit with a period of 9.20205d. The planet radius is 1.19 \\Rjup\\ and the mass is 0.72 \\Mjup. It is a gas-giant planet with a ''normal'' internal structure of mainly H and He. CoRoT-Exo-4b has the second longest period of the known transiting planets. It is an important discovery since it occupies an empty area in the mass-period diagram of transiting exoplanets.} {} ",
        "introduction": "\\label{sec:intro}  The masses and radii of extrasolar planets are fundamental parameters for constraining the physics at play in their interiors. A combination of photometric and Doppler techniques is presently the only way to indirectly derive these parameters. However, this possibility is limited to planets that transit in front of their stars, hence a sample of objects that is highly biased towards close-in orbits.  The planet search programme conducted from space by CoRoT \\citep{baglin06} is currently monitoring tens of thousands of stars and is regularly discovering new transiting planets \\citep{barge08,alonso08,deleuil08b}. CoRoT has the ability to discover planets with smaller radii and/or longer periods than what is routinely achieved by ground-based transit surveys. The evolution of fundamental parameters of extrasolar planets will thus be explored as a function of the distance to the star and as related characteristics like irradiation.  This has important implications for our understanding of these planets. For example, it is a way to test whether tidal damping (highly dependent on orbital distance), increased opacities, or alternative theories are the explanation for the anomalously large sizes of some extrasolar planets \\citep[e.g.][]{guillot06,burrows07,chabrier07}.  With the current detection threshold of the first CoRoT run (45 day duration), 26 transiting events were detected. Two of them turned out to be planets (CoRoT-Exo-1b and 4b, \\citet{barge08, aigrain}, 2 are grazing eclipsing binaries, 7 are low-mass binaries, and 10 are background eclipsing binaries. Four events remain unsolved cases, either because they are fast rotators (no possible radial velocity follow-up, the cross-correlation function is too wide or undetected with HARPS) or because they showed no radial-velocity variations. These cases are still in the list of systems to be observed at the next visibility season with higher accuracy. The statistics achieved on this run is compatible with expectations and the simulations on CoRoT detection, with its current threshold and follow-up performances.  The present paper focuses on the fourth planet discovered by CoRoT. The detection and light curve analysis is described in a companion paper \\citep[ hereafter Paper~IV]{aigrain}. With its 9.20205d period, the planet CoRoT-Exo-4b has the second longest period known so far in the family of transiting planets, after HD~17156b \\citep{barbieri07}.  We conducted a series of complementary observations from the ground to establish the planetary nature of CoRoT-Exo-4b and to derive the fundamental parameters of both the star and the planet. ",
        "conclusions": "The planet CoRoT-Exo-4b has the second longest period of the transiting systems known today and is found in a region of the mass-period parameter space that was previously empty (see Fig.~\\ref{fig:fig6} available on-line).  The mean density of $0.525$\\,\\gcmmm\\ makes CoRoT-Exo-4b a gas-giant with a ``normal'' internal structure dominated by H and He. It is located in the bulk of Hot Jupiters in the mass-radius diagram. We estimate the age of the parent star to be about 1\\,Gyr; thus, the star is likely to be fairly active and its age and rotation rate (Paper IV) are compatible.  New discoveries made with CoRoT will likely populate the mass-period diagram in the period range 5--50 days and allow further investigations of the importance of star-to-planet distance on the internal structure of giant planets.  This will complement the output from current radial-velocity surveys, allow us to better understand the mechanisms that govern the contraction of these planets, and provide constraints for planet formation models.    \\begin{table} \\begin{minipage}[t]{\\columnwidth} \\begin{center} \\caption{The orbit parameters and the fundamental parameters of CoRoT-Exo-4b and its host star.} \\renewcommand{\\footnoterule}{}  % \\begin{tabular}{llll} \\hline\\hline $V_0$ [\\kms]        & $32.301   \\pm 0.005     $&  $K$ [\\ms]           & $63      \\pm 6         $ \\\\ $e$                 & $0.0       \\pm 0.1       $ & $a$ [AU]            & $0.090  \\pm 0.001    $ \\\\ $P$ [d] & 9.20205 $\\pm$ 0.00037& & \\\\ $M_p$ [\\Mjup]       & $0.72    \\pm 0.08       $ &  $R_p$ [\\Rjup]       & $1.190_{-0.05}^{+0.06}$ \\\\ $\\rho_p$ [g/cm$^3$] & $0.525   \\pm 0.15      $ & $T_{\\rm eq}$ [K]    & $1,074    \\pm 19       $ \\\\ \\hline $M_s$ [\\Msun]       & $1.16_{-0.02}^{+0.03}  $ & $R_s$ [\\Rsun]       & $1.17_{-0.03}^{+0.01}  $ \\\\ $v\\sin i$ [km/s]    & $6.4     \\pm 1.0       $ & $T_{\\rm eff}$ [K]   & $6190    \\pm 60        $ \\\\ $\\log g$            & $4.41    \\pm 0.05      $ & Age  [Gy] &  $1_{-0.3}^{1.0}$\\\\ \\hline \\hline \\label{tab:param} \\end{tabular} \\end{center} \\end{minipage} \\end{table}    \\onlfig{5}{ \\onecolumn \\begin{figure} \\begin{center} \\epsfig{file=MPcorot4.ps,width=0.8\\linewidth,angle=0} \\caption{Mass-period diagram of the 35 known transiting planets showing the location of CoRoT-Exo-4b (diamond).} \\label{fig:fig6} \\end{center} \\end{figure} }  \\onltab{2}{ \\begin{table}% \\begin{center} \\caption[]{Radial-velocity measurements of CoRoT-Exo-4 with SOPHIE (lines 1 to 7) and with HARPS (lines 8 to 20). \"BIS\" is the slope of the bisector span, measured in the cross-correlation function.  } \\begin{tabular}{lllll} \\hline\\hline \\small BJD &  RV & Uncertainty & BIS & err(BIS)\\\\ - 2400000 &      [km~s$^{-1}$] & [km~s$^{-1}$]&[km~s$^{-1}$]&[km~s$^{-1}$] \\\\ \\hline \\hline 54405.6545 &  32.370&0.024&  $ -0.021$ & 0.048 \\\\ 54408.6667 &  32.278&0.018&  $ -0.004$ & 0.036 \\\\ 54410.6344 &  32.258&0.027&  $ -0.025$ & 0.054 \\\\ 54413.6784 &  32.343&0.034&  $ -0.113$ & 0.068 \\\\ 54416.6908 &  32.294&0.027&  $ -0.097$ & 0.054 \\\\ 54494.3741 &  32.259&0.019&  $ +0.020$ & 0.036 \\\\ 54507.3778 &  32.320&0.020&  $ +0.067$ & 0.040 \\\\ \\hline 54423.8430&   32.352&0.022&   $ +0.020$ & 0.044 \\\\ 54424.8273&   32.344&0.019&   $ +0.012$ & 0.040 \\\\ 54428.8091&   32.234&0.014&   $ -0.017$ & 0.028 \\\\ 54442.7618&   32.381&0.012&   $ +0.032$ & 0.024 \\\\ 54443.7161&   32.361&0.021&   $ -0.050$ & 0.042 \\\\ 54445.7748&   32.270&0.012&   $ -0.012$ & 0.024 \\\\ 54447.7813&   32.239&0.013&   $ +0.015$ & 0.026 \\\\ 54448.7990&   32.268&0.014&   $ -0.005$ & 0.028 \\\\ 54449.7871&   32.272&0.017&   $ +0.002$ & 0.034 \\\\ 54450.7729&   32.373&0.013&   $ -0.072$ & 0.026 \\\\ 54481.6668&   32.287&0.016&   $ +0.040$ & 0.032 \\\\ 54482.6719&   32.287&0.015&   $ +0.040$ & 0.030 \\\\ 54486.6509&   32.290&0.015&   $ -0.002$ & 0.030 \\\\ \\hline\\hline \\label{tab:rv} \\end{tabular} \\end{center} \\end{table} }"
    },
    "0807/0807.0606_arXiv.txt": {
        "abstract": "{[ABSTRACT ABRIDGED] IC10 X-1 has recently been confirmed as a black hole (BH) + Wolf-Rayet (WR) X-ray binary, and NGC300 X-1 is thought to be. IC10 X-1 and NGC300 X-1 have  similar X-ray properties, with  luminosities $\\sim$10$^{38}$erg s$^{-1}$, and  orbital periods $\\sim$30 hr.} {We investigate similarities between these systems, as well as differences between them and the known Galactic BH binaries.} {We have examined XMM-Newton observations of NGC300 X-1 and IC10 X-1. For each observation, we extracted lightcurves and spectra; power density spectra  (PDS) were constructed from the lightcurves, and the X-ray emission spectra were modeled.} { Each source exhibits PDS  that are characteristic of  turbulence in wind accretion or high state  X-ray binaries (XBs). In this state, Galactic XBs with known BH primaries have soft, thermal emission; however the emission spectra of NGC300 X-1 and IC10 X-1  in the XMM-Newton observations are  predominantly non-thermal.  The remarkable similarity between the behaviour of NGC300 X-1 in Observation 1 and that of IC10 X-1 lends strong evidence for NGC300 X-1 being a BH+WR  binary. } {The unusual spectra of NGC300 X-1 and IC10 X-1 may be  due to these systems existing in a persistently high state, whereas all known BH LMXBs are transient. BH XBs  in a persistent high state could retain their corona, and hence exhibit a large non-thermal component. LMC X-1 is a  BH XB that has only been observed in the high state, and its spectrum is remarkably similar to those of our targets.  We therefore classify  NGC300 X-1, IC10 X-1 and perhaps LMC X-1 as a new breed of BH XB, defined by their persistently high accretion rates and consequent stable disc configuration and corona. { This scenario may also explain the lack of ultraluminous X-ray sources in the canonical soft state.}} ",
        "introduction": "Characterised by high temperatures and large mass loss rates, Wolf Rayet (WR) stars represent the final stage of stellar evolution prior to supernova for massive stars \\citep[$>$25 M$_{\\odot}$;][]{review}. Their spectra are characterised by strong, broad emission lines, which reveal their chemically evolved nature due to, successively, the CNO (WN stars) and triple-$\\alpha$ (WC stars) processes. These products of nuclear burning become visible as the  hydrogen rich mantle of the O star progenitor is stripped away. For single stars this is accomplished via the stellar wind of the progenitor; however binary interaction may also contribute to this process. Indeed, if a stellar merger may be avoided, interaction between the two components of a high mass X-ray binary (OB star and compact object) is thought to lead to a short period WR+compact object binary (WR+co).  Empirically, such systems  are exceedingly rare; to date we know of only three possible WR+co systems: Cygnus X-3 in our Galaxy \\citep{lom05}, IC 10 X-1 \\citep{cc04, bb04}, and most recently, NGC 300 X-1 \\citep{carp07a,cch07}.  Population synthesis carried out by \\citet{lom05} suggests that $\\sim$1 WR + BH binary, and $\\sim$1 WR + neutron star (NS) binary,  may exist  as bright X-ray sources in a galaxy like our own; this is certainly in keeping with the rarity of candidates.   \\citet{prest07}  re-analysed archival Gemini Multi-Object Spectrograph spectra of the WR star  associated with IC10 X-1, and  found evidence for a  black hole (BH) primary, with  a  plausible BH mass range of 23--34 M$_{\\odot}$. This result was subsequently confirmed   by \\citet{sf08}, who obtained spectra over five nights of observations with the Keck-I 10m telescope. They report a BH mass of 20--35 M$_{\\odot}$ for the compact object,  making it the most massive known stellar BH.  NGC300 X-1 has not yet been optically identified as a WR+BH binary; however, in this work we find striking similarities between the X-ray properties of NGC300 X-1 and the confirmed WR+BH binary IC10 X-1. We will show that NGC300 X-1 is likely to be a WR+BH binary also.      WR+co binaries are just one  subset of high mass X-ray binaries (HMXBs), and it is possible that their temporal and spectral characteristics closely resemble other types of HMXBs. Hence it is instructive to  briefly review the   pertinent properties of HMXBs in general. HMXBs are well classified by the nature of the donor star \\citep[see e.g.][ and references within]{wnp95}.  Around 60\\% of Galactic HMXBs include a Be donor \\citep{lvv07}; these have elliptical orbits lasting tens to hundreds of days, only accreting near periastron. As a result, Be HMXBs are transient sources, generally  powered by Bondi-Hoyle accretion \\citep{bh44} from a decretion disc around the Be star \\citep{on01}. Several  Bondi-Hoyle accreting   HMXBs with a  NS accretor  exhibit rapid aperiodic variability in X-ray intensity with amplitude $\\sim$10--30\\%;  their power density spectra (PDS) resemble a power law with  spectral index $\\gamma$ $\\sim$1 \\citep[{ power $\\propto$ $\\nu^{-\\gamma}$ where $\\nu$ is the frequency}, see e.g.][]{nag89,bh90,vdk95}. These works suggested that such variability could be present in all accretion-powered pulsars. However,  an inspection of the catalogues of HMXBs in the Galaxy and Magellanic clouds \\citep{lvv05,lvv07} revealed that only 16 out of 242 HMXBs have published records of  aperiodic variability.  HMXBs that have exhibited  the variability described above  have power law emission spectra with spectral index, $\\Gamma$, $\\sim$0.5--1.8.  There are no known BH HMXBs that accrete in this way.    A further $\\sim$32\\% of Galactic HMXBs contain supergiant (SG) donors \\citep{lvv07}; the orbital periods range over 1.4--41 days, and the donor stars are (or are close to) filling their Roche lobe \\citep[see][ for a review]{kap04}. The compact object in SG HMXBs is continuously accreting from the donor's wind and  can be either wind-fed  via Bondi-Hoyle accretion (with luminosities $\\sim$10$^{35}$--10$^{36}$ erg s$^{-1}$) or disc-fed, reaching $\\sim$10$^{38}$ erg s$^{-1}$  \\citep{kap04}.  Studies of Galactic disc-accreting XBs have shown that their spectral and variability properties depend as much on the accretion rate as the nature of the accretor (NS or BH). At low accretion rates, the emission spectra and temporal variability of disc-accreting binaries are remarkably similar \\citep{vdk94}. However, at certain, higher, accretion rates there are characteristic differences that may allow one to determine the nature of the primary; the PDS are well described by a power law with  index $\\gamma$ $\\sim$1, but the emission spectra for NS and BH XBs are very different. For a disc-accreting XB with a NS primary, the emission corresponding to  the power law PDS is non-thermal, while disc-accreting BH XBs exhibiting the same PDS have a thermal X-ray emission spectrum \\citep{vdk94}. Indeed, \\citet{don04} examined $\\sim$ 1 Tbyte of RXTE archival data on Galactic XBs, and found a region in colour-colour space that is uniquely associated with BH XBs in this high, soft state.  Hence, we may learn a great deal about disc-accreting binaries from their X-ray variability and emission spectra.       \\begin{table}[!b] \\centering \\caption{ Journal of XMM-Newton observations of NGC300 X-1 and IC10 X-1. For each observation, we provide the target, date, revolution number, exposure and good time interval. }\\label{journal} \\begin{tabular}{cccccc} \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} Object &  Date  & Rev  &  Exp&GT\\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} NGC300 X-1  & 2000 Dec 26 & 0192 &  32 ks& 25 ks \\\\ NGC300 X-1  & 2001 Jan 01 & 0195 & 40 ks & 40 ks \\\\ NGC300 X-1  & 2005 May 22 & 0998 & 35 ks & 25 ks \\\\ NGC300 X-1 & 2005 Nov 25 & 1092 & 35 ks & 35 ks \\\\ IC10 X-1 & 2003 Jul 03 & 0653 & 44 ks & 28 ks \\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} \\end{tabular} \\end{table}  NGC300 X-1 and IC10 X-1 both have orbital periods $\\sim$30 hr and X-ray luminosities $\\sim$10$^{38}$ erg s$^{-1}$ \\citep{carp07a,prest07,sf08}. We examined all existing XMM-Newton observations of NGC300 X-1 and IC10 X-1, in order to compare their properties. In particular, we examined the PDS from the X-ray lightcurves of these systems for the first time. In Sect.~\\ref{obs} we describe the observations and data analysis, then provide the results from NGC 300 X-1 and IC10 X-1 in turn in Sect.~\\ref{res}. We discuss our findings in Sect.~\\ref{discuss},  and consider four scenarios for NGC300 X-1 and IC10 X-1: wind accretion onto a neutron star, disc accretion onto a neutron star, wind accretion onto a black hole and disc accretion onto a black hole. { We then compare the properties of  these sources with those of the ultraluminous X-ray sources (ULXs), many of which are thought to be HMXBs.} Finally we  draw our conclusions in Sect.~\\ref{conc}. ",
        "conclusions": "\\label{conc}  IC10 X-1 is the first confirmed BH+WR system, and NGC300 X-1 is thought to be one also. They are known to have strikingly similar X-ray periods and luminosities, so we studied  them in detail. We have examined for the first time the PDS of NGC300 X-1, in four XMM-Newton observations, and also of IC10 X-1, in one XMM-Newton observation. We find that they are  well described by a simple power law with spectral index $\\simeq$1, over the 0.002--0.1 Hz range; such variability is characteristic of disc-accreting XBs in their high state,  or of some HMXB pulsars that are powered by Bondi-Hoyle accretion.  We also find  a striking resemblance between the Obs. 1 NGC300 X-1 spectrum and the IC10 X-1 spectrum. Since IC10 X-1 is a confirmed WR+BH system, our results provide strong support for NGC300 X-1 being one too.   We have considered four accretion scenarios:  disc accretion or  Bondi-Hoyle accretion,  onto a BH or NS. We favour disc accretion onto a BH for both systems, but cannot rule out Bond-Hoyle accretion onto a BH. The observed X-ray spectra of NGC300 X-1 allowed  us to rule out Bondi-Hoyle accretion onto a NS,  but do not exclude disc accretion onto a NS+WR binary, or onto a NS LMXB that is merely coincident with the WR.    If  both NGC300 X-1 and IC10 X-1  are indeed WR+BH binaries, then they are spectrally distinct from the known BH XBs. We propose that this difference may be due to these systems being persistently in the high state, allowing them to keep their disc corona,  like LMC X-1; contrariwise, BH LMXBs enter the high state only in violent outbursts, where the corona may be ejected. NGC300 X-1 and IC10 X-1 (and possibly LMC X-1) would then comprise a new class of BH binary, defined by their high mass transfer rates and consequent stable accretion disc configurations. { This stable disc scenario may also explain why only $\\sim$10\\% of ULXs exist in the thermally-dominated high soft state. A fair comparison between the ULXs and NGC300 X-1 and IC10 X-1 requires detailed analyis of the variability of ULXs in addition to spectral analysis; we will carry this out in a future paper.}"
    },
    "0807/0807.2603_arXiv.txt": {
        "abstract": "We calculate the shear viscosity $\\eta$ and thermal conductivity $\\kappa$ of a nuclear pasta phase in neutron star crusts.  This involves complex non-spherical shapes.  We use semiclassical molecular dynamics simulations involving 40,000 to 100,000 nucleons.  The viscosity $\\eta$ can be simply expressed in terms of the height $Z^*$ and width $\\Delta q$ of the peak in the static structure factor $S_p(q)$.  We find that $\\eta$ increases somewhat, compared to a lower density phase involving spherical nuclei, because $Z^*$ decreases from form factor and ion screening effects.  However, we do not find a dramatic increase in $\\eta$ from non-spherical shapes, as may occur in conventional complex fluids. ",
        "introduction": "\\label{intro} The viscosity of dense matter can dampen the oscillations of neutron stars.  High multipolarity p-mode oscillations may impact the pulse shapes of some radio pulsars \\cite{pulsars}.  For p-modes, the primary restoring force is the pressure, and the modes may be damped by the shear viscosity of the neutron star crust \\cite{chugunov}.  Inertial or r-modes of a rotating neutron star may radiate gravitational waves and could limit the spin period \\cite{rmodes}.  For these modes the coriolis force provides the primary restoring force and they may be damped by the shear, and or bulk, viscosities of dense matter through out the star.  Flowers and Itoh \\cite{flowers} and Nandkumar and Pethick \\cite{nandkumar} have calculated the shear viscosity of neutron star crust matter in the liquid phase.  While recently, Chugunov and Yakovlev have calculated the shear viscosity of both the solid and liquid phases \\cite{chugunov}.  They assume the viscosity is dominated by the momentum carried by degenerate electrons and that this is limited by electron-ion scattering.   At high densities in the neutron star core, the shear viscosity comes from momentum carried by nucleons and this is limited by NN scattering.  See for example Ref. \\cite{benhar}.  The thermal conductivity of the neutron star crust helps determine the temperature profile of the star and is very important for the cooling time of the crust.  Recently, crust cooling times have been observed for two neutron stars after extended outbursts \\cite{cooling}.  We have calculated the thermal conductivity of the outer crust including the possible role of impurities \\cite{thermalcon}.  The thermal conductivity could also be important for setting the ignition conditions for carbon superbursts \\cite{superbursts}.  At intermediate densities, around 10$^{14}$ g/cm$^3$ just below nuclear density, matter may form complex nuclear pasta phases \\cite{pasta}.  Competition between short range attractive nuclear, and long range repulsive Coulomb, interactions can lead to clusters with many different non-spherical shapes including long rods or flat plates.  Because pasta may form at high densities, it could represent as much as half the mass of the neutron star crust.  The complex shapes in nuclear pasta have sizes of tens of Fermis.  This is comparable to the wavelength of neutrinos in Supernova explosions.  Therefore, coherent scattering from the nuclear pasta shapes may impact the opacity of Supernova neutrinos \\cite{pastaopacity}.  In the laboratory, complex fluids with large non-spherical molecules can have shear viscosities that are much larger than for simple fluids.  How might the complex pasta shapes influence the shear viscosity of nuclear pasta?  Is there a drastic increase in viscosity that is similar to that seen for laboratory complex fluids?  We are not aware of any previous calculations of the shear viscosity of nuclear pasta.  Although Chugunov and Yakovlev have calculated the shear viscosity for similar densities \\cite{chugunov}, they apparently assumed spherical nuclei.  In this paper, we calculate the shear viscosity and thermal conductivity using a semiclassical model for the nuclear pasta \\cite{pasta1} and molecular dynamics (MD) simulations with from 40,000 to 100,000 nucleons \\cite{pasta2}.  Our goal is to gain some simple insight into how transport properties, such as the shear viscosity or thermal conductivity, depend on the sizes and shapes of the clusters.  Furthermore, we wish to know how transport properties might change as one goes from isolated spherical nuclei at low densities, to complex pasta phases, and then to uniform nuclear matter at high densities.  One important advantage of our semiclassical model is that it can be directly applied in all of these density regimes.  We wish to emphasize the importance of simple qualitative insight.   Astrophysicists request very detailed transport properties, such as the shear viscosity, even though there are large remaining uncertainties in much more basic dense matter properties.  It may be unrealistic to predict the exact sizes and shapes of the pasta clusters.  Instead, we hope to gain qualitative insight into how pasta shapes may impact the shear viscosity or other properties.  The manuscript is organized as follows. In Sec.~\\ref{formalism} we review the simple semiclassical model, explain the calculation of the shear viscosity, and provide some details of the molecular dynamics simulations.  In Sec.~\\ref{results} we present results for the static structure factors and use these to calculate the shear viscosity and thermal conductivity.  Finally, conclusions and future directions are presented in Sec.~\\ref{conclusions}. ",
        "conclusions": "\\label{conclusions} In this paper we calculate the shear viscosity $\\eta$ and thermal conductivity $\\kappa$ of nuclear pasta using molecular dynamics (MD) simulations of a semiclassical model.  Our model includes Coulomb repulsion and reproduces nuclear saturation.  Furthermore, it is directly applicable at low densities in a phase with spherical nuclei as well as for complex pasta phases.  Finally, the model makes no assumptions about the pasta shapes.  Instead the MD simulations could produce any shapes.  We assume the shear viscosity is dominated by the momentum carried by electrons and that this is limited by electron-proton scattering.  We calculate the static structure factor $S_p(q)$ of the protons from our MD trajectories for simulations involving 40,000 to 100,000 nucleons.  We find a peak in $S_p(q)$ that can be characterized by a height $Z^*$, a width $\\Delta q$ and a position $q^*$.  The height $Z^*$ represents the effective number of protons that an electron can scatter from.  This is reduced from the total number of protons in a nucleus or pasta cluster because of form factor and ion screening effects.  The shear viscosity can be approximated \\begin{equation} \\eta\\approx \\frac{\\pi n_e}{20 \\alpha^2} \\bigl(\\frac{q^*}{\\Delta q\\, Z^*}\\bigr)\\, , \\end{equation} with $n_e$ the electron density.  We find that $\\eta$, for nuclear pasta, is somewhat increased over that for spherical nuclei, because $Z^*$ is much less than the total number of protons in a piece of pasta.  However $\\eta$ is not increased by orders of magnitude just because of the non-spherical pasta shapes.  This is in contrast to conventional complex fluids where large non-spherical molecules can dramatically increase $\\eta$.  For nuclear pasta, the non-spherical shapes only impact $\\eta$ through their influence on $S_p(q)$ and the parameters $Z^*$, $\\Delta q$, and $q^*$.  Different effective interactions may lead to different sizes and shapes for the nuclear pasta.  Indeed it may be very difficult to predict these accurately.  However one can calculate transport properties such as $\\eta$, the thermal conductivity $\\kappa$ and the electrical conductivity in terms of $S_p(q)$ and the parameters $Z^*$, $\\Delta q$, and $q^*$.  In future work we hope to develop a simple way to estimate these parameters, given a pasta configuration.  In addition we plan to calculate other mechanical properties of nuclear pasta such as the shear modulus, and extend our MD pasta simulations to additional values of proton fraction, temperature and density."
    },
    "0807/0807.2573_arXiv.txt": {
        "abstract": "We found an unusual complex of narrow blue filaments, bright blue knots, and H$\\alpha$-emitting filaments and clouds, which morphologically resembled a complex of ``fireballs,'' extending up to 80 kpc south from an E+A galaxy RB199 in the Coma cluster. The galaxy has a highly disturbed morphology indicative of a galaxy--galaxy merger remnant. The narrow blue filaments extend in straight shapes toward the south from the galaxy, and several bright blue knots are located at the southern ends of the filaments. The $R_{\\rm C}$ band absolute magnitudes, half light radii and estimated masses of the bright knots are $\\sim -12 - -13$ mag, $\\sim 200-300$ pc and $\\sim 10^{6-7}$ M$_{\\odot}$, respectively. Long, narrow H$\\alpha$-emitting filaments are connected at the south edge of the knots. The average color of the fireballs is $B - R_{\\rm C}\\approx 0.5$, which is bluer than RB199 ($B - R = 0.99$), suggesting that most of the stars in the fireballs were formed within several times $10^8$ yr. The narrow blue filaments exhibit almost no H$\\alpha$ emission. Strong H$\\alpha$ and UV emission appear in the bright knots. These characteristics indicate that star formation recently ceased in the blue filaments and now continues in the bright knots. The gas stripped by some mechanism from the disk of RB199 may be traveling in the intergalactic space, forming stars left along its trajectory. The most plausible fireball formation mechanism is ram pressure stripping by high-speed collision between the galaxy and the hot intra-cluster medium. The fireballs may be a snapshot of diffuse intra-cluster population formation, or halo star population formation in a cluster galaxy. ",
        "introduction": "Clusters of galaxies are ideal experimental laboratories for studying environmental effects on galaxy evolution (e.g. Bosselli \\& Gavazzi 2006). Cluster galaxy density varies widely from edge to core, the intergalactic space is often filled with X-ray emitting hot gas, and the cluster's large mass forms a very deep gravitational potential around its center. These properties give us observational means to investigate how galaxy density, ambient gas--galaxy interaction, and strong gravitational drag force affect galaxy evolution.  Several lines of observational evidence suggest that drastic galaxy evolution occurred in clusters of galaxies from the redshift $z \\sim 1$ to the present. Early type galaxies form a dominant population in the inner region of local rich clusters, and in particular, S0 galaxies are preferentially found in and around the cluster core region \\citep{Dressler1980,Dressler1994,Dressler1997,Postman1984,Goto2003,VanDerWel2007}. This relationship---the morphology--density relationship---is also observed at high redshift ($z \\sim 1$). However, this enhancement of the fraction of S0 galaxies around a cluster core is less remarkable in high redshift clusters than in local clusters \\citep{Smith2005,Postman2005}. In addition, a higher blue galaxy fraction has been established for high-redshift clusters \\citep{Butcher1978,Butcher1984}. Furthermore, a rapid rise in the luminosity function (LF) of member galaxies toward the faint end has been reported in several local rich clusters \\citep{Popesso2005,Trentham2005,Popesso2006,Milne2007,Adami2007,Yamanoi2007}. However, the LFs of some rich clusters at $z \\approx 0.3$ may exhibit no such rapid increase at those faint ends \\citep{Harsono2007}. These observational facts suggest that in terms of morphology, color, and population, rapid evolution of galaxies occurred in clusters from $z \\sim 1$ to the present.  Several environmental effects that explain the evolution of cluster galaxies have been proposed, including galaxy--galaxy interaction \\citep{Farouki1981,Icke1985}, galaxy harassment \\citep{Moore1996}, tidal interaction with cluster potential \\citep{Henriksen1996}, ram pressure stripping \\citep{Gunn1972,Vollmer2001,Kron2008}, turbulent viscous stripping \\citep{Nulsen1982}, and galaxy starvation \\citep{Larson1980,Bekki2002}. In addition, many observational studies have been made to identify the dominant environmental processes in clusters (cf., Bosselli \\& Gavazzi 2006 and references therein). However, the kinds of physical processes that are dominant in rapid galaxy evolution in clusters are still unclear.  Recently, \\citet{Cortese2007} (C07) found two peculiar galaxies near the central regions of rich clusters Abell 1689 and Abell 2667 at $z\\approx 0.2$. These galaxies have disturbed morphology and are associated with many small blue blobs extending toward one side of the galaxies. These blobs have absolute magnitudes of $M_B \\approx -11 - -12$, corresponding to masses of $10^{6-8}$ M$_{\\odot}$. C07 detected [OII] emission from some of the blobs and between the blobs and the host galaxies. The brightness and the size of the blobs are similar to those of dwarf elliptical galaxies or ultracompact dwarf galaxies \\citep{Drinkwater2004}. C07 suggested that the phenomena they found may be snapshots of the transformation of spiral galaxies to S0 galaxies, and may give insight into the origin of cluster faint-end populations.  Here we report the discovery of a strange complex of narrow blue filaments and knots extending toward one side of a disturbed E+A galaxy, RB199, in the Coma cluster. New deep optical broadband and narrowband (H$\\alpha$) imaging observations of the central region of the Coma cluster revealed this structure. The morphology of the structure is very similar to that of the phenomena reported in C07, and the sizes and the luminosities of the knots in the structure are comparable to C07's blobs. This structure may be a nearby counterpart of C07's phenomena.  We assumed the cosmological parameters ($h_0$, $\\Omega_m$, $\\Omega_\\lambda$) = (0.73, 0.24, 0.72) and the distance modulus of the Coma cluster to be 35.05 \\citep{Yagi2007}. The linear scale at the Coma cluster is 474 pc arcsec$^{-1}$ under this assumption. ",
        "conclusions": "We found a strange complex (the ``fireballs'') of blue filaments and H$\\alpha$ clouds that extends up to 80 kpc toward the south from the merging galaxy RB199 in the Coma cluster. The narrow blue filaments are extended in straight shapes and several bright blue knots are located at the south end of the blue filaments. Strong H$\\alpha$ emission is associated with the knots. In addition, faint, narrow H$\\alpha$-emitting filaments extend farther from the southern edge of the knots.  The total $B$ and $R_{\\rm C}$ magnitudes of the fireballs are 19.87 and 19.33, respectively. The luminosities of the H$\\alpha$ emission associated with the bright knots are $ 1-6 \\times 10^{38}$ erg s$^{-1}$, which correspond to a star-formation rate of $\\sim 10^{- 3}$ M$_\\odot$ yr$^{-1}$, assuming all the H$\\alpha$ comes from star-forming activity. The average color of the fireballs is $B - R_{\\rm C}\\approx 0.5$, which is bluer than RB199 ($B - R_{\\rm C} = 0.99$), and the blue filaments and the bright knots grow bluer with distance from the nucleus of the galaxy. These observational facts suggest that the stars of which the fireballs are composed were formed within several times $10^8$ yr and that the stellar population grows younger farther from the galaxy.  Almost no H$\\alpha$ emission is observed in the blue filaments but strong H$\\alpha$ and UV emissions are associated with the bright knots. These characteristics indicate that star formation has already ceased in the blue filaments and is now in progress in the bright knots. The gas stripped by some mechanism from the disk of RB199 may be traveling in the intergalactic space, forming stars and leaving the formed stars along its trajectories. The $R_{\\rm C}$-band absolute magnitudes, the half light radii, and the estimated masses of bright knots in the fireballs are $\\sim -12 - -13$ mag, $\\sim 200-300$ pc and  $\\sim 10^{6-7}$ M$_{\\odot}$, respectively. These values are similar to those of faint dwarf spheroidals in the Virgo cluster or the faint end of local dwarf galaxies. This suggests that the bright knots are gravitationally self-bound systems.  The most plausible formation mechanism of the fireballs is ram pressure stripping by the hot ICM. Tidal stripping by galaxy--galaxy interaction or galaxy--cluster interaction would not be a primary mechanism for the formation of the fireballs. A galaxy--galaxy merger, however, is probably important to make subsequent ram pressure stripping of the ICM more effective in forming this strange structure.  Although whether the stars in the fireballs are all gravitationally bound to RB199 is not known, some structures close to the galaxy may fall back to the galaxy in the future, oscillating between the two sides of the galaxy disk and eventually forming the halo star population of RB199. Some part of the fireballs may be stripped from RB199 in the future by tidal interaction with nearby galaxies or the cluster gravitational potential, even if they are currently bound to the galaxy. If this is the case, the stripped stars may form a diffuse intra-cluster population \\citep{Gregg1998,Okamura2002,Murante2004,Krick2006,Gerhard2007}.   The fireballs we found in the Coma cluster resemble the compact young blobs found by C07 and Sun et al. (2007) around some cluster galaxies. C07 argued that formation of these compact blobs is rather rare in clusters at $z \\approx 0.2$. They found only two examples out of 130 cluster spiral galaxies at $z \\approx 0.2$. They also predicted that these features would be more remarkable at $z \\geq 0.2$ because the galaxy infalling rate and the gas content of each galaxy are higher in high- $z$ clusters than in local clusters. However, our findings show that formation of compact stellar systems like C07's blobs around cluster galaxies is still ongoing in the local universe. The Coma cluster has a large substructure of the ICM and a group of galaxies associated with this substructure, indicating that this cluster itself is a merging system \\citep{Colless1996}. RB199 is located near the border between the main body and the substructure of the ICM \\citep{Poggianti2004}. The special conditions of the Coma cluster and RB199 may produce a special environment around RB199,namely, high-ensity ambient gas, high-peed infall, and/or frequent galaxy--galaxy interactions"
    },
    "0807/0807.2059_arXiv.txt": {
        "abstract": "Broad \\feii\\ emission is a prominent feature of the optical and ultraviolet spectra of quasars.  We report on a systematical investigation of optical \\feii\\ emission in a large sample of 4037 $z<0.8$ quasars selected from the Sloan Digital Sky Survey.  We have developed and tested a detailed line-fitting technique, taking into account the complex continuum and narrow and broad emission-line spectrum.  Our primary goal is to quantify the velocity broadening and velocity shift of the \\feii\\ spectrum in order to constrain the location of the \\feii-emitting region and its relation to the broad-line region.  We find that the majority of quasars show \\feii\\ emission that is redshifted, typically by $\\sim 400$~ \\kms\\ but up to 2000~\\kms, with respect to the systemic velocity of the narrow-line region or of the conventional broad-line region as traced by the \\hb\\ line.  Moreover, the line width of \\feii\\ is significantly narrower than that of the broad component of \\hb.  We show that the magnitude of the \\feii\\ redshift correlates inversely with the Eddington ratio, and that there is a tendency for sources with redshifted \\feii\\ emission to show red asymmetry in the \\hb\\ line.  These characteristics strongly suggest that \\feii\\ originates from a location different from, and most likely exterior to, the region that produces most of \\hb.  The \\feii-emitting zone traces a portion of the broad-line region of intermediate velocities whose dynamics may be dominated by infall. ",
        "introduction": "Fe emission contributes significantly to the optical and ultraviolet (UV) spectra of most active galactic nuclei (AGNs), both in terms of wavelength coverage and flux.  The properties of the \\feii-emitting clouds may provide important clues to the underlying physics in the broad-line region (BLR). First, \\feii\\ emission can be used to constrain the covering factor of BLR clouds from energy budget considerations.  Second, the ratio of the equivalent width (EW) of \\feii\\ to that of \\hbb\\ strongly varies with statistical measures of AGN correlations, such as the so-called Eigenvector 1 derived from principal component analysis, which is believed to be driven by some fundamental property such as mass accretion rate \\citep[e.g.,][]{bg92,sul00a,sul00b,marzi01,marzi03a}.  Third, Fe abundance derived from \\feii\\ emission can be used to study the cosmological evolution of AGNs and possibly chemical enrichment of their hosts and environment \\citep[e.g.,][]{wil85,wheel89,die02,die03,maio03}.  Careful measurement of the properties of \\feii\\ emission in a large sample of AGNs will clearly have a considerable impact on our understanding of these systems.  The origin of the optical/UV Fe emission has been hotly debated for more than two decades. Thousands of Fe emission lines blend together to form a pseudo-continuum, which, when combined with Balmer continuum emission, results in the ``small blue bump'' around 3000 \\ang\\ \\citep{grandi82,wil85}. Previously, theoretical calculations of photoionized clouds in the BLR encountered difficulties reproducing the observed strength of strong \\feii\\ emission, prompting many authors to propose additional physical mechanisms \\citep{netzer83,wil85,joly87,joly91,collin88,sigut98}.  Recently, a more sophisticated calculation by \\citet{bal04} revealed that the predicted shape and EW of the 2200--2800\\AA\\ \\feii\\ UV bump can only be made consistent with observed values if either microturbulence of hundreds of \\kms\\ or another collisionally excited component is included in the model.  Despite its significance, current observations of \\feii\\ emission provide poor constraints on its origin.  Some studies suggest that \\feii\\ emission originates from the same region as the other broad emission lines.  For example, \\citet{phil77}, \\citet{bg92}, \\citet{laor97b}, and \\citet{veron04} observed similar line widths and profiles for \\feii\\ and \\hbb, and \\citet{maoz93} found that both \\feii\\ emission and the Balmer continuum have comparable variation amplitudes. But there is a growing debate on this issue. \\citet{marzi03b} found that \\hbb\\ is systematically broader than \\feii\\ for sources with \\fwhmhbb\\ $>$ 4000 \\kms.  Recent studies of \\feii\\ emission variability have shown that \\feii\\ emission responds to variations in continuum flux but that the variability amplitude of \\feii\\ is not the same as that of \\hb\\ (\\citealt{mv05} and references therein; \\citealt{wang05}; \\citealt{kuehn08}). In fact, the upper limit on the time lag between \\feii\\ emission and the continuum exceeds the lag of any other observed emission lines obtained \\citep{mv05}.  This implies that \\feii\\ may be emitted from further out in the BLR than any other broad emission line. \\citet{kuehn08} suggested that \\feii\\ may be produced from a region between the BLR and the dust sublimation radius. The study of \\citet{matsuoka07}, based on measurements of the O {\\sc i} and Ca {\\sc ii} emission lines, supports the notion that \\feii\\ emerges from the outer portion of the BLR. In a recent three-dimensional spectroscopic study of Mrk 493, \\citet{popovic07} found that the \\feii\\ emission region is extensive and that the line width of \\feii\\ is only 1/3 of that of \\hbb, leading them to suggest that \\feii\\ emission originates in an intermediate-line region.  If these findings can be confirmed, \\feii\\ emission can be a probe of the intermediate-line region, which may be the transition from the torus to the BLR and accretion disk.  Quasar emission lines often exhibit considerable velocity shifts with respect to each other.  However, the relative velocity of \\feii\\ emission with respect to other lines has not been well studied, especially in systematically for a large sample of objects (only a few measurements of individual sources have been published; e.g., \\citealt{veron04}).  In fact, almost all broad quasar emission lines show blueward velocity shifts \\citep{gastell82,carswell91}, with the exception of \\mgii\\ (no shift; \\citealt{junk89}) and \\hb\\ (redward shift in some studies; e.g., \\citealt{mcin99}).  Most previous studies simply {\\it assumed}\\ that \\feii\\ has no shift with respective to \\oiii\\ \\citep[e.g.,][]{bg92,marzi96,mclure02,die03,greene05b,kim06,woo06} and that it has the same line width as the broad component of \\hb\\ \\citep[e.g.,][]{netzer07,salv07}.  The goal of this study is to test this assumption.  The present paper presents the first detailed investigation of the velocity shift and width of optical \\feii\\ emission\\footnote{Unless otherwise noted, the \\feii\\ emission in this paper refers to the optical band.} in a large sample of quasars selected from the Sloan Digital Sky Survey (SDSS; \\citealt{york00}).  Our primary motivation is to determine the physical location and origin of the \\feii-emitting region. We describe selection of the sample in \\S \\ref{sample} and spectral analysis in \\S \\ref{fitting}.  We test the reliability of the measurements and the errors using Monte Carlo simulations (\\S \\ref{mcsim}), and then check how significantly our method improves the spectral fit and how our method affects the measurements of other emission-line parameters (\\S \\ref{compare}). Section \\ref{result} discusses the results we obtained, including the distribution of \\feii\\ emission shifts and widths, correlations with other parameters, and also an analysis of the composite spectra. The implications of our results are discussed in \\S \\ref{discuss}, with conclusing remarks given in \\S \\ref{summary}.  Throughout this work, we adopt the following cosmological parameters: $H_0=70~{\\rm km~s^{-1}}$ Mpc$^{-1}$, $\\Omega_{m}=0.3$, and $\\Omega_\\Lambda=0.7$ \\citep{spergel06}. ",
        "conclusions": "\\label{discuss}  \\subsection{Redshifted \\hb\\ Component: OQ 208-like Quasars} \\label{shifthb} In our sample, we find a class of sources whose \\hb\\ profile can be fitted well including an additional, substantially redshifted Gaussian component. Usually this additional component has a velocity width intermediate between that of the broad and narrow components. We call this the intermediate component.  An interesting and possibly highly significant fact is that the width and velocity shift of this additional \\hb\\ component is consistent with those of the \\feii\\ emission. A prototype of this kind of sources is OQ 208 (Mrk 668), which was first studied in detail by \\citet{marzi93}. They pointed out that in OQ 208 the \\feii\\ lines at 4924 and 5018 \\ang\\ have the same peak displacement as the red peak of \\hb.  \\begin{figure*} \\centering \\noindent\\includegraphics[angle=-90,width=0.9\\textwidth]{f14.eps} \\caption{ Spectrum of SDSS J094603.94+013923.6, an example of a OQ 208-like source. ({\\it a}) Continuum decomposition. ({\\it b}) Fe model and the observed spectrum after subtracting the power law and Balmer continuum.  The spectra and the marks plotted in panel ({\\it a}) and ({\\it b}) are as the same as those in the top and bottom panels of Fig. \\ref{fig-conti1}.  ({\\it c}) Fitting of \\hb\\ and \\oiii\\ emission lines. The multiple Gaussian components are in blue, and the sum of them is in red. The two blue dashed lines mark the rest-frame wavelength of \\hb\\ and \\oiii\\ $\\lambda$5007. ({\\it d}) Fitting of \\oii\\ emission line.  The blue dashed line marks the rest-frame wavelength of \\oii\\ $\\lambda$3727.  Note that the redshifts and widths of \\oii, \\oiii, and \\hbn\\ are almost the same.  Note also that we added an additional redshifted Gaussian line to fit \\hb. The width and velocity shift of this redshifted \\hb\\ component are consistent with those of \\feii\\ emission. See text and Table \\ref{tab-redhb} for more details.} \\label{fig-redhb} \\end{figure*}  Figure \\ref{fig-redhb} shows the spectrum of SDSS J094603.94+013923.6, an example of a OQ 208-like source.  Panel ({\\it a}) shows the continuum decomposition, following the same convention as used in the top panel of Figure \\ref{fig-conti1}. Panel ({\\it b}) shows the emission-line spectrum after continuum subtraction.  The red solid line is our \\feii\\ model.  The blue dashed line marks the position of the peak of \\feii\\ $\\lambda$4924 at zero velocity shift. The shift of the \\feii\\ spectrum is obvious.  Panel ({\\it c}) illustrates the detailed fitting of the \\hb\\ and \\oiii\\ emission lines.  The profile of \\hbn\\ is fixed to that of \\oiii.  The two blue dashed lines mark the rest-frame wavelength of \\hb\\ and \\oiii\\ $\\lambda$5007; both \\hbn\\ and \\oiii\\ share the same velocity.  A prominent, intermediate-width redshifted component is clearly required to fit \\hb.  Panel ({\\it d}) shows the fit for \\oii\\ $\\lambda$3727, which also has the same redshift as \\oiii.  This means that this source is not an \\oiii\\ ``blue outlier'' \\citep{boroson05}. Table \\ref{tab-redhb} lists the FWHMs and velocity shifts of \\feii, \\oiii, \\oii, and each component of \\hb. The velocity shifts are all with respect to \\oiii. The width and velocity shift of the redshifted \\hb\\ component are consistent with those of \\feii\\ emission.  The emission-line spectrum of SDSS J094603.94+013923.6 can be divided into three systems: (1) \\hb, \\oiii, and \\oii\\ emission lines with narrow (FWHM $\\approx 500$ \\kms) widths and no velocity shift; (2) a normal ``broad'' \\hb\\ component with FWHM $\\approx 5700$ \\kms, approximately at rest with respect to the narrow lines; and (3) \\feii\\ emission and \\hb\\ of intermediate width (FWHM $\\approx 1500$ \\kms) redshifted by $\\sim 1500-1700$ \\kms.  \\begin{deluxetable*}{ccccccc} \\tablewidth{0pt} \\tablecolumns{7} \\tablecaption{Emission-line Properties of SDSS J094603.94+013923.6 \\label{tab-redhb}} \\tablehead{ \\colhead{} & \\colhead{Fe} & \\multicolumn{3}{c}{\\hb} & \\colhead{\\oiii\\tablenotemark{a}} & \\colhead{\\oii} \\\\ \\cline{3-5} \\colhead{} & \\colhead{} & \\colhead{Redshifted} & \\colhead{Broad} & \\colhead{Narrow} & \\colhead{} & \\colhead{} } \\startdata FWHM (\\kms)  & 1543(55) & 1428(18) & 5730(58) & \\nodata\\tablenotemark{b} & 505(6) & 568(18) \\\\ Velocity shift (\\kms) & 1533(24) & 1756(7)  & 373(25)  & 49(10) & \\nodata\\tablenotemark{c} & $-$22(8) \\enddata \\tablecomments{The number in parenthesis is the error.} \\tablenotetext{a}{Refers to the line core component used to derive the redshift.} \\tablenotetext{b}{The width of the \\hbn\\ is fixed to the width of \\oiii.} \\tablenotetext{c}{The redshift of \\oiii\\ is used as the systemic redshift.} \\end{deluxetable*}  OQ 208-like sources can offer an unique view to understand the origin of \\feii\\ emission and the structure of the BLR.  A detailed study of this class of objects is beyond the scope of this paper.  We will analyze a sample of such sources discovered in our work in a forthcoming publication. For the purposes of the present discussion, we simply note: \\feii\\ emission is not associated with the conventional broad component of \\hb\\ but instead originates from the same region emitting the intermediate component of \\hb. This is consistent with our findings in \\S \\ref{resultshift} that the width of \\feii\\ is systematically narrower than that of \\hbb.  \\subsection{Where is the \\feii\\ Emission Region?} \\label{feorig}  This paper has demonstrated that \\feii\\ emission most likely does not originate from the same location that produces the broad component of \\hb. The evidence comes from the systematic redshifts and narrower line widths of \\feii\\ compared to \\hbb.  As discussed below, the simplest interpretation is that \\feii\\ emission originates from an inflow that is located at the outer parts of the BLR.  The \\hbb\\ emission line itself shows no systematic velocity shift, and any velocity shift itself is also small. This suggests that \\hbb\\ emission region is well virialized and that the width of \\hbb\\ is dominated by gravity, making this line suitable for estimating BH masses.  Our finding that \\fwhmfe\\ $\\approx 3/4$ \\fwhmhbb\\ suggests that the \\feii\\ emission region is farther from the central BH than the \\hbb\\ emission region. If the \\feii\\ emission region is also virialized, so that $R\\propto v^{-2}$, then it is about 2 times farther from the center than the \\hbb\\ region.  On the other hand, the systematic redshift of \\feii\\ indicates that the assumption of virialization may be incorrect.  The redward shift of \\feii\\ and the inverse correlation between \\vfe\\ and \\er\\ favor a scenario in which \\feii\\ emission emerges from an inflow. To explain the systematic redshift, the inflow on the back side of the accretion disk must be obscured by the accretion disk and the torus, so that we can only observe the redshifted part on the nearer side.  To explain the inverse correlation between \\vfe\\ and \\er, we speculate that the inflow is driven by gravity toward the center and decelerated by the radiation pressure.  An increase in \\er\\ would enhance the radiation pressure and lead to a decrease of the inward velocity of the inflow.  Previous studies of the \\civ\\ line have shown that it tends to be systematically blueshifted \\citep[e.g.,][]{gastell82,wilkes84,marzi96,sul00a,richards02b}. This finding has led to suggestions that high-ionization lines and low-ionization lines originate from distinct regions \\citep[see also the result of reverberation mapping][]{peterson99}. High-ionization lines such as \\civ\\ may be emitted from some kind of outflowing disk wind \\citep[e.g.,][and references therein]{leighly04a,leighly04b,baskin05,shang07}, whereas low-ionization lines such as \\hb\\ are anchored to a more disk-like configuration.  This paper adds an additional element to this picture.  We suggest that in addition to a disk and a wind, the BLR has yet another component, one associated with inflowing material that produces the \\feii\\ emission.  Interestingly, \\citet{welsh07} recently reanalyzed the spectral variability of the well-studied Seyfert 1 galaxy NGC 5548 and suggested, based on the differential lag between the red and blue wings of the \\hb\\ profile, that the BLR in this object contains an inflowing component. Such an inflow is likely to develop from the inner edge of the dusty torus, which may connect with the accretion disk. If this picture is correct, the redshifted \\feii\\ emission can be used as a probe of the transition region from the dusty torus to the BLR or accretion disk. More detailed theoretical modelling and observations of the \\feii\\ emission region are clearly required to test this picture."
    },
    "0807/0807.2868_arXiv.txt": {
        "abstract": "We develop a model for the physical origin and redshift evolution of spheroid scaling relations. We consider spheroid sizes, velocity dispersions, dynamical masses, profile shapes (Sersic indices), stellar and supermassive black hole masses, and their related scalings. Our approach combines advantages of prior observational constraints in halo occupation models and libraries of high-resolution hydrodynamic simulations of galaxy mergers. This allows us to separate the relative roles of dissipation, dry mergers, formation time, and evolution in progenitor properties, and identify their impact on observed scalings at each redshift. We show that, at all redshifts, dissipation is the most important factor determining spheroid sizes and fundamental plane scalings, and can account (at $z=0$) for the observed fundamental plane tilt and differences between observed disk and spheroid scaling relations. Because disks (spheroid progenitors) at high redshift have characteristically larger gas fractions, this predicts more dissipation in mergers, yielding systematically more compact, smaller spheroids. In detail, this gives rise to a mass-dependent evolution in the sizes of spheroids of a given mass, that agrees well with observations. This relates to a subtle weakening of the tilt of the early-type fundamental plane with redshift, important for a number of studies that assume a non-evolving stellar mass fundamental plane. This also predicts evolution in the black hole-host mass relations, towards more massive black holes at higher redshifts. Dry mergers are also significant, but only for large systems which form early -- they originate as compact systems, but undergo a number of dry mergers (consistent with observations) such that they have sizes at any later observed redshift similar to systems of the same mass formed more recently. Most of the observed, compact high-redshift ellipticals will become the cores of present BCGs, and we show how their sizes, velocity dispersions, and black hole masses evolve to become consistent with observations. We also predict what fraction might survive intact from early formation, and identify their characteristic $z=0$ properties. We make predictions for residual correlations as well: e.g.\\ the correlation of size and fundamental plane residuals with formation time of a given elliptical, that can be used as additional tests of these models. ",
        "introduction": "\\label{sec:intro}  Understanding the scaling relations between the photometric and kinematic properties of galaxy spheroids -- their masses, sizes, velocity dispersions, and luminosities -- is fundamental to explaining their origin. Any model which attempts to account for these correlations must, of course, also determine how they evolve as a function of redshift. In a $\\Lambda$CDM universe, objects grow hierarchically, implying that ellipticals do not comprise a monolithic population formed at an early epoch, but have evolved and grown by mergers from $z\\gtrsim6$ to $z=0$.  Observations are beginning to probe the populations of spheroids at these epochs, and reveal that there is indeed evolution in such scaling relations \\citep[see e.g.][]{trujillo:size.evolution,mcintosh:size.evolution, zirm:drg.sizes,vandokkum:z2.sizes}, generally in the sense of massive ellipticals being more compact at high redshifts.  Some of these systems, in fact, appear sufficiently compact \\citep[see e.g.][]{zirm:drg.sizes,vandokkum:z2.sizes} that there may be no similar $z=0$ analogs whatsoever, leading to considerable debate regarding the fate of such objects.  Ultimately, the causes of such evolution and its implication for e.g.\\ the fundamental plane of ellipticals at those redshifts (and today) are not well understood.  Moreover, observations now demonstrate that essentially all massive spheroids host a supermassive black hole at their centers \\citep{KormendyRichstone95}. The mass of this black hole is tightly correlated with a variety of structural parameters: the mass \\citep{magorrian}, velocity dispersion \\citep{FM00,Gebhardt00}, and concentration/Sersic index \\citep{graham:concentration,graham:sersic}.  Recently, \\citet{hopkins:bhfp.obs} and \\citet{aller:mbh.esph} demonstrated that these correlations could be understood in terms of a ``fundamental plane''-like relation, essentially one where black hole (BH) mass tracks the binding energy of the stellar bulge. \\citet{hopkins:bhfp.theory} showed that this is the natural expectation of theories where some form of feedback from accretion leads to self-regulated growth of the BH.  This has important consequences for e.g.\\ quasar lightcurves \\citep{hopkins:faint.slope} and lifetimes \\citep{hopkins:lifetimes.letter}, as well as for the effects of ``feedback'' on the host galaxy, possibly necessary for ``quenching'' of star formation to create red elliptical galaxies at $z=0$ \\citep{croton:sam,hopkins:red.galaxies,hopkins:clustering,hopkins:groups.ell}.  Observations suggest that there may be evolution in the black hole host-galaxy scalings \\citep[e.g.][]{peng:magorrian.evolution}, but the details remain ambiguous \\citep[see e.g.][]{woo06:lowz.msigma.evolution, salviander:msigma.evolution,lauer:mbh.bias}.  As pointed out in \\citet{hopkins:bhfp.theory}, at least some such evolution should be a natural consequence of evolution in host spheroid scaling relations, provided that there is some fundamental quantity or property traced by black hole mass. Again, if true, this is of basic importance for models of black hole growth, AGN and quasar activity, star formation, and galaxy formation and evolution.  A number of important questions therefore arise: If indeed the observations are correct, why are these high-redshift spheroids different from their analogs at $z=0$?  What can this tell us about the process of spheroid formation? Can we use them to constrain and test models for spheroid formation, quenching, and black hole growth? Are such objects typical, or are they the result of complex selection effects? What happens to them by $z=0$? Can we find the remnants or left-overs of the population today?  Are evolution in spheroid structural properties related to evolution in their hosted black holes? And how is this connected to the cosmological process of quasar fueling and black hole growth?  In a series of papers, \\citet{hopkins:cusps.mergers,hopkins:cusps.ell, hopkins:cores,hopkins:cusps.fp} (hereafter \\paperone-\\paperfour), we combined libraries of hydrodynamic simulations of galaxy merger remnants and observations of nearby ellipticals spanning the largest available dynamic range in e.g.\\ spatial scale, surface brightness, and galaxy properties, and developed a methodology by which we could empirically separate spheroids into their two dominant physical components. First, a dissipationless component -- i.e.\\ an ``envelope,'' formed from the violent relaxation/scattering of stars that are already present in merging stellar disks that contribute to the final remnant. Because disks are very extended, with low phase-space density (and collisionless processes in a merger cannot raise this phase-space density), these stars will necessarily dominate the profile at large radii, hence the envelope, with a low central density.  Second, a dissipational component -- i.e.\\ a dense central relic of starbursts from gas that has been brought into the remnant, lost its angular momentum, and turned into stars in a compact central starburst similar to those observed in e.g.\\ local ULIRGs (although these would represent the most extreme cases).  Because gas can radiate, it can collapse to very high densities, and will dominate the profile within radii $\\sim0.5-1\\,$kpc, accounting for the high central densities in ellipticals. The gas will probably reflect that initially brought in from merging disks, but could in principle also stem from new cooling or stellar mass loss in the elliptical \\citep[see e.g.][]{ciottiostriker:recycling}.  In subsequent mergers the two components will act dissipationlessly (they are just stars and dark matter), but the segregation between the two is sufficient that they remain distinct even after multiple major dissipationless or ``dry'' re-mergers: i.e.\\ one can still, in principle, distinguish the dense central stellar component that is the remnant of the combined dissipational starburst(s) from the less dense outer envelope that is the remnant of low-density disk stars.  In \\paperone\\ we developed a methodology to empirically separate these components in observed systems with data of sufficient quality, and tested this on observed samples of nearby merger remnants from \\citet{rj:profiles}. Using e.g.\\ comparison with stellar populations, colors, and other properties as well as direct comparison of simulations with observed surface brightness profiles, galaxy shapes, and kinematics, we confirmed that this could reliably extract the dissipational component of the galaxy \\citep[see also][and references therein]{jk:profiles}. In \\papertwo\\ we extended this to observed ellipticals with central ``cusps'' and in \\paperthree\\ to those with central ``cores'' (believed to have undergone subsequent dry mergers, but as we argue, still retaining evidence for the initial dissipational/dissipationless stellar components), from \\citet{jk:profiles} and \\citet{lauer:bimodal.profiles}. In \\paperfour, we used these results to argue that the relative mass fraction in the dissipational component -- i.e.\\ the total mass fraction of the elliptical built up dissipationally, as opposed to brought in as stars in stellar disks, was the most important parameter determining the sizes and densities of ellipticals at $z=0$. We showed that this fraction depends systematically on mass: low mass ellipticals experienced more dissipation than high-mass ones (presumably reflecting the well-established observational fact that low-mass disks are more gas-rich than higher-mass disks). We then empirically demonstrated that this systematic dependence is both necessary and sufficient to explain e.g.\\ the difference between the size-mass relation of ellipticals and that of their (ultimate) progenitor disks, and to explain the observed ``tilt'' in the fundamental plane relation of spheroids.  There are, however, limitations in our ability to theoretically model these processes using idealized simulations and to understand them observing only $z=0$ ellipticals. Many galaxies are expected to have had complex merger histories (perhaps forming early in a rapid series of multiple, highly dissipational mergers, rather than a single disk-disk merger at late times), which might, in principle, ``smear our'' or bias some of these comparisons or systematically alter some predictions \\citep[see e.g.][]{kobayashi:pseudo.monolithic,naab:etg.formation}. Moreover, the gas content in pre-merger disks (and possibly in ellipticals themselves) is both expected and observed \\citep[e.g.][]{erb:lbg.gasmasses} to evolve strongly as a function of redshift. The disks building up ellipticals at $z\\gtrsim2$ are much more gas rich, so their remnant ellipticals at these times would necessarily be more dissipational, and as such might have different structural properties and obey different scaling relations from their $z=0$ counterparts. It is therefore of particular interest to combine our empirical understanding of the role of dissipation from local samples, and our inferences from idealized simulations, with fully cosmological models that can account for the evolution in elliptical progenitors and merger histories with redshift.  Some initial attempts at studying these correlations have been made using cosmological simulations \\citep{naab:etg.formation,dimatteo:cosmo.bhs, sijacki:radio,croft:cosmo.morph.density}. It is, unfortunately, prohibitively expensive to simultaneously resolve the $\\sim100\\,$pc scales necessary to meaningfully predict e.g.\\ spheroid scaling lengths, black hole masses, and central velocity dispersions in a fully cosmological simulation which would also contain a volume capable of modeling a large number of galaxies up to masses $\\sim10^{12}\\,\\msun$. Moreover, disk formation remains an unsolved problem in such simulations, with the uncertain physics of e.g.\\ star formation and feedback probably important to form disks with the appropriate scale lengths, thicknesses, and bulge-to-disk ratios for their masses \\citep{robertson:cosmological.disk.formation, donghia:disk.ang.mom.loss,governato:disk.formation,ceverino:cosmo.sne.fb, zavala:cosmo.disk.vs.fb}. If initial disks do not have the appropriate observed structural properties, then even a perfect model for spheroid formation will have severe systematic inaccuracies.  In order to get around some of these limitations, alternative attempts have been made to instead predict scaling resolutions from semi-analytic models of galaxy formation \\citep{khochfar:size.evolution.model, almeida:durham.ell.scalings}. Although they could not calibrate the details of their predictions with numerical simulations, \\citet{khochfar:size.evolution.model} used this to illustrate, with very simple assumptions, that increasing dissipational fractions in ellipticals at higher redshift leads to the expectation that spheroids should be more compact. \\citet{almeida:durham.ell.scalings} explicitly ignored dissipation, and as a result found that none of the models they considered could reproduce the observed $z=0$ fundamental plane correlations of spheroids (nor their redshift evolution).  Such models, however, potentially suffer from the same problems regarding the nature of disks (the ``initial conditions''): uncertainties related to modeling spheroid formation are difficult to disentangle from prescriptions for e.g.\\ cooling, disk growth, star formation, and feedback.  Moreover, while a few properties can easily be estimated (in a mean sense) from {\\em a priori} analytic arguments, many important properties and physical factors affecting spheroid evolution -- the role of dissipation in changing the shape and velocity structure of a spheroid, the detailed density profiles (especially at small radii and in extended envelopes) of ellipticals, and the dark matter to stellar or baryonic mass ratios within various radii in the galaxy -- require detailed, high-resolution numerical simulations to robustly model.  One way of circumventing the problems in these cosmological models is to adopt a halo occupation distribution (HOD) approach, in which (beginning from a halo+subhalo distribution determined in numerical simulations) halos are populated with galaxies according to empirical constraints on their clustering properties. The success of such models at reproducing the observed distributions at a variety of redshifts, as a function of various galaxy properties, has now been demonstrated \\citep[see e.g.][]{conroy:monotonic.hod}.  This allows us to begin with the empirical knowledge of where galaxies lie, and what their gas properties are, without reference to some (potentially incomplete) {\\em a priori} model of e.g.\\ star formation and feedback (which we are not, in any case, interested in testing here). Using this approach to identify and follow galaxies and their subsequent mergers, we can then reference each merger to an idealized hydrodynamic simulation with similar properties (e.g.\\ initial morphologies, gas fractions, sizes, masses, etc.), and use this to predict what the properties of the remnant should be. Constructing a sample of such remnant spheroids at any given redshift is straightforward, and allows us to compare with observed scaling relations in a direct manner.  In this paper, we adopt this method to combine the advantages of high-resolution simulations, semi-analytic models, and observational constraints on spheroid progenitors in order to develop robust predictions for the evolution of spheroid structure and the correlation between black hole mass and host properties. We study how e.g.\\ the evolution in disk gas fractions, structural properties, and the interplay between gas-rich and gas-poor mergers and multiple mergers in a hierarchical cosmology relate to observed correlations from $z=0-4$.  In \\S~\\ref{sec:model}, we describe our methodology, including our cosmological models of merger histories (\\S~\\ref{sec:model:cosmology}), empirical models for progenitor disks (\\S~\\ref{sec:model:hod}), and library of simulations which we use to determine remnant properties (\\S~\\ref{sec:model:sims}). In \\S~\\ref{sec:z0}, we consider and compare observations to the resulting predicted scalings of spheroids at $z=0$, and specifically highlight the effects of dissipation (\\S~\\ref{sec:z0:diss}) and dry mergers (\\S~\\ref{sec:z0:dry}).  In \\S~\\ref{sec:evol} we make a number of predictions and comparisons to observations regarding how these correlations evolve with redshift, and again illustrate the effects of dissipation (\\S~\\ref{sec:evol:diss}) and dry mergers (\\S~\\ref{sec:evol:dry}).  In \\S~\\ref{sec:track} we follow the history of individual systems forming at different redshifts to compare them at later redshifts, and discuss the fates of early-forming galaxies which (initially) reflect these evolved correlations. Finally, in \\S~\\ref{sec:track}, we summarize our results and outline future observational tests.  Throughout, we adopt a WMAP5 $(\\Omega_{\\rm M},\\,\\Omega_{\\Lambda},\\,h,\\,\\sigma_{8},\\,n_{s}) =(0.274,\\,0.726,\\,0.705,\\,0.812,\\,0.960)$ cosmology \\citep{komatsu:wmap5}, and normalize all observations and models shown to these parameters. Although the exact choice of cosmology may systematically shift the associated halo masses (and high-redshift evolution) at a given galaxy mass (primarily scaling with $\\sigma_{8}$), our comparisons (in terms of stellar mass) are for the most part unchanged (many of these differences are implicitly normalized out in the halo occupation approach). Repeating our calculations for a ``concordance'' $(0.3,\\,0.7,\\,0.7,\\,0.9,\\,1.0)$ cosmology or the WMAP1 $(0.27,\\,0.73,\\,0.71,\\,0.84,\\,0.96)$ and WMAP 3 $(0.268,\\,0.732,\\,0.704,\\,0.776,\\,0.947)$ results of \\citet{spergel:wmap1} and \\citet{spergel:wmap3} has little effect on our conclusions. We also adopt a \\citet{salpeter64} stellar initial mass function (IMF), and convert all stellar masses and mass-to-light ratios accordingly. Again, the choice of the IMF systematically shifts the normalization of stellar masses herein, but does not substantially change our comparisons. All magnitudes are in the Vega system, unless otherwise specified.   \\breaker ",
        "conclusions": "\\label{sec:discuss}  We have combined the results from a large library of numerical simulations with simple, observationally constrained models for the halo occupation (and corresponding merger history) of galaxies in order to predict how spheroid scaling laws evolve. We demonstrate that, despite a variety of progenitor properties, complex merger histories, and evolution in e.g.\\ progenitor sizes and gas content with redshift, the most important property driving the structural scaling relations of spheroids at each redshift (and, in particular, making these scaling relations {\\em different} from those obeyed by disk galaxies) is the amount of dissipation involved in the formation of the spheroid. At the highest masses, dry mergers at relatively late times are also important, but are few, consistent with observational constraints. Together, these predict a significant mass-dependent evolution in spheroid scaling laws, which has important consequences for the fundamental plane, spheroid densities, and the formation histories of massive galaxies.  At $z=0$, a simple model accounting for the effects of dissipation in mergers is able to explain the size-mass, velocity dispersion-mass, and fundamental plane (dynamical-stellar mass) correlations, the relation between galaxy profile shape (Sersic index) and mass, the correlation between black hole mass and various host properties (bulge binding energy, mass, velocity dispersion, and profile shape), the relative mass fractions in compact dissipational components in ellipticals, and the abundance of cuspy/disky/rapidly rotating versus cored/boxy/slowly rotating ellipticals \\citep[see][]{hopkins:groups.ell}. Dissipation in the initial disk-disk mergers that form the spheroid progenitors is the most important parameter controlling these relations.  In particular, at high masses, where disks are gas-poor, the sizes, densities, and dynamical mass to stellar mass ratios of ellipticals reflect those of their progenitor disks. At lower masses, where disks are gas-rich, mergers involve a significant amount of new star formation in dissipational nuclear starbursts triggered by loss of energy and angular momentum in the gas in the merger, building up a compact central stellar distribution (the ``dissipational'' component) that yields a smaller, more dense remnant. This mass-dependence in the dissipational content of progenitor disks makes the size-mass relation of ellipticals steeper than that of disks and gives the observed tilt in the fundamental plane (essentially the statement that low-mass disks, being more dominated by this nuclear dissipational component, have higher baryon-to-dark matter fractions within their {\\em stellar} effective radii). Black holes are formed with a mass that, in a simple feedback-regulated model, corresponds to the binding energy of the dissipational component of bulge at the time of this dissipation, which the black hole must work against in order to self-regulate its growth. This matches the detailed observations of a black hole fundamental plane in \\citet{hopkins:bhfp.obs} and gives rise to the secondary (indirect) correlations between black hole mass and spheroid velocity dispersion, mass, size, and Sersic index (all of which agree with those observed).  To lowest order, the dissipational fractions that give rise to these trends at $z=0$ reflect the observed gas fractions of disk galaxies at moderate redshifts $z\\sim0-2$. This is because most ellipticals are formed in a relatively small number of major mergers at low redshifts in this range. It is only when one considers the most massive populations of e.g.\\ BCGs and massive cored, boxy, slowly rotating ellipticals that dry mergers become important. Their effect is to puff up these systems, producing extended envelopes and raising their outer Sersic indices. This is sufficient to explain the observed weak dependence of outer Sersic index on mass in elliptical galaxies, following the detailed multi-component decompositions of observed systems in \\papertwo\\ and \\paperthree. When systems are fitted to a single Sersic index, the result is a steeper dependence that reflects a complex combination of e.g.\\ the outer Sersic index (primarily driven by merger history), the mass fraction in a dissipational component (yielding rising densities at small radius), the shape of the two components where they are comparable, and the dynamic range of the fit (since real ellipticals are not perfect Sersic profiles, there is some dependence, and in particular when fitting the entire profile to a single Sersic law, the dependence can be significant).  The detailed interplay of dissipation and dry mergers also gives rise to important residual trends that can be used as tests of these models. For example, at a given stellar mass, the oldest systems will have the largest effective radii and dynamical masses. This is driven by two effects: first, disks which have high star formation efficiencies (earlier star formation times) will be more gas-poor at the time of their mergers -- i.e.\\ will have less mass in a dissipational component. Second, systems which may be gas-rich but form very early (i.e.\\ have a high redshift of their last gas-rich merger) tend to do so because they are in dense environments, and will therefore experience subsequent mergers with both more gas-poor disks and other spheroids, increasing their size and dynamical-to-stellar mass ratios.  At higher redshifts, disk gas fractions are expected and observed to be systematically larger. As a result, spheroids observed at these redshifts are expected to have formed in more gas-rich mergers and to have a larger mass fraction in their dissipational component, making them more compact. We predict that elliptical sizes evolve in a mass-dependent fashion: since low-mass disks are still quite gas-rich at $z=0$, there is little room for them to become more so at high redshift, so they cannot evolve much. High mass disks, on the other hand, could in principle be much more gas rich than their observed $z=0$ $\\lesssim10\\%$ values. As well, at high masses, dry mergers become important in the predicted size evolution. Consequently, the highest mass ellipticals are expected to evolve to become a factor $\\sim2-3$ smaller at $z\\gtrsim2$. In terms of the mean effective radius of spheroids of a given stellar mass, relative to that at $z=0$, we find evolution of the form $(1+z)^{-\\beta}$ where $\\beta\\approx0.23\\,\\log{(M_{\\ast}/10^{9}\\,\\msun)}$ (of course enforcing $\\beta>0$ or whatever minimum evolution is set by observed evolution in progenitor disk sizes at these masses).  This is consistent with observations from e.g.\\ \\citet{trujillo:size.evolution, mcintosh:size.evolution,zirm:drg.sizes,vandokkum:z2.sizes}, although we highlight important biases arising from stellar population gradients which are relevant when attempting to estimate the sizes of the most massive galaxies at $z\\gtrsim2$. We demonstrate that this size evolution owes almost entirely to the change in dissipation with redshift, although dry mergers are of some importance for the most massive systems. Whether we explicitly include it or not, the evolution cannot be explained by simple scaling of disk sizes with redshift \\citep[which, in any case, is observed to be weak; see e.g.][and references therein]{somerville:disk.size.evol}.  This evolution has important implications for the other fundamental plane correlations. Because the fundamental plane tilt (at $z=0$) arises owing to systematically higher dissipational fractions in low-mass systems (making them more compact and yielding lower dark matter fractions within the stellar $R_{e}$), at high redshift (where even high-mass systems must be highly dissipational, since disk gas fractions are much higher) we expect the tilt to be weaker. Note that we refer here to the {\\em stellar mass} fundamental plane, i.e.\\ $M_{\\rm dyn}$ versus $\\mstar$ -- there will of course be stellar population effects introducing normalization and tilt changes in various observed bands. The predicted change is within the $z=0$ uncertainties at $z\\le1$, but by $z\\gtrsim3$, the tilt is predicted to largely disappear. This is important both as a test of these models, but also as a caution to observational attempts to use the fundamental plane as a detailed stellar population probe. It is certainly much less significant than the expected evolution in mass-to-light ratios in optical bands, but could be important for detailed stellar population probes such as that in \\citet{vandokkum:imf.evol} which assume a fixed stellar mass fundamental plane (although selection of e.g.\\ satellite galaxies in clusters, which are less likely to undergo significant future merging relative to central galaxies in the field or groups, may be sufficient to mitigate these evolutionary effects).  We argue that velocity dispersions will change weakly with redshift, despite the evolution in effective radii, because systems with higher dissipational fractions (smaller $R_{e}$) are more baryon-dominated in their stellar effective radii (i.e.\\ the velocity dispersion includes a weaker contribution from the halo). This is important, because velocity dispersion is in general preserved or increases in subsequent re-mergers. If high-redshift massive galaxies had extreme velocity dispersions $\\gtrsim600\\,{\\rm km\\,s^{-1}}$ (corresponding to the simple assumption that $\\sigma$ is inversely proportional to just the stellar effective radius), then the appropriate number of systems with equal or higher velocity dispersions would have to exist at $z=0$ \\citep[and they do not, even in survey volumes of the SDSS;][]{bernardi:most.massive}. However, accounting for the interplay between dark matter and baryonic mass, our predicted evolution in velocity dispersions is completely consistent with observational constraints.  Furthermore, because of the increasing dissipational components and deepening potential wells at high redshift, we expect that, in any feedback-regulated model of black hole growth, black holes must become at least somewhat more massive (relative to their host spheroid stellar mass) at high redshift. We predict a similar mass-dependent evolution in $M_{\\rm BH}/M_{\\ast}$, with little or no significant evolution at $M_{\\ast}\\lesssim10^{10}\\,\\msun$ (corresponding to black holes with masses $\\lesssim10^{7}\\,\\msun$), and factor $\\sim2$ increase to $z\\sim2-3$ in the masses of the black holes in the most massive systems ($M_{\\ast}\\gg10^{11}\\,\\msun$, corresponding to $\\mbh\\gtrsim10^{9}\\,\\msun$). Likewise, we predict the evolution in black hole mass relative to host velocity dispersion, which exhibits a more complex behavior. The physical details of driving these trends are discussed in more detail in \\citet{hopkins:bhfp.theory}, but here we embed those models in a more fully cosmological context. The evolution predicted is consistent with recent observations \\citep{peng:magorrian.evolution} and with indirect constraints based on the relative evolution of black hole mass density and spheroid mass density \\citep{hopkins:old.age, merloni:magorrian.evolution}, however it is still relatively weak. Selection effects \\citep{lauer:mbh.bias} would be expected to explain the difference between some observational estimates that infer very strong \\citep[e.g.][]{walter04:z6.msigma.evolution} or rapid low-redshift \\citep{woo06:lowz.msigma.evolution} evolution and our predictions.  We outline the history of these massive systems which are different when first formed at high redshift (where they are very dissipational) and observed at low redshift. Most of these compact, early forming ellipticals will undergo a significant number of dry mergers, or mergers with lower-redshift, more gas-depleted and larger disks, which will serve to puff up the remnant, building up some stellar mass and an extended stellar envelope (although changing the central velocity dispersion by a relatively small amount). This raises their effective radii more rapidly than their stellar mass, in the sense of moving them up in the size-mass relation. Interestingly, the evolution in disk gas fractions and the effects of such dry mergers almost exactly offset one another in massive systems. In other words, a $\\sim10^{12}\\,\\msun$ galaxy which first formed from a merger at $z=3$ (which had a high gas content, leading to a very compact $\\gtrsim10^{11}\\,\\msun$ initial galaxy with $R_{e}\\sim1$\\,kpc) will have almost the same size at $z=0$ as a galaxy of the same observed stellar mass which formed at much later time, say $z=1$, from a less dissipational merger (because disk gas fractions have decreased in that time interval). The system which formed earlier was initially more compact, but experienced more mergers in the intervening time, so for a given stellar mass, the scatter in $R_{e}$ is not large. Likewise, velocities dispersions and other fundamental plane properties have small scatter at a given stellar mass at any observed redshift because of this cancellation effect.  This is not to say that most systems will undergo many dry mergers -- the overall density of dry mergers predicted here is low, consistent with observational constraints from e.g.\\ \\citet{bell:dry.mergers,vandokkum:dry.mergers,lin:mergers.by.type}, and as we noted above most (especially $\\lesssim L_{\\ast}$) ellipticals today formed in a few (or just one) gas-rich mergers at relatively low redshifts. However, the systems which form earliest do so because they live in the most dense environments, which evolve the most rapidly, and are expected to undergo the most subsequent merging activity. We predict, as a result, that most of these compact systems observed at high redshifts will end up constituting some of the central dense stellar mass in observed brightest cluster or group galaxies at $z=0$; this is consistent with the observed sizes \\citep{vonderlinden:bcg.scaling.relations}, surface brightness profile shapes (extended envelopes build up through their subsequent dry merging) \\citep[see e.g.\\ \\paperthree\\ and][]{gallagherostriker72, naab:dry.mergers}, and masses and velocity dispersions \\citep{bernardi:bcg.scalings} of these $z=0$ systems, and also expected given the clustering properties of high redshift massive systems \\citep{quadri:highz.color.density}.  We estimate what fraction, as a function of stellar mass, of the massive galaxy population today may have survived from these early times un-remerged and thus reflect the smaller sizes and higher dissipational content characteristic of systems formed at that time, finding that e.g.\\ $\\sim1-5\\%$ of $z=0$ $\\sim10^{11}\\,\\msun$ systems may have survived since $z\\sim2$. This is small, but not negligible, and by identifying compact ellipticals with old stellar population ages, it should be possible to locate the surviving remnants of this population. We predict that these will be massive, compact (with other indicators of gas-rich origin and high inferred dissipational fractions from their stellar profiles), with old stellar population ages. Because they still generally would initially form in high-density environments, but have avoided merging since then, we would roughly expect them to be in dense environments (rich groups and clusters) today, but {\\em not} to be the central galaxy (again, most early-forming systems will be the central galaxy -- but if they are all the way until $z=0$, it is unlikely that they will have avoided further mergers, and thus would not be un-remerged). There may be some observed candidate members of this population in the so-called ``compact elliptical'' class, galaxies similar to NGC 4486B in Virgo which appears to satisfy these criteria. Further observations to test this hypothesis could prove extremely valuable as nearby probes of high-redshift galaxy evolution.  Future work should extend these models to include the (at present more uncertain) role of particularly minor mergers (where e.g.\\ satellite disruption may be important to explain the extended halos in massive BCGs) and other sources of dissipation, such as stellar mass loss \\citep[see e.g.][]{ciottiostriker:recycling}. However, we show that other effects, e.g.\\ evolution in the sizes of disk galaxies, dry mergers alone, and evolution in galaxy host halo properties, are insufficient to explain effects like those predicted here, in particular to account for e.g.\\ the tilt of the fundamental plane, the relative sizes of ellipticals (versus those of disks) as a function of mass, the scaling of elliptical galaxy profile shapes with mass, and the redshift evolution in the fundamental plane correlations. Therefore, these predictions are important tests of any model in which dissipational star formation is a significant influence on galaxy formation.  Because dissipation is necessary in a basic sense to reconcile the densities of ellipticals and spirals, this represents an important test of the merger hypothesis itself, as well as a test of our understanding of galaxy merger histories and the evolution in different galaxy components with redshift. In addition, we provide a unified theoretical lens through which to interpret a number of observations of spheroid and black hole-host scalings at both $z=0$ and high redshift, in which spheroids are fundamentally multi-component objects (with dissipational components originally formed in central starbursts triggered in the initial gas-rich interactions that formed the system, and dissipationless components representing the scattered, violently relaxed stars from the pre-initial merger stellar disks). As multi-component objects, it is primarily the relative mass fraction in the dissipational component, reflecting the gas content of spheroid progenitors, that drives their structural properties and the evolution of those structural properties with redshift."
    },
    "0807/0807.1108_arXiv.txt": {
        "abstract": "The recent robust and homogeneous analysis of the world's supernova distance-redshift data, together with cosmic microwave background and baryon acoustic oscillation data, provides a powerful tool for constraining cosmological models.  Here we examine particular classes of scalar field, modified gravity, and phenomenological models to assess whether they are consistent with observations even when their behavior deviates from the cosmological constant $\\Lambda$.  Some models have tension with the data, while others survive only by approaching the cosmological constant, and a couple are statistically favored over $\\lcdm$.  Dark energy described by two equation of state parameters has considerable phase space to avoid $\\Lambda$ and next generation data will be required to constrain such physics. ",
        "introduction": "}  A decade after the discovery of the acceleration of the cosmic expansion \\citep{perl99,riess98} we still understand little about the nature of the dark energy physics responsible.  Improved data continues to show consistency with Einstein's cosmological constant $\\lam$, and in terms of a constant equation of state, or pressure to density, ratio $w$, the best fit to the data is $w=-0.969^{+0.059}_{-0.063}({\\rm stat})^{+0.063}_{-0.066} ({\\rm sys})$, where $\\Lambda$ has $w=-1$ \\citep{union}. However, the magnitude of $\\lam$ required and the coincidence for it becoming dominant so close to the present remain unexplained, and an abundance of motivated or unmotivated alternative models fills the literature.  Using the latest, most robust data available we examine the extent to which data really have settled on the cosmological constant.  The vast array of models proposed for dark energy makes comparison of every model in the literature to the data a Sisyphean task.  Here we select some dozen models with properties such as well defined physical variables, simplicity, or features of particular physical interest. These embody a diversity of physics, including scalar fields, phase transitions, modified gravity, symmetries, and geometric relations. While far from exhaustive, they provide roadmarks for how well we can say that current data have zoomed in on $\\lam$ as the solution.  For such comparisons it is critical to employ robust data clearly interpretable within these ``beyond $\\lam$'' cosmologies.  Geometric probes from the Type Ia supernovae (SN) distance-redshift relation, cosmic microwave background (CMB) acoustic peak scale shift parameter, and baryon acoustic oscillations (BAO) angular scale serve this essential role.  Equally important is confidence in the error estimates, incorporating systematics as well as statistical uncertainties. This has been studied in detail in the recent unified analysis of the world's published heterogeneous SN data sets -- the Union08 compilation \\citep{union}.  This SN compilation includes both the large data samples from the SNLS and ESSENCE survey, the compiled high redshift SNe observed with the Hubble Space Telescope, a new sample of nearby SNe, as well as several other, small data sets.  All SNe have been analyzed in a uniform manner and have passed a number of quality criteria (such as having data available in two bands to measure a color, and sufficient lightcurve points to make a meaningful fit). The samples have been carefully tested for inconsistencies under a blinded protocol before combining them into a single final data set comprising 307 SNe, the basis for this analysis. In this work the SNe data will be combined with the constraints obtained from the baryon acoustic oscillation scale \\citep{eisenstein05} and from the five year data release of WMAP and ground based CMB measurements \\citep{komatsu08}.  In Section~\\ref{sec:models} we describe the general method for cosmological parameter estimation and present a summary table of the various models considered and the $\\chi^2$ statistics of the fit. Sections~\\ref{sec:const}--\\ref{sec:grownu} then briefly describe the dark energy models, their parameters, and show the likelihood contours.  The concluding discussion occurs in Section~\\ref{sec:concl}. ",
        "conclusions": "}  We have considered a wide variety of dark energy physics quite different from the cosmological constant.  These include a diversity of physical origins for the acceleration of the expansion: from dynamical scalar fields to dark energy that will eventually cause deceleration and collapse, to gravitational modifications arising from extra dimensions or from quantum phase transitions, to geometric or kinematic parametrization of the acceleration, to dark energy that may have influenced the early universe and that may have its magnitude set by the neutrino mass. The comparison to $\\lcdm$ and constant $w$ cases covers 5 one-parameter and 5 two-parameter dark energy equation of state models. (\\citet{lh05} detail how even next generation data will not generically be able to tightly constrain more than two such parameters.)  Two key results to emphasize are that current data 1) are consistent with $\\Lambda$, and 2) are also consistent with a diversity of other models and theories, even when we restrict consideration to those with at least modest physical motivation or justification.  As explicitly shown by the mirage model, any inclination toward declaring $\\lam$ the answer based on consideration of a constant $w$ has an overly restricted view. The need for next generation observations with far greater accuracy, and the development of precision growth probes, such as weak gravitational lensing, is clear.  All major classes of physics to explain the nature of dark energy are still in play.  However there are already quite hopeful signs of imminent progress in understanding the nature of dark energy.  For example, for the braneworld model tight control of systematics would decrease the goodness of fit to $\\Delta\\chi^2=+15$, even allowing for spatial curvature, diminishing its likelihood by a factor 2000 na{\\\"\\i}vely, effectively ruling out the model.  For the doomsday model, improving errors by 30\\% extends our ``safety margin'' against cosmic collapse by 10 billion years -- a nonnegligible amount!  Every improvement in uncertainties pushes the limits on the neutrino mass within the growing neutrino model closer toward other astrophysical constraints -- plus this model essentially guarantees a deviation from $w=-1$ of $0.1\\,(m_{\\nu 0}/{\\rm eV})$, excitingly tractable.  Terrestrial neutrino oscillation bounds already provide within this model that $1+w>0.005$.  As points of interest, we note that the model with noticeably positive $\\Delta\\chi^2$ relative to $\\Lambda$, and hence disfavored, is completely distinct from the cosmological constant, i.e.\\ the braneworld model has no limit within its parameter spaces equivalent to $\\Lambda$.  This does not say that no such model could fit the data -- the $R_{\\rm low}$ model is also distinct from $\\Lambda$ but fits as well as many models.  Certainly many successful models under current data do look in some averaged sense like a vacuum energy but this does not necessarily point to static dark energy.  Two serious motivations to continue looking for deviations are that physicists have failed for 90 years to explain the magnitude required for a cosmological constant, and that the previous known occurrence of cosmic acceleration -- inflation -- evidently involved a dynamical field not a cosmological constant.  To guide further exploration of the possible physics, we highlight those models which do better than $\\Lambda$: the geometric dark energy and algebraic thawing approaches.  One of the sole models where adding a degree of freedom is justified (albeit modestly) by the resulting reduction in $\\chi^2$ is the ${\\rm R_{high}}$ model directly studying deviations of the spacetime curvature from the matter dominated behavior.  This has one more parameter than the constant $w$ EOS approach, but improves in $\\chi^2$ by 1.  In addition, it has a built-in test for the asymptotic de Sitter fate of the future expansion.  We recommend that this model be considered a model of interest for future fits.  The other model improving by at least one unit of $\\chi^2$ is the algebraic thawing model, performing better than the other thawing models, with a general parametrization explicitly incorporating the physical conditions imposed by matter domination on the scalar field dynamics.  The diversity of models also illustrates some properties of the cosmological probes beyond the familiar territory of vanilla $\\lcdm$. For example, for the algebraic thawing and other such evolutionary models, the premium is on precision of $w_0$ and $w_a$ much more than the averaged or pivot EOS value $w_p$.  Not all models possess the wonderful three-fold complementarity of the probes seen in the constant $w$ case; for many of the examples BAO and CMB carry much the same information as each other.  However, we clearly see that for every model SN play a valuable role, complementary to CMB/BAO, and often carries the most important physical information: such as on the doomsday time or the de Sitter fate of the universe or the Planck scale nature of the PNGB symmetry breaking.  The diversity of physical motivations and interpretations of acceptable models highlights the issue of assumptions, or priors, on how the dark energy should behave.  For example, in the $R_{\\rm low}$ model should priors be flat in $r_0$, $r_1$ or in $\\om$, $w_0$; in the PNGB model should they be flat in $f$, $\\phi_i/f$ or in $\\om$, $w_0$, etc.? Lacking clear physical understanding of the appropriate priors restricts the physical meaning of any Bayesian evidence one might calculate to employ model selection; the $\\chi^2$ goodness of fit used here does not run into these complications that can obscure physical interpretation.  We can use our diversity of models for an important consistency test of our understanding of the {\\it data\\/}.  If there would be systematic trends in the data which do not directly project into the $\\lcdm$ parameter space (i.e.\\ look like a shift in those parameters), then one might expect that one of the dozen models considered might exhibit a significantly better fit. The fact that we do not observe this can be viewed as evidence that the data considered here is not flawed by significant hidden systematic uncertainties.  The data utilize the Union08 compilation of uniformly analyzed and crosscalibrated Type Ia supernovae data, constituting the world's published set, with systematics treated and characterized through blinded controls.  The data are publicly available at http://supernova.lbl.gov/Union, and will be supplemented as further SN data sets become published; the site contains high resolution figures for this paper as well.  However, to distinguish deeply among the possible physics behind dark energy requires major advances in several cosmological probes, enabling strong sensitivity to the time variation of the equation of state.  This is especially true for those models that are now or were in the past close to the cosmological constant behavior.  We are getting our first glimpses looking beyond $\\Lambda$, but await keen improvements in vision before we can say we understand the new physics governing our universe."
    },
    "0807/0807.1752_arXiv.txt": {
        "abstract": "We report the detection in CO of the far-side counterpart of the well-known expanding 3-Kpc Arm in the central region of the Galaxy. In a CO longitude-velocity map at $b = 0^{\\circ}$ the Far 3-Kpc Arm can be followed over at least $20^{\\circ}$ of Galactic longitude as a faint lane at positive velocities running parallel to the Near Arm. The Far Arm crosses $l = 0^{\\circ}$ at +56 \\kms, quite symmetric with the $-53$ \\kms expansion velocity of the Near Arm. In addition to their symmetry in longitude and velocity, we find that the two arms have linewidths ($\\sim21$ \\kms), linear scale heights ($\\sim103$ pc FWHM), and H$_2$ masses per unit length ($\\sim4.3$ x $10{^6}$ \\msun kpc$^{-1}$) that agree to 26\\% or better. Guided by the CO, we have also identified the Far Arm in high-resolution 21 cm data and find, subject to the poorly known CO-to-H$_2$ ratio in these objects, that both arms are predominately molecular by a factor of 3--4.  The detection of these symmetric expanding arms provides strong support for the existence of a bar at the center of our Galaxy and should allow better determination of the bar's physical properties. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0807/0807.4177_arXiv.txt": {
        "abstract": "We investigate the implications of the intergalactic opacity for the evolution of the cosmic ultraviolet luminosity density and its sources. Our main constraint follows from our measurement of the Lyman-$\\alpha$ forest opacity at redshifts $2\\leq z\\leq4.2$ from a sample of 86 high-resolution quasar spectra. In addition, we impose the requirements that intergalactic HI must be reionized by $z=6$ and HeII by $z\\approx3$, and consider estimates of the hardness of the ionizing background from HI to HeII column density ratio measurements. The derived hydrogen photoionization rate is remarkably flat over the \\Lya~forest redshift range covered. Because the quasar luminosity function is strongly peaked near $z\\sim2$, the lack of redshift evolution indicates that star-forming galaxies likely dominate the photoionization rate at $z\\gtrsim3$, and possibly at all redshifts probed. Combined with direct measurements of the galaxy UV luminosity function, this requires only a small (emissivity weighted) fraction $f_{\\rm esc}\\sim0.5\\%$ of galactic hydrogen ionizing photons to escape their source for galaxies to solely account for the entire ionizing background. Under the assumption that the galactic UV emissivity traces the star formation rate (which is the case if the escape fraction and dust obscuration are constant with redshift), current state-of-the-art observational estimates of the star formation rate density, which peak similarly to quasars at $z\\sim2$, appear to underestimate the total photoionization rate at $z\\approx4$ by a factor $\\approx4$, are in tension with the most recent determinations of the UV luminosity function, and fail to reionize the Universe by $z=6$ if extrapolated to arbitrarily high redshift. A star formation history peaking earlier, as in the theoretical model of Hernquist \\& Springel, fits the \\Lya~forest photoionization rate well, reionizes the Universe in time, and is in better agreement with the rate of $z\\approx4$ gamma-ray bursts observed by \\emph{Swift}. Quasars suffice to doubly ionize helium by $z\\approx3$, provided that most of the HeII ionizing photons they produce escape into the intergalactic medium, and likely contribute a non-negligible and perhaps dominant fraction of the hydrogen ionizing background at their $z\\sim2$ peak. ",
        "introduction": "\\label{introduction} The opacity of the Lyman-$\\alpha$ (\\Lya) forest is set by a competition between hydrogen photoionizations and recombinations \\citep[][]{1965ApJ...142.1633G} and can thus serve as a probe of the photonization rate \\cite[e.g.][]{1997ApJ...489....7R}. The hydrogen photoionization rate $\\Gamma$~is a particularly valuable quantity as it is an integral over all sources of ultraviolet (UV) radiation in the Universe, \\begin{equation} \\Gamma(z) = 4\\pi \\int_{\\nu_{\\rm HI}}^{\\infty} \\frac{d\\nu}{h \\nu} J_{\\nu}(z) \\sigma(\\nu), \\end{equation} where $J_{\\nu}$ is the angle-averaged specific intensity of the background, $\\sigma(\\nu)$ is the photoionization cross section of hydrogen, and the integral is from the Lyman limit to infinity. As such, it bears a signature of cosmic stellar and quasistellar activity that is not subject to the completeness issues to which direct source counts \\citep[e.g.,][]{1999ApJ...519....1S, 2004ApJ...606L..25B, 2006ApJ...642..653S, 2007ApJ...670..928B, 2006ApJ...653..988Y, 2007ApJ...654..731H} are prone. In addition, unlike the radiation backgrounds directly observed on Earth, the \\Lya~forest is a \\emph{local} probe of the high-redshift UV background, as only sources at approximately the same redshift contribute to $\\Gamma$~at any point in the forest \\citep[][]{1996ApJ...461...20H}.  The hydrogen photoionization rate is not only a powerful observational tracer of cosmic radiative history, but is also a key ingredient in cosmological simulations and theoretical models of galaxy formation. Photoionization heating can, for example, suppress the formation of dwarf galaxies \\citep[][]{1992MNRAS.256P..43E, 1996ApJ...465..608T, 1997ApJ...477....8W}. While the photoionization rate of hydrogen by itself provides limited information on the shape of the background spectrum, it can be combined with the photoionization rates of other species (helium, for example) to provide a direct measure of the background hardness \\citep[e.g.,][]{1997ApJ...488..532C, 2003MNRAS.340..473S, 2006MNRAS.366.1378B}. The background spectrum sets the temperature of the gas and so it must be specified in any hydrodynamical simulation of structure formation \\citep[e.g.,][]{1996ApJ...457L..51H, 1999ApJ...511..521D, 2003MNRAS.339..312S}. Knowledge of the ionizing background spectrum is also necessary to infer the mass density of metals from observations of absorption by particular ions (``ionization corrections'') \\citep[e.g.,][]{2003ApJ...596..768S, 2004ApJ...602...38A, 2007arXiv0712.1239A}.  Recent results make it especially worthwhile to reconsider the implications for the ionizing sources, usually assumed to be a combination of quasars and star-forming galaxies \\citep[e.g.,][]{1999ApJ...514..648M}, of the \\Lya~forest opacity in the light of new data. The comoving star formation rate (SFR) density as estimated from the observed luminosity of galaxies, for instance, is in tension with the independently derived evolution of the stellar mass density \\citep[e.g.,][]{2006ApJ...651..142H}. This suggests problems with observational estimates of the cosmic SFR, to which $\\Gamma$~is sensitive. The abundance of intergalactic metals as traced by CIV appears consistent with being constant from $z\\sim2$ to perhaps $z\\sim6$ \\citep[][]{2001ApJ...561L.153S, 2003ApJ...594..695P, 2006MNRAS.371L..78R}, suggestive of an early period of nucleosynthetic activity that may be at odds with estimates of the SFR density that decline above $z\\sim2-3$ \\citep[][]{2004ApJ...615..209H, 2006ApJ...651..142H}. Stellar mass density measurements at $z\\gtrsim5$ \\citep[e.g.,][]{2006ApJ...649L..67L, 2006ApJ...651...24Y, 2007MNRAS.374..910E, 2007ApJ...659...84S} also suggest significant early star formation. In addition, deep $z\\sim6-10$ searches utilizing magnification by foreground galaxy clusters \\citep[][]{2006A&A...456..861R, 2007ApJ...663...10S} have yielded surprisingly large abundances of low-luminosity galaxies, in tension with blank field surveys and perhaps even requiring an upturn in their abundance with increasing redshift. At lower redshifts, \\cite{2007arXiv0711.1354R} discovered a large population of faint \\Lya~emitters at $2.67\\leq z \\leq 3.75$ in an extremely deep 92-hour exposure with VLT/FORS2, most of which would not have been selected in existing Lyman-break surveys. This underscores the definite possibility that photometric-selection surveys may presently miss a substantial fraction of the cosmic luminosity density originating in low-luminosity objects.  Here, we use improved measurements of the intergalactic opacity at $z\\gtrsim2$ to obtain a measurement of the hydrogen photoionization rate $\\Gamma$. We present a detailed comparison with direct source counts of quasars and star-forming galaxies and ask whether these can account for the measured $\\Gamma$. Finally, we consider the implications for the estimates of the cosmic SFR density and the escape fraction of ionizing photons from galaxies, which relates the soft UV emission to the photons that actually contribute to photoionizations.  The observational constraints we use are described in \\S \\ref{measurements}. In \\S \\ref{gamma}, we derive the hydrogen photoionization rate implied by our recent measurement of the \\Lya~forest opacity \\citep[][]{2008ApJ...681..831F} and explore potential systematic effects. In \\S \\ref{photoionizations to emissivity}, we describe how to convert the photoionization rate to a UV luminosity density. We investigate the implications for the cosmic sources of UV radiation, quasars and star-forming galaxies in \\S \\ref{cosmic sources}. We also consider the constraints imposed by the reionizations of hydrogen and helium in this section. We revisit the star formation history in \\S \\ref{madau diagram} before discussing our results in \\S \\ref{discussion} and summarizing our conclusions \\S \\ref{conclusions}.  Throughout, we assume a cosmology with $(\\Omega_{\\rm m},~\\Omega_{\\rm b},~\\Omega_{\\Lambda},~h,~\\sigma_{8})=(0.28,~0.046,~0.72,~0.70,~0.82)$, as inferred from the \\emph{Wilkinson Microwave Anisotropy Probe} (WMAP) five-year data in combination with baryon acoustic oscillations and supernovae \\citep[][]{2008arXiv0803.0547K}. Unless otherwise stated, all error bars are $1\\sigma$. Some of the results presented here were already reported in a concise companion \\emph{Letter} \\citep[][]{2008ApJ...682L...9F}. They are described here with the full details of the analysis and underlying assumptions, as well as supporting arguments.  \\begin{deluxetable}{lcccccl} \\tablewidth{0pc} \\tablecaption{\\Lya~Effective Optical Depth and Derived Hydrogen Photoionization Rate\\label{our measurement}} \\tabletypesize{\\footnotesize} \\tablehead{\\colhead{$z$} & \\colhead{$\\tau_{\\rm eff}$}\\tablenotemark{a} & \\colhead{$\\sigma_{\\tau_{\\rm eff}}$}\\tablenotemark{b} & $\\Gamma$\\tablenotemark{c} & $\\sigma_{\\Gamma}$\\tablenotemark{d} \\\\ \\colhead{} & \\colhead{} & \\colhead{} & 10$^{-12}$ s$^{-1}$ & 10$^{-12}$ s$^{-1}$} \\startdata 2.0 & 0.127 & 0.018 & 0.64 & 0.18 \\\\ 2.2 & 0.164 & 0.013 & 0.51 & 0.10 \\\\ 2.4 & 0.203 & 0.009 & 0.50 & 0.08 \\\\ 2.6 & 0.251 & 0.010 & 0.51 & 0.07 \\\\ 2.8 & 0.325 & 0.012 & 0.51 & 0.06 \\\\ 3.0 & 0.386 & 0.014 & 0.59 & 0.07 \\\\ 3.2 & 0.415 & 0.017 & 0.66 & 0.08 \\\\ 3.4 & 0.570 & 0.019 & 0.53 & 0.05 \\\\ 3.6 & 0.716 & 0.030 & 0.49 & 0.05 \\\\ 3.8 & 0.832 & 0.025 & 0.51 & 0.04 \\\\ 4.0 & 0.934 & 0.037 & 0.55 & 0.05 \\\\ 4.2 & 1.061 & 0.091 & 0.52 & 0.08 \\\\ \\enddata \\tablenotetext{a}{Measurement reported by \\cite{2008ApJ...681..831F} corrected for continuum bias binned in $\\Delta z=0.2$ intervals and with the contribution of metals to absorption in the \\Lya~forest subtracted based on the results of \\cite{2003ApJ...596..768S}.} \\tablenotetext{b}{Statistical only.} \\tablenotetext{c}{Derived as in \\S \\ref{gamma}.} \\tablenotetext{d}{Includes contributions from the statistical uncertainties in both $\\tau_{\\rm eff}$ and $T_{0}$.} \\end{deluxetable} ",
        "conclusions": "\\label{conclusions} We have investigated the implications of the evolution of the intergalactic opacity for the sources responsible for its photoionization. The principal constraint is our measurement of the \\Lya~effective optical depth over the redshift range $2\\leq z \\leq4.2$ from a sample of 86 high-resolution quasar spectra \\citep[][]{2008ApJ...681..831F}. We also imposed the requirements that intergalactic HI must be reionized by $z=6$ and HeII by $z\\approx3$, and used estimates of the hardness of the ionizing background from HI to HeII column density ratio measurements. Our main conclusions are as follows: \\begin{enumerate} \\item The intergalactic hydrogen photoionization rate is remarkably constant over the redshift range $2 \\leq z \\leq 4.2$, with a value $\\Gamma \\approx (0.5\\pm0.1)\\times 10^{-12}$ s$^{-1}$, subject to remaining systematic uncertainties in normalization. \\item Because the quasar luminosity function is strongly peaked near $z\\sim2$, the lack of redshift evolution implies that star-forming galaxies are likely to dominate the photoionizing background at $z\\gtrsim3$, and possibly at all redshifts probed. \\item Although our arguments are robust to normalization systematics, fiducial assumptions regarding the mean free path of ionizing photons, the spectral energy distribution of quasars, and an escape fraction of ionizing photons from quasars of unity imply that quasars alone overproduce the total ionizing background near their peak. This puzzle suggests that these assumptions warrant more scrutiny and highlights the uncertainties in calculating the ionizing background by integrating luminosity functions. \\item Only a small escape fraction of ionizing photons $f_{\\rm esc}\\sim0.5\\%$ is needed for galaxies to solely account for the entire UV luminosity density implied by the \\Lya~forest, with the scaling with uncertain parameters given in equation \\ref{fesc scaling}. This small fraction might be reconciled with higher measurements from the direct detection of Lyman-continuum photons from high-redshift galaxies \\citep[][]{2001ApJ...546..665S, 2006ApJ...651..688S} if it increases with galaxy luminosity. \\item The state-of-the-art observational fit to the cosmic star formation history by \\cite{2006ApJ...651..142H} (peaking near $z\\sim2$ similarly to quasar activity) appears to underestimate the total photoionization rate by almost an order of magnitude at $z\\approx4$ if the escape fraction and dust obscuration are constant with redshift for $z\\gtrsim2$, and is also in tension with the most recent high-redshift determinations of the galaxy UV luminosity function. \\item Normalizing the ionizing emissivity predicted by the best-fit star formation history of \\cite{2006ApJ...651..142H} to our \\Lya~forest measurement at $z=3$, their star formation history fails to reionize the Universe by $z=6$. \\item A star formation history peaking at a higher redshift $z=5-6$, like the theoretical model of \\cite{2003MNRAS.341.1253H}, fits the $2\\leq z\\leq 4.2$ \\Lya~forest well, reionizes the Universe in time, and is in better agreement with the rate of GRBs observed by \\emph{Swift} at $z\\approx4$. \\item Quasars alone suffice to doubly ionize helium by $z\\approx3$ provided that most of the HeII ionizing photons they produce escape into the IGM. \\item The ratio the HI to the HeII column densities in the post-HeII reionization epoch $2\\lesssim z \\lesssim3$ suggests, with a significant uncertainty arising from the large fluctuations observed, that quasars contribute a non-negligible fraction ($\\gtrsim20$\\%) and perhaps dominate the hydrogen ionizing background at their $z\\sim2$ peak. \\end{enumerate}  Although the present uncertainties in the redshift evolution of UV dust corrections, the escape fraction of ionizing photons from galaxies, and possible evolution of the stellar initial mass function however prohibit definitive conclusions about the cosmic SFR to be reached at this time, we have illustrated how the IGM represents a powerful integral probe of the cosmic luminosity density free of the completeness issues that affect the direct detection of its sources. In this sense, its study is very complementary to that of luminosity functions.  In this work, we have strived to make the simplest assumptions consistent with our present body of knowledge. For instance, we have assumed that the escape fraction of ionizing photons from star-forming galaxies is constant with redshift and found evidence that the high-redshift ($z\\gtrsim3$) SFR density is currently underestimated by direct galaxy counts. Of course, many of our arguments could be logically reversed and, for example, instead favor an evolving escape fraction on the basis of the present estimates of the star formation history. More work is certainly needed in order to break such degeneracies, but a robust conclusion is that the simplest picture contains inconsistencies, the resolution of which is poised to teach us about early star formation, AGN activity, and perhaps more.  Better constraints on the thermal history of the IGM and on the abundance of the Lyman-limit systems which determine the mean free path of ionizing photons will be key in reducing the systematic uncertainties in the arguments presented here. In addition, improved determinations of the spectral indices of quasars and star-forming galaxies, especially shortward of the Lyman limit, will be useful in refining our calculations. In the same vein, an updated calculation of the full spectrum of the ionizing background, like that of \\cite{1996ApJ...461...20H}, is clearly warranted in the light of our improved empirical knowledge of the physical state of the IGM and of the luminosity functions of both quasars and galaxies (Faucher-Gigu\\`ere et al., in prep.)\\nocite{ionspectrum}."
    },
    "0807/0807.1145.txt": {
        "abstract": "We provide a comprehensive analytic study (rigorous and qualitative) of the conditions for the existence of a a positive kinematic back reaction term $\\QQ>0$, in the context of Buchert's scalar averaging formalism applied to spherically symmetric Lema\\^{\\i}tre--Tolman--Bondi (LTB) dust solutions in which averaging domains are given as spherical comoving regions containing a symmetry center. We introduce proper volume and quasi--local average functionals and functions in order to examine the conditions for $\\QQ\\geq 0$, and in the process we also explore the relation between back reaction, spatial curvature and binding energy for a wide variety of LTB configurations. The back reaction term is positive for all ``hyperbolic'' regular domains with negative spatial curvature, either in the full radial range or in the radial asymptotic range.  This result is also valid if these domains contain an inner ``elliptic'' region with positive curvature undergoing local collapse.  For some cases in which positive spatial curvature decreases asymptotically, the conditions for a positive back reaction can still be met but seem to be more restrictive. Since $\\QQ>0$ is a necessary condition for a positive ``effective'' acceleration that would mimic the effect of dark energy (in the context of Buchert's formalism), we examine this issue in LTB models in a follow up paper (part II). ",
        "introduction": "The possibility that observations could be influenced by diferent ways of coarse graining and averaging of dust inhomogeneities provides a number of popular alternative explanations for cosmic acceleration found in the literature~\\cite{InhObs1,InhObs2,InhObs3,InhObs4,InhObs5,InhObs6}. Among these theoretical proposals, the spatial averaging formalism developed by  Buchert~\\cite{buchert} and co--workers considers the ``effective'' acceleration that mimics a cosmological constant, which might arise from the so--called ``back--reaction'' terms that emerge by rewriting scalar evolution equations in terms of spatial averages of matter--energy density and the expansion scalar. See \\cite{ave_review} for a comprehensive review of this formalism and \\cite{IW} for further discussion. For alternative proposals on averaging inhomogeneities, see \\cite{zala,colpel}.  Unfortunately, whether based on Buchert's formalism or not, it is technically challenging to calculate in practice coarse grained observational parameters or back--reaction terms for general non--linear inhomogeneus or ``realistic'' conditions. Hence, besides perturbative approach~\\cite{pert} and idealized spacetimes~\\cite{wiltshire}, proposals that examine cosmic acceleration without dark energy have been discussed or tested for the spherically symmetric Lema\\^\\i tre--Tolman--Bondi (LTB) class of dust solutions~\\cite{InhObs1,InhObs2,InhObs5,ltbmods,LTBave1,LTBave2,LTBave3}. These models~\\cite{LTB,kras} are simple, but general enough, inhomogeneous spacetimes and so are ideal to test the effects of inhomogeneity.  See \\cite{ltbstuff1,ltbstuff2, ltbstuff3, ltbstuff,suss02} for further discussion on regularity of LTB models and \\cite{kras} for a comprehensive review.  Several articles among the references in \\cite{InhObs1,InhObs2,InhObs5,ltbmods,LTBave1} have already considered the application of Buchert's formalism to LTB models.  More recently, Paranjape and Singh~\\cite{LTBave2} utilized asymptotic or late time approximations, while Chuang, Gu and Hwang \\cite{LTBave3} relied on particular cases of exact LTB solutions. These articles apparently signal that a positive effective acceleration can occur under certain conditions (negative curvature and low density). In the present article and its continuation (part II) we enhance and complement this existing literature by looking at this theoretical issue in in full general analytic form, without specializing for particular LTB models, and without perturbations or approximations. Specifically: we examine in detail sufficient conditions for a non--negative back--reaction term, $\\QQ\\geq 0$, which become necessary conditions for the existence of a positive effective  acceleration $A_{{\\rm{eff}}}\\geq 0$ (applied in part II to those cases previously found in this paper to comply with $\\QQ\\geq 0$). We have summarized these results in Table 1, for the benefit of those readers who wish to see them before going into the technical detail. The relation between back--reaction, spatial curvature and a binding energy functional, which can be defined in a covariant manner~\\cite{sussQL}, appears in the process of discussing the conditions for $\\QQ\\geq 0$. This relation, summarized  in section 14,  provides a solid, though model--dependent, support for the arguments suggested by Wiltshire~\\cite{wiltshire2} on back--reaction, quasi--local energy and binding energy.  We summarize now the section by section contents of the article. We present in section 2 the basic parameters and properties of LTB models. In section 3 we examine the natural and covariant time slicing afforded by the normal geodesic 4--velocity, discussing geometric properties of the hypersurfaces $\\T(t)$ orthogonal to $u^a$, which provide the spatial slices in which the proper volume averages associated with Buchert's formalism will be defined and computed.  Section 4 provides a rigorous definition of the proper volume spatial averaging functional, as well as its associated local valued function, while in section 5 we introduce the quasi--local average functional and function, which will be needed in subsequent sections. The evolution equations for averaged scalars in Buchert's formalism, including the back--reaction term $\\QQ$ and its associated effective acceleration $A_{{\\rm{eff}}}$, are given in section 6, while in section 7 we present the generic sufficient conditions for $\\QQ\\geq 0$. We discuss the  conditions for the existence of zeros of radial gradients of scalars  (``turning values'' or TV's) in section 8. The sufficient conditions for $\\QQ\\geq 0$ are then applied to domains in various LTB configurations: ``parabolic'' (section 9), ``hyperbolic'' (section 10), ``elliptic'' with open and closed topology (sections 11 and 12) and mixed configurations in which an inner elliptic region (enclosing a center) is surrounded by a hyperbolic or parabolic exterior (section 13). In section 14 we provide an overview and final discussion on the interrelation between back reaction, spatial curvature and binding energy. The discussion on properties of LTB models given in sections 2 and 3 is complemented with a summary of regularity conditions (Appendix A), the relation between radial coordinate and proper radial length  (Appendix B) and the fluid flow evolution equations~\\cite{ellisbruni89,BDE,1plus3,LRS} (Appendix C). Appendices D and E provide the detailed proof of Propositions 7 and 8.  The most relevant result in sections 9--13 is the fact that averaging domains always exist in which $\\QQ\\geq 0$ holds for all regular LTB models whose radial asymptotic range has, either negative spatial curvature ($\\RR<0$ hyperbolic), or radially decaying positive curvature ($\\RR>0$ elliptic models with open $\\mathbb{R}^3$ topology) with sufficiently large gradients $\\RR'<0$. In a follow up paper (part II) we explore conditions for which $A_{{\\rm{eff}}}\\geq 0$ can hold for those configurations in which $\\QQ\\geq 0$ holds, leading to a robust analytic proof on the compatibility between specific LTB models and the existence of a positive effective acceleration in the context of Buchert's formalism. ",
        "conclusions": "We have provided in this article a fully comprehensive discussion of the back--reaction term $\\QQ$ for LTB models in the context of Buchert's averaging formalism. All possible regular configurations admitting at least a symmetry center were examined in ten propositions that where rigorously proven for parabolic, hyperbolic and elliptic domains, regions and models, and for turning values (TV') of $\\Theta,\\,\\FF$ and $R$ along hypersurfaces $\\T(t)$. A summary of the conditions for $\\QQ\\geq 0$ is provided in Table 1.  The form of $\\C$ for all domains $\\eta[r]$ leads to a determined sign of $\\QQ$ for parabolic ($\\QQ=0$ Proposition 1) and for hyperbolic ($\\QQ\\geq 0$) or elliptic ($\\QQ\\leq 0$) models without TV's (Propositions 2 and 4). In all other configurations, the sufficient conditions for the occurrence of $\\QQ\\geq 0$ depend on the selected domain $\\eta[r]$ and on the existence of TV's of $\\Theta$ or $\\FF$. These conditions are not too restrictive for domains and models (hyperbolic or elliptic) whose hypersurfaces $\\T(t)$ have topology $\\mathbb{R}^3$ and have a full radial asymptotic range (Propositions 3, 5--7 and 9--10). For elliptic models whose $\\T(t)$ have spherical topology (TV of $R$, Proposition 8), these conditions are much more restrictive.  From the proofs of the Propositions it is evident that the profiles of $\\FF,\\,\\FF_p$ and their radial gradients determine the conditions for $\\C\\geq 0$.  However, as discussed in several parts of the article (for example, see equations (\\ref{KTVa})--(\\ref{KTVb}) in section 8.2, as well as sections 10 and 13), these radial profiles are themselves closely related to the spatial curvature $\\RR$, its average quasi--local function $\\RR_q$ and their gradients $\\RR',\\,\\RR_q'$. Hence, statements about configurations favoring $\\QQ\\geq 0$ can also be given as statements on $\\RR$ in specific domains. Thus, fulfillment of $\\QQ\\geq 0$ (and possibly $A_{\\rm{eff}}\\geq 0$) is compatible with the following domains % \\begin{itemize} % \\item $\\RR\\leq 0$ for all $r$ (``pure'' hyperbolic models of section 10). \\item $\\RR\\leq 0$ for large $r$, even if these domains contain elliptic inner regions with $\\RR>0$ (hyperbolic models enclosing an elliptic region, section 13). \\item $\\RR\\geq 0$, with sufficiently large gradients $\\RR'<0$ and $\\RR\\to 0$ asymptotically (elliptic domains with open topology and a TV of $\\FF$, section 11) % \\end{itemize} % This connection between the sign of $\\QQ$ and the sign and profile of the spatial curvature (specially negative curvature) fits very well with the dynamical importance of this curvature due to the presence of large cosmic voids in the pattern of cosmic large scale structure~\\cite{buch08}.  However, $\\FF$ is an invariant scalar in LTB models~\\cite{sussQL} and can also be associated with a covariant characterization (valid for spherical symmetry) of ``binding energy'', which has been proposed in \\cite{sussQL,MSQLM,hayward1,hayward2,szab} by means of the comparison between the quasi--local or ``effective'' mass--energy function $\\MM_q$ (the function $M$ in (\\ref{fieldeq1})--(\\ref{fieldeq2})) and the proper mass--energy $\\MM_p$  function: % \\bse\\ba 2\\MM_p(r) =\\frac{2G}{c^2}\\int_{\\DD[r]}{\\rho\\,\\dd\\VV_p}=\\kappa\\int_0^r{\\rho\\,\\FF^{-1}\\,R^2\\,R'\\,\\dd x},\\\\ 2\\MM_q(r) = 2M(r) =\\frac{2G}{c^2}\\int_{\\DD[r]}{\\rho\\,\\FF\\dd\\VV_p}=\\kappa\\int_0^r{\\rho\\,R^2\\,R'\\,\\dd x},\\ea\\ese % where $\\kappa=8\\pi G/c^2$. A binding energy integral functional can be defined as % \\ba\\fl\\B[r] &=& 2\\left[\\MM_q-\\MM_p\\right]=\\frac{2G}{c^2}\\int_{\\DD[r]}{\\rho\\,(\\FF-1)\\,\\dd\\VV_p}=\\kappa\\int_0^{r}{\\rho\\,\\frac{(\\FF-1)\\,R^2\\,R'}{\\FF}\\,\\dd x}\\nonumber\\\\ \\fl &=& \\kappa\\left\\langle \\left(\\FF-1\\right)\\,\\rho\\right\\rangle\\, \\VV,\\label{Bdef}\\ea % which for every domain is a conserved quantity along the 4--velocity flow: $\\dot\\B[r]=0$. Bearing in mind that the radial profiles of $\\FF$ and $\\FF_p$ are qualitatively analogous and that $\\B[r]$ in (\\ref{Bdef}) is not a local but an integral quantity associated to a whole domain $\\eta[r]$, all statements about a given asymptotic behavior of $\\FF_p$ can be given as statements on the behavior and profiles of this binding energy in different domains. Therefore, fulfillment of condition $\\QQ\\geq 0$ can be associated with domains in which % \\begin{itemize} % \\item $\\B[r]\\geq 0$ for all $r$ (``pure'' hyperbolic models of section 10). \\item $\\B[r]\\geq $ for large $r$, even if these domains contain elliptic inner regions with $\\B[r]<0$ (hyperbolic models enclosing an elliptic region, section 13). \\item $\\B[r]\\leq 0$, where the gradient $\\B'<0$ becomes less negative for large $r$ and $\\B\\to 0$ asymptotically (elliptic domains with open topology and a TV of $\\FF$, section 11) % \\end{itemize} % This connection between binding energy and back--reaction has been highlighted by Wiltshire~\\cite{wiltshire2}.  It is important to notice that, regardless of the TV's of $\\Theta$ or $\\FF$, or changes of signs of spatial curvature or binding energy and other local complexities of the radial profiles, the conditions for positive $\\C$ (sufficient for $\\QQ\\geq 0$) are basically determined by the asymptotic radial behavior of the incumbent scalars. The dominance of the asymptotic behavior is particularly evident in the case of the elliptic region with a hyperbolic exterior discussed in section 13 and illustrated by figures \\ref{fig7} and \\ref{fig8}: a local feature (like the elliptic region or a TV of $\\Theta$) has no effect on the occurrence of $\\C\\geq 0$, as long as we select domains that lie sufficiently far away from the center. This is so, not only in the external hyperbolic region, but even in the transition zone still inside the elliptic region (see figure \\ref{fig8}). The fact that the asymptotic behavior is the key factor is expected, as back--reaction is a non--local effect, and so local features can always be ``coarse grained'', even if they exhibit critical conditions like a collapsing singularity in the elliptic region.  The Propositions that we have proved provide information on the existence of a positive $\\QQ$, but not on how large this term can be.  It is possible to argue qualitatively that, while  back--reaction can be positive in elliptic and hyperbolic domains, the latter are likely to yield a larger value. This follows from comparing the integrands of $\\Phi$ and $\\Psi$ in (\\ref{CC2})--(\\ref{psi_def}) for each type of domain. Assuming the gradient $\\Theta'$ to be of similar magnitude in each case, the function $\\varphi$ in (\\ref{phi_def}) is likely to be a small quantity in most of the asymptotic radial range of an elliptic domain compatible with $\\C>0$. If we look at the profile of $\\FF_p$ displayed by figure \\ref{fig4}: as $r/y\\gg 1$, we have $\\FF\\to 1$ but $\\FF<1$ for all $x$. Since $\\FF_p<\\FF$, then $\\FF_p\\to 1$ (asymptotically small curvature). But then, $\\FF_p(x)\\approx\\FF_p(r)\\approx 1$, which makes $\\varphi$ in (\\ref{phi_def}) close to zero for a long asymptotic radial range (notice also that $\\FF$ and $\\FF'$ are strongly constrained by (\\ref{FFell1})). In comparison, $\\FF_p$ is not restricted to remain below $\\FF_p=1$ in the asymptotic range of hyperbolic domains or regions (see figures \\ref{fig7} and \\ref{fig8}), and so this function can reach much larger values than in elliptic regions, which makes the integrand of $\\Phi$ much less restricted.  The magnitude of the back--reaction term $\\QQ$ is also related to the magnitude of the radial gradients of various scalars. Using (\\ref{Vprime}), (\\ref{prop2}), (\\ref{Vqprime}), (\\ref{propq2}), (\\ref{CCdef}) we can write $\\C$ as % \\begin{equation}\\fl \\C = \\left(\\frac{R}{3R'}\\right)^2\\,\\left[\\left(\\Theta_p'\\frac{\\FF}{\\FF_p}\\right)^2-\\Theta_q'{}^2\\right]=\\left(\\frac{R}{3R'}\\right)^2\\,\\left[\\Theta_p'\\frac{\\FF}{\\FF_p}-\\Theta_q'\\right]\\left[\\Theta_p'\\frac{\\FF}{\\FF_p}+\\Theta_q'\\right].\\label{CCgrads}\\end{equation} % However, the gradients $\\Theta_p'$ and $\\Theta_q'$ can be related by means of (\\ref{Hamcon}) and (\\ref{ave_eveqs_ef3}) to the gradients of $\\rho_p,\\,\\rhoav,\\,\\RR_q$ and $\\RRav$, while gradients of $\\RR_q$ relate to gradients of $\\FF$, then it is evident that $\\C$ will be larger in radial ranges where these gradients are significant ({\\it{i.e.}} in regions where inhomogeneity is not negligible). As a consequence of (\\ref{CCgrads}) and the qualitative arguments in the previous paragraph, it seems that hyperbolic models that allow for larger gradients of $\\FF$ and $\\FF_p$ (in a wider radial range) are much more likely to provide a back--reaction $\\QQ$ that is sufficiently large for $\\ACal>0$ in (\\ref{efeacc}) to hold for density profiles that are not too restrictive, though  even if the resulting $\\C>0$ could be small in elliptic models an effective acceleration could still occur if $\\rho$ is sufficiently small.  Since $\\Theta_p$ (or $\\Thetaav$) and $\\Theta_q'$ in (\\ref{Hamcon}) and (\\ref{ave_eveqs_ef3}) play the role of ``kinetic energy'' terms, the relation between $\\C$  and the gradients $\\Theta_p'$ and $\\Theta_q'$ given by (\\ref{CCgrads}) provides theoretical and contextual support to Wiltshire's interpretation of back--reaction~\\cite{wiltshire2}. Given the LTB configurations which we have proven to be compatible with $\\QQ\\geq 0$, we examine in a follow up paper (part II) the magnitude and radial variation of the back--reaction term, as well as boundary conditions that are needed to find the conditions for the fulfillment of (\\ref{efeacc}) and/or (\\ref{efeacc2}), which lead to the existence of a positive effective acceleration that could mimic the effects of dark energy.  Another line for future work is to extend the results presented here and in part II to the more general Szekeres models~\\cite{kras,bolejko1}. These tasks are being undertaken in separate continuing articles.    %\\setcounter{section}{1} \\appendix"
    },
    "0807/0807.2567_arXiv.txt": {
        "abstract": "s{Sources of gravitational waves (GW) and emitters of high energy (HE) neutrinos both involve compact objects and matter moving at relativistic speeds. Coincidences between \\vo~and \\anta~would give a unique insight on the physics of the most powerful objects in the Universe. The feasibility of such GW/HE $\\nu$ coincidences is analysed. }  \\noindent The forthcoming years should be very exciting both in GW astronomy and HE $\\nu$ astronomy. The \\vo~interferometer~\\cite{virgo}, currently closed down for upgrade, should be taking data with an improved injection system in 2009, whereas the \\anta~collaboration has completed the deployment and connections of its 12 lines~\\cite{antares} early in June 2008, starting the operation of the first {\\it undersea} neutrino telescope. The two \\lo~interferometers (ITF) are in operation~\\cite{ligo}, and {\\sc IceCube} is deploying its 1~km$^3$ neutrino telescope in the ice of the South Pole, already having 40 lines in data taking mode \\cite{icecube}. Both GW sources and HE $\\nu$ emitters involve compact objects and matter moving at relativistic speeds. As a result, time coincidences between \\vo~and \\anta~can be envisaged. Some astrophysical objects, invisible in electromagnetic channels, may be observable only {\\it via} their GW and HE $\\nu$ emissions. Finally, in many quantum gravity (QG) models~\\cite{qg}, the propagation velocity of a particle depends on the energy: measuring a non-zero time delay between GW bursts and HE $\\nu$ transients would allow to probe QG effects. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0807/0807.0879_arXiv.txt": {
        "abstract": "A claim for evidence of dark matter interactions in the DAMA experiment has been recently reinforced. We employ a new type of germanium detector to conclusively rule out a standard isothermal galactic halo of Weakly Interacting Massive Particles (WIMPs) as the explanation for the annual modulation effect leading to the claim. Bounds are similarly imposed on a suggestion that dark pseudoscalars might lead to the effect. We describe the sensitivity to light dark matter particles achievable with our device, in particular to Next-to-Minimal Supersymmetric Model candidates. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0807/0807.1728_arXiv.txt": {
        "abstract": "{}{Molecular emission lines originating in the nuclei of luminous infra-red galaxies are used to determine the physical properties of the nuclear ISM in these systems. }{A large observational database of molecular emission lines is compared with model predictions that include heating by UV and X-ray radiation, mechanical heating, and the effects of cosmic rays.}{The observed line ratios and model predictions imply a separation of the observed systems into three groups: XDRs, UV-dominated high-density ($n$$\\geq$$10^5$cm$^{-3}$) PDRs, and lower-density ($n$$=$$10^{4.5}$cm$^{-3}$) PDRs that are dominated by mechanical feedback.}{The division of the two types of PDRs follows naturally from the evolution of the star formation cycle of these sources, which evolves from deeply embedded young stars, resulting in high-density ($n$$\\geq$$10^5$cm$^{-3}$) PDRs, to a stage where the gas density has decreased ($n$$=$$10^{4.5}$cm$^{-3}$) and mechanical feedback from supernova shocks dominates the heating budget.} ",
        "introduction": "\\label{sec:introduction} Dense ($n$$>$$10^4{\\rm cm}^{-3}$) molecular gas plays an important role in the physics of (Ultra-) Luminous Infrared Galaxies. It provides the materials from which stars are formed and fuels possible active galactic nuclei.  The radiation originating in molecular gas in the nuclei of these galaxies provides information about the physical properties of the nuclear environment, such as the gas density and chemical composition, and the dominant type of radiation field. There are two important types of radiation: UV radiation (6-13.6\\,eV) in star-forming regions, which generates Photon Dominated Regions \\citep[PDRs; e.g.][]{1999RvMP...71..173H}; and X-rays (1-100\\,keV) emanating from accreting black holes, which produce in X-ray Dominated Regions \\citep[XDRs; e.g.][]{1996A&A...306L..21L, 1996ApJ...466..561M}.  Analysis of the state of molecular gas in a significant number of galaxies allows systematic effects to be identified, which can provide insight into the processes influencing the (gas in the) nuclei of these galaxies: the star-formation rate, the evolutionary stage of the system, and feedback processes.  \\cite{2008A&A...477..747B} presented data for the CO, HCN, HNC, \\HCOP, CN, and CS line emission of 37 infrared luminous galaxies and 80 additional sources taken from the literature \\cite[e.g.,][]{NguyenEA1992, SolomonDR1992, AaltoPHC2002, GaoS2004a, GraciaCarpioGPC2006}. They suggested that the sources can be divided into PDRs and XDRs, based on the ratios of emission lines.  We expand on their analysis and explore in more detail the different physical and chemical processes responsible for the diversity in line ratios found in the \\cite{2008A&A...477..747B} data, by comparing their data with the predictions of PDR and XDR chemistry models first presented by \\cite{MeijerinkS2005}.    \\begin{figure*}[htbp] \\sidecaption \\includegraphics[width=12cm]{0327fig1.eps} \\caption{Diagnostic diagrams using integrated line ratios of HCN, HNC, and \\HCOP\\ versus each other. {\\bf a)}~Integrated \\HCOP/HCN versus HNC/\\HCOP\\ ratios, {\\bf b)}~Integrated \\HCOP/HCN versus HNC/HCN ratios, and {\\bf c)}~Integrated HNC/HCN versus HNC/\\HCOP\\ ratios. The symbols X and P and the dotted lines mark the regions of the line ratios of XDRs and PDRs. Squares indicate reliable values and triangles upper and lower limits.  In addition, filled symbols indicate OH or H$_2$CO megamasers. Colored lines indicate results for PDR and XDR modeling, as function of column density (see sect.~\\ref{sec:pdrs-xdrs}), where the dot denotes the highest column. See Table~\\ref{tab:shock-params} and legend for details about the models; n denotes the density in cm$^{-3}$, and Fuv and Fx, respectively, the UV and X-ray flux in erg s$^{-1}$ cm$^{-2}$.} \\label{3luik-no-shocks} \\end{figure*} ",
        "conclusions": "\\label{sec:conclusions} Although PDR, XDR or CR chemistry models can explain the general behavior of the HCN, HNC and \\hcop\\ line ratios  observed in nearby star-forming galaxies \\citep{2008A&A...477..747B}, they cannot explain the low HNC/HCN line ratios observed in a large number of systems. The addition of heating by mechanical feedback provides a solution to this problem.  Our simulations in Figs.~\\ref{3luik-no-shocks} and \\ref{3luik-shocks} show that the observed sources can be split into three main groups, based on their HNC/HCN line ratios. First, there are a few sources that exhibit XDR chemistry (log(HNC/HCN)$>$$0$).  The other (PDR) sources can be divided into two groups: a small group (composed mostly of the OH-MM and ULIRG systems) of sources that can be modeled with standard PDR chemistry and high ($n$$\\geq$$10^5$cm$^{-3}$) densities and a larger group of sources with lower densities ($n$$\\sim$$10^{4.5}$cm$^{-3}$) that are heated by mechanical feedback.  All evidence leads to the natural conclusion that the division between the PDR sources is a result of the evolution of the star-formation cycle of a galaxy, during which the dominant heating process changes and density of the gas diminishes.  The high-density ($n$$\\geq$$10^5$cm$^{-3}$) PDRs (with HNC/HCN ratios around unity and weak \\hcop\\ lines) represent the early stage of star-formation where the stars are forming in the densest clouds and the stellar UV radiation dominates the chemistry in those clouds. During this period most of the gas will be converted into stars and the increasing amount of UV radiation of the stars will gradually penetrate a larger volume of the (lower-density) gas.  The resulting lower average density leads to an increase in \\hcop\\ with respect to HCN and HNC. After about 10\\,Myr, the first massive stars become SNe and their shocks begin to dominate the heating budget, leading to lower HNC/HCN line ratios. These two effects alter the properties of the galaxies to those of a second group consisting of mechanical-feedback-dominated lower-density PDRs. Because of the short duty-cycle of the first stage compared to the second, most sources in our sample are observed to be in the later stage of their evolution.  It therefore appears that mechanical feedback plays a crucial role in the time dependent physical state of the dense star-forming molecular gas in luminous infrared galaxies."
    },
    "0807/0807.4611_arXiv.txt": {
        "abstract": "The Phase-Referenced Imaging and Micro-arcsecond Astrometry (PRIMA) facility is scheduled for installation in the Very Large Telescope Interferometer observatory in Paranal, Chile, in the second half of 2008. Its goal is to provide astrometric accuracy in the micro-arcsecond range. High precision astrometry can be applied to explore the dynamics of the dense stellar cluster. Especially models for the formation of stars near super massive black holes or the fast transfer of short-lived massive stars into the innermost parsec of our galaxy can be tested. By measuring the orbits of stars close to the massive black hole one can probe deviations from a Keplerian motion. Such deviations could be due to a swarm of dark, stellar mass objects that perturb the point mass solution. At the same time the orbits are affected by relativistic corrections which thus can be tested. The ultimate goal is to test the effects of general relativity in the strong gravitational field. The latter can be probed with the near infrared flares of SgrA* which are most likely due to accretion phenomena onto the black hole. We study the expected performance of PRIMA for astrometric measurements in the Galactic Center based on laboratory measurements and discuss possible observing strategies. ",
        "introduction": "\\label{sec:intro}  %  During the past three to four decades sub-milliarcsecond radio astrometry with intercontinental long baseline interferometry (VLBI) led to the discoveries of super-luminal motions in relativistic radio jets of distant quasars and of molecular gas outflows in star forming regions. VLBI also yielded precision measurements of the dynamics of a warped accretion disk around the supermassive black hole (SMBH) in the external galaxy NGC 4258 and provided primary distance measurements throughout the Milky Way and reaching several external galaxies. Hipparcos astrometry of stars in the solar neighborhood revolutionized our knowledge of the dynamics of the Milky Way. Speckle and adaptive optics, near-infrared (NIR) astrometry of stars with the European Southern Observatory (ESO) New Technology Telescope (NTT) and Very Large Telescope (VLT) demonstrated that SgrA*, the compact radio source at the center of the Milky Way, must be a SMBH of about 4 million solar masses, beyond any reasonable doubt. NIR and optical interferometry with the Very Large Telescope Interferometer (VLTI) will open a new era of high resolution, narrow angle precision astrometry. It will allow phase-referenced imaging of faint sources at milliarcsecond (mas) resolution, and $10 - 100$ microarcsecond ($\\mathrm{\\mu as}$) precision astrometry, all at the exquisite sensitivity provided by the large collecting area of the VLT (2~m diameter Auxiliary Telescopes, ATs, and 8~m diameter Unit telescopes, UTs).  In this paper we will discuss some of the exciting possibilities for dynamical measurements that PRIMA will offer in the Galactic Center research. Studies of stars and gas in the immediate vicinity of the event horizon of the massive SMBH in the Galactic Center (GC), with the ultimate goal of testing General Relativity (GR) in its strong field limit will be viable in only a few years' time. ",
        "conclusions": "The Phase-Referenced Imaging and Micro-arcsecond Astrometry (PRIMA) facility at the VLTI will allow astrometric observations with a precision of up to 10~$\\mathrm{\\mu as}$. It will first be available for observations with the ATs (aimed for limiting magnitudes of $m_{K}$ of 8/15 for the reference/science object) and somewhat later for the UTs with limiting magnitude aims of $m_{K} = 10/18$. The Galactic Center star cluster contains a bright star, IRS~7, which will allow the use PRIMA already with the ATs.  With PRIMA one can determine stellar orbits of many of the early-type stars with time baselines of only a few years, much faster than with existing instrumentation like NACO. This will help to better understand the dynamics of the central star cluster of our Galaxy and to constrain models about star formation near supermassive black holes.  Using PRIMA together with the UTs, influence of General Relativity on the orbits of the S-stars in the central arcsecond may be observed. However, the needed integration time of 30 minutes may be too large to time-resolve the infrared flare events, which originate in the immediate vicinity of the supermassive black hole. This observation may first be done with the planned GRAVITY interferometer (Eisenhauer\\cite{Eisenhauer2005} et~al.(2005)), which will use all four UTs."
    },
    "0807/0807.5084_arXiv.txt": {
        "abstract": "We investigate effects of D-term contributions to the anomaly mediated supersymmetry breaking scenario. If we introduce an E$_6$ GUT motivated D-term, it is possible that slepton mass squared is positive at weak scale even if there are no other D-term contributions. Moreover, as a consequence of additional D-term contributions to scalar masses, we obtain various soft supersymmetry breaking mass spectra, which are different from those obtained in the conventional anomaly mediation scenario. Then there would be a distinct signature of this scenario at the LHC. For example, in several cases, there exist mass splittings between the right-handed sfermion and the left-handed sfermion. We also discuss some characteristic features of the neutralino dark matter in this model. ",
        "introduction": "Supersymmetry (SUSY) extension is one of the most promising way to solve the gauge hierarchy problem in the standard model (SM) \\cite{SUSY}. Since any superpartners have not been observed in current experiments, SUSY should be broken at low energies. Furthermore, soft SUSY breaking terms are severely constrained to be almost flavor blind and CP invariant. Thus, the SUSY breaking has to be mediated to the visible sector not to induce too large CP and flavor violation effects. Some mechanisms to achieve such SUSY breaking mediation have been proposed \\cite{Luty:2005sn}.  The anomaly mediated supersymmetry breaking (AMSB) scenario \\cite{AMSB1, AMSB2, AMSB3} is one of the most attractive scenario due to its flavor-blindness and ultraviolet (UV) insensitivity for the resultant soft SUSY breaking terms. The pattern of SUSY breaking does not depend at all on physics at higher energy scales. On the eve of the Large Hadron Collider (LHC) operation at CERN, which start this year, there are several studies in the aspects of collider physics to discriminate the AMSB scenario from the other SUSY breaking mediation \\cite{Barr:2002ex, Datta:2002vy, Asai:2007sw}. Despite the appeal of the AMSB, the original version of the AMSB is excluded. % The slepton squared masses become negative at the weak scale, and hence the theory would break U(1)$_{\\rm em}$. There have been many attempts to solve this problem by incorporating additional positive contributions to slepton squared masses at tree level \\cite{AMSB1, tree1, tree2, tree3, tree4, tree5, tree6, Hodgson:2007kq} or at quantum level \\cite{PR, quantum}.  Among them, adding some D-terms such as ${\\rm U}(1)_Y$ and ${\\rm U}(1)_{\\rm B-L}$ may be the most interesting possibility \\cite{tree5}, because U(1)$_Y$ and U(1)$_{\\rm B-L}$ as local symmetries are included in SO(10) grand unified theory (GUT) \\footnote{For the group theoretical aspects of {\\rm SO}(10), see for example, \\cite{FIKMO}.}. The idea of GUT's bears several profound features. Perhaps the most obvious one is that GUT's have the potential to unify the diverse set of particle representations and parameters found in the SM into a single, comprehensive, and hopefully predictive framework. For example, by unifying all U(1) generators within a non-Abelian theory, GUT's would provide an explanation for the quantization of electric charge. By combining GUT's with SUSY, we hope to unify the attractive features of GUT's simultaneously with those of SUSY into a single theory, SUSY GUT's \\cite{Sakai:1981gr}. The apparent gauge couplings unification of the minimal supersymmetric standard model (MSSM) is strong circumstantial evidence in favor of the emergence of a SUSY GUT near $M_{\\rm GUT} \\sim 10^{16}$ [GeV] \\cite{Amaldi:1991cn, unification}. While there are {\\it a priori} many choices for possible GUT gauge group, the list can be narrowed down by requiring groups of rank $\\geq 4$ that have complex representations. The smallest groups satisfying these requirements are SU(5), SU(6), SO(10), and E$_6$. Among these choices, SO(10) is particularly attractive \\cite{Fritzsch:1974nn}, because SO(10) is the smallest simple Lie group for which a single anomaly-free irreducible representation (namely the spinor ${\\bf 16}$ representation) can accommodate the entire SM fermion content of each generation. In addition, the recent progress in neutrino physics \\cite{Mohapatra:2005wg} makes SO(10) GUT's \\cite{so10} a favorite candidate for GUT's because it naturally incorporates the see-saw mechanism \\cite{see-saw} that can naturally explain the lightness of the light neutrino masses.  While it is interesting to add SO(10) motivated D-terms to the AMSB scenario, an extension of the grand unification group of SO(10) to E$_6$ is also worthy of attention \\cite{E6, Bando:1999km, tree3}. In the E$_6$ GUT, the fundamental representation ${\\bf 27}$ includes ${\\bf 10}$ as well as ${\\bf 16}$ of SO(10) in the following manner: \\bea {\\bf 27} &\\to& {\\bf 16}_{+1} + {\\bf 10}_{-2} + {\\bf 1}_{+4} \\nonumber\\\\ &\\to& [{\\bf 10}_{-1} + \\overline{\\bf 5}_{+3} + {\\bf 1}_{-5} ] + [{\\bf 5}_{+2} + \\overline{\\bf 5}_{-2} ] + {\\bf 1}_0 \\label{decom} \\eea under $E_6 \\supset {\\rm SO}(10) \\times {\\rm U}(1)_V,~{\\rm SO}(10) \\supset {\\rm SU}(5) \\times {\\rm U}(1)_W$. Hereafter we take the following linear combination of those two U(1) charges to match with the convention in \\cite{Martin}. \\be S = -(5V+W)/12,~~X=W/3 \\;. \\ee The symmetry breaking pattern of SO(10) down to the SM gauge group can be found in the literature \\cite{FIKMO}. A linear combination of U(1)$_X$ charge and U(1)$_{\\rm B-L}$ charge gives a correct SM hypercharge normalization under the following relation (see, for example, \\cite{FIKMO}). \\be Y = \\frac{1}{4} \\left[3 X + 5 (B-L) \\right] \\;. \\ee While there are several studies on the extra gauge boson, $Z^\\prime$, which is associated with an extra U(1) gauge symmetry, potentially arisen from E$_6$ GUT, the effect of the U(1)$^\\prime$ D-term has not been studied well in the literature. Hereafter, we label the U(1)$^\\prime$ gauge symmetry which is originated from E$_6$ GUT as U(1)$_S$. In this paper, we specifically consider the effect of such an extra U(1)$_S$ D-term, in addition to the ${\\rm U}(1)_Y$ and the ${\\rm U}(1)_{\\rm B-L}$ D-terms to the soft SUSY breaking mass terms. Since the effects of this extra U(1)$_S$ D-term to the sfermion mass squared gives common signature contributions, the tachyonic slepton problem can simply be solved in this framework.  This paper is organized as follows. In section 2, we begin with a brief review of AMSB and explain a model of D-term assisted anomaly mediation in E$_6$ GUT motivated models. In section 3, we show numerical results for the calculation of sparticle mass spectra and a set of allowed parameter spaces of $(D_Y, D_{\\rm B-L}, D_{\\rm S})$. After calculating all the sparticle masses by using {\\tt ISAJET 7.75} \\cite{ISAJET}, the relic density of the cold dark matter (CDM) in this scenario is also estimated by using {\\tt micrOMEGAs} \\cite{Belanger:2006is}. The last section is devoted to summary and discussions. ",
        "conclusions": "Anomaly mediation of supersymmetry breaking (AMSB) is very attractive because the resultant soft supersymmetry breaking parameters at a given energy scale are determined only by physics at that energy scale (UV insensitivity) and hence is highly predictive (only one parameter, $F_\\phi$). However, there is tachyonic slepton problem. It is known that adding some D-terms such as ${\\rm U}(1)_Y$ and ${\\rm U}(1)_{\\rm B-L}$ is the well-motivated solution to this problem.  In this paper, we have considered the effects of additional D-terms, which might be originated from the E$_6$ GUT models. We have evaluated the soft SUSY breaking terms and obtained various sparticle mass spectra for various input values of ($\\alpha_Y$, $\\alpha_{\\rm B-L}$, $\\alpha_{\\rm S}$), that are different from those obtained in the conventional anomaly mediation. It has been found that even if we add only one extra ${\\rm U}(1)_{\\rm S}$ D-term (without U(1)$_Y$ and U(1)$_{B-L}$ D-terms), it can become a solution to the tachyonic slepton problem, so it is more economical. Since there could have some amounts of mass splittings between the right-handed sfermion and the left-handed one, this scenario can have a very distinct signature at the LHC.  We have also evaluated the dark matter relic density in our scenario, and we have shown that there exist parameter space, which is consistent with the WMAP observational data. In the WMAP consistent region, the mass of the thermal relic DM is required to be heavy since the lightest neutralino is Wino-like. However, if one consider the non-thermal production of the neutralino DM, there is a possibility to keep the neutralino DM light."
    },
    "0807/0807.2492_arXiv.txt": {
        "abstract": "By adopting the chemical evolution model of the Milky Way disk, we have studied the star formation and chemical evolution history for M31 galaxy disk. We mainly concentrated on the global properties of the M31 disk. The model has been scaled to the related disk parameters, mainly the disk scale length and total disk baryon mass of M31, which we have adopted to be $r_d$=5.5kpc and $M_{tot} = 7\\times 10^{10}$ \\ms. It is found that, when the classical Kennicutt star formation law was applied, the obtained radial profiles of gas surface density and star formation rate (SFR) have great difference from the observed results in M31 disk. Then we have adopted modified SFR as we did for the Milky Way galaxy, that is the SFR is radial dependent. Detailed calculations show that by adjusting the star formation efficiency, it is possible to get reasonable gas and abundance profiles, but the total disk SFR is a factor of 2-3 higher than that estimated from observations. And also the predicted SFR radial profile is also much higher than what GALEX observed in the outer part. Possible reasons could be that the M31 disk has been interacted by other factors which seriously altered the star formation history, or the observed SFR is underestimated due to inappropriate extinction correction. ",
        "introduction": "As one of the three disk galaxies in the Local Group, Andromeda galaxy (M31, or NGC224) provides an unique opportunity for testing theory of galaxy formation and evolution (Widrow et al. 2003; Renda et al. 2005; Widrow \\& Dubinski 2005; Brown et al. 2006; Tamm et al. 2007; Tempel et al. 2007). In order to understand how star formation has influenced the evolution of M31 disk, Williams (2003a, 2003b) has measured star formation history in several regions of the M31 disk from the KPNO/CTIO Local Group Survey. It is found that the total mean star formation rate for the disk is about 1 \\ms yr$^{-1}$. By deep HST photometry, Bellazzini et al. (2003) have studied the metallicity distributions and star formation history in many locations of the disk (see also Ferguson \\& Johnson 2001; Sarajedini \\& van Duyne 2001; Olsen et al. 2006).  However, some basic properties of the Milky Way and M31 galaxies are still uncertain despite years of detailed observations and analysis. For example, the total masses of two galaxies could be varied by a factor of few based on the current published literatures in which the mass of galaxy is mostly derive from kinematic analysis of the galaxy and its neighbors. And there also exists disagreement over which of the two galaxies should be more massive (Evans \\& Wilkinson 2000; Gottesman et al. 2002).  In this work, we apply an infall model, which is similar to what has been used for the Milky Way galaxy, to M31 galaxy. We will mainly concentrate on the global properties of the M31 disk based on the currently available data. We would like to know whether the similar model could explain both the Milky Way and M31 disk observations. If not, what is the main difference in the formation history of those two disks. We would like to demonstrate and hope to understand the similarities and differences of star formation and chemical evolution history between M31 and the Milky Way disks. ",
        "conclusions": "In this work, we have studied the chemical evolution for the M31 disk based on the model similar to the Milky Way disk (Boissier \\& Prantzos 1999, 2000). The main observational constrains for M31 disk are the profiles of gas and star formation rate, the abundance gradients. The star formation profile of the M31 disk is from recent UV data of GALEX (Boissier et al. 2007).  We tried to use a unified chemical evolution model to predict the disk properties of both M31 and Milky Way by using the modified KS SFR law. When the results are scaled to disk scale lengths, we find that the M31 need a higher star formation efficiency than the Milky Way disk. And the model always predicts a higher SFR in the outer part of the M31 disk. It is also found that the adopted different infall time scales are not very sensitive to the radial profiles of gas, SFR and abundances. The inside-out disk formation is mainly controlled by the radial dependence of star formation rate instead of the time scale."
    },
    "0807/0807.0497_arXiv.txt": {
        "abstract": "AstraLux is the Lucky Imaging camera for the Calar Alto 2.2-m telescope, based on an electron-multiplying high speed CCD. By selecting only the best 1-10\\% of several thousand short exposure frames, AstraLux provides nearly diffraction limited imaging capabilities in the SDSS {\\em i'} and {\\em z'}\\ filters over a field of view of 24$\\times$24 arcseconds. \\\\ By choosing commercially available components wherever possible, the instrument could be built in short time and at comparably low cost. We present the instrument design, the data reduction pipeline, and summarise the performance and characteristics\\footnote{Based on observations collected at the Centro Astron\\'omico Hispano Alem\\'an (CAHA) at Calar Alto, operated jointly by the Max-Planck Institut f\\\"ur Astronomie and the Instituto de Astrof\\'isica de Andaluc\\'ia (CSIC)}. ",
        "introduction": "The recovery of the full theoretical resolution of large ground-based optical telescopes in the presence of atmospheric turbulence has been a major goal of technological developments in the field of astronomical instrumentation in the past decades. The de-facto standard nowadays is adaptive optics, the active real-time correction of the incoming distorted wavefronts by means of deformable mirrors. While diffraction limited image quality can now be routinely achieved in the near infrared at large telescopes, this is only possible with considerable technical effort.  Even before the age of adaptive optics, speckle interferometric techniques  gave access to the diffraction limit of large telescopes by applying auto- or cross-correlation methods in Fourier space to large sets of atmospherically degraded short-time exposures\\cite{Labeyrie:1970,Labeyrie:1974,Knox:1973,Knox:1974,Lohmann:1983}.  An even more simple approach has been employed especially by amateur astronomers in order to produce sharp images of the moon or bright planets: instead of using all images of a large set of short-time exposures, only those showing little image distortion due to the variable strength of atmospheric turbulence are combined to a high-resolution result. Like in the speckle interferometric approach, this requires only a detector with fast readout capability and moderate computational effort. The application of this technique to fainter astronomical targets, e.g. low-mass double stars, was limited by the readout noise of the available detectors at visible wavelengths until a few years ago. At typical exposure times of few ten milliseconds -- necessary to ``freeze'' the effects of atmospheric turbulence -- photon noise limited detectors are obligatory to use this so-called ``Lucky Imaging'' technique on targets fainter than $\\approx$10\\,mag. Since the image quality of each single frame has to be determined, e.g. by measuring the Strehl ratio of a reference star, the readout noise of the detector sets stringent limits on the minimum brightness of this reference.  The advent of electron multiplying CCDs (EMCCD)\\cite{Jerram:2001,Hynecek:2002,Mackay:2001} led to considerable interest in the Lucky Imaging technique on the part of professional astronomers. First experiments at the Nordic Optical Telescope (NOT) with LuckyCam\\cite{Baldwin:2001,Tubbs:2002,Tubbs:2003,Law:2006} proved that Lucky Imaging is a very promising alternative to adaptive optics, allowing diffraction limited imaging at telescopes in the 2--3\\,m class at wavelengths below 1$\\mu$m. These results triggered the development of a similar instrument for the 2.2-m telescope at Calar Alto by our group. The Lucky Imaging camera AstraLux was built in less than 5 months thanks to the availability of most parts as off-the-shelf equipment.  AstraLux routinely reaches Strehl ratios of 15--20\\% in the SDSS {\\em i'} and {\\em z'} passbands and has become the standard tool for double star observations at Calar Alto. Small technical overheads, an easy to use software interface and the availability of pipeline results in near real-time allow the observation of up to 20 targets per hour.  In the following we describe the instrument design and data processing pipeline and summarise the instrument's performance and key characteristics. Interested readers looking for more comprehensive information are referred to the diploma thesis\\cite{Hormuth:2007a}\\footnote{Available online at \\texttt{http://www.mpia.de/ASTRALUX/}}, which covers AstraLux and its performance in full detail. ",
        "conclusions": "Within less than one year it was possible to design, build, and characterise a Lucky Imaging instrument for the Calar Alto 2.2-m telescope. Beyond evaluation of the observing technique and data reduction strategies, scientific observations commenced right in the first night at the telescope. As a common user instrument from 2007 on, AstraLux has become the standard tool for diffraction limited imaging at the Calar Alto observatory. It is currently mostly used for binarity surveys among stars and minor planets.  The encouraging results regarding both development speed as well as scientific output triggered the start of the project ``AstraLux Sur'', an almost identical copy of the Calar Alto version to be used as visitor instrument at ESO's New Technology Telescope (NTT) at La Silla, Chile. AstraLux Sur is currently on its way to Chile after a development time of less than 2 months and will have first light in mid-July 2008.  Apart from Lucky Imaging, AstraLux can be used as a high speed photometer with single photon detection capability. In combination with the custom GPS timing hardware MicroLux\\cite{Hormuth:2008} it has been successfully used for observations of the optical pulse profile of the Crab pulsar.  We are currently investigating possible improvements of the optical design, namely a telecentric magnification optics and a tunable atmospheric dispersion corrector. However, the strength of the current system is its simple design, leading to short development times, low hardware costs, and high stability. \\\\ An ambitious project for the future would be the combination of AstraLux with a low-order adaptive optics system, currently developed for the Calar Alto 2.2-m telescope. This could enable observations at shorter wavelengths or increase the resulting Strehl ratio under seeing conditions normally unfavourable for Lucky Imaging. It has recently been proven at the 5-m Palomar telescope\\cite{Law:2008} that this approach actually works and can boost the performance of the Lucky Imaging technique."
    },
    "0807/0807.1435.txt": {
        "abstract": " ",
        "introduction": "We study the quantum dynamics of fermions in a homogeneous but temporally varying background field including collisions with a thermal background. These conditions are appropriate to model for example the particle production during preheating at the end of inflation or during cosmological phase tranistions~\\cite{brandenbeger,Prokopec_partnumber}, as well as baryogenesis during preheating~\\cite{shapo}, or coherent baryogenesis~\\cite{cohbaryog}. They are relevant also for neutrino oscillations in the early universe~\\cite{EKT}, or just for generic studies of thermalisation of quantum systems~\\cite{generic_therm}. The formalism we will be developing here can also be modified for a treatment of static problems with planar symmetry~\\cite{HKR1,HKRfuture}. In this form the resulting kinetic equations will be of interest for problems involving quantum reflection such as electroweak baryogenesis~\\cite{CKV,CJK}.  For related studies of quantum transport equations for electroweak baryogenesis starting from similar theoretical foundation see {\\em e.g.}~\\cite{KPSW,PSW}.  By use of the CTP-formalism~\\cite{Schwinger-Keldysh,CTP} we will set up quantum kinetic equations for the fermionic 2-point function $G^<$. We find that these equations admit a rich structure of spectral solutions including the expected mass-shell solutions for particle and antiparticle excitations, but also a new class of solutions living in $k_0=0$-shell in the phase space~\\cite{HKR1}. We interpret  the arbitrary weight functions on these shells as describing the out-of-equilibrium particle and antiparticle numbers (mass shells) and the quantum coherence between particles and antiparticles ($k_0=0$-shell). New coherence solutions are shown to be excluded from the spectral function by the spectral sum-rule, which indicates that they are not part of the kinematic phase space, although they necessarily occur in the dynamical function $G^<$.  These solutions are also eliminated from $G^<$ by the KMS-conditions in the thermal limit. As a consequence of the singular shell structure, an integration procedure is needed in order to define a physical density matrix in terms of the original 2-point function. We show how this procedure necessarily involves specifying precisely the amount of information on the system, and derive an evolution equation for a density matrix relevant for a spatially homogenous case, including quantum coherence. We then introduce the interactions and show that the spectral structure for the phase space including the coherence shell solutions survives in the quasiparticle limit with the interactions.  We set up dynamical equations for the physical density matrix including the interactions and compute explicit expressions for the collision integrals in the case of a simple model Lagrangian describing decays and inverse decays.  The interaction terms are shown to contain the usual collision terms that push the mass-shell functions towards thermal limit, but also other collision terms that tend to bring the quantum coherence functions to zero. The loss of coherence does not happen instantly, like a collapse of a wave function, but smoothly over a characteristic time scale set by the strength of the interactions at the shell $k_0=0$, in close analogy to the damping of coherence in the case of neutrino oscillations~\\cite{EKT}.  In our approach we {\\em define} the fermionic number density in the same way as in thermal field theory, that is, as an dimensionless real-valued function living on the positive mass-shell of the spectral form 2-point correlator. With this definition we see that in thermal limit the number density is indeed the standard Fermi-Dirac distribution. We will show that our definition for the particle number agrees with that of ref.~\\cite{Prokopec_partnumber}, where it was derived using the operator methods and Bogolyubov transformation to diagonalize the fermionic Hamiltonian. We also derive expressions for the energy density and the pressure. The latter is shown to differ from the statistical pressure, but the statistical pressure is retrieved for any realistic measurement that averages out the quantum oscillations. As applications of our formalism we calculate particle number production during fermionic preheating including finite decay width for the heavy particles produced. We will see that the decoherence induced by the decays can have dramatic effect on the particle number evolution. We also show explicitly how an initially highly correlated out-of-equilibrium density matrix relaxes to a thermal equilibrium as a result of collisions.  This paper is organized as follows: In section \\ref{sec:formalism} we give a brief intoduction to Shwinger-Keldysh CTP-formalism in order to write down the fermionic Kadanoff-Baym (KB) transport equations for 2-point correlation functions. In section \\ref{sec:free} we study the noninteracting KB-(or Dirac) equation and find out the nontrivial phase-space shell structure including the new $k_0=0$-shell. In section \\ref{sec:weighted} we define a nonsingular physical density matrix, as a weighted integral over the initial singular 2-point function and show that its evolution is heavily dependent on the extrenous information we have on the system. The material in the sections \\ref{sec:free}-\\ref{sec:weighted} partly overlaps with the derivation in the companion paper~\\cite{HKR1}, but we include a shortened discussion here for completeness. In section \\ref{sec:physical} we will compute the particle number density, energy density and pressure in terms of the density matrix. In section \\ref{sec:interactions} we generalize our kinetic equations to include collisions. We study the particle production at the preheating and the decoherence phenomenon in sections \\ref{sec:preheating} and \\ref{sec:decoherence}, and finally section \\ref{sec:discussion} contains our conclusions.  % ",
        "conclusions": "\\label{sec:discussion} In this paper we have developed a quantum transport formalism, simple enough to be used in practical calculations, for treating the coherent evolution of fermionic system in the presence of decohering collisions. In our derivation we used the CTP-formalism for the quantum field theory in the out-of-equilibrium conditions. The key element in the derivation was the observation that in the mean field limit and in the quasiparticle approximation in an interacting theory, the 2-point function has a singular shell structure, that in addition to the usual mass-shell solutions contain new $k_0=0$-solutions, which carry the information about the quantum coherence between particles and antiparticles. Thus the actual 2-point function can be understood as a dynamical phase space weight function, or measure, with a number of unknown on-shell coefficient functions that carry the physical information about the system. These can be interpreted as the out-of-equlibrium particle numbers on mass shells and as functions measuring quantitatively the level of quantum coherence on $k_0=0$-shell. We showed that the new coherence solutions are excluded from the spectral function by the spectral sum-rule, and that they are also eliminated from $G^<$ by the KMS-condition in the thermal equlibrium limit. The fact that coherence solutions appear only in the dynamical function makes sense, because they describe correlation between mass-shell states, and no coherence should be present without interfering physical states (Solitary coherence functions can however, in some cases describe pairs of localized virtual states~\\cite{HKR1}.)  We then proceeded to show how a sensible physical density matrix can be defined by an integration procedure which involves specifying the amount of external information on the system. This information affects both the definition of the physical density matrix and its eventual equation of motion. Finally we derived an explicit dynamical equation of motion for a physical density matrix corresponding to a spatially homogenous system, where the 3-momentum and helicity are good quantum numbers.  We also computed the collision integrals in the thermal limit for a model interaction Lagrangian, showing in particular how different shells pick up different collision terms in the integration procedure, corresponding to the different external momentum configurations at different shells. Finally, we applied our formalism to the case of coherent particle production during the preheating stage of the inflation including a finite decay width to the produced heavy particles. We showed that including collisions can have dramatic effect on the evolution of the particle number, in particular when one includes a nontrivial evolution of the inflaton field. We stress that our method can be used to a quantitative analysis of the particle production during inflation if one models the inflaton oscillation and evolution numerically, and includes a network of transport equations also for the daughter states in the decay process. Finally we pointed out that the method is also suitable for a quantitative analysis of coherent baryogenesis and a computation of a baryon- or lepton number generation during decay of the preheating produced heavy states. In our last example we showed explicitly how an initially highly correlated out-of-equilibrium density matrix relaxes to a thermal equilibrium.  We believe that the current formalism will find applications in many different contexts, also outside the cosmology. In particular we are currently extending the formalism also to the case of relativistic bosonic fields and also to nonrelativistic systems. It is also interesting to see how the usual fermionic flavour oscillation pattern arises in the present context when the formalism is extended to the case with many mixing fermion flavours~\\cite{HKR1,HKRfuture}."
    },
    "0807/0807.0518_arXiv.txt": {
        "abstract": "We present a method we developed for the correction of the beam divergence in finite distance sources imaging through coded mask instruments. We discuss the defocusing artifacts induced by the finite distance showing two different approaches to remove such spurious effects. We applied our method to one-dimensional coded-mask systems, although it is also applicable to 2D systems as well. We provide a detailed mathematical description of the adopted method and of the systematics introduced in the reconstructed image (e. g. the fraction of source flux collected in the reconstructed peak counts). The accuracy of this method was tested by simulating point-like and extended sources at finite distance with the instrumental set-up of the SuperAGILE experiment, the one dimensional coded-mask X-ray imager on-board the AGILE mission. We obtained reconstructed images of good quality and high source location accuracy. Finally, we show the results obtained by applying this method to real data collected during the calibration campaign of SuperAGILE. Our method was demonstrated to be a powerful tool to investigate the imaging response of the experiment, particularly, the absorption due to the materials intercepting the line of sight of the instrument and the conversion between detector pixel and sky direction. ",
        "introduction": "``Coded mask aperture imaging technique'' is a spatial multiplexing technique to image X-ray sources over large fields of view, at X-ray energies where the absorption is dominated by the photoelectric effect, typically below 100-150 keV. Although the new multi-layer and long focal length telescopes are now starting being used to focus X-rays between 10 and 70-80 keV, this can be only done over a very narrow field of view (below 1 degree) and the coded mask technique is not surpassed yet for imaging large solid angles.  Coded aperture systems encode the sky radiation through a \\textit{mask} of opaque and transparent elements before detector records it. The illuminated mask pattern projects a ``shadowgram'' on the detection plane shifted with respect to the central position by an amount that depends on the direction towards the point source. The source position in the field of view is reconstructed at a level of accuracy depending on the angular resolution of the system and the number of counts. The experimental parameters of a coded mask system can be optimized to fit the scientific requirements. For example, the angular resolution can be optimized at design level by a proper choice of the size of mask elements and a mask-detector separation. On the other hand coded mask systems have the disadvantage of no direct imaging that implies the need of decoding procedures to reconstruct the sky image. Another disadvantage with respect to direct imaging is that the Poisson noise from any source in the sky is induced in every other position of the sky. We refer the reader to \\cite{izand94} for an extensive and general discussion of the design and properties of coded mask systems. Here we limit our discussion to the properties relevant to the algorithm we have developed.  The image of a source in the sky is obtained by correlating the image accumulated by detector with the proper decoding patterns. Let $D$ be the detector image of the sky, then D is given by the convolution of the true sky image $S$ with the mask pattern $M$ plus the background $B$:  \\begin{equation}D = S \\otimes M + B.\\end{equation}  It is possible to reconstruct the sky image $S'$ by cross correlating the detector image with the reconstruction matrix  R, which has to be opportunely defined:  \\begin{equation}S' = D\\odot R= S \\otimes (M \\odot R) + B \\odot R.\\end{equation}  The reconstruction is obtained when  $M\\odot R$ is the identity matrix. This can be achieved if the mask is based on a pseudo-random URA (Uniformly Redundant Array) pattern, built using cyclic difference sets for which the auto-correlation function is a $\\delta$ function (see \\cite{gupo76}). If the mask pattern presents an open fraction of 0.5 (the number of opaque elements equal to the number of the transparent ones) the reconstruction matrix is given by $R{ij}=1,-1$ for transparent and opaque elements, respectively. It is evident from Eq. (2) that this kind of matrix allows for the  background subtraction in the assumption of its uniformity. Generally speaking, if the complete mask pattern is projected on the detector then $S'=S$. For several mask-detector systems this is true only for a limited fraction of the field of view, while the remaining  sky is only partially coded due to the collimator shadow. In these regions of the reconstructed images a source induces false peaks (the so-called ``coding noise\"). The removal of these false peaks occurs by using an Iterative Removal of the Sources in the sky (IROS). We refer the reader to \\cite{izand92} for details about the removal of such effect from sky images.  \\indent Imaging source at finite distance (i. e. divergent beams) is more complicated because the size of the shadow in the detector plane is magnified depending on the distance of the source to the plane itself. In this paper we will call the effect of the beam divergence on the optical system mask-detector  as a ``defocusing'' effect, although this term is not fully appropriate.  Using simple geometrical considerations the magnification $m$ of the mask pattern size is related to the source to the mask plane distance (here $z$) and to the mask-detector separation $f$:  \\begin{equation}m = (z+f)/z. \\end{equation} The standard deconvolution procedure does not allow the source peak reconstruction because the beam divergence defocuses the incident radiation. However we can consider the detector image as the result of convolution of the source with a virtual mask, i. e. the magnified version of the true mask, or alternatively we can consider a virtual detector (the \\textit{reduced } version of the true one) as the convolution with the true mask. Both approaches consist in a sort of a ``collimation'' of the modulated radiation thus suppressing the defocusing due to the beam divergence. In the case of a virtual mask, the focusing requires the replacement of the reconstruction matrix $R$ in Eq. (2) with a matrix $R'$ obtained by expanding R with a factor of $m$ so that $R'{ij}$ will have values between -1 and 1, defined by interpolation. Instead, the use of a virtual detector offers the advantage to leave $R$ unchanged. In both cases, the open fraction value ($0.5$) is altered, because only a fraction of the illuminated mask pattern is projected on the detector.   Thus, there is no strong argument in support of the choice of either procedure. However, in order to have the same configuration of mask-detector system as for parallel beams (i. e. the same correspondence between mask element and detector pixel) and to decode the image with the same reconstruction matrix $R$, we decided to develop a reconstruction method based on a virtual detector ``reduced'' by a factor equal to $1/m$.  The reduction procedure (hereafter ``shrinking procedure'') has been checked on the 1-D coded mask system of a true experiment, SuperAGILE, however this procedure, mainly based on simple geometrical considerations, may also be used for 2-D systems.  In the next section we present the reduction procedure together  with a detailed analytical description of the systematics related to the method for the case of 1-D coded mask-instruments.  In section 3 we show the reliability of the method by means of computer simulations, assuming the instrumental set-up that was used to validate our method by means of experimental tests. The results of these tests are presented in section 4, where our method was applied to the ground calibrations of the SuperAGILE experiment (\\cite{fer07}). ",
        "conclusions": "We described the method developed to correct for the beam divergence in coded mask imaging of sources placed at finite distance from the detector plane. The method has been studied and applied on 1-D coded-mask instrument, although its extension to the 2-D instruments is feasible. It works to collimate the modulated radiation from the X-ray source with reference to one direction within the beam divergence. The choice of the direction is determined by the best suppression of defocusing effects of the image, associated to the magnification of the source shadowgram in the detector.   This requires the knowledge of the source distance to the mask plane and the mask-detector separation. If these two values are not known or known within some uncertainty, the best focusing may be found searching for the highest signal to noise ratio in the reconstructed image. This is a function of both the source distance to the detector plane (and then of the magnification factor) and the direction (within the beam divergence) with respect to which photons will be collimated. We argued that there is only a value of the magnification factor corresponding to the maximum of the signal to noise ratio, while  there are several (depending on the distance) favorite directions to collimate photons, going to one possibility for the infinite distance case. A detailed study of the systematics introduced by our method on the reconstructed source image has been done. We found a reduction of the peak counts compared with the infinite distance, which amount depends on the source distance and the beam divergence (neglecting the source spatial extension).    To check the reliability of the method we performed a set of Monte Carlo simulations of X-rays sources imaged by the SuperAGILE experiment.  We simulated point-like sources as well as extended sources at finite distance. We found that the flux attenuation with respect to the parallel beam case mainly depends on the geometric extension of the source. The dependence on the source distance is quite negligible due to the particular configuration of the coded-mask system of SuperAGILE.  Then we found  a flux attenuation of $\\sim 17\\%$ compared with the parallel beam case for a point-like source. This clearly quantifies the widening of the Point Spread Function due to the systematics of our method. The flux reduction increases  to $\\sim 40\\%$ if e. g. a source with a disk shape of $5$ mm in diameter is considered.  By means of these simulations, we estimated the statistical errors on both orthogonal and lateral displacements of the source obtained with our method. For a signal to noise ratio of  $~ 300$, we find that the 1-sigma error for the orthogonal displacement is  0.1 cm, while it is negligible  for the lateral displacement. These estimations of the method accuracy have been used to improve the ground calibration set-up of SuperAGILE. Searching for a good trade-off between the integration time (at the required $S/N$ ratio) and the physical quantities needed to be calibrated, we fixed a distance of $200$ cm of the source to the detector plane. We imaged Cd$^{109}$ and I$^{125}$ at $\\sim 200$ cm. In particular, tests with Cd$^{109}$ were done from -50 up to 50 degrees with a step of 10 degrees in the experiment field of view. In this case we succeeded to obtain a good reconstruction of the source peak from -40 up to 50 degrees off-axis. The values of source distance from the detection plane as well as its angular displacement inferred from the shrinking procedure are in good agreement with the metrological measurements performed with a high precision laser tracker ($\\sim 100 \\mu$m accuracy). This allowed us to calibrate the conversion between detector pixel and sky degree within 2 arc minutes uncertainty (to be compared with SuperAGILE pixel size of 6 arc minutes) and to study the transparency of materials surrounding SuperAGILE."
    },
    "0807/0807.0243_arXiv.txt": {
        "abstract": "We present a revaluation of the optical morphology for 549 galaxies from the Catalog of Isolated Galaxies in the Northern Hemisphere (CIG; Karachentseva 1973) that are available in the Sloan Digital Sky Survey (SDSS; DR6). Both the high resolution and high dynamic range of the SDSS images and our semi--automatic image processing scheme, allow for a major quality and uniform morphological analysis. The processing scheme includes (1) sky--subtracted, cleaned, and logarithmic scaled $g-$band images; (2) filtered--enhanced versions of the images in (1); and (3) the corresponding RGB images available in the SDSS database. Special attention was paid for distinguishing between E, S0, and Sa candidates through an additional analysis of (4) the surface brightness, $\\epsilon$, $PA$ and $A_{4}/B_{4}$ Fourier series expansion profiles. An atlas of mosaics containing [(1), (2) and (3)] images for Sab--Sm/Irr types and [(1), (2), (3), (4)] images for E/S0/Sa types was produced and is available on the web. The median type in the sample corresponds to Sbc, with 65\\% of the sample being of this type or later. A scarce population of E (3.5\\%) and S0 (5\\%) galaxies amounting to 8.5\\%, and a non-negligible 14\\% fraction of early--type SaSab spirals are identified. We compare our results against a previous reclassification of the CIG catalog based on the digitized POSS II images (Sulentic et al. 2006).  We calculate also the $gri$ absolute magnitudes corrected by Galactic and internal extinctions and present the $g-i$ color distribution and the color--magnitude diagram. Among the spirals, we find tentative fractions of strong and suspected bars of 65.8\\% and of 33.3\\% of rings. A detailed image analysis of the E galaxies (18) allows us to find a richness of distinct substructure in their isophotal shape and also of morphological distortions. At least 78\\% of the E galaxies show some kind of morphological distortion (shells, dust lanes,  diffuse halos, etc.), suggesting that these galaxies suffered late dry mergers. The isophotes of 42\\% (37\\%) of the E galaxies are boxy (disky). Among 4 blue E's, 3 are disky. Finally, we calculate for all the sample the $CAS$ (concentration, asymmetry and clumpiness) structural parameters in the $ugriz$ bands.  We analyze the loci of these galaxies in different projections of the $CAS$ volume diagram and discuss some trends of the $CAS$ parameters with the color band, as well as with the morphological type and the galaxy color. ",
        "introduction": "\\label{S1}   A galaxy is considered as isolated if it has not suffered any interaction with another normal galaxy or with a group/cluster environment over a Hubble time or at least since approximately one half of its mass was assembled. In this sense, isolation offers a unique opportunity to identify the intrinsic physical and evolutive processes of galaxies, and then to confront the results with theoretical predictions of galaxy evolution (e.g., Avila-Reese et al. 2008). Isolated galaxies are also important as comparison objects in studies of the morphological mix as a function of environment, and the environmental effects on galaxy properties. However, the definition of isolated galaxy from an observational point of view is not an easy task. Karachentseva (1973) introduced an operational criterion for this definition and constructed the Catalog of Isolated Galaxies in the Northern Hemisphere (hereafter CIG), one of the most extended and used samples for these objects. The sample contains 1050 galaxies of all morphological types and $m_B<15.7$ mag, identified in the photographic Palomar Observatory Sky Survey (POSS) plates. A careful revision of the morphology and general properties of these galaxies in the light of recent, much more detailed observations is highly desirable.  A first step in this direction has been given by Sulentic et al. (2006; hereafter Sul06), who revisited the morphology of the CIG galaxies by using the digitalized POSS II data, and found some fraction of galaxies having a morphological type different to those reported in the original CIG.   In this paper, we are aimed to analyze the CIG galaxies in common with the huge galaxy sample of the Sloan Digital Sky Survey (SDSS), a sample that we will call CIG$\\cap$SDSS. The advantages of the SDSS CCD images over previous image surveys in similar regions of the sky are notorious. Among these we mention: (1) the uniformity in both the instrumental set and photometric conditions, (2) the homogeneity in multicolor $ugriz$ imaging and photometry, and (3) the high resolution  (0.4 $\\arcsec$/pix) of the images.  These advantages are crucial for studying galaxy morphology and structural properties.   Because of the SDSS large numbers, uniformity, and completeness down to a certain magnitude, it is an ideal database for the search of fundamental relations between physical properties of galaxies and their dependencies on local and larger--scale environments. For example, recently Park et al. (2007,2008) used the SDSS data for studying the environmental dependence of the observed morphology --and other properties-- at different density smoothing scales. Based on their results, they concluded that galaxy morphology depends primarily on the small scale environment, which may reflect the influence of later evolutionary effects such as galaxy--galaxy interactions, rather than on large--scale environment in which the galaxy initially formed.  A relevant goal of the present paper is the identification of the isolated galaxies that may show evidence of morphological disturbances. It is possible that the isolated galaxies showing signs of distortions are the recent merger products, at a given large-scale background density, as predicted by Park, Gott \\& Choi (2008).  An important question concerning morphological classification is the correct definition of what is an elliptical (E), lenticular (S0) or early--type spiral (Sa) galaxy. These types are the most affected in their definition by saturation, poor resolution, etc. Saucedo--Morales \\& Bieging (2001) estimated that one--half of E galaxies could be misclassified in the POSS plates--based CIG catalog. The existence of true isolated early--type galaxies is in general an important problem in galaxy formation. Our finding of detailed structure in the E and S0 isolated galaxies is expected to provide new tests for the various hypotheses of their formation. Have the isolated E galaxies systematic differences from those ones found in groups and clusters? Have they assembled in the same way in both cases? Similar questions could be asked about the isolated S0 galaxies, whose detailed formation histories are poorly understood. A few decades back, Gott \\& Thuan (1976) argued that E galaxies were produced by larger initial--density fluctuations, which would have higher density at turnaround and where star formation (SF) would be completed before collapse. Such larger initial--density fluctuations would be more likely in a region that would later turn into a high density environment. In the modern context of the $\\Lambda$ Cold Dark Matter ($\\Lambda$CDM) hierarchical model, the formation of E/S0 galaxies is also favored in the dense environments. The models predict a small probability of finding isolated E/S0 galaxies and they should have younger stars on average than their cluster counterparts, as well as assembled their present--day masses by late major mergers (Kauffmann 1996; de Lucia et al. 2006).  A complementary analysis of galaxy morphology and gross photometric properties has been introduced some years ago, namely the $CAS$ (concentration, asymmetry and clumpiness) system. This system has been proposed to distinguish galaxies at different stages of evolution (Conselice 2003, and references therein) and it provides the possibility to carry out a physical and quantitative classification based on measurable structural parameters. We present here the measured $CAS$ parameters in the $ugriz$ bands for the 549 CIG$\\cap$SDSS isolated galaxies.  The outline of the paper is as follows. Section 2 summarizes the relevant past work on CIG morphologies. In \\S 3 the image processing scheme and considerations used to reevaluate morphology for the CIG$\\cap$SDSS galaxies are presented. In \\S 4 a description of our Atlas of mosaic images for the 549 CIG$\\cap$SDSS galaxies is presented.  The results of our classification and morphological analysis, and the calculation of the $CAS$  parameters in the $ugriz$ bands are presented. Section 5 is a discussion containing  a comparison of our morphological classification to previous works (\\S 5.1), a discussion on the bar and ring fractions in spiral galaxies (\\S 5.2), the morphological distortions and formation of isolated E galaxies (\\S 5.3), and the $CAS$ parameters in different bands and their correlations with other galaxy properties (\\S 5.4). Finally, in \\S 6 a summary and the main conclusions of this paper are presented. ",
        "conclusions": "\\label{S5}  The building up of a well--defined sample of local isolated galaxies with uniformly--derived and detailed morphological information is of great relevance because it provides a fair database for comparison with model predictions as well as with observed samples of galaxies in other environments and at higher redshifts. In this paper we reported (i) a detailed morphological revaluation of a sample of 549 CIG galaxies that were available in the SDSS(DR6) database, (ii) looked for morphological distortions in the analyzed isolated galaxies, (iii) calculated corrected magnitudes and colors, as well as the structural $CAS$ parameters for all the galaxies, and (iv) discussed some preliminary conclusions based on the obtained results.  We have taken advantage of the improved scale and dynamic range of the SDSS database and proposed an image processing scheme that enhances the detection of low/high surface brightness morphological and structural components in these CIG galaxies. We have used logarithmic scaled images to look for internal/external details; $g$ band filtered--enhanced images to look for internal structure in the form of star forming regions, bars, rings and/or structure embedded into dusty regions; and used complementary RGB color images from the SDSS database to visualize the spatial distribution of the SF and other components like dust (blue colors for recent SF and red colors for older populations/dusty components).  In the case of E/S0/Sa galaxy candidates, we have estimated and used additional information in the form of a surface brightness profile and the corresponding geometric profiles (ellipticity $\\epsilon$, Position Angle $PA$ and $A_{4}/B_{4}$ coefficients of the Fourier series expansions of deviations of a pure ellipse) from the $r$ band images in order to provide evidence of boxiness/diskiness and other structural details that help to discriminate morphology in these galaxies.  The main results and conclusions of the paper are as follow:  --The median morphological type in the present CIG$\\cap$SDSS sample is Sbc ($T=4$); 35\\% of the galaxies are earlier than Sbc, and 65\\% are Sbc or later. This new revaluation gives somewhat different morphological fractions with respect to previous reports, mainly in the E, S0, Sa and Sab types. While Sul06 found 5.7\\% of E types, 6.6\\% of S0 types and 1.4\\% of transition E/S0 types, amounting to 13.7\\% of E/S0 galaxies, our results indicate 3.5\\% of E types and 5\\% of S0 types, amounting to a fraction of only 8.5\\% of E/S0 galaxies. This emphasizes the difficulty of finding E/S0 galaxies in environments typical to the isolated galaxies. For early--type S galaxies, while Sul06 found 1.3\\% of Sa types and 5.1\\% of Sab types, we accounted for 6.5\\% of Sa types and 7.8\\% of Sab types. In general, for 74\\% of the galaxies in common with Sul06 both classifications have a difference $|\\Delta T | < 2$.  --For the 539 CIG$\\cap$SDSS galaxies we have calculated the absolute magnitudes in the bands $gri$ and corrected by both the Galaxy and internal extinction. The average and 1$\\sigma$ dispersions of the $g-i$ color for the E/S0, Sa--Sb, and Sbc--Sm/Irr morphology ranges are $1.18 \\pm 0.26$, $0.89 \\pm 0.21$, and $0.66 \\pm 0.55$, respectively. The color--magnitude diagram for the E/S0's is quite flat and narrow, while for the disk galaxies, a trend of redder galaxies as they get brighter, but with a large scatter is appreciated.  --We reported a tentative $r-$band fraction of bars in disk galaxies of 65.8\\% (defined only by an eye inspection and excluding objects with $i> 60^{\\circ}$); $35\\%$ are clearly barred (SB), and $30.8\\%$ show evidence of a weak bar (SAB). The reported fraction of bars in the present CIG isolated disk galaxies is larger than others reported in the literature for samples related to field and cluster/group environments. The fraction of barred galaxies does not change with the morphological type, but it correlates significantly with the light concentration parameter $C_{80/20}$. We have also found 33.3\\% of the disk galaxies showing ring structures.   --Our image processing scheme of the SDSS data allowed us to find a richness of distinct substructure in the isophotal shapes of the E galaxies as well as a series of morphological distortions. Out of the 18 E galaxies in the sample, 8 (44\\%) show boxy isophotes and 7 (38\\%) show disky isophotes. Among the 4 blue E's ($g-i < 1$), 3 are disky and 1 boxy, but the latter shows evidence of an inner disk. The fraction of E's with one or more kinds of morphological distortion amounts to 78\\%: fine stuctures (50\\%), dust lane structures (17\\%), diffuse halo (39\\%), and inner disk (39\\%). Our results suggest that most of the isolated E's suffered recent dry major mergers (E--E). Only a few of them show pieces of evidence of mergers of galaxies with gas, and/or of late gas accretion and intense SF.  --In order to complement our morphological study, we have estimated the $ugriz$ band $CAS$ parameters of the 549 CIG$\\cap$SDSS galaxies as part of a model--independent morphological system. The $CAS$ averages and standard deviations in the $r-$band for the E/S0 galaxies are $<C> = 3.81 \\pm 0.46$, $<A> = 0.09 \\pm 0.03$, $<S> = 0.08 \\pm 0.10$, for the SaSb galaxies are $<C> = 3.58 \\pm 0.52$, $<A> = 0.14 \\pm 0.06$, $<S> = 0.16 \\pm 0.16$, and for the SbcSm galaxies are $<C> = 2.97 \\pm 0.49$, $<A> = 0.20 \\pm 0.07$, $<S> = 0.25 \\pm 0.18$. These values are in rough agreement with those reported previously for smaller samples of isolated galaxies in the $R$-band Johnson-Cousins system (Hern\\'andez-Toledo et al. 2006;2008).  The mean asymmetries reported here are definitively lower than the typical ones for interacting disk galaxies (Hern\\'andez--Toledo et al. 2005). While the $C$ parameter systematically increases from bluer to redder bands, both $A$ and $S$ significantly decrease. The $CAS$ parameters present more robust trends with the morphological type $T$ and the total $(u-r)$ color in the redder bands, suggesting that the basic structure of galaxies is revealed better in the redder bands. In any case, the mentioned trends are very noisy, the less scattered being with $C$."
    },
    "0807/0807.4710_arXiv.txt": {
        "abstract": "{Toroidal obscuration is a keystone of AGN unification. There is now direct evidence for the torus emission in infrared, and possibly water masers. Here I summarize the torus properties, its possible relation to the immediate molecular environment of the AGN and present some speculations on how it might evolve with the AGN luminosity. ",
        "introduction": "Although there are numerous classes of active galactic nuclei (AGN), a unified scheme has been emerging steadily \\citep[e.g.,][]{Ski93, Urry95}. The nuclear activity is powered by a supermassive (\\about\\E6--\\E{10} \\Mo) black hole and its accretion disk. This central engine is surrounded by a dusty torus, which can be considered an acronym for Toroidal Obscuration Required by Unification Schemes: much of the observed diversity is simply explained as the result of viewing this axisymmetric geometry from different angles. Because of the anisotropic obscuration of the central region, sources viewed face-on are recognized as ``type 1'', those observed edge-on are ``type 2''. From basic considerations, \\cite{Krolik88} concluded that the obscuration is likely to consist of a large number of individually very optically thick dusty clouds. ",
        "conclusions": ""
    },
    "0807/0807.2650_arXiv.txt": {
        "abstract": "{ We present lunar occultation (LO) observations obtained in August 2006 with the recently demonstrated burst mode of the ISAAC instrument at the ESO VLT. The results presented here follow the previously reported observations carried out in March 2006 on a similar but unrelated set of sources. } { Interstellar extinction in the inner regions of the galactic bulge amounts to tens of magnitudes at visual wavelengths. As a consequence, the majority of sources in that area are poorly studied and large numbers of potentially interesting sources such as late-type giants with circumstellar shells, stellar masers, infrared stars, remain excluded from the typical investigations which are carried out in less problematic regions. Also undetected are a large numbers of binaries. By observing LO events in this region, we gain the means to investigate at least a selected number of sources with an unprecedented combination of sensitivity and angular resolution. } { The LO technique permits to achieve milliarcsecond resolution with a sensitivity of K$\\approx$12\\,mag at a very large telescope. We have used the opportunity of a favorable passage of the Moon over a crowded region in the general direction of the Galactic Center to observe 78 LO events of heavily reddened stellar sources. } {We have detected six new binary and one triple star, with typical projected separation of $\\approx 0\\farcs01$. We have also detected the compact circumstellar emission around one maser and one central star of a planetary nebula. Additionally we have measured the diameter and/or circumstellar shell of two carbon stars and other IR sources. } {We have used the upper limits on the size of about 60 unresolved or marginally resolved sources to establish the performance of the method. In agreement with our previous result, we conclude that lunar occultations in fast read-out mode on a detector subwindow at an 8\\,m-class telescope can achieve an angular resolution close to $0\\farcs001$ with a sensitivity K$\\approx 12$\\,mag. } ",
        "introduction": "This paper follows closely the scientific rationale, method, observations and results already presented in Richichi et al. (\\cite{PaperI}, Paper~I hereafter), where we introduced the technique of lunar occultations (LO) at the ESO Very Large Telescope using a fast readout (burst mode) of the ISAAC instrument. For our application, the burst mode allows to sample subarrays of 32 or 64 squared pixels with sampling times of 3.2 or 6.4\\,ms, respectively. We demonstrated that in this way it is possible to record LO data with the highest quality ever achieved in terms of signal-to-noise ratio (SNR), thanks also to the strong reduction of scintillation made possible by the 8.2\\,m mirror in the near-IR. In Paper~I we characterized in detail the performance, which approached $\\approx 1$ milliarcsecond (mas) in angular resolution, K$\\approx 12$\\,mag in sensitivity, and $\\Delta$K$\\approx 8$\\,mag in dynamic range on scales smaller than the Airy disk of the telescope. Notwithstanding the obvious limitations of LO with respect to the choice of the sources and the repeatability, this combined performance is superior to that of any other technique presently available for high angular resolution in the near-IR.  In Paper~I we took advantage of a recently developed data pipeline which employs an automated generation of masks and light curve extraction  and generates  the first-guess parameters for the fitting (AWLORP, Fors et al. \\cite{fors08}). This pipeline allows us to perform a quick preliminary inspection of large volume of LO data. Subsequently, a more detailed interactive analysis is performed on a selected number of sources using both model-dependent and model-independent procedures (ALOR and CAL respectively, Richichi et al. \\cite{richichi96} and Richichi \\cite{CAL}).  In the present paper we report on LO observations carried out in the night of August 5, 2006, during which 78 events were observed. We descrive the observations and the list of sources in Sect.~\\ref{data}, without recalling the details of the method and of the data analysis which can be found in Paper~I and references therein. Subsequently, we discuss the results in Sect.~\\ref{results}. These include several new detections of binary and triple stars, as well as angular diameters and circumstellar shells of diverse objects such as a maser, two carbon stars, and the central star of a planetary nebula. ",
        "conclusions": "We have reported on additional lunar occultation observations, after those reported in Paper~I, using the burst mode of the ISAAC instrument at the ESO VLT. During a passage of the Moon in the central regions of the galactic bulge, we have recorded 78 LO events (72 with a positive detection) over 8.5 hours. Our results include the new detection of six binaries and one triple star, with typical projected separations of $\\approx 10$\\,mas, i.e. five times less than the diffraction limit of the 8.2\\,m telescope. Brightness ratios were as low as 1:100. We have also detected and measured the extended emission arising from circumstellar dust around one maser star, two carbon stars, the central star of a planetary nebula and other IR stars.  Our sources have no or very few known counterparts, and other information such as fluxes or spectra is also missing or very scarce. Follow-up studies are necessary in order to better characterize the energy distribution and spectral characteristics of these sources. Our measurements have also shown a significant amount of variability in many of the sources, with differences between our determination and the 2MASS flux which approached 2\\,mag in some cases.  As shown in this paper and in Paper~I, the LO technique using a fast readout on a detector subwindow at a large telescope  provides an unprecedented combination of sensitivity and high angular resolution, namely K$\\approx$12\\,mag and $0\\farcs001$ respectively. Fast-photometry on a detector subwindow is now a well-characterized observational mode, publicly offered at the VLT."
    },
    "0807/0807.2699_arXiv.txt": {
        "abstract": "We investigate the variation of bar strength with central velocity dispersion in a sample of barred spiral galaxies. The bar strength is characterized by $Q_g$, the maximal tangential perturbation associated with the bar, normalized by the mean axisymmetric force. It is derived from the galaxy potentials which are obtained using near-infrared images of the galaxies. However, $Q_g$ is sensitive to bulge mass. Hence we also estimated bar strengths from the relative Fourier intensity amplitude ($A_{2}$) of bars in near-infrared images. The central velocity dispersions were obtained from integral field spectroscopy observations of the velocity fields in the centers of these galaxies; it was normalized by the rotation curve amplitude obtained from HI line width for each galaxy. We found a correlation between bar strengths (both $Q_g$ and $A_{2}$) and the normalized central velocity dispersions in our sample. This suggests that bars weaken as their central components become kinematically hotter. This may have important implications for the secular evolution of barred galaxies. ",
        "introduction": "In past decades there has been a large amount of discussion regarding the secular evolution of bars in spiral galaxies. Numerical simulations show that gas inflow, star formation and the subsequent buildup of a central mass concentration (CMC) can considerably alter the shape of a bar and even result in the dissolution of the bar itself (Norman, Sellwood \\& Hasan 1996; Shen \\& Sellwood 2004; Kormendy \\& Kennicutt 2004; Hozumi \\& Hernquist 2005; Athanassoula, Lambert \\& Dehnen 2005; Bournaud, Combes \\& Semelin 2005). However, the fraction of bars remains fairly high out to a redshift of z$\\sim$0.7 to 1 (Sheth et al. 2003; Elmegreen, Elmegreen \\& Hirst 2004; Jogee et al. 2004; Sheth et al. 2007) which suggests that either bars are long-lived or can reform in galaxy disks. Simulations suggest that bars can dissolve but form again once the disk is cooled through gas accretion (Bournaud \\& Combes 2002) or through tidal interactions (Berentzen et al. 2004).  Although there has been a considerable amount of numerical work on bar formation and dissolution, there is not much observational evidence. Molecular gas has a higher central concentration in barred galaxies compared to non-barred galaxies which suggests that bars are important for building up a CMC (Sakamoto et al. 1999; Sheth et al. 2005). The bar length has been found to correlate with bar intensity contrast and central luminosity densities (Elmegreen et al. 2007). A more direct indicator of how bars change with mass concentration is the correlation of CMC and bar ellipticity (Das et al. 2003) which shows that bar ellipticity declines with increasing dynamical CMC. An enhanced CMC also results in the deepening of the central potential and gives rise to a dynamically hotter galaxy center. This will lead to a higher observed central stellar velocity dispersion in a galaxy ($\\sigma_{v}$). This effect has been seen in numerical simulations where a dynamically hotter center results when a bar dissolves (Friedli \\& Benz 1993; Hasan, Pfenniger \\& Norman 1993) or weakens (Athanassoula, Lambert \\& Dehnen 2005). In this paper we investigate whether there is any observational evidence for this heating effect by examining the correlation between the central velocity dispersion in bars and the bar strength. ",
        "conclusions": "The main result of this paper are shown in Figures 1 \\& 2. There are 26 galaxies in the plots of which about half are intermediate type spirals and the remaining a mixture of early and late type spirals. Since the number of galaxies in each Hubble type is not very large, we cannot investigate trends within the different Hubble types. But from Figures 1 \\& 2, it appears that early type spirals have relatively lower central dispersions; this may be because they have larger bulges where rotational velocity is comparatively higher and the central velocity dispersion lower compared to the later Hubble types.  Since the correlation is significant but not very strong we examined the two galaxies that define the higher and lower limits of $Q_g$ and $A_2$. We looked at them closely to see if they are odd in some way, and not characteristic of the rest of the sample. (i)~NGC~3162 ($\\sigma_{e}/v_g$=1.13)~:~This is an intermediate type spiral galaxy with a weak bar and prominent bulge. There may also be a ring in the center. Though the spiral arms are somewhat asymmetric, the nucleus is fairly undisturbed. (ii)~NGC~4314 ($\\sigma_{e}/v_g$=0.18)~:~This is a bright, early type galaxy with a strong bar and large bulge with a LINER type nucleus. There may be significant rotation in the nuclear region which may lower the $\\sigma_{e}$. Also, the mass may be more widely distributed over the bucle and hence not as centrally concentrated as in NGC~3162. There appears to be a ring of star formation in the nucleus as well (Gonzalez Delgado et al. 1997). Both galaxies are thus fairly normal and not unusually different from the rest of the sample galaxies.  Figures~1 \\& 2 suggest that galaxies with dynamically hotter nuclei have weaker bars. It is now well established that a galaxy's central black hole mass and bulge velocity dispersion are correlated (Ferrarese \\& Merritt 2000; Gebhardt et al. 2000). Later results show that the nuclear mass also correlates with the overall mass of a galaxy (Ferrarese et al. 2006; Hopkins et al. 2007). These results all indicate that the nuclear mass in a galaxy is intimately connected to the dynamics of its disk and halo. If so, then it is not suprising that $\\sigma_{e}$ in our sample of barred galaxies is correlated with $Q_g$ or $A_{2}$. It suggests that the growth of a central mass and evolution of the bar for spiral galaxies may be closely linked.  Our results have important implications for the secular evolution of barred galaxies. Simulations suggest that there may be several factors responsible for the dissolution of a bar. One is the CMC growth (Hasan \\& Norman 1990; Friedli \\& Pfenniger 1991) which weakens the bar-supporting $x_1$ orbits and increases the fraction of chaotic orbits in the galaxy center. Second is the inflow of gas towards the galaxy center which results in the transfer of angular momentum to the bar wave which then weakens the bar itself (Bournaud, Combes \\& Semelin 2005). The buckling instability is an important bar thickening mechanism that results in a boxy/peanut bulge which can temporarily weaken a bar (Raha et al. 1991; Berentzen et al. 1998; Athanassoula \\& Misiriotis 2002; Martinez-Valpuesta, Shlosman \\& Heller 2006; Debattista et al. 2006). All these effects result in a more massive and dynamically hotter central component and a weaker bar. The correlations that we see in Figures 1 and 2 may be observational indications of this ongoing evolution. In particular, while the apparent drop of $Q_g$ with central velocity dispersion might be an artifact caused by the bulge dilution effect (see Laurikainen, Salo \\& Buta 2004; a bias could follow since nuclear velocity dispersion and bulge mass are strongly correlated), the similar correlation between the bar intensity contrast $A_2$ suggest that the effect is real. The found correlation between $A_2$ and $\\sigma_e/v_g$ is also consistent with Das et al. (2003) who found that the bar ellipticity (closely related to $A_2$) drops with central mass concentration.  The evolution of bars by secular processes in galaxies is an issue which is expected to gain more attention in the near future. Recent observational evidence shows that the fraction of strong bars in bright galaxies increases from under 10\\% at redshift z=0.84 to about 30\\% in the local universe (Sheth et 2007). Also, it has been shown by (Laurikainen et al. 2007) that among the early-type barred galaxies the bulge-to-total flux ratios are on average smaller than in the non-barred galaxies. These results, together with ours, may indicate that bars evolve with their parent galaxies."
    },
    "0807/0807.0180_arXiv.txt": {
        "abstract": "A model of multi-component hybrid inflation, dubbed multi-brid inflation, in which various observable quantities including the non-Gaussianity parameter $f_{NL}$ can be analytically calculated was proposed recently. In particular, for a two-brid inflation model with an exponential potential and the condition that the end of inflation is an ellipse in the field space, it was found that, while keeping the other observational quantities within the range consistent with observations, large non-Gaussianity is possible for certain inflationary trajectories, provided that the ratio of the two masses is large. One might question whether the resulting large non-Gaussianity is specific to this particular form of the potential and the condition for the end of inflation. In this paper, we consider a model of multi-brid inflation with a potential given by an exponential function of terms quadratic in the scalar field components. We also consider a more general class of ellipses for the end of inflation than those studied previously. Then, focusing on the case of two-brid inflation, we find that large non-Gaussianity is possible in the present model even for the equal-mass case. Then by tuning the model parameters, we find that there exist models for which both the non-Gaussianity and the tensor-to-scalar ratio are large enough to be detected in the very near future. ",
        "introduction": "The primordial non-Gaussianity has been one of the hottest topics in cosmology in recent years. The conventional, single-field slow-roll inflation predicts that the curvature perturbation is Gaussian to an extremely high accuracy~\\cite{Salopek:1990jq}. In other words, if any primordial non-Gaussianity is detected, it strongly indicates that the dynamics of inflation is not as simple as we thought it to be.  The primordial non-Gaussianity is conveniently represented by a parameter denoted by $f_{NL}$~\\cite{Komatsu:2001rj}. Roughly, it is the ratio of the 3-point correlation function (or the bispectrum) to the square of the 2-point correlation function (or the square of the spectrum). It is expected that near-future experiments such as those of PLANCK will be able to detect $f_{NL}$ at a level as small as 5~\\cite{Babich:2004yc}.  Finding even a small deviation from Gaussianity will have profound implications on the theory of the early universe. Consequently, numerous types of inflationary models that produce detectable non-Gaussianity have been proposed and studied~\\cite{NGmodels1,NGmodels2,Bernardeau:2002jf, Dvali:2003em,Kofman:2003nx,Lyth:2005qk,Alabidi,BarnabyCline,Chambers:2007se}. In terms of the nature of non-Gaussianities, most of these models can be classified into two categories; those with non-Gaussianities arising intrinsically from the quantum fluctuations, and those with non-Gaussianities due to nontrivial classical dynamics on superhorizon scales. A typical example of the former is the DBI inflation, in which the slow-roll condition can be fully violated~\\cite{Silverstein:2003hf}. In this case, the equilateral $f_{NL}$ (denoted by $f_{NL}^{\\rm equil}$) representing the amplitude of the bispectrum of the equilateral configurations, is found to play an important role~\\cite{DBIinflation}. On the other hand, in the latter case where non-Gaussianities are produced on superhorizon scales, by causality the local $f_{NL}$ (denoted by $f_{NL}^{\\rm local}$) characterizes the level of the non-Gaussianity. It is defined in terms of the coefficient in front of the second order curvature perturbation~\\cite{Komatsu:2001rj}, \\begin{eqnarray} \\Phi=\\Phi_L+f_{NL}^{\\rm local}\\Phi_L^2\\,, \\label{fNLlocaldef} \\end{eqnarray} where $\\Phi$ is the curvature perturbation on the Newtonian slice and $\\Phi_L$ is its linear, Gaussian part.  In this paper, we focus on the latter case, that is, we consider models that may produce a large value of $f_{NL}^{\\rm local}$, for example 10--100. More specifically, we consider hybrid inflation with multiple inflaton fields, dubbed multi-brid inflation~\\cite{Sasaki:2008uc}. The inflaton fields are assumed to follow the slow-roll equations of motion, and their fluctuations are assumed to be Gaussian. In this case, the $\\delta N$ formalism is most useful for the evaluation of the curvature perturbation and non-Gaussianity~\\cite{Sasaki:1995aw,Sasaki:1998ug,Wands:2000dp, Lyth:2004gb,Lyth:2005fi}.  As in the conventional hybrid inflation, the inflaton fields are coupled to a water-fall field, and inflation ends when the inflaton fields satisfy a certain condition that triggers the instability of the water-fall field. However, unlike the case of a single inflaton field in which there is essentially no degree of freedom in the condition for the end of inflation, there is a substantial increase in the degree of freedom at the end of inflation in multi-brid inflation and it widens the viable range of the parameter space considerably and leads to the possibility of generating large non-Gaussianity.  As a model of multi-brid inflation, an analytically solvable two-brid inflation model was recently investigated in detail~\\cite{Sasaki:2008uc}, where the potential was assumed to be exponential with the exponent given by a linear combination of the inflaton fields. In this paper, we consider a two-brid model with again an exponential potential but with the exponent given by a quadratic function of the inflaton fields. The potential has point symmetry about the origin of the field space, in contrast to the case of the linear exponent which has no symmetry. Thus by investigating the quadratic case, we will be able to see if the generation of large non-Gaussianity is a generic feature of multi-brid inflation or if it is due to the lack of symmetry that leads to large non-Gaussianity in the linear exponent case.  In passing, we mention that the possibility of large non-Gaussianity from loop correction terms in the perturbative expansion has recently been discussed by Cogollo et al.~\\cite{Cogollo:2008bi}. Although such a case certainly needs further investigation, in this paper we concentrate on the case in which leading order (tree) terms dominate over loop correction terms.  This paper is organized as follows. In \\S 2, we describe our model and derive basic formulas to be used in the proceeding sections. In \\S 3, using the formulas derived in \\S 2, we analytically compute the spectrum of the curvature perturbation ${\\cal P}_S(k)$, its spectral index $n_S$, the tensor-to-scalar ratio $r$, and the non-Gaussianity parameter $f_{NL}^{\\rm local}$. In \\S 4, as a couple of special cases of our model, we analyze in detail the equal-mass case as well as the case of large mass ratio. We find that a large $f_{NL}^{\\rm local}$ is possible in both cases. In particular, in the equal-mass case, by tuning the parameters to some extent, we find that it is possible to have both $r$ and $f_{NL}^{\\rm local}$ large enough to be detected. We conclude the paper in \\S 5. Some computational details are described in Appendix~\\ref{app:cal}. For comparison, we also summarize the result for the case of the linear exponent two-brid model in Appendix~\\ref{app:linear}. We use the Planck units where $M_{pl}^{-2}=8\\pi G=1$. ",
        "conclusions": "\\label{sec:conclusion}  We analytically investigated the curvature perturbation and its non-Gaussianity in a model of multi-field hybrid inflation, dubbed {\\it multi-brid\\/} inflation. The model we considered is a two-field hybrid inflation (two-brid inflation) model with the potential mimicking conventional quadratic potentials. The new ingredient of the model is the generalization of the condition for the end of inflation. We considered a very general coupling of the two inflaton fields to a water-fall field.  Then, using the $\\delta N$ formula, we derived an analytical expression for the curvature perturbation. Based on this expression, we obtained the curvature perturbation spectrum ${\\cal P}_S$, the spectral index $n_S$, the tensor-to-scalar ratio $r$, and the non-Gaussian parameter $f_{NL}^{\\rm local}$. We found that a large positive $f_{NL}^{\\rm local}$ is possible in this model. Then, at least for a certain limited range of the parameters, we explicitly showed that it is possible to have large non-Gaussianity while keeping the values of the other quantities consistent with those of the observation. In particular, we showed that when the two inflaton masses are equal, the parameters can be tuned so that they lead to a fairly large tensor-to-scalar ratio, $r\\sim 0.1$, as well as a large non-Gaussian parameter, $f_{NL}^{\\rm local}\\sim 50$. These values will be at a detectable level in the very near future. On the other hand, interestingly, we found that having a large mass ratio in the present model does not help in producing both $r$ and $f_{NL}^{\\rm local}$ large enough to be detected. This is in contrast to the model studied in~\\cite{Sasaki:2008uc}.  The standard lore has been that $f_{NL}^{\\rm local}$ is too small for models with large $r$ or vice versa. We have shown, in this paper, not be the case, particularly in this model of spontaneously symmetry breaking at the end of inflation. This may be the most important conclusion of this work. At the moment, we have no clear physical explanation for this result. We hope we will be able to answer this question in the near future."
    },
    "0807/0807.0713_arXiv.txt": {
        "abstract": "A robust second order, shock-capturing numerical scheme for multi-dimensional special relativistic magnetohydrodynamics on computational domains with adaptive mesh refinement is presented. The base solver is a total variation diminishing Lax-Friedrichs scheme in a finite volume setting and is combined with a diffusive approach for controlling magnetic monopole errors. The consistency between the primitive and conservative variables is ensured at all limited reconstructions and the spatial part of the four velocity is used as a primitive variable. Demonstrative relativistic examples are shown to validate the implementation. We recover known exact solutions to relativistic MHD Riemann problems, and simulate the shock-dominated long term evolution of Lorentz factor 7 vortical flows distorting magnetic island chains. ",
        "introduction": "Relativistic flows of magnetized plasmas are observed in a wide variety of astrophysical objects. At galactic scales, accretion of matter around black holes in Active Galactic Nuclei (AGN) in addition results in well collimated jets emitted along the rotation axis. To explain their observed superluminal plasma motion at the parsec length scales, flows with Lorentz factors of more than 10 are needed. The observed synchrotron emission indicates that these jets are pervaded by magnetic fields. Even more powerful processes are at play in the highly relativistic blast waves associated with Gamma Ray Bursts (GRBs). Here, the plasma flows easily reach Lorentz factors of 100 or even higher. The morphology and time evolution of these and other astrophysical objects often involve strong shocks and complex magnetic field topologies. To compute this kind of challenging relativistic plasma dynamics, the combination of Adaptive Mesh Refinement (AMR) with a robust, shock-capturing numerical method is therefore indispensible.  With growing interest in relativistic astrophysical phenomena, various efforts are ongoing to develop numerical special relativistic hydrodynamic and magnetohydrodynamic codes. Significant progress was achieved in the last decade with the development of conservative shock-capturing schemes, which use either exact Riemann solvers or approximate Riemann solvers, or more robust central type schemes, in relativistic hydrodynamics~\\cite{Eulderink&Melemma94, Fontetal94, Donatetal, Aloyetal99, delzanna02, Marti&Muller03, Mignone&Bodo05} (for a contemporary review see \\cite{Marti&Muller03}) and in relativistic magnetohydrodynamics \\cite{komis, balsara, koldoba, delzanna, jeroen, leis, Mignone&Bodo06, mignone2007, hllc07}. The study of relativistic hydrodynamic fluids currently starts to benefit also from using spatial and temporal adaptive techniques, or Adaptive Mesh Refinement (AMR) \\cite{Zhang&MacFadyen06, Melianiet07, Wangetal07}. To date, AMR was incorporated only in some works, both in special relativistic magnetohydrodynamics \\cite{jeroen} and in general relativistic magnetohydrodynamics \\cite{Anninosetal05,Andersonetal06,giacom2}. At the same time, various authors, including \\cite{Melianiet07}, have shown that AMR is imperative to simulate extreme astrophysical phenomena, such as those encountered in GRBs.  In order to handle GRB and other extreme relativistic flow regimes, the employed solver needs to be very robust under a wide variety of plasma conditions. Therefore, we decided to use a conservative discretization which does not fully exploit the detailed solution knowledge of the Riemann problem at each cell interface, defined by two constant states in contact. This contrasts with a true Godunov scheme, which would compute the flux across cell interfaces from the exact solution to the local Riemann problem. While fairly recently~\\cite{giacom}, an exact Riemann problem solver for the RMHD equations became available, the nonlinear iteration involved would make it a fairly computationally costly ingredient for a multidimensional code. As pointed out further, we already face a similar Newton-Raphson procedure to deduce primitive from conservative variables in every grid point. In fact, we follow the trend away from using exact or approximate Riemann problem solvers in most codes in use to date. In the works cited above, Komissarov~\\cite{komis} already used a linearized (approximate) Riemann solver, and even fell back on an HLLE variant~\\cite{hll83, einf88} which only uses the fastest wave speeds. Koldoba et al.~\\cite{koldoba} and Balsara~\\cite{balsara} also employ an approximate, linearized Riemann solver and presented numerical solutions for stringent 1D test problems. All other works mentioned~\\cite{delzanna, jeroen, leis, Mignone&Bodo06, hllc07} use central-type schemes, of either the local Lax-Friedrichs, HLL, or HLLC variant, suitably adopted to the relativistic MHD regime. We will use the simplest of these in our implementation, since in combination with automated grid adaptivity, we rely on convergence and robustness of the underlying discretization, while accuracy is most easily gained by raising the effective resolution employed.   The macroscopic dynamics of astrophysical objects with relativistic plasma flows are governed by the conservation of particle and momentum-energy, along with the Maxwell equations to account for the involvement of the magnetic field. In Sect.~\\ref{sec-eq}, we recast the RMHD equations in conservation form. Details on the implementation of the high-resolution shock-capturing scheme and the adaptive mesh refinement in our code can be found in Sects.~\\ref{sec-scheme}-\\ref{sec-amr}. The code is validated against several known 1D analytical solutions in Sect.~\\ref{sec-tests} and two multi-dimensional relativistic flow simulations are performed. In Sect.~\\ref{sec-summ}, a summary of the paper containing our main conclusions and an outlook for future applications of the code is given. ",
        "conclusions": "\\label{sec-summ} We provided details on our relativistic grid-adaptive MHD code, and tested it using recently available Riemann problem solutions, as well as multi-dimensional setups. The latter include a new long-term simulation demonstrating shock-shock interactions and reconnection events, and this could be of interest to benchmark existing shock-capturing algorithms on long-term shock-dominated relativistic magnetized flow problems, such as those recently presented in~\\cite{komis2}.  In future work, the RMHD code AMRVAC will be used to model AGN jet propagation~\\cite{jetsrmhd} and GRB dynamics. The magnetic field is suggested as possible ingredient to achieve the high Lorentz factor flows, such as reached by the GRB fireball, and plays a crucial role in both AGN and GRB flow collimation. Other future projects concentrate on implementing a more realistic equation of state to investigate the launching and the propagation of relativistic jet. We will discuss these astrophysical applications with more detail in future publications."
    },
    "0807/0807.3526_arXiv.txt": {
        "abstract": "We present a two-dimensional analysis of the bright nearby galaxy group NGC 5044 using the currently available \\chandra\\ and \\xmm\\ data. In the inner 10 kpc a pair of cavities are evident together with a set of bright X-ray filaments. If the cavities are interpreted as gas displaced by relativistic plasma inflated by an AGN, even in the absence of extended 1.4 GHz emission, this would be consistent with a recent outburst as also indicated by the extent of dust and H$\\alpha$ emission. The soft X-ray filaments coincident with H$\\alpha$ and dust emission are cooler than the ones which do not correlate with optical and infrared emission. We suggest that dust-aided cooling contributes to form warm ($T \\sim 10^4$ K) gas, emitting H$\\alpha$ radiation. At 31 kpc and 67 kpc a pair of cold fronts are present, indicative of sloshing due to a dynamical perturbation caused by accretion of a less massive group, also suggested by the peculiar velocity of the brightest galaxy NGC 5044 with respect to the mean group velocity. ",
        "introduction": "\\label{Introduction} The current X-ray observatories, \\chandra\\ and \\xmm, have revolutionized our understanding of the cores of relaxed galaxies, groups and clusters, which show a highly peaked X-ray emission from a hot interstellar medium whose radiative cooling time is less than 1 Gyr \\citep[for recent determinations of cooling times in these objects see for example][]{Voigt.ea:04,Sanderson.ea:06,Jetha.ea:07}. In the absence of heating, a cooling flow is established, in which the gas cools, condenses and flows toward the center, accreting onto the central galaxy \\citep{Fabian:94}. However the mass sink for all this supposedly cooling and condensing gas has never been found \\citep{Donahue.ea:04}. X-ray observations with {\\em Chandra} and {\\em XMM} have established that there is little evidence for emission from gas cooling below $\\sim$ \\tvir/3: just when gas should be cooling most rapidly it appears not to be cooling at all \\citep[see the recent review by][and references therein]{Peterson.ea:06}.  A compensating heat source must therefore resupply the radiative losses, and many possibilities have been proposed, including thermal conduction \\citep[e.g.,][]{Narayan.ea:01}, energy released by mergers \\citep[e.g.,][]{Motl.ea:04} or by supernovae \\citep[e.g.,][]{Silk.ea:86}. However, feedback from the central AGN has become the most appealing solution to the problem \\citep[see the recent review by][and references therein]{McNamara.ea:07}. There is, in fact, clear observational evidence for AGN heating as the majority of brightest central galaxies of clusters and groups host a radio loud AGN \\citep[e.g.,][]{Burns:90,Best.ea:07} and, following the launch of \\chandra, in an increasing number of objects such disturbances as shocks, ripples and cavities have been found in the central atmospheres of clusters, groups and elliptical galaxies \\citep[e.g.,][]{Fabian.ea:06,Birzan.ea:04,Dunn.ea:06,Forman.ea:05,Vrtilek.ea:02,Jetha.ea:07,Croston.ea:08,Finoguenov.ea:01*2,Allen.ea:06}. The cavities, which appear as X-ray surface brightness depressions, have been interpreted as bubbles of low density relativistic plasma inflated by radio jets, displacing the thermal gas causing $PdV$ heating \\citep[e.g.,][]{Churazov.ea:02}, although \\citet{Mathews.ea:08} pointed out that cavity formation contributes thermal energy that may offset radiative cooling only by injecting ultra-hot but non relativistic gas; X-ray cavities formed solely from relativistic gas have a global cooling effect. In most cases, the energy introduced by the AGN is more than sufficient to counteract putative cooling flows \\citep{McNamara.ea:07}, although the physical process of the coupling of the feedback energy with the ambient medium is not well understood. While many of the observed X-ray cavities are filled with plasma emitting in the radio at 1.4 GHz, some are undetected at this frequency and have been referred to as ``ghost cavities''. These may result from the aging of the relativistic particle population and be the signature of a previous AGN outburst. Observations at low radio frequency of several rich clusters with ghost cavities show that they are indeed filled with relativistic plasma \\citep[e.g., A 2597,][]{Clarke.ea:05}, although even low frequency emission cannot be easily detected.  However, some net cooling is probably occurring, as the galaxies at the center of cool cores show properties  not shared by typical elliptical galaxies \\citep{Crawford:04}: 1) the presence of strong, low-ionization emission line nebulae both in clusters \\citep[e.g.,][]{Crawford.ea:99,Conselice.ea:01} and in groups \\citep[e.g.][]{Macchetto.ea:96}; 2) the galaxies at the center of clusters that display these nebular emission show a component of excess ultraviolet/blue continuum associated with young stars \\citep[e.g.,][]{Cardiel.ea:98,Crawford.ea:99}. This excess blue light is again extended on scales of several kpc and there is evidence that some (spatially extended) star formation has been triggered by interaction with the radio source \\citep{Mcnamara:04} 3) CO emission lines have been detected in several cool core clusters at millimeter wavelengths \\citep[e.g.,][]{Edge:01,Salome.ea:03}, implying the presence of a substantial amount of warm molecular gas ($10^{9-11.5}$ M$_\\odot$) within a 50 kpc radius of the central galaxy. Less massive warmer molecular regions have been observed in H$_2$ lines, which are often spatially associated with the stellar UV and H$\\alpha$ emission extending over $\\sim 20$ kpc \\citep{Jaffe.ea:05}.  The origin and excitation mechanism of the nebular H$\\alpha$ emission have been much debated and they are still poorly known. The nebulae require a constant and distributed heating source \\citep[e.g.][]{Johnstone.ea:88} which could be stellar \\citep[e.g. massive OB stars,][]{Allen.ea:92} or the ICM \\citep[e.g. conduction,][]{Sparks.ea:89}. Deep \\chandra\\ X-ray imaging of Perseus \\citep{Fabian.ea:03*1} and M87 \\citep{Sparks.ea:04} have renewed interest in conduction because of the spatial coincidence of filamentary H$\\alpha$ and soft X-ray emission, which can be due to conduction and mixing of the cold gas with the ICM. The disposition of some of the optical filaments in Perseus strongly suggests that the filaments are due to the buoyant radio bubbles drawing out the cold gas from a central reservoir  \\citep{Fabian.ea:03*1} and it reveals the possibility of tracing the flow of the rising bubbles \\citep{Hatch.ea:06}.  But the question is still open if galaxy-galaxy interactions are required to stimulate the central optical nebulae and if the optical filaments are trails of galaxies punching through a molecular hydrogen reservoir as suggested by the examples in \\citet{Wilman.ea:06}. The excitement caused by the shift in the cooling flow paradigm is stimulating also a renewed theoretical interest in the quest for the origin and excitation of molecular \\citep{Ferland.ea:08} and H$\\alpha$ emission \\citep{Nipoti.ea:04,Pope.ea:08,Revaz.ea:08}. In particular \\citet{Nipoti.ea:04} have shown how the H$\\alpha$ nebulae can be stable against thermal evaporation only for the conditions of temperature and pressure found in cool cores, explaining their association with that type of environment.  High resolution \\chandra\\ observations have also revealed another interesting and unexpected feature in the ICM of relaxed clusters: the presence of cold fronts. Many clusters have been found to exhibit sharp arc-shaped jumps in their gas density and temperature which, unlike shock fronts, have the gas on the dense side cooler, so the pressure is continuous across the front \\citep[e.g.,][]{Markevitch.ea:00,Vikhlinin.ea:01}. In merging systems they were interpreted as contact discontinuities between gases from different sub-clusters \\citep{Markevitch.ea:00}. However cold fronts are present in the centers of many, if not most, relaxed clusters with cool cores \\citep[e.g.,][]{Mazzotta.ea:01,Markevitch.ea:03*1,Dupke.ea:03,Ascasibar.ea:06,Ghizzardi.ea:07}. The widely accepted scenario is that cold fronts in these systems are due to sloshing of the cool gas in the central gravitational potential, which is set off by minor mergers/accretions; the only necessary condition is a steep entropy profile as observed in relaxed clusters \\citep{Ascasibar.ea:06}.  Whereas a growing number of clusters and elliptical galaxies have deep enough multi-wavelength data (X-rays, radio and optical) to study the rich phenomenology of cool cores in a detailed spatially resolved fashion, only an handful of groups with such coverage exists, as for example the GEMS objects \\citep{Osmond.ea:04} presented in \\citet{Rasmussen.ea:07}, and therefore, ``unfortunately, AGN heating is not as well studied in groups as in clusters'' \\citep{McNamara.ea:07}. Examination of AGN feedback at the mass scale of groups is valuable because, although the scale of outbursts in groups is less energetic and often on a smaller spatial scale than in clusters, the impact can be even more dramatic than in rich clusters due to the shallower group potential. Statistical studies examining the impact of AGN on groups are starting to address the points raised above \\citep{Croston.ea:05,Jetha.ea:07}. Cold fronts have not been investigated in detail so far in relaxed groups of galaxies.  In this paper we show with more detail the currently available \\chandra\\ and \\xmm\\ data for the galaxy group NGC 5044, one of the brightest groups in X-rays: it is one of the only 5 objects with ${\\rm{k}}T < 2$ keV in the HIFLUGCS sample \\citep{Reiprich.ea:02}. Our estimate, using the new \\chandra\\ and \\xmm\\ data, for the bolometric (0.1-100 keV) X-ray luminosity within $r_{500}=443 h_{70}^{-1}$ kpc \\citep{Gastaldello.ea:07*1} is $1.05\\pm0.06 \\times 10^{43} h_{70}^{-2}$ \\ergsec, in good agreement with the \\rosat\\ estimate presented in \\citet{Reiprich.ea:02} of $2.46 \\times 10^{43} h_{50}^{-2}$ within $560 h_{50}^{-1}$ kpc. In our previous analysis, mainly focused on radial, azimuthally averaged properties of the system \\citep{Buote.ea:03,Buote.ea:03*1}, interesting features like holes and filaments within 10 kpc, a sharp edge, resembling a cold front, at 67 kpc, and unusually low iron abundances at large radii \\citep{Buote.ea:04} were already pointed out. An H$\\alpha$ nebula is present in the core of NGC 5044 showing an extended filamentary structure \\citep{Goudfrooij.ea:94*1,Caon.ea:00}. NGC 5044 is also remarkable because {\\em Spitzer} data show extended cold dust emitting at 70 $\\mu m$ \\citep{Temi.ea:07}. Moreover, \\citet{Temi.ea:07*1} showed extended 8 $\\mu m$ excess (likely arising from PAH, polycyclic aromatic hydrocarbon, molecules) extending out to several kpc and spatially coincident with the H$\\alpha$ emitting nebulosity and the brightest soft X-ray emission. As proposed in \\citet{Temi.ea:07*1}, current evidence is consistent with an internal origin of this dust, which has been buoyantly transported from the galactic core out to several kpc into the hot X-ray emitting gas following an AGN outburst.  All distance-dependent quantities have been computed assuming \\ho = 70 \\kmsmpc, \\omegam = 0.3 and \\omegalambda = 0.7. At the redshift of $z=0.009$ 1\\arcmin\\ corresponds to 11.1 kpc. All the errors quoted are at the 68\\% confidence limit. ",
        "conclusions": "\\label{conclusion}  We have presented results of a two-dimensional analysis of the currently available \\chandra\\ and \\xmm\\ data for the bright nearby galaxy group \\source. The results can be summarized as follows:  \\begin{itemize} \\item A pair of X-ray cavities have been detected, further confirming the recent outburst indicated by the extent and morphology of H$\\alpha$ and dust emission. \\item The presence of cooler filamentary X-ray emission co-spatial with H$\\alpha$ and dust emission strengthens our previous suggestion that dust-aided cooling could contribute to the H$\\alpha$ emission. \\item The presence of a set of two cold fronts together with a peculiar velocity of the central galaxy NGC 5044 suggests a disturbance of an overall relaxed system by an off-axis merger with a smaller satellite. \\end{itemize}  The detection of such a rich phenomenology has been possible due to \\source\\ being bright and nearby, much like on bigger mass scales has been possible with M87 and Perseus. As for these objects up-coming deeper X-ray observations with \\chandra\\ \\citep{David:07} and \\xmm\\ \\citep{Kaastra:07} will shed further light on the properties of \\source\\ and groups of galaxies in general. Multi-wavelength observations, i.e. in H$\\alpha$ and radio bands, are needed for a deeper understanding of the physical processes in the core."
    },
    "0807/0807.4254_arXiv.txt": {
        "abstract": "{ AGILE is an Italian Space Agency mission dedicated to the observation of the gamma-ray Universe. The AGILE very innovative instrumentation combines for the first time a gamma-ray imager (sensitive in the energy range 30 MeV - 50 GeV), a hard X-ray imager (sensitive in the range 18-60 keV) together with a Calorimeter (sensitive in the range 300 keV - 100 MeV) and an anticoincidence system. AGILE was successfully launched on April 23, 2007 from the Indian base of Sriharikota and was inserted in an equatorial orbit with very low particle background.} {AGILE  provides crucial data for the study of Active Galactic Nuclei, GRBs, pulsars, unidentified gamma-ray sources, Galactic compact objects, supernova remnants,  TeV sources, and fundamental physics by microsecond timing.} {An optimal angular resolution (reaching 0.1-0.2 degrees in gamma-rays and 1-2 arcminutes in hard X-rays) and very large fields of view (2.5 sr and 1 sr, respectively) are obtained by the use of Silicon detectors integrated in a very compact instrument.} {AGILE surveyed the gamma-ray sky and detected many Galactic and extragalactic sources during the first months of observations. Particular emphasis is given to multifrequency observation programs of extragalactic and galactic objects.} { AGILE is a successful high-energy gamma-ray mission which reached its nominal scientific performance. The AGILE Cycle-1 pointing program started on 2007 December 1, and is open to the international community through a Guest Observer Program. } ",
        "introduction": "\\label{}    Gamma-ray astrophysics above 100 MeV is an exciting field of astronomical sciences  that received a strong impulse in recent years. Detecting cosmic gamma-ray emission in the energy range from a few tens of MeV to a few tens of GeV is possible only from space instrumentation, and in the past 20 years several space missions confronted the challenge of detecting  cosmic gamma-rays. Gamma-ray emission from cosmic sources at these energies is intrinsically non-thermal, and provides a diagnostic of particle acceleration and  radiation processes in extreme conditions. gamma-ray sources with the most energetic phenomena or objects of our Universe including: Galactic compact objects and related sources (e.g., pulsars, SNRs, pulsar wind nebulae, binary systems with accreting neutron stars and black holes), molecular clouds shining in gamma-rays because of interaction with energetic cosmic rays, extragalactic massive black holes residing in Active Galactic Nuclei (AGN), exploding stars originating Gamma-Ray Bursts (GRBs).   The history of gamma-ray astronomy is a topic beyond the scope of this paper. We briefly summarize here the pioneering efforts and experiments that were carried out over the years. The first detection of cosmic energetic radiation of energy well above the electron's rest mass started in the late Sixties with the detection of gamma-rays above 50 MeV by the OSO-3 satellite \\citep{oso-3}. The first gamma-ray telescope with an angular resolution of a few degrees was launched on November, 1972, by the Malindi site in Kenya as the second Small Astronomy Satellite (SAS-2). A clear concentration of gamma-ray emission from the Galactic plane was firmly established together with a handful of pointlike gamma-ray sources \\citep{fichtel-1975,bignami-1979}. On August 8, 1975 the European mission COS-B was launched and, after 7 years of observations, several tens of gamma-ray sources were studied and catalogued \\citep{mayer-1979,swanenburg-1981}.  The Compton Gamma-Ray Observatory (CGRO) (active between 1991-2000) substantially increased our knowledge of the gamma-ray Universe and provided a wealth of data on a large variety of sources as well as unsolved puzzles. In particular, CGRO hosted the Energetic Gamma-Ray Experiment Telescope (EGRET) operating in the energy range 30 MeV-30 GeV which carried out a complete sky survey detecting hundreds of gamma-ray sources \\citep{fichtel-1997,hartman-1999,th0,th1,th2,th3}. This scientific inheritance is the starting point for any high-energy astrophysics mission.  The new generation of high-energy space missions has to address some of the fundamental issues that were left open (or unresolved) by EGRET. For a mission following CGRO, it  is important to make substantial progress in the instrumentation and mission concept to achieve the following goals: (1) improving the gamma-ray angular resolution near 100 MeV by at least a factor of 2-3 compared to EGRET; (2) obtaining the largest possible field of view (FOV) at 100 MeV reaching 2.5-3 sr; (3) drastically reducing the deadtime for gamma-ray detection from the EGRET value of 100 ms; (4) obtaining broadband spectral information possibly including a simultaneous detection capability in the MeV and X-ray energy ranges; (5) carrying out a rapid quicklook analysis of the gamma-ray data and a fast dissemination of results and alerts; (6) stimulating efficient multifrequency programs. The AGILE Mission is aimed at meeting all these goals with an unprecedented configuration for space gamma-ray astrophysics: a Small Scientific Mission with optimized on-board and ground-segment resources.   The space program AGILE ({\\it Astro-rivelatore Gamma a Immagini LEggero}) is a high-energy astrophysics Mission supported by the Italian Space Agency (ASI) with scientific and programmatic participation by INAF, INFN, CNR, ENEA and several Italian universities \\citep{tavani1,tavani-a,barbiellini-a,tavani-texas,tavani-2}. The main industrial contractors include Carlo Gavazzi Space, Thales-Alenia-Space (formerly Laben), Rheinmetall Italia (formerly Oerlikon-Contraves), Telespazio, Galileo Avionica, and Mipot.  The main scientific goal of the AGILE program is to provide a powerful and cost-effective mission with excellent imaging capability simultaneously in the 30 MeV-50 GeV and 18-60 keV energy ranges with a large field of view that is unprecedented in high-energy astrophysics space missions.    AGILE was successfully launched by the Indian PSLV-C8 rocket from the Sriharikota base on April 23, 2007 (see Fig.~\\ref{lancio}). The launch and orbital insertion were nominal, and a quasi-equatorial orbit was obtained with the smallest inclination (2.5 degrees) ever achieved by a high-energy space mission (e.g., SAS-2 and Beppo-SAX had inclinations of $\\sim4$ degrees). The satellite commissioning phase was carried out during the period May-June, 2007. The scientific Verification phase and the in-orbit calibration (based on long pointings at the Vela and Crab pulsars) were carried out during the period July-November 2007. The nominal scientific observation phase (AGILE Cycle-1, AO-1) started on December 1, 2007. AGILE is the first gamma-ray Mission successfully operating in space after the long temporal hiatus of almost ten years since the end of EGRET operations. It will operate together with the GLAST mission that was launched on June~11,~2008 \\citep{glast}.     \\begin{figure}[t!] \\begin{center} \\includegraphics [height=12cm]{agile-lancio-10.jpg} \\caption { {\\itshape The launch of the AGILE satellite by the Indian PSLV-C8 rocket from the Sriharikota base on April 23, 2007. } } \\end{center} \\label{lancio} \\end{figure}   \\begin{figure}[t!] \\begin{center} \\includegraphics [height=9cm]{agile-IABG-2.jpg} \\caption { {\\itshape The integrated AGILE satellite in its final configuration being covered by the thermal blanket during the qualification tests in IABG (Munich), July 2006. The total satellite mass is equal to 350 kg. } \\label{sat-1}} \\end{center} \\end{figure}     The AGILE instrument design is very innovative and based on solid state Silicon detector technology and state-of-the-art electronics  and readout systems developed in Italian laboratories \\citep{barbiellini-1,barbiellini-2,nina}. The  instrument is light ($\\sim 100$~kg) and very compact (see Fig.~\\ref{PL-1}). The total satellite mass is about 350 kg (see Fig.~\\ref{sat-1}).  AGILE is expected to substantially advance our knowledge in several research areas including the study of Active Galactic Nuclei and massive black holes, Gamma-Ray Bursts (GRBs), the unidentified gamma-ray sources, Galactic transient and steady compact objects, isolated and binary pulsars, pulsar wind nebulae (PWNae), supernova remnants, TeV sources, and the Galactic Center while mapping the overall gamma-ray emission from our Galaxy. Furthermore, the fast AGILE electronic readout and data processing (resulting in detector deadtimes smaller than $\\sim $ 200 $\\mu$sec) allow to perform, for the first time, a systematic search for sub-millisecond gamma-ray/hard X-ray transients that are of interest for both Galactic compact objects (searching outburst durations comparable with the dynamical timescale of $\\sim 1 \\, M_{\\odot}$ compact objects) and quantum gravity studies in extragalactic sources.   The AGILE Science Program is aimed at providing a complete sky coverage during its first year while allowing a prompt response to gamma-ray transients and alert for follow-up multiwavelength observations. AGILE provides important  information  complementary to several high-energy space missions (Chandra, INTEGRAL, RXTE,  XMM-Newton, SWIFT, {\\ch Suzaku}, GLAST) and  supports ground-based investigations in the radio, optical, and TeV bands. Part of the AGILE Science Program is open for Guest Investigations on a competitive basis. Quicklook data analysis and fast communication of new transients has been implemented as an essential part of the AGILE Science Program. ",
        "conclusions": ""
    },
    "0807/0807.4638_arXiv.txt": {
        "abstract": "{ We present a study of \\ion{N}{v} absorption systems at $1.5 \\lesssim z \\lesssim 2.5$ in the spectra of 19 QSOs, based on data obtained with the VLT/UVES instrument. Our analysis includes both the absorbers arising from the intergalactic medium, as well as systems in the vicinity of the background quasar.  We construct detailed photoionization models to study the physical conditions and abundances in the absorbers and to constrain the spectral hardness of the ionizing radiation.  The rate of incidence for intervening \\ion{N}{v} components is $d\\mathcal{N}/dz = 3.38 \\pm 0.43$, corresponding to $d\\mathcal{N}/dX = 1.10 \\pm 0.14$. The column density distribution function is fitted by the slope $\\beta = 1.89 \\pm 0.22$, consistent with measurements of \\ion{C}{iv} and \\ion{O}{vi}. The narrow line widths ($b_\\ion{N}{v} \\sim 6\\,\\mathrm{km\\,s}^{-1}$) imply photoionization rather than collisions as the dominating ionization process. The column densities of \\ion{C}{iv} and \\ion{N}{v} are correlated but show different slopes for intervening and associated absorbers, which indicates different ionizing spectra. Associated systems are found to be more metal-rich, denser, and more compact than intervening absorbers. This conclusion is independent of the adopted ionizing radiation. For the intervening \\ion{N}{v} systems we find typical values of $\\mathrm{[C/H]} \\sim -0.6$ and $n_{\\element{H}} \\sim 10^{-3.6}\\,\\mathrm{cm}^{-3}$ and sizes of a few kpc, while for associated \\ion{N}{v} absorbers we obtain $\\mathrm{[C/H]} \\sim +0.7$, $n_{\\element{H}} \\sim 10^{-2.8}\\,\\mathrm{cm}^{-3}$ and sizes of several 10\\,pc. The abundance of nitrogen relative to carbon [N/C] and $\\alpha$-elements like oxygen and silicon [N/$\\alpha$] is correlated with [N/H], indicating the enrichment by secondary nitrogen. The larger scatter in [N/$\\alpha$] in intervening systems suggests an inhomogeneous enrichment of the IGM. There is an anti-correlation between [N/$\\alpha$] and [$\\alpha$/C], which could be used to constrain the initial mass function of the carbon- and nitrogen-producing stellar population. ",
        "introduction": "Studying elemental abundances in various environments is important for understanding the enrichment history and chemical evolution of astronomical objects. Of particular interest are metal-poor objects, which are thought to be very old and to conserve the abundances pattern of early enrichment. In the local Universe these are, for example, extremely metal-poor halo stars (see \\citet{beerschristlieb2005} for a review). At higher redshifts, the abundance pattern of damped Ly$\\alpha$ (DLA) systems can be studied in detail \\citep[reviewed by][]{wolfeetal2005}. Metal-rich objects like quasars at very high redshift suggest that a substantial amount of metals must have been produced at very early epochs, implying star formation prior to the redshift of observation \\citep[see the review by][]{hamannferland1999}. Measured abundances can then be compared to predicted chemical yields of different stellar populations \\citep[e.g.][]{woosleyweaver1995, iwamotoetal1999, hegerwoosley2002, limongichieffi2003, chieffilimongi2004, umedanomoto2005, kobayashietal2006} optionally weighted by an initial mass function (IMF) to constrain the production site of the heavy elements.  An important tracer of the chemical evolution is the abundance of nitrogen. Nitrogen is created by the conversion of carbon and oxygen during the CNO cycle in hydrogen-burning zones. If the seed carbon and oxygen atoms are produced in the star itself, the nitrogen is called primary. Secondary nitrogen is produced if carbon and oxygen have already been present in the gas that formed the star. Since $\\alpha$-elements (elements built up of $^4$He-cores, e.g.\\ oxygen and silicon) are mainly produced by Type II supernova explosions of massive stars, they are released into the surrounding medium at earlier epochs than carbon and nitrogen, which are mainly created during nuclear burning within massive and intermediate-mass stars. Stars with $M > 8\\,M_{\\sun}$ dominate the production of carbon at early epochs, while stars with lower masses are the main production sites for nitrogen released with a given time lag. Because of this time delay, \\citet{edmundspagel1978} first proposed to use [N/O] as an age indicator. The sensitivity for such an age indicator depends on the duration of the time lag. \\citet{pettinietal2002b} estimate that the $\\sim 250\\,\\mathrm{Myr}$ inferred by \\citet{henryetal2000} are too short to explain all the DLA systems with low [N/$\\alpha$] abundances. However, the time delay increases if stellar rotation is taken into account shifting the production site of primary nitrogen to stars with lower masses and longer evolutionary time scales \\citep{meynetmaeder2002, meynetpettini2004}.  Furthermore, [N/$\\alpha$] traces the production process of nitrogen. While the presence of only primary nitrogen leads to constant [N/$\\alpha$] as found in low-metallicity DLA systems \\citep[e.g.][]{prochaskaetal2002, centurionetal2003, petitjeanetal2008, pettinietal2002b, pettinietal2008}, [N/$\\alpha$] increases with metallicity if secondary nitrogen contributes. The early release of $\\alpha$-elements also leads to an $\\alpha$-enhancement with respect to carbon in the intergalactic medium \\citep[IGM; $\\mathrm{[Si,\\,O/C]} > 0$, e.g.][]{aguirreetal2004, aguirreetal2008, simcoeetal2004, simcoeetal2006}, indicating enrichment by SN II explosions of massive stars \\citep[e.g.][]{matteuccicalura2005, qianwasserburg2005}. Such estimates are based on the detection of \\ion{C}{iv}, \\ion{Si}{iv}, and \\ion{O}{vi} features in QSO absorption spectra whose measured column densities can be transformed into abundances applying appropriate photoionization corrections. However, due to the sparse detection rate of nitrogen features the nitrogen abundance of the high-redshift IGM is rather unknown.  Though absorption features of \\ion{N}{i} and/or \\ion{N}{ii} are usually observed in DLA systems, nitrogen is rarely detected in optically thin (i.e.\\ $N_{\\ion{H}{i}} \\lesssim 10^{17}\\,\\mathrm{cm}^{-2}$) intergalactic absorption systems. Features of \\ion{N}{v} are often identified in absorption systems proximate to the background quasar \\citep[e.g.\\ the recent work of][]{foxetal2008}. Because of its rather high ionization potential ($77.5\\,\\mathrm{eV} = 5.7\\,\\mathrm{Ryd}$ to be created, $97.9\\,\\mathrm{eV} = 7.2\\,\\mathrm{Ryd}$ to be destroyed) this ion is generated by the hard radiation of the close-by QSO. Moreover, quasars are generally metal-rich \\citep[e.g.][]{dietrichetal2003} and about solar [N/H] has been measured in associated systems \\citep[e.g.][]{dodoricoetal2004} making \\ion{N}{v} easily detectable in proximate absorbers. Together with other highly ionized species such as \\ion{O}{vi}, \\ion{N}{v} absorption is also detected in the galactic halo and high-velocity clouds \\citep{savageetal1997, indebetouwshull2004b}. \\citet{foxetal2007} found 3 DLA systems exhibiting weak \\ion{N}{v} features in a sample of \\ion{O}{vi}-bearing DLAs \\citep[see also][]{lehneretal2008}. Recently, \\citet{prochaskaetal2008} report on the detection of \\ion{N}{v} absorption close to long-duration gamma-ray bursts. The generally narrow line widths of the features imply photoionization of the gas by the gamma-ray burst afterglow.  In the IGM the metallicity and therefore the nitrogen content is lower and the ionizing UV background is substantially softer than unfiltered QSO radiation. Thus, \\ion{N}{v} features are expected to be weak and will arise mainly in the Ly$\\alpha$ forest making the detection and identification difficult. One might expect to detect \\ion{N}{v} absorption preferentially when the absorbing material is exposed to a locally hardened radiation field. This may occur in the vicinity of a foreground QSO close to the line of sight generating a so-called transverse proximity effect. Due to the locally higher level of ionization this effect should produce a decrease in the density of neutral hydrogen, which is, however, rarely detected \\citep[e.g.][]{liskewilliger2001, croft2004, schirberetal2004, galleranietal2008}. Yet, combined analyses considering also \\ion{He}{ii} \\citep{jakobsenetal2003, worsecketal2007} and metal line absorption \\citep{goncalvesetal2008} demonstrate the existence of the transverse proximity effect in spectral hardness. Studying the line of sight towards HE~2347-4342 \\citet{worsecketal2007} find a foreground QSO at $z=2.282$ with a transverse proper distance of $1.76\\,\\mathrm{Mpc}$ in the vicinity of the intergalactic \\ion{N}{v} system at $z=2.2753$.  In this paper we present a systematic search for \\ion{N}{v} at $z \\sim 2$ in QSO absorption spectra aiming to probe the nitrogen abundance in the IGM and to test whether \\ion{N}{v} is a tracer for spectral hardness of the ionizing radiation. Therefore we select both intervening \\ion{N}{v} absorbers and systems associated to the background quasar. Their statistical and physical properties are studied and the characteristics of the different absorber classes are extensively compared. After presenting our selection procedure and the resulting sample in Sect.\\ \\ref{sample}, we investigate the statistical properties of \\ion{N}{v} absorption lines like number density and column density distribution in Sect.\\ \\ref{statistics}. In order to estimate elemental abundances we compute photoionization models as described in Sect.\\ \\ref{mod_procedure}. The results are presented in Sect.\\ \\ref{results} and the implications for the enrichment of the IGM and for the presence of hard radiation sources close to \\ion{N}{v} absorption systems are discussed in Sect.\\ \\ref{discussion}. Our conclusions are presented in Sect.\\ \\ref{conclusions}. Throughout the paper we use a cosmology with $\\Omega_{\\mathrm{M}} = 0.3$, $\\Omega_{\\mathrm{\\Lambda}} = 0.7$, and $H_0 = 70\\,\\mathrm{km\\,s}^{-1}\\,\\mathrm{Mpc}^{-1}$. Abundances are given in the notation $\\mathrm{[X/Y]} = \\log (\\mathrm{X/Y}) - \\log (\\mathrm{X/Y})_{\\sun}$ with solar abundances taken from \\citet{asplundetal2005}. ",
        "conclusions": "We have performed a survey for \\ion{N}{v} absorption systems in the spectra of 19 QSOs observed with VLT/UVES. The \\ion{N}{v} systems are selected from 198 intervening and 30 associated metal line systems exhibiting \\ion{C}{iv} features, where absorbers within $|\\Delta v| < 5000\\,\\mathrm{km\\,s}^{-1}$ from the QSO emission redshift are classified as associated.\\\\  (1) In total, 21 intervening systems with 61 individual components and 11 associated systems with 46 individual components are identified in the redshift range $1.5 \\lesssim z \\lesssim 2.5$. Thus, the fraction of \\ion{C}{iv} systems exhibiting \\ion{N}{v} features is roughly 11\\,\\% and 37\\,\\% for intervening and associated absorbers, respectively. The rate of incident for intervening \\ion{N}{v} systems is $d\\mathcal{N}/dz = 3.38\\pm0.43$, corresponding to $d\\mathcal{N}/dX = 1.10\\pm0.14$.\\\\  (2) The column density distribution function can be described by the slope $\\beta = 1.89\\pm0.22$. This slope is consistent with measurements for \\ion{C}{iv} and \\ion{O}{vi}, suggesting a common origin of these species, which is photoionized gas. Further evidence for photoionization is provided by the narrow line width of the \\ion{N}{v} components ($b_{\\ion{N}{v}} \\sim 6\\,\\mathrm{km\\,s}^{-1}$ for the intervening systems), implying temperatures of $T \\sim 3\\cdot10^{4}\\,\\mathrm{K}$ too low for collisional ionization that produces \\ion{N}{v} at much higher temperatures of $T \\sim 2\\cdot10^{5}\\,\\mathrm{K}$. The \\ion{N}{v} features in the individual systems are detected at the same position as \\ion{C}{iv} while \\ion{O}{vi} is shifted in velocity space to slightly higher values, typically by $v_{\\ion{O}{vi}} - v_{\\ion{C}{iv}} \\simeq +4\\,\\mathrm{km\\,s}^{-1}$. Therefore, \\ion{N}{v} is apparently in the same gas phase as \\ion{C}{iv} despite its higher ionization potential.\\\\  (3) The column densities of the \\ion{N}{v} and \\ion{C}{iv} components are closely correlated with different slopes for intervening and associated absorbers where intervening systems exhibit generally weaker \\ion{N}{v} features than associated systems. The different slopes indicate a different spectral energy distribution of the ionizing radiation. While associated absorbers are exposed to the hard spectrum of the background QSO, the intervening systems are ionized by the substantially softer radiation of the intergalactic UV background \\citep[][HM01]{haardtmadau2001}. In conclusion, intervening \\ion{N}{v} features arise in rather \"normal\" intergalactic metal absorption systems and do not trace locally hardened radiation. This interpretation is supported by extensive photoionization modeling testing several spectral energy distributions. No indications for ionizing spectra substantially harder than the HM01 background are found in case of intervening systems.\\\\  (4) Detailed photoionization models are derived for each system in order to estimate the elemental abundances of the absorbers. We estimate a systematic error of roughly $0.7\\,\\mathrm{dex}$ for the individual metallicities of the associated absorbers due to the uncertain spectral energy distribution of the background QSO. The resulting metallicities for intervening systems are typically low with a median of $\\mathrm{[C/H]} = -0.63$. $\\alpha$-elements like oxygen and silicon are enhanced yielding median values $[\\alpha/\\element{C}] = +0.41$ and $[\\element{N}/\\alpha] = -0.58$, respectively. The typical density is $n_{\\element{H}} \\simeq 10^{-3.6}\\,\\mathrm{cm}^{-3}$ and sizes of a few to several 10\\,kpc with a median $\\sim 5\\,\\mathrm{kpc}$ are found. In contrast, associated systems are more metal-rich independently of the adopted radiation field, having solar metallicity or even higher (median $\\mathrm{[C/H]} = +0.66$) and nearly solar relative abundances ($[\\alpha/\\element{C}] = -0.29$ and $[\\element{N}/\\alpha] = -0.05$, respectively). These absorbers are denser, yielding $n_{\\element{H}} \\gtrsim 10^{-2.8}\\,\\mathrm{cm}^{-3}$, and about one order of magnitude smaller with typical sizes of 10 to a few 100\\,pc.\\\\  (5) For both intervening and associated systems [N/C] and [N/$\\alpha$] are correlated with metallicity where nitrogen is more abundant in metal-rich absorbers. The increase of [N/$\\alpha$] with [N/H] suggests that the nitrogen content is dominated by secondary nitrogen. The best-fit slope is $\\sim 0.5$ consistent with abundance measurements in extragalactic \\ion{H}{ii} regions and metal-poor halo stars. The values inferred for intervening systems show a larger scatter around this relation, indicating more inhomogeneous enrichment of the IGM. With our sample we only probe systems with $\\mathrm{[N/H]} \\gtrsim -2.0$. In order to detect nitrogen in lower-metallicity, intervening systems \\ion{N}{iv} has to be surveyed whose main transition would arise in the UV at $z=2$.\\\\  (6) The relative abundances [$\\alpha$/C] and [N/$\\alpha$] trace the overall metallicity in the sense that metal-poor systems show enhanced $\\alpha$-abundances compared to carbon and nitrogen. [N/$\\alpha$] is anti-correlated with [$\\alpha$/C]. A linear fit yields $[\\element{N}/\\alpha] = -(1.04 \\pm 0.13)\\cdot [\\alpha/\\element{C}] - (0.15 \\pm 0.08)$. This correlation is consistent with abundance measurements in metal-poor halo stars and extragalactic \\ion{H}{ii} regions, whereas the most metal-poor DLA systems do not fit the relation, probably reflecting the difference between primary and secondary nitrogen enrichment. Since the investigated sample is selected by \\ion{N}{v}, the systems are expected to be rather nitrogen-rich and therefore dominated by secondary nitrogen. The relation between [N/$\\alpha$] and [$\\alpha$/C] could be used in future studies to constrain the initial mass function of the carbon-and nitrogen-producing stellar population.\\\\"
    },
    "0807/0807.3829_arXiv.txt": {
        "abstract": "We derive  the Fundamental  Plane (FP) relation  for a sample  of 1430 early-type galaxies in  the optical (r band) and  the near-infrared (K band),  by  combining  SDSS  and  UKIDSS data.   With  such  a  large, homogeneous dataset, we are able to assess the dependence of the FP on the  waveband.   Our  analysis  indicates  that  the  FP  of  luminous early-type  galaxies  is essentially  waveband  independent, with  its coefficients increasing at most by  $8\\%$ from the optical to the NIR. This finding fits well into a consistent picture where the tilt of the FP  is not  driven  by  stellar populations,  but  results from  other effects, such as non-homology.  In this framework, the optical and NIR FPs require more massive galaxies  to be slightly more metal rich than less massive ones,  and to have highly synchronized  ages, with an age variation per decade in mass smaller than a few percent. ",
        "introduction": "\\label{sec:INTR}  Early-type galaxies  (ETGs) populate a two-dimensional  surface in the space  of parameters  that  reflect size  (effective radius),  density (mean   surface  brightness),   and   kinetic  temperature   (velocity dispersion) \\citep{George92}.  A key  feature of the Fundamental Plane (FP;   \\citealt{Dressler87};  \\citealt{George87})  is   its  deviation (`tilt')  from the  virial  theorem,  which may  be  interpreted as  a variation  of the  M/L ratio  along the  sequence of  ETGs  and/or the breaking  of  homology  assumption,  i.e.   the  fact  that,  for  all galaxies, the  observed parameters have the  same power-law dependence on the corresponding physical quantities (namely, the central velocity dispersion on kinetic energy,  the effective radius on `gravitational' radius,  and the  effective surface  brightness on  the  overall light profile,  see e.g.~\\citealt{DjS93}).   Despite  all the  observational efforts, the origin of the tilt  is still under debate.  The change of the  M/L ratio  can be  explained by  a change  in either  the stellar population (e.g.~\\citealt{PrS96})  or dark matter  content with galaxy mass  (\\citealt{CLR96}). Both  structural  and dynamical  non-homology have also been  invoked as physical explanations of  the observed tilt (see            e.g.~\\citealt{HjM95},~\\citealt{CdC95},~\\citealt{GrC97}, and~\\citealt{BCC97}).  Recently, ~\\citet[from now TBB04]{Tru04} showed that  the   tilt  is  mostly   driven  by  dynamical   and  structural non-homology,  while  stellar populations  account  for  only a  small fraction  of it.~\\citet{BBT07}  argued that  the tilt  is  more likely because of  a variation  of the dark  matter content with  mass, still favoring a picture where stellar populations play a minor role.  Since the contribution of different stellar populations to galaxy luminosity is  expected to  be  wavelength-dependent, while  other effects  (e.g. non-homology) are not, the dependence of the FP on wavelength directly informs on how properties of the stellar populations change with mass, which is a crucial point to understand galaxy formation and evolution.  Previous  studies of  the wavelength  dependence led  to contradictory results.~\\citet{PDdC98b}   and~\\citet{SCO98}   found   the   tilt   to significantly decrease  from optical to  NIR wavelengths, interpreting this  result  as  an  increase  of  age  and  metallicity  with  mass. ~\\citet{MGA99} and~\\citet{ZGS02}  found only  a small decrease  of the tilt with wavelength, with the  FP still being significantly tilted in the NIR.  However, ~\\citet{BER03b},  deriving the FP for ETGs observed in the Sloan  Digital Sky Survey (SDSS), found evidence  for the FP to be  wavelength independent  from the  $g$ to  the $z$  bands.  Several different  effects  can produce  this  puzzling  picture.   The FP  in different  wavebands has often  been derived  for small  samples, with inhomogeneous measurements  of galaxy parameters,  different selection criteria (e.g.  galaxy samples  spanning different ranges in magnitude and/or velocity dispersion), and  with different fitting methods.  The FP by~\\citet{BER03b} avoided all  these problems by analyzing the same sample of galaxies  at different wavebands, but it  was limited to the short wavelength baseline provided by  the SDSS.  In the present work, for  the first  time,  we derive  the  FP by  using  the same,  large, homogeneous sample of ETGs  over the wide wavelength baseline provided by the r- and K-band data of  the SDSS and the UKIRT Infrared Deep Sky Survey (UKIDSS).  The layout of  the paper is the following. In  Sec.~2, we describe the selection of the sample, while Sec.~3 details how we obtain the r- and K-band    structural   parameters,    and    the   central    velocity dispersions. Sec.~4  deals with  the comparison of  the r-  and K-band FPs. In  Sec.~5, we  show how  the optical and  NIR FPs  constrain the variation of stellar population  parameters along the galaxy sequence. The discussion follows  in Sec.6.  Throughout the paper,  we adopt the cosmology  $\\rm H_0  \\!   = \\!   75  \\, km  \\,  s^{-1} \\, Mpc^{-1}$, $\\Omega_{\\rm m} \\!  = \\!  0.3$, and $\\Omega_{\\Lambda} \\!  = \\!  0.7$. ",
        "conclusions": "\\label{sec:DISC}  This work presents the waveband  dependence of the FP by comparing the optical and NIR  FPs for a large sample  of galaxies, with homogeneous measurements of structural  parameters and velocity dispersions.  This is allowed,  for the  first time, thanks  to the availability  of both SDSS r-band photometry and  spectroscopy, and UKIDSS K-band photometry for the same sample of ETGs.  Such a dataset, together with the use of the same  fitting procedure in  both bands, makes our  study virtually free from any methodological effect  on the waveband dependence of the FP.  Our analysis shows that the FP does not change significantly from the  optical to  the  NIR, bringing  interesting  questions about  the nature of the sequence of ETGs.  In  Sec.~5, we  have shown  how the  r- and  K-band FPs  constrain the variation of  the stellar population properties  (age and metallicity) with stellar mass, and how  such a constraint is strongly dependent on the fraction,  $f$, of the  FP tilt resulting from  stellar population effects.  Previous studies of the color-magnitude (CM) relation and of line-strength indices of ETGs might  help us to solve this dependency, by deriving  the proper value  for $f$. \\citet{Kod98} showed  that the little  redshift evolution  of the  CM  relation implies  (i) all  the (luminous) ETGs to be equally old and (ii) more massive galaxies to be more metal  rich, with a metallicity  change of $\\delta  (\\log Z) \\sim 0.22$dex per  decade in stellar  mass.  This finding  is qualitatively consistent  with  that  of~\\citet{THO05},  who  found  absorption-line indices consistent with a metallicity  change of $\\delta (\\log Z) \\sim 0.12$dex  per  mass decade.   They  also  found  evidence for  an  age gradient along the sequence of ETGs, with $\\delta (\\log t) = +0.05 \\pm 0.07$  (see their  eq.3).   Fig.~\\ref{fig:DLOGT_DLOGZ} compares  these values of $\\delta (\\log Z)$ with those derived by the optical--NIR FP. For the BrC03  (Vaz96) model, the maximum value  of $\\delta (\\log Z)$, which is consistent with the  FP, amounts to $0.06 $($0.09$) $\\pm 0.04 $($0.04$)dex for $f=0$. This value  is $4 \\sigma$ ($3.7 \\sigma$) lower than that  derived by the CM  relation, while it is  only $1.5 \\sigma$ ($0.8  \\sigma$)  lower  than  that  found  by~\\citet{THO05}.   As  $f$ increases, the FP requires the value of $\\delta (\\log Z)$ to decrease, making  the  above differences  even  larger.  Therefore,  reconciling previous estimates  of $\\delta (\\log Z)$  with our results  leads to a scenario where  $f=0$, which means that  the FP tilt is  not driven by stellar  populations.    We  have   to  remark,  however,   that  this interpretation is  troublesome, since  galaxy colors and  line indices are always measured within a given fixed aperture, and the presence of internal population gradients in galaxies can significantly affect the inferred  values   of  $\\delta  (\\log   Z)$  and  $\\delta   (\\log  t)$ (e.g.~\\citealt{Scodeggio:01}).  Further constraints come from  previous works addressing the origin of the  FP tilt  itself.   Performing a  detailed  dynamical analysis  of twenty-five galaxies,  \\citet{Cap06} derived \\ml  \\, ratios consistent with  those obtained  from the  virial  theorem in  the assumption  of homology, concluding  that structural and orbital  non-homology have a negligible role in the tilt  of the FP (see also~\\citealt{ZZG08}).  In support  to this  view, they  also showed  that the  variation  of the dynamical \\ml \\, is  correlated with the H$_\\beta$ line-strength, thus ascribing   most   of   the   tilt   to   stellar   population   (age) effects. However, as the  authors notice, this result strictly applies to their measurement of the velocity dispersion as the average over an aperture of radius equal to  $r_e$, a fact that alone might compensate part of the dynamical non-homology.  Moreover, most of the galaxies in their sample (68\\%) are fast  rotators, while five of them (20\\%) have low  velocity  dispersion  ($\\sigma$=60-85  km  s$^{-1}$).   As  found by~\\citet{ZGZ06}  and~\\citet{DOF08}, bright  and faint  spheroids have different FPs, with  the tilt becoming larger for  galaxies having low velocity dispersion.  Hence, the different selection of our sample and that   of~\\citet{Cap06}   prevents   a   straightforward   comparison. \\citet{BBT07} showed  that by  replacing mean surface  brightness with mass density,  the FP relation  closely approaches the  virial theorem expectation, implying that  most of the tilt is  caused by a variation of dark  matter content with galaxy mass.   However, the uncertainties on their  FP coefficients  and the possible  biases introduced  by the gravitational-lens   selection  do   not   definitively  exclude   the contribution of non-homology  to the FP tilt.  Our  result agrees with \\citet{BBT07} regarding  the minor role played  by stellar populations on  the tilt.   TBB04,  agreeing with~\\cite{BCC97}  and~\\citet{GrC97}, found that structural and  dynamical non-homology can account for more than  two-thirds  of  the  FP  tilt, with  the  remaining  part  being explained by  stellar population effects.   In particular, restricting the analysis to the  magnitude-complete subsample, they found that the contribution of  stellar populations  to the tilt  becomes negligible. Notice also  that their FPs  are derived from different  sources, with significantly different coefficients in the optical and NIR wavebands.  To understand  how our  results may be  affected by the  assumption of homology,  we followed  an approach  similar to  that of  TBB04, using spherical,  isotropic,  non  rotating,  one-component models  of  ETGs following the  Sersic law (see~\\citealt{LBC05}).  For  each galaxy, we considered  the  model with  the  corresponding  Sersic  index in  the r-band, and used that model to correct the central velocity dispersion to the  quantity $\\sigma_{\\rm t}$  (defined as the square-root  of the total  specific   kinetic  energy),   the  effective  radius   to  the gravitational radius,  $r_{\\rm g}$, and to calculate  the mean surface brightness within the gravitational radius, $ <\\!  \\mu \\!  >_{\\rm g}$. Applying the orthogonal  fit, we obtain the following  equation of the FP in the K band: $\\log  r_{\\rm g} \\propto (2.3 \\pm 0.2) \\log \\sigma_t + (0.4 \\pm 0.02) <\\!  \\mu \\!  >_{\\rm g}$, which is remarkably close to the virial theorem expectation, implying that non-homology may account for the entire tilt.  To  explore how this result would affect stellar population properties, we normalized the r- and K-band FP coefficients in such  a way  to match the  virial theorem expectation  ($a=2.0$ and $b=0.4$) in the K band.   Then, we derived the corresponding values of \\dz \\,  and \\dt  \\, (see Sec.~\\ref{sec:ANALYSIS}).   As we can  see in Fig.~\\ref{fig:DLOGT_DLOGZ}, accounting  for non-homology leads  to the same \\dz \\, and \\dt \\, values as those derived under the assumption of homology if the tilt of the FP is not due to stellar populations (i.e. $f=0$).   As discussed above,  this is  consistent with  what expected from the color-magnitude relation and absorption-line indices of ETGs.   In summary, our  analysis suggests a consistent picture  where (i) the tilt of the FP does  not originate from stellar population effects but is due  to other effects,  such as non-homology; (ii)  the SDSS-UKIDSS FPs require  more massive galaxies to  be mildly more  metal rich than less massive  systems, and to  have extremely synchronized  ages, with the age variation per mass decade being smaller than few percent."
    },
    "0807/0807.1298_arXiv.txt": {
        "abstract": "We discover a large-scale shell G53.9+0.2 around the Crab-like pulsar wind nebula (PWN) G54.1+0.3 with 1420 MHz continuum VLA observations. This is confirmed by a new infrared image at 8 $\\mu$m from the GLIMPSE Legacy Project, which reveals an intriguing infrared shell just surrounding the large radio shell. We analyze the 21 cm HI absorption spectra and $^{13}$CO emission spectra towards PWN G54.1+0.3 and bright sources on both radio and IR shells. Continuous HI absorption up to the tangent point and absence of negative HI absorption features imply that PWN G54.1+0.3 has a distance beyond the tangent point but within the Solar circle, i.e. 4.5 to 9 kpc. G54.1+0.3 is likely at distance of $\\simeq$6.2 kpc due to the morphological association between the PWN and a CO molecular cloud at velocity of $\\simeq$53 km/s, as revealed by high-resolution $^{13}$CO images. Based on the HI absorption spectrum and recombination line velocity ($\\simeq$40 km/s) of the bright HII region G54.09-0.06, which is on the IR shell, the IR shell is likely located at a distance of $\\simeq$7.3 kpc, which is also the distance of the associated large-scale radio shell. At this distance, the radio shell has a radius of $\\sim$30 pc. The radio shell  may be thermal and lack IR emission due to dust destruction, or it may be nonthermal and part of a newly found old SNR. In either case it is located at a different distance than PWN G54.1+0.3. ",
        "introduction": "PWN G54.1+0.3 is a Crab Pulsar Wind Nebula (PWN) analog, with a young central pulsar PSR J1930+1852 spinning down to energize its PWN (Velusamy \\& Becker, 1988). The Crab Nebula, so far, despite intensive searches (Frail et al. 1995) shows no evidence for an outer shock (i.e. supernova remnant (SNR) shell) from the supernova explosion which produced the pulsar. No satisfactory explanation has been proposed for this lack of an outer SNR shell. Like the Crab, G54.1+0.3 also, as yet, has no identified outer SNR shell. Here we search for such a shell for G54.1+0.3. We find an SNR-like shell around G54.1+0.3 and then study its properties to determine if it is associated with G54.1+0.3.  The distance of PWN G54.1+0.3 is important for determination of the luminosity and energetics of the pulsar and its PWN and for comparison with other PWN. The distance has been estimated by several times previously. Based on Chandra X-ray observations, Lu et al. (2002) estimated an absorption column density of about half the total Galactic absorption in the direction to G54.1+0.3, and therefore suggested a distance of $\\sim$5 kpc. From an HI  absorption spectrum, Weisberg et al. (2008) found PSR J1930+1852, which is the central pulsar of PWN G54.1+0.3, is at a distance of $>$5 kpc. Using an improved free-electron density/distance model (Cordes \\& Lazio 2002), Camilo et al. (2002) gave a DM distance of $\\le$8 kpc to PSR J1930+1852. % The distance uncertainty on the system is thus still large, and improved distance determination to the PWN system is needed. Recently, we have revised distances to several SNR/PWN or SNR/HESS source or SNR/CO cloud interaction systems with good precision by analyzing HI+CO spectra of these system (Leahy \\& Tian 2008; Tian, Leahy \\& Wang 2007; Tian et al. 2008). Here, we use these effective methods to constrain the distance to the PWN G54.1+0.3/ PSR J1930+1852 system, and to revise its X-ray luminosity. ",
        "conclusions": "Here we have searched for a SNR shell counterpart to the PWN G54.1+0.3 using radio continuum, HI and CO line emission, and infrared observation. In addition we have used HI absorption and CO emission data to constrain distances to PWN G54.1+0.3 and other nearby sources. We determine a distance of $\\simeq$6.2 kpc to G54.1+0.3 using an associated CO cloud, consistent with the HI absorption spectrum which limits G54.1+0.3 to be between the tangent point and the far side of the solar circle. In addition, we find a SNR-like shell in 1420 MHz continuum emission, G53.9+0.2, which is offset from G54.1+0.3. The radio shell has no thermal IR emission from its shell but is surrounded by an IR shell with an embedded HII region and associated CO clumps, both of which are located at 7.3 kpc distance. The radio shell may be thermal and lack IR emission due to dust destruction, or it may be nonthermal and part of a newly found old SNR. In either case it is located at a different distance than PWN G54.1+0.3."
    },
    "0807/0807.1251_arXiv.txt": {
        "abstract": "The recurrent nova RS Ophiuchi, one of the candidates for Type Ia supernova progenitors, underwent the sixth recorded outburst in February 2006. We report a complete light curve of supersoft X-ray that is obtained for the first time. A numerical table of X-ray data is provided. The supersoft X-ray flux emerges about 30 days after the optical peak and continues until about 85 days when the optical flux shows the final decline. Such a long duration of supersoft X-ray phase can be naturally understood by our model in which a significant amount of helium layer piles up beneath the hydrogen burning zone during the outburst, suggesting that the white dwarf mass is effectively growing up.  We have estimated the white dwarf mass in RS Oph to be $1.35 \\pm 0.01~ M_\\odot$ and its growth rate to be  about $(0.5-1) \\times 10^{-7} M_\\odot$ yr$^{-1}$ in average. ",
        "introduction": "RS Oph is one of the well-observed recurrent novae and is suggested to be a progenitor of Type Ia supernova. It has undergone its sixth recorded outburst on 2006 February 12 \\citep{nar06} and many observational results are reported (see other papers in this proceedings). It's $y$ magnitude light curve has been obtained throughout the outburst \\citep{hac06} [the numerical table is provided in \\citet{hac08} in this volume], which shows a mid-plateau phase that lasts 45-75 days from the optical peak followed by a quick decrease (Figure 1).  RS Oph has also been observed with X-ray satellites. We analyzed the {\\it Swift} XRT observations available in the HEAS-ARC' database and extracted the count rate in the energy band of 0.3-0.55 keV binned in 2000 s. \\citep[see][for more details]{hac07}. The supersoft X-ray (SSX) light curve is plotted in Figure 1 (see also Table 1). The light curve rises at about 30 days after the optical peak and shows a long plateau phase that lasts as long as about 50 days corresponding to a long mid-plateau phase of optical light curve. ",
        "conclusions": "We summarize our results as follows;  \\noindent 1. The WD mass is estimated to be $1.35 \\pm 0.01~ M_\\odot$ from the optical and X-ray light curve fittings.  \\noindent 2. The accreted matter during  21 years before the outburst is estimated from the envelope mass at the optical peak to be $4 \\times 10^{-6}M_{\\odot}$, $50-70\\%$ of which is ejected by the wind and the rest $30-50\\%$ accumulates on the WD. Therefore the net growth rate of the WD is $(0.5-1) \\times 10^{-7}M_{\\odot}$yr$^{-1}$.  \\noindent 3. The durations of the mid-plateau phase of optical and the peak plateau phase of SSX suggest the presence of a helium layer which accumulates on the WD. Therefore, the WD mass of RS Oph is now growing.  \\noindent 4. The hydrogen content is expected to decrease by $\\Delta X = 0.2-0.4$ during the outburst, because convection may quickly descend and unable to mix the helium ash into hydrogen-rich material.   \\begin{table}[!ht] \\label{fonts} \\caption{Supersoft X-ray count rates} \\smallskip \\begin{center} {\\tiny \\begin{tabular}{llllllllllll} \\tableline \\noalign{\\smallskip}  Time &Count &Time& Count & Time& Count & Time &Count& Time &Count& Time &Count  \\\\  & rate& &  rate& & rate&  & rate&  & rate & & rate \\\\  [day]  & [s$^{-1}$]  &  [day]  & [s$^{-1}$] &[day]  & [s$^{-1}$]&[day]  & [s$^{-1}$]&[day]  & [s$^{-1}$]& [day] & [s$^{-1}$]\\\\ \\noalign{\\smallskip} \\tableline \\noalign{\\smallskip} 3.53 &   -0.20 &  36.74 &    1.64 & 40.52 &    1.56 & 47.09 &    2.07 & 52.67 &    2.10 & 64.59 &    2.00\\\\ 11.33 &   -0.69 &  36.77 &    1.59 & 40.54 &    1.47 & 47.16 &    2.10 & 52.74 &    2.08 & 66.12 &    2.00\\\\ 11.40 &   -0.95 &  36.79 &    1.54 & 40.59 &    1.52 & 47.37 &    2.12 & 52.81 &    2.06 & 66.19 &    2.00\\\\ 11.47 &   -0.78 &  36.81 &    1.47 & 40.61 &    1.53 & 47.42 &    2.08 & 52.86 &    2.09 & 66.26 &    1.98\\\\ 13.94 &   -0.93 &  36.84 &    1.50 & 40.68 &    1.72 & 47.44 &    2.14 & 52.93 &    1.99 & 67.12 &    1.99\\\\ 15.96 &   -0.86 &  36.86 &    1.49 & 41.07 &    1.12 & 47.49 &    2.10 & 52.97 &    2.09 & 67.14 &    1.99\\\\ 18.53 &   -1.02 &  36.91 &    1.43 & 41.12 &    1.62 & 47.51 &    2.12 & 52.99 &    2.11 & 67.18 &    1.99\\\\ 18.57 &   -0.80 &  36.93 &    1.46 & 41.14 &    1.70 & 47.83 &    2.10 & 53.04 &    2.11 & 67.21 &    1.95\\\\ 26.33 &   -0.51 &  36.98 &    1.22 & 41.19 &    1.74 & 47.90 &    2.12 & 53.06 &    2.12 & 67.25 &    1.99\\\\ 26.35 &   -0.51 &  37.00 &    1.20 & 41.21 &    1.74 & 48.02 &    2.09 & 53.18 &    2.10 & 67.28 &    1.94\\\\ 29.34 &    1.07 &  37.05 &    1.11 & 41.26 &    2.00 & 48.04 &    2.09 & 53.20 &    2.11 & 67.32 &    1.97\\\\ 29.36 &    1.06 &  37.12 &    1.37 & 41.28 &    2.02 & 48.09 &    2.09 & 53.25 &    2.09 & 67.35 &    1.86\\\\ 30.22 &    0.48 &  37.14 &    1.39 & 41.33 &    2.04 & 48.11 &    2.12 & 53.27 &    2.09 & 67.51 &    2.02\\\\ 30.24 &    0.59 &  37.18 &    1.39 & 41.35 &    2.04 & 48.16 &    2.08 & 53.32 &    2.12 & 67.53 &    2.02\\\\ 32.35 &    1.21 &  37.21 &    1.41 & 41.40 &    2.04 & 48.18 &    2.13 & 53.34 &    2.14 & 68.67 &    1.99\\\\ 32.37 &    1.19 &  37.25 &    1.58 & 41.42 &    2.07 & 48.23 &    2.14 & 53.39 &    1.79 & 68.74 &    1.96\\\\ 33.16 &    1.78 &  37.32 &    1.52 & 41.47 &    1.97 & 48.30 &    2.12 & 53.41 &    1.80 & 69.66 &    1.96\\\\ 33.25 &    1.86 &  37.39 &    1.64 & 42.46 &    2.09 & 48.37 &    2.11 & 53.46 &    2.10 & 69.68 &    1.96\\\\ 33.30 &    1.95 &  37.42 &    1.56 & 42.53 &    2.03 & 48.43 &    2.10 & 53.48 &    2.10 & 69.73 &    1.97\\\\ 33.32 &    1.88 &  37.46 &    1.81 & 43.13 &    1.90 & 48.50 &    2.07 & 53.53 &    2.09 & 69.75 &    1.95\\\\ 33.37 &    1.88 &  37.49 &    1.79 & 43.20 &    1.87 & 48.57 &    2.13 & 53.55 &    2.06 & 70.59 &    1.95\\\\ 33.39 &    1.94 &  37.53 &    1.81 & 43.27 &    1.98 & 48.90 &    2.10 & 53.60 &    2.04 & 70.61 &    1.97\\\\ 33.43 &    1.94 &  37.60 &    1.84 & 43.34 &    1.97 & 49.04 &    2.10 & 53.62 &    1.97 & 70.66 &    1.96\\\\ 33.46 &    1.94 &  37.67 &    1.92 & 43.41 &    1.92 & 49.11 &    2.09 & 53.67 &    2.05 & 70.68 &    1.96\\\\ 33.50 &    1.85 &  37.74 &    1.85 & 43.48 &    1.78 & 49.18 &    2.09 & 53.69 &    1.99 & 72.35 &    1.90\\\\ 33.53 &    1.88 &  37.79 &    1.88 & 43.55 &    1.94 & 49.22 &    2.08 & 53.74 &    1.88 & 72.37 &    1.87\\\\ 33.57 &    1.92 &  37.81 &    1.50 & 43.94 &    1.69 & 49.24 &    2.10 & 53.78 &    2.03 & 72.42 &    1.89\\\\ 33.60 &    1.86 &  38.13 &    2.10 & 44.01 &    1.91 & 49.29 &    2.05 & 53.80 &    2.01 & 72.49 &    1.88\\\\ 33.64 &    1.18 &  38.32 &    2.13 & 44.08 &    1.91 & 49.31 &    2.07 & 53.85 &    1.79 & 73.30 &    1.91\\\\ 33.67 &    0.96 &  38.34 &    2.11 & 44.15 &    1.81 & 49.36 &    2.07 & 53.87 &    2.05 & 73.34 &    1.89\\\\ 33.71 &    0.37 &  38.39 &    2.14 & 44.20 &    1.67 & 49.38 &    2.08 & 53.92 &    2.09 & 73.37 &    1.87\\\\ 33.74 &    0.34 &  38.41 &    2.11 & 44.22 &    1.76 & 49.43 &    2.06 & 53.94 &    2.03 & 73.41 &    1.89\\\\ 33.78 &    0.27 &  38.46 &    2.10 & 44.27 &    1.91 & 49.45 &    2.04 & 53.99 &    2.10 & 73.43 &    1.86\\\\ 33.90 &    1.34 &  38.48 &    2.07 & 44.29 &    1.90 & 49.50 &    2.06 & 54.01 &    2.11 & 74.22 &    1.82\\\\ 33.97 &    1.66 &  38.53 &    1.97 & 44.34 &    1.43 & 49.52 &    2.06 & 54.06 &    2.08 & 74.29 &    1.79\\\\ 34.04 &    1.94 &  38.60 &    1.92 & 44.36 &    1.73 & 49.82 &    2.10 & 54.13 &    1.91 & 74.36 &    1.80\\\\ 34.11 &    1.98 &  38.67 &    2.02 & 44.41 &    1.93 & 49.85 &    2.12 & 54.18 &    2.08 & 74.43 &    1.87\\\\ 34.18 &    1.99 &  38.74 &    2.05 & 44.43 &    1.97 & 50.03 &    2.12 & 54.20 &    2.12 & 74.48 &    1.87\\\\ 34.20 &    2.01 &  38.78 &    2.02 & 44.48 &    1.95 & 50.05 &    2.13 & 54.24 &    1.78 & 74.50 &    1.88\\\\ 34.24 &    1.68 &  38.85 &    1.54 & 44.50 &    1.97 & 50.10 &    2.09 & 54.27 &    1.78 & 75.17 &    1.87\\\\ 34.27 &    1.56 &  38.92 &    0.84 & 44.55 &    1.99 & 50.17 &    2.08 & 54.31 &    2.09 & 75.22 &    1.86\\\\ 34.31 &    1.28 &  38.99 &    0.92 & 45.15 &    1.92 & 50.24 &    2.07 & 54.34 &    2.10 & 75.24 &    1.85\\\\ 34.34 &    1.45 &  39.06 &    1.33 & 45.17 &    1.90 & 50.31 &    2.04 & 54.41 &    2.10 & 77.23 &    1.81\\\\ 34.36 &    1.57 &  39.13 &    1.70 & 45.22 &    2.09 & 50.38 &    2.09 & 54.45 &    2.11 & 77.25 &    1.79\\\\ 34.38 &    1.69 &  39.34 &    1.55 & 45.29 &    2.11 & 50.52 &    2.16 & 54.48 &    2.12 & 77.30 &    1.81\\\\ 34.43 &    1.90 &  39.41 &    1.01 & 45.36 &    1.98 & 51.51 &    2.12 & 57.60 &    2.03 & 77.37 &    1.81\\\\ 34.45 &    1.95 &  39.48 &    0.88 & 45.43 &    2.12 & 51.54 &    2.11 & 57.62 &    2.05 & 79.15 &    1.74\\\\ 34.50 &    2.07 &  39.52 &    1.19 & 45.47 &    1.90 & 51.84 &    2.14 & 58.62 &    2.03 & 79.18 &    1.74\\\\ 34.52 &    2.04 &  39.55 &    1.20 & 45.49 &    1.85 & 51.91 &    2.11 & 58.64 &    1.99 & 80.52 &    1.67\\\\ 35.05 &    1.86 &  39.59 &    1.24 & 45.56 &    2.14 & 51.93 &    2.12 & 59.15 &    2.11 & 80.56 &    1.68\\\\ 36.26 &    1.51 &  39.62 &    1.27 & 45.82 &    2.09 & 51.98 &    2.12 & 59.87 &    2.07 & 80.59 &    1.68\\\\ 36.30 &    1.51 &  39.66 &    1.16 & 45.89 &    2.00 & 52.05 &    2.11 & 59.89 &    2.08 & 80.63 &    1.67\\\\ 36.33 &    1.52 &  39.68 &    1.08 & 46.03 &    2.09 & 52.12 &    2.12 & 61.63 &    2.06 & 80.66 &    1.68\\\\ 36.37 &    1.54 &  39.75 &    1.41 & 46.07 &    2.02 & 52.18 &    2.10 & 61.70 &    2.05 & 85.15 &    1.36\\\\ 36.40 &    1.53 &  39.82 &    1.33 & 46.10 &    2.02 & 52.25 &    2.12 & 62.16 &    2.01 & 85.19 &    1.35\\\\ 36.44 &    1.59 &  39.85 &    1.58 & 46.14 &    2.09 & 52.30 &    2.09 & 62.18 &    2.00 & 85.22 &    1.35\\\\ 36.47 &    1.62 &  39.87 &    1.61 & 46.17 &    2.09 & 52.32 &    2.12 & 62.23 &    2.05 & 87.14 &    1.17\\\\ 36.51 &    1.50 &  39.99 &    1.56 & 46.21 &    2.08 & 52.37 &    2.12 & 62.25 &    2.04 & 87.16 &    1.18\\\\ 36.54 &    1.43 &  40.05 &    1.66 & 46.24 &    2.08 & 52.39 &    2.14 & 62.58 &    2.08 & 87.21 &    1.17\\\\ 36.58 &    1.38 &  40.12 &    1.55 & 46.28 &    2.08 & 52.44 &    2.12 & 62.65 &    2.07 & 87.28 &    1.17\\\\ 36.61 &    1.20 &  40.19 &    1.44 & 46.30 &    2.09 & 52.46 &    2.13 & 63.18 &    2.02 & 87.35 &    1.15\\\\ 36.65 &    1.47 &  40.26 &    1.11 & 46.35 &    1.97 & 52.53 &    2.08 & 63.25 &    2.04 & 91.03 &    0.70\\\\ 36.68 &    1.41 &  40.33 &    0.87 & 46.37 &    1.97 & 52.58 &    2.10 & 64.18 &    2.04 & 91.10 &    0.74\\\\ 36.70 &    1.70 &  40.40 &    1.01 & 46.42 &    2.09 & 52.60 &    2.09 & 64.24 &    1.83 & 93.57 &    0.43\\\\ 36.72 &    1.65 &  40.47 &    1.53 & 46.49 &    2.06 & 52.65 &    2.10 & 64.52 &    1.98 & 93.90 &    0.36\\\\ \\noalign{\\smallskip} \\tableline\\ \\end{tabular} } \\end{center} \\end{table}   \\normalsize \\begin{table}[!ht] \\label{fonts2} \\caption{Evolutional data of the outburst model} \\smallskip \\begin{center} {\\tiny \\begin{tabular}{lllllll} \\tableline \\noalign{\\smallskip} Time &$\\log T_{\\rm ph}$&$\\log R_{\\rm ph}$& $\\log$ (mass loss rate)&$\\log V$  & $\\log L$ &  $\\log g$ \\\\  [day]& [K]&[cm] &[$M_{\\odot}~$yr$^{-1}$] & [cm~s$^{-1}$]&  [erg~s$^{-1}$] &[cm~s$^{-2}$]\\\\ \\noalign{\\smallskip} \\tableline \\noalign{\\smallskip} 9.8 &  4.86 & 11.04 &   -4.754 &  7.990 &  38.38 &  4.181 \\\\ 22.1 &  5.09 & 10.60 &   -5.287 &  8.087 &  38.40 &  5.060 \\\\ 27.3 &  5.16 & 10.46 &   -5.507 &  8.096 &  38.41 &  5.342 \\\\ 32.6 &  5.29 & 10.20 &   -5.778 &  7.971 &  38.42 &  5.846 \\\\ 34.8 &  5.37 & 10.05 &   -5.952 &  7.815 &  38.43 &  6.151 \\\\ 38.2 &  5.50 &  9.79 &   -6.490 &  7.405 &  38.43 &  6.679 \\\\ 38.9 &  5.55 &  9.69 &   -6.820 &  7.154 &  38.44 &  6.873 \\\\ 39.1 &  5.56 &  9.67 &   -6.921 &  7.067 &  38.44 &  6.918 \\\\ 39.5 &  5.63 &  9.53 &    0.000 &  0.000 &  38.44 &  7.191 \\\\ 40.5 &  5.82 &  9.15 &    0.000 &  0.000 &  38.44 &  7.948 \\\\ 44.9 &  5.95 &  8.89 &    0.000 &  0.000 &  38.44 &  8.471 \\\\ 50.7 &  6.03 &  8.73 &    0.000 &  0.000 &  38.43 &  8.796 \\\\ 60.6 &  6.10 &  8.58 &    0.000 &  0.000 &  38.41 &  9.093 \\\\ 71.7 &  6.13 &  8.51 &    0.000 &  0.000 &  38.39 &  9.239 \\\\ 78.1 &  6.14 &  8.44 &    0.000 &  0.000 &  38.30 &  9.368 \\\\ 80.8 &  6.13 &  8.43 &    0.000 &  0.000 &  38.24 &  9.399 \\\\ 85.0 &  6.12 &  8.41 &    0.000 &  0.000 &  38.16 &  9.425 \\\\ 89.5 &  6.09 &  8.40 &    0.000 &  0.000 &  38.02 &  9.445 \\\\ 94.4 &  6.04 &  8.40 &    0.000 &  0.000 &  37.82 &  9.460 \\\\ 97.5 &  6.01 &  8.39 &    0.000 &  0.000 &  37.70 &  9.466 \\\\ 100.3 &  5.99 &  8.39 &    0.000 &  0.000 &  37.58 &  9.469 \\\\ 116.4 &  5.85 &  8.39 &    0.000 &  0.000 &  37.03 &  9.480 \\\\ \\noalign{\\smallskip} \\tableline\\ \\end{tabular} } \\end{center} \\end{table}    \\begin{figure}[!ht] \\plottwo{ kato_fig1a.epsi}{kato_fig1b.epsi} \\caption{(a) The observational light curves of optical ($y$ band) and supersoft X-ray (0.3-0.55 keV) as well as our best fit model of $1.35 M_{\\odot}$ WD. (b) The evolution of physical values in our best fit model. The wind mass-loss rate, wind velocity at the photosphere, photospheric temperature, photospheric radius, total luminosity, and the surface gravity $g \\equiv GM/r^2$. The scale of the total luminosity $\\log L$ is 37.0 (bottom) and 39.25 (top). All the values are plotted in the logarithmic scale. The X-ray count rates and evolution data are tabulated in Tables 1 and 2. The origin of time is set to be JD 2453779. } \\label{bestfitmodel} \\end{figure}"
    },
    "0807/0807.4642_arXiv.txt": {
        "abstract": "Measuring the 3D distribution of mass on galaxy cluster scales is a crucial test of the $\\Lambda$CDM model, providing constraints on the behaviour of dark matter.  Recent work investigating mass distributions of individual galaxy clusters (e.g. Abell 1689) using weak and strong gravitational lensing has revealed potential inconsistencies between the predictions of structure formation models relating halo mass to concentration and those relationships as measured in massive clusters.  However, such analyses employ simple spherical halo models while a growing body of work indicates that triaxial 3D halo structure is both common and important in parameter estimates.  Though lensing is sensitive only to 2D projected structure and is thus incapable of independently constraining 3D models, the very strong assumptions about the symmetry of the lensing halo implied with circular or perturbative elliptical NFW models are not physically motivated and lead to incorrect parameter estimates with significantly underestimated error bars.  We here introduce a Markov Chain Monte Carlo (MCMC) method to fit fully triaxial models to weak lensing data that gives parameter and error estimates that fully incorporate the true uncertainty present in nature.  Using weak lensing data alone, the fits are sensitive to the Bayesian priors on axis ratio; we explore the impact of various general and physically motivated priors, and emphasise the need for future work combining lensing data with other data types to fully constrain the 3D structure of galaxy clusters.  Applying the MCMC triaxial fitting method to a population of NFW triaxial lenses drawn from the shape distribution of structure formation simulations, we find that including triaxiality cannot explain a population of massive, highly concentrated clusters within the framework of $\\Lambda$CDM, but easily explains rare cases of apparently massive, highly concentrated, very efficient lensing clusters.  Our MCMC triaxial NFW fitting method is easily expandable to include constraints from additional data types, and its application returns model parameters and errors that more accurately capture the true (and limited) extent of our knowledge of the structure of galaxy cluster lenses. ",
        "introduction": "Galaxy clusters are ideal laboratories in which to study dark matter, being the most massive bound structures in the universe and dominated by their dark matter component ($\\sim90\\%$). Constraining the clustering properties of dark matter is crucial for refining structure formation models that predict both the shapes of dark matter halos and their mass function (e.g. \\cite{navarro}; \\cite{bahcall}; \\cite{dahle}).  Several methods are used to measure galaxy cluster dark matter profile shapes and halo masses on a range of scales, including X-ray and Sunyaev-Zeldovich (SZ) studies, dynamical analyses, and gravitational lensing.  However, all of these methods require simplifying assumptions to be made regarding the shape and/or dynamical state of the cluster in order to derive meaningful constraints from available data.  Crucially, while most parametric methods typically assume spherical symmetry of the halo, observed galaxy clusters often exhibit significant projected ellipticity and halos in CDM structure formation simulations (e.g. \\cite{bett} (using the Millennium simulation); \\cite{shaw}) show significant triaxiality in cluster-scale halos, with axis ratios between minor and major axes as small as 0.4.  Understanding and accurately incorporating the impact of this physical reality on cluster mass and parameter estimates is crucial for accurate comparisons between measured cluster properties and model predictions.  In addition to the determination of masses, most cluster profile fits are carried out in the hope of either supporting or refuting the universality of the NFW profile and thus testing the CDM paradigm.  The NFW profile is typically parameterised by an approximate virial mass $M_{200}$ and a concentration parameter, $C$, and simulations predict a strong correlation between the two.  For a cluster of $M=10^{15}$ M$_{\\odot}$, $C \\sim 4$.  However, several authors (e.g. \\cite{limousin}; \\cite{kneib}; \\cite{gavazzia}) have recently reported gravitational lensing results in the very low probability tail of the predicted distribution; notably, in a combined weak and strong lensing analysis of Abell 1689, \\cite{broadhurst} report a concentration parameter of $C\\sim10$, when $C\\sim4$ is expected.  While a few such results are not damning, especially given the very complex, likely not relaxed, structures of the lensing halos (see e.g. the work of \\cite{lokas} on A1689), it is nonetheless of interest to investigate how possible future discrepancies between observations and the predictions of $\\Lambda$CDM should be interpreted.  Efforts to understand the impacts of triaxiality in gravitational lensing and its potential role in explaining apparent discrepancies with CDM began when \\cite{ogurib} applied a fully triaxial NFW model to the shear map of Abell 1689 to find that it is consistent with $6\\%$ of cluster-scale halos. These continued as \\cite{gavazzib} showed that a triaxial NFW can reconcile parameter values derived from observations of the cluster MS2137-23 to predictions from N-body simulations.  More generally, \\cite{corl} demonstrated in the weak lensing regime that neglecting halo triaxiality in parameterised fits of NFW models to NFW haloes with axis ratios significantly less than one can lead to over- and underestimates of up to 50$\\%$ and a factor of 2 in halo mass and concentration, respectively.  While extreme cases of triaxiality are rare, such haloes can be much more efficient lenses than their more spherical counterparts, especially when in configurations that hide most of the triaxial shape along the line of sight.  Further, even haloes with less extreme axis ratios are inaccurately fit by spherical models.  We expect the vast majority of galaxy cluster scale dark matter haloes, and so it is important to include this expectation in model fitting.  In the past, triaxial models have not generally been fit to lensing data because they cannot be well-constrained.  A simple NFW triaxial halo model has six free parameters (concentration, mass, two axis ratios, and two orientation angles) as opposed to the two parameters of the spherical model (mass and concentration only), and even combined weak and strong lensing data is not enough to meaningfully constrain so many parameters.  This is sometimes due to the limited depth of currently available observational data, but more importantly to the intrinsic limitations of lensing.  Because lensing is affected only by the the projected surface density and shear of the underlying mass distribution, it is inherently impossible to fully constrain a three-dimensional structure without imposing strong priors on the shape of the halo or supplementing lensing data with other data types more sensitive to line-of-sight halo structures, such as dynamical studies.   Despite these difficulties, since triaxiality has been convincingly demonstrated to be an important factor in model parameter estimation in lensing analyses, it is necessary to find a way to directly include it in NFW fits to galaxy clusters. For individual clusters it is crucial, if claims regarding the validity of the $\\Lambda$CDM paradigm based on NFW parameter estimates are to be meaningfully evaluated.  It is also of importance across populations to obtain more accurate distributions of galaxy cluster parameters, and especially in measuring the galaxy cluster mass function -- if the mass function is signficantly sloped as expected, even the best-case scenario of a symmetric scatter of mass estimates due to neglected triaxiality would lead to an asymmetric shift in the calculated mass function, as there are more low mass halos to shift up in mass than there are high mass halos to shift down.  To that end, we present here a Bayesian Markov Chain Monte Carlo method to fit fully triaxial NFW halos to weak lensing data. This method flexibly combines weak lensing data with prior probability functions on the model halo parameters to return parameter and error estimates that reflect the true uncertainties of the problem. Importantly, it also allows for the straightforward and statistically-robust inclusion of additional constraints from other data sources such as observations from SZ, X-ray and spectroscopic surveys. ",
        "conclusions": "Some may argue that it is foolhardy to fit a density profile model that fundamentally cannot be fully constrained with available data. Indeed, in many cases calculating the Bayesian evidence would likely favour a Spherical over a Flat prior, due to a significant decrease in available parameter space coupled with a relatively small decrease in likelihood values under the Spherical prior. However, unlike in cases of fitting multiple halos to account for substructure, or in another context, adding additional parameters to cosmological models, we know {\\it a priori} from physical observations that a triaxial model is a better model than a spherical model -- we see clearly in non-parametric lensing mass maps that galaxy clusters have significant levels of ellipticity and are not spherical.  The parameter estimates and uncertainties resulting from spherical models, or 2D elliptical models, that assume a $\\delta$-function prior on at least one axis ratio, therefore intrinsically claim a level of certainty we know we do not possess.  Making our best choice of prior on the axis ratios and fitting triaxial models is the more physically true approach, and gives the best estimate of parameter values and errors that can be obtained from lensing data alone.  The large size of these contours emphasises the need to combine lensing data with other data sources, be they SZ, dynamical, or X-ray, because the  dependence of all lensing effects -- strong, weak, intermediate -- on projected mass fundamentally constrains the precision and accuracy with which galaxy cluster masses and concentrations can be obtained.  Further, there is no way to know for a single lensing cluster how important triaxiality is without first carrying out a full triaxial analysis and comparing its results to those from a spherical fit; Figure~\\ref{fig:plot14} helps illustrate why this is the case.  The four panels plot the distributions of two observable properties of the triaxial halos of the lensing population of Section \\ref{ssec:pop} for which the 68\\% and 95\\% confidence contours obtained fitting an NFW under a Spherical prior {\\it do not} contain the true halo values.  These observables are the axis ratio of the projected isodensity contours $q$ of the lens and the maximum likelihood value of the NFW fit under a Spherical prior.  It might be hoped that either a high visible ellipticity (low $q$) or a low likelihood value would indicate that the spherical model is a bad fit; however, as seen in the figure, the distributions show no bias compared to the distribution of the entire lensing population, and thus can provide no indication of the adequacy of the spherical fit.  This can be understood by example by referring back to the extreme halo shapes and orientations of Section \\ref{ssec:ex}: the highly triaxial prolate halo in a Line of Sight orientation will have both a low ellipticity ($q\\sim 1$) and a very high maximum likelihood value (it is an efficient lens with circular projected isodensity contours and so will be very well fit by a spherical model), but the most-probable parameter values and their errors obtained from fitting a spherical NFW are very inaccurate.  Thus, it is necessary to fit a triaxial NFW using the best prior available to every lens, as there is no way to determine beforehand how triaxial any given halo might be.  \\begin{figure} \\epsfig{file=f13.eps,scale=0.3} \\caption{The left panel shows the $M_{200}-C$ 68$\\%$ and 95$\\%$ confidence contours, under various priors, for a triaxial NFW fit to a triaxial halo typical of those found in simulations, with $M_{200}=10^{15}M_{\\odot}$ and $C=4$, $\\{a=0.76, b=0.85\\}$.  The right panel plots the $M_{200}-C$ distribution shaded by the minor axis ratio $a$ for the same halo under a Flat prior.} \\label{fig:plot13} \\end{figure}  A great advantage of this MCMC method is the ease with which other data sets and constraints may be incorporated.  Whether through extra terms in the likelihood function (though such inclusions require a careful analysis, not seen yet in the literature for the combination of many data types, of the proper weights to give data types with very different scales and sources of errors), or through a prior (e.g. a prior on the cluster mass taken from the best estimate and errors of a dynamical study).  The future of understanding cluster structure lies in the combination of various data types, as all available methods are limited either by fundamental constraints, such as lensing, or by the availability of data, as in dynamics, in which highly simplified dynamical models are required by the relatively small numbers of position-velocity pairs.  Methods such as this, and those currently being developed by e.g., \\cite{jullo} and \\cite{feroz}, are a crucial element of the next generation analysis toolbox for galaxy cluster studies.   \\begin{figure} \\epsfig{file=f14.eps,scale=0.4} \\caption{The right (left) top and bottom panels show, respectively, the distributions of the projected axis ratios $q$ and the maximum likelihood values obtained from fitting a spherical NFW, for the halos of the triaxial lensing population of Section \\ref{ssec:pop} for which the 68\\% (95\\%) confidence contours recovered fitting an NFW under a Spherical prior {\\it do not} enclose the true triaxial halo values.} \\label{fig:plot14} \\end{figure}  \\subsection{Selecting better priors} Placing well-motivated priors on the shapes of galaxy clusters is crucial to future work modelling the most massive structures in the universe, and understanding both their characteristics as individuals and a population.  While structure formation simulations provide a good starting point for such priors, using their predictions does of course bias all results towards agreement with those simulations.  One good alternative is to use a distribution of 2d axis ratios observed in a sample of galaxy clusters, preferably from mass-sensitive methods (e.g. lensing mass reconstructions) and from this construct a 3d shape prior.  For this, even the simplest elliptical lensing models would be adequate, as CK07 showed that even a Singular Isothermal Ellipsoid model consistently recovered projected axis ratios accurately, even if the true lensing profile was NFW (this is not true for other model parameters).  While the necessity and importance of these priors may be discomforting, using standard techniques and simpler models is simply disguising the problem, as such models contain highly-restrictive hidden priors.  Bayesian techniques such as the one presented in this paper are simply tools that allow us to better understand the true constraints we can place on physical models, not solutions in themselves to physical problems.  The broad posterior probability distributions for the triaxial NFW profiles fit to simulated lensing data in this paper indicate the weakness of our current constraints, and emphasise the need for focus on the careful combination of complementary data types to further constrain galaxy cluster structure models. To this end, this Bayesian MCMC triaxial NFW fitting  method provides, through the prior probability functions, a statistically robust and straightforward way to combine constraints from data types with very different error properties.  \\subsection{Application} This MCMC triaxial NFW fitting method, using fully tested look-up tables to significantly speed up the calculation of the triaxial NFW lensing quantities, can be implemented for a standard weak lensing data set on a single fast processor in about 1 day, making it fully feasible for use across even the largest existing lens surveys.  To begin this process, in a companion paper Corless, King, \\& Clowe (in preparation), we apply this method employing a range of priors to galaxy clusters Abell 1689, Abell 1835 and Abell 2204.  In applying this method to observational data it will be important to account for associated and line-of-sight structures in the lensing field.  Galaxy clusters form at the intersections of filaments, so it is not surprising that spectroscopic studies often reveal correlated structures in their fields(e.g. \\cite{lokas}) that can bias our estimation of cluster parameters if neglected (\\cite{kingb}).  Unaccounted for large scale structure along the line of sight can also have a severe impact on lensing analyses (\\cite{hoekstra}),leading to a factor of $\\sim 2$ increase in errors on cluster mass. As demonstrated by \\cite{dodel}, this can be somewhat mitigated by including a noise term for large scale structure in the lensing analysis.  Building on the greater coherence scale of large scale structure noise compared with the noise associated with the intrinsic ellipticity dispersion of galaxies, Dodelson notes that the errors on cluster mass can be reduced by $\\sim 50\\%$ for wide-field data.  This MCMC triaxial fitting method will be very useful in the analysis of current and future surveys such as LoCuSS, PAN-STARRS and DES.  \\subsection{Summary} We here present an MCMC method for fitting triaxial NFW models to weak lensing data that \\begin{itemize} \\item consistently and accurately returns parameter and error estimates representing the true uncertainties of the problem; \\item includes Bayesian priors on the halo geometry, chosen from simulations or observations; \\item allows great flexibility in the inclusion of constraints from different observations and analyses, including strong lensing, dynamical studies, and S-Z and X-ray derived mass models; \\item demonstrates that while triaxiality can explain rare cases of overconcentrated galaxy cluster halos within the $\\Lambda$CDM paradigm, it cannot explain a population-wide trend of over-concentration; \\item should be widely applied in weak lensing analyses to better determine cluster parameters and errors that represent the true (and limited) extent of our knowledge of the distribution of matter in galaxy cluster lenses. \\end{itemize}"
    },
    "0807/0807.1067_arXiv.txt": {
        "abstract": "{} {Pointed observations with XMM-Newton provide the basis for creating catalogues of X-ray sources detected serendipitously in each field. This paper describes the creation and characteristics of the 2XMM catalogue.} {The 2XMM catalogue has been compiled from a new processing of the XMM-Newton EPIC camera data. The main features of the processing pipeline are described in detail. } {The catalogue, the largest ever made at X-ray wavelengths, contains 246,897 detections drawn from 3491 public XMM-Newton observations over a 7-year interval, which relate to 191,870 unique sources. The catalogue fields cover a sky area of more than 500\\,deg$^2$. The non-overlapping sky area is $\\sim\\!360$\\,deg$^2$ ($\\sim\\!1$\\% of the sky) as many regions of the sky are observed more than once by XMM-Newton. The catalogue probes a large sky area at the flux limit where the bulk of the objects that contribute to the X-ray background lie and provides a major resource for generating large, well-defined X-ray selected source samples, studying the X-ray source population and identifying rare object types. The main characteristics of the catalogue are presented, including its photometric and astrometric properties } {} ",
        "introduction": "\\label{intro}  Surveys play a key role in X-ray astronomy, as they do in other wavebands, providing the basic observational data  characterising the underlying source populations. {\\sl Serendipitous} X-ray sky surveys, based on the field data from individual pointed observations, take advantage of the relatively wide field of view afforded by typical X-ray instrumentation. Such surveys have been pursued with most X-ray astronomy satellites since the Einstein Observatory. The resulting serendipitous source catalogues (e.g., EMSS: Gioia et al. 1990, Stocke et al. 1991; WGACAT: White et al. 1994; ROSAT 2RXP: Voges et al.\\ 1999; ROSAT 1RXH: ROSAT Team 2000; ASCA AMSS: Ueda et al., 2005) have been the basis for numerous studies and have made a significant contribution to our knowledge of the X-ray sky and our understanding of the nature of the various Galactic and extragalactic source populations.  The XMM-Newton observatory provides unrivalled capabilities for serendipitous X-ray surveys by virtue of the large field of view of the EPIC cameras and the high throughput afforded by the heavily nested telescope modules. This capability guarantees that each XMM-Newton observation provides a significant harvest of serendipitous X-ray sources in addition to data on the original target. In addition, the extended energy range of XMM-Newton ($\\sim 0.2-12$ keV) means that XMM-Newton detects significant numbers of obscured and hard-spectrum objects that are absent in many earlier soft X-ray surveys.  This paper describes the Second XMM-Newton Serendipitous Source Catalogue (2XMM) which has been created from the serendipitous EPIC data from from 3491 XMM-Newton pointed observations made over a $\\sim\\!7$-year interval since launch in 1999. The XMM-Newton serendipitous source catalogues are produced by the XMM-Newton Survey Science Centre (SSC), an international consortium of ten European institutions, led by the University of Leicester, as a formal project activity performed on behalf of ESA. The catalogues are based on the EPIC source lists produced by the scientific pipe-line used by the SSC for the processing of all the XMM-Newton data. The first serendipitous source catalogue, 1XMM, was released in 2003 (Watson et al. 2003a; XMM-SSC 2003). The current 2XMM catalogue incorporates a wide range of improvements to the data processing, uses the most up-to-date instrument calibrations and includes a large number of new parameters. In parallel, the 2XMM catalogue processing also produces a number of additional data products, for example time-series and spectra for the brighter individual X-ray sources. A pre-release version of the current catalogue, 2XMMp (XMM-SSC 2006), was made public in 2006. This includes $\\sim\\!65$\\% of the fields and $\\sim\\!75$\\% of the sky area covered by 2XMM, while $\\sim\\!88$\\% of all 2XMMp sources appear in the 2XMM catalogue. Around 56\\% of all 2XMM sources  already have an entry in the 2XMMp catalogue.  The 2XMM catalogue provides an unsurpassed sky area for serendipitous science and reaches a flux limit corresponding to the dominant extragalactic source contribution to the cosmic X-ray background. The catalogue is part of a wider project to explore the source populations in the XMM-Newton serendipitous survey (the XID project; Watson et al. 2001; Watson et al. 2003b) through optical identification of well-defined samples of serendipitous sources (e.g., Barcons et al. 2002, 2007; Della Ceca et al. 2004, Caccianiga et al. 2008, Motch et al. 2002; Schwope et al. 2004; Page et al. 2007; Yuan et al. 2003; Dietrich et al. 2006). Indeed these identification programs were effectively based on less mature versions of the XMM-Newton catalogue data processing. XMM-Newton serendipitous survey results have also been used to study various statistical properties of the populations such as X-ray spectral characteristics, source counts, angular clustering, and luminosity functions (Severgnini et al. 2003; Mateos et al. 2005; Carrera et al. 2007; Caccianiga et al. 2007; Mateos et al. 2008; Della Ceca et al. 2008; Ebrero et al. 2008). Other projects based on XMM-Newton serendipitous data include the HELLAS2XMM survey (Baldi et al. 2002; Cocchia et al. 2007).  The 2XMM serendipitous catalogue described here is complementary to ``planned\" XMM-Newton surveys which provide coverage of much smaller sky areas, but often with higher sensitivity, thus exploring the fainter end of the X-ray source population. The deepest such surveys, such as the Lockman Hole (Hasinger et al. 2001; Brunner et al. 2008) and the CDF-S (Streblyanska et al. 2004), cover essentially only a single XMM-Newton field of view, have total integration times $\\sim\\!300-1000$\\,ks and reach fluxes $\\sim$ few $\\times 10^{-16}\\rm\\,erg\\,cm^{-2}\\,s^{-1}$, close to the confusion limit. XMM-Newton has also carried out contiguous surveys of various depths covering much larger sky areas utilising mosaics of overlapping pointed observations to achieve the required sensitivity and sky coverage. Currently the largest contiguous XMM-Newton survey is the XMM-LSS (Pierre at al. 2007) covering $\\sim\\!5$\\,deg$^2$ with typical exposure time 10\\,--\\,20\\,ks per observation. Other medium-deep surveys of 1\\,--\\,2\\,deg$^2$ regions include the SXDS ($\\sim\\!1.1$\\,deg$^2$, 50\\,--\\,100\\,ks exposures; Ueda et al. 2008), the COSMOS surveys ($\\sim\\!2$\\,deg$^2$, $\\sim\\!80$\\,ks exposures; e.g., Cappelluti et al. 2007; Hasinger et al.\\ 2007), and the Marano field survey (Krumpe et al. 2007). These larger area surveys typically reach limiting fluxes of $10^{-14}$ to $<10^{-15}\\rm\\,erg\\,cm^{-2}\\,s^{-1}$.  We also note that Chandra observations have been used to compile a serendipitous catalogue including $\\sim\\!7000$ point sources (the ChaMP catalogue; Kim et al. 2007) and plans are underway to compile a serendipitous catalogue from all suitable Chandra observations (Fabbiano et al. 2007).  The paper is organised as follows. Section~\\ref{instr} introduces the XMM-Newton observatory. Section~\\ref{catobs} presents the XMM-Newton observations used to create the catalogue and the characteristics of the fields. Section~\\ref{processing} outlines the XMM-Newton data processing framework and provides a more detailed account of the EPIC data processing, focusing in particular on source detection and parameterisation, astrometric corrections and flux computation. Section~\\ref{ssp} provides an account of the automatic extraction of time-series and spectra for the brighter sources, while Sect.~\\ref{extcatcross} outlines the external catalogue cross-correlation undertaken. Section~\\ref{qual_eval} describes the quality evaluation undertaken and some recommendations on how to extract useful sub-samples from the catalogue. Section~\\ref{catcompile} describes additional processing and other steps taken to compile the catalogue including the identification of unique sources.  The main properties and characterisation of the catalogue is presented in Sect.~\\ref{catchar}. Section~\\ref{cataccess} summarises access to the catalogue and plans for future updates to 2XMM, and Sect.~\\ref{conclusion} gives a summary. ",
        "conclusions": "The 2XMM catalogue contains more than 20\\,000 entries of extended detections. The reliable detection and parameterisation of extended sources is significantly more demanding than for point-like sources because there are many more degrees of freedom in the parameter space. The relatively simple analysis approach used in the creation of the catalogue (Sect.~\\ref{extdsrc}) means that the catalogue contains a significant number of extended object detections that are either spurious or at least uncertain (cf.\\ Sects.~\\ref{srcdetproblems} and~\\ref{autoflag}). The most common causes of problems with extended sources are summarised below and illustrated in Fig.~\\ref{examplesextdfig}.  \\begin{description}  \\item{\\bf Spurious detections near bright point sources:} These are mostly due to inaccuracies of the PSF models, leading to inaccurate modelling of the internal background by the source fitting routine.  \\item{\\bf Confusion of point sources:} Pairs or multiples of point sources can be detected as one extended source since only up to two point sources are modelled simultaneously by the fitting algorithm.  \\item{\\bf Insufficient background subtraction:} Some spatial variations of the intrinsic background are poorly modelled by the spline map. In regions where the background is underestimated, spurious detections of extended sources are possible. (In many cases the extent parameter of these sources is at or near the maximum of the allowed range, $80\\arcsec$.)  \\item{\\bf Multiple detections of extended sources:} The surface brightness distribution of extended sources is generally more complex than the fitted $\\beta$-model. This can lead to additional detections in the wings of extended sources. The most extreme cases are observations of complex, bright extended sources (e.g., Galactic supernova remnants), leading to the detection of numerous extended sources in one field. Also, extended emission following the fitted $\\beta$-model, but with an extent greater than the maximum allowed in the fit, tends to be broken up into multiple detections.  \\item{\\bf Instrumental artefacts:} OOT events of piled-up sources, single reflections arcs, and scattered light from the RGA (cf.\\ Fig.~\\ref{examplesfieldfig}\\,c) can cause both point-like and extended spurious detections.  \\end{description}  \\begin{figure*}[ht!] \\resizebox{\\hsize}{!}{\\includegraphics{example_extd_sq.ps}} % \\caption[]{Examples of extended source detections. Green circles mark point source detections. In panels (i) -- (vi) the magenta and yellow circles mark real and spurious extended detections respectively, plotted with their fitted extent (i.e., core radius, see Sect.~\\ref{extdsrc}). In panels (vii) \\& (viii) the yellow ellipses indicate the position of spurious extended source detections.\\\\ {\\sl Top row}: (i) a compact extended source with a small core radius; (ii) a large, low surface brightness extended source at the edge of the FOV with low likelihood but high flux (see Fig.~\\ref{extsrc}); (iii) an object with a point-like core detected both as a point source as well as an extended source; (iv) a clearly extended source with a spurious detection nearby (yellow circle) which is smaller and fainter (by a factor of 45) and which therefore does not significantly affect the parameters of the real source.\\\\ {\\sl Bottom row}: (v) a SNR in the LMC where intrinsic structure is detected as point sources (note that the core radius is not representative as the extended emission does not follow the $\\beta$-model fitted); (vi) a bright extended source with multiple spurious detections around the centre: the core radii of these spurious detections are comparable to or greater than the extent of the real source and will thus significantly affect the parameters of the real source (note that the maximum core radius allowed in the fitting is $80\\arcsec$); (vii) a faint filamentary structure broken up into several extended detections where the parameters have little meaning (due to the circularly symmetric nature of the fit); (viii) a crowded region where several point sources are detected as extended due to source confusion (the algorithm is restricted to fitting at most two confused sources simultaneously). } \\label{examplesextdfig} \\end{figure*}  The catalogue contains extensive detection flags (Sect.~\\ref{qual_eval}) which can be used to produce much cleaner extended-source samples, albeit at the expense of removing some genuine extended objects. (This is the case as the flagging scheme only provides warnings about generic problems with the analysis or the data rather than a specific assessment of the reality of each detection.) In particular the automated quality Flags 4, 5, and 6 (see Table~\\ref{autoflagtab}) are set to warn about possible spurious detections of extended sources. The combined Flag~7 for extended sources is set if one of the Flags $4-6$ is set. This flag is set for 9\\,882 out of 20\\,837 detections, indicating a potential spurious fraction of about $50$\\%. However, the rate of spurious detections is distributed very unevenly over the catalogue observations as is discussed below.  Figure~\\ref{extsrc} illustrates some of the main features of the extended source detections in the catalogue. The plot shows that there is, as expected, an overall correlation of extent likelihood with EPIC flux. The considerable scatter in the plot has three origins: (i) the observations from which the detections are drawn have a considerable range of exposure times and background values; (ii) source extent: sources with larger spatial extent have lower likelihoods at the same integrated flux; (iii) the presence of significant numbers of spurious detections. The detections with Flag~7 set show, as expected, a broader distribution than those without this flag, and a much broader distribution than for the detections with `best' summary flags (i.e., summary flag $<\\!2$). This is, of course, due to the fact that spurious detections will often have implausible likelihoods for the fitted flux or correspond to very large source extent which is rare in genuine detections.  \\begin{figure}[ht] \\resizebox{\\hsize}{!}{\\includegraphics[angle=270]{extsrc.ps}} % \\caption[]{Distribution of extent likelihood as a function of total-band EPIC flux for the extended source detections in the 2XMM catalogue. Red dots are potentially spurious detections with Flag 7=T, yellow dots are detections with Flag 7=F, black dots are the `best' sample detections with summary flag $<\\!2$. Green stars mark the targets of the XMM-Newton observations classified as extended object types and blue squares targets which are object types classified as point-like. The vertical concentrations of target points at flux $\\sim\\!3\\times 10^{-11}$ and $\\sim\\!2\\times 10^{-10}\\rm\\,erg\\,cm^{-2}\\,s^{-1}$ are real, being due to multiple detections of two different SNRs used as XMM-Newton calibration targets. } \\label{extsrc} \\end{figure}  Based on the sample with `best' summary flags it is clear that there are very few reliable extended source detections with extent likelihood above $\\sim\\!1000$ or flux above $\\sim\\!4\\times 10^{-13}\\rm\\,erg\\,cm^{-2}\\,s^{-2}$, highlighting the problems that the detection algorithm has with bright objects.\\footnote{We also note that this is what is expected from the source counts of clusters of galaxies which are expected to dominate the extended detections, at least at high Galactic latitudes.} Indeed the majority of reliable extended objects in this region of the diagram are the XMM-Newton targets themselves (but note that many of these have Flag~7 set which would otherwise indicate potentially spurious detections). At the highest fluxes a large fraction of the detections relate to very bright point-like targets that are incorrectly parameterised as being extended due the deficiencies of the fitting algorithm noted above.  We have investigated a small subset of the extended detections at high Galactic latitudes covered by SDSS DR6 (excluding targets). We selected detections with extent likelihood $>\\!100$ and no warning flags set (i.e., summary flag 0) and evaluated their validity by examining the X-ray images visually and by searching for matches with catalogued objects. We find that less than 5\\% of these may be spurious extended source detections, around 40\\% are clearly associated with catalogued clusters or groups of galaxies and a few percent are associated with single nearby galaxies.  For a further $\\sim\\!30$\\% of the detections we find convincing evidence of a previously uncatalogued cluster or group of galaxies at the X-ray source location from visual inspection of the SDSS DR6 images.  These results demonstrate that the overall reliability of the `best' extended source sample is high, at least at higher likelihoods, and that, as expected, the extended source sample is dominated by groups and clusters of galaxies. We have not carried out a similar exercise systematically at low Galactic latitudes but checks of selected detections demonstrate the expected associations with SNRs, HII regions, and discrete extended features in the Galactic Centre region.   \\label{conclusion}  We have presented the 2XMM catalogue, described how the catalogue was produced and discussed the main characteristics of the catalogue. Table~\\ref{sumtable} provides a summary of its main properties, bringing together information presented elsewhere in this paper.  \\begin{table}[bt] \\caption{Summary of 2XMM catalogue characteristics} \\label{sumtable} \\small \\centering \\tabcolsep 1.5mm \\begin{tabular}{llrl} \\hline \\multicolumn{2}{l}{Energy range (keV; set by EPIC cameras)} & 0.2\\,--\\,12.0 \\\\ \\hline Observations & {\\sl total} & 3491 \\\\ & pn data     & 2674 \\\\ & MOS1 data   & 3384 \\\\ & MOS2 data   & 3394 \\\\ \\hline Time interval   & \\multicolumn{2}{r}{Feb 2000\\,--\\,Mar 2007} \\\\ \\hline Detections & {\\sl total}            & 246897 \\\\ & {\\sl  total} L$\\ge 10$ & 201275 \\\\ & {\\sl  total} sum flag 0& 199359 \\\\ & point-like             & 226060 \\\\ & extended               & 20837 \\\\ & with products          & 38320 \\\\ \\hline Unique sources & {\\sl total}   & 191870 \\\\ & point-like    & 173066 \\\\ & extended      & 18804 \\\\ \\hline Sky area \\ (deg$^2$ ) & {\\sl total}\\,$^{a}$ & $\\sim\\!560$ \\\\ & net MOS1/2          & $\\sim\\!355$ \\\\ & net pn              & $\\sim\\!330$ \\\\ \\hline Median exposure time & MOS1/2 & $\\sim\\!16000$ s\\\\ \\ \\ \\ \\ \\ \\ (per observation)    & pn     & $\\sim\\!12500$ s\\\\ \\hline Flux limit (pn) at  $\\sim 10$\\% sky & 0.5\\,--\\,2.0 keV   & $\\sim\\!2$ \\\\ \\ coverage \\  ($10^{-15} \\rm\\ erg\\ cm^{-2}\\ s^{-1}$ )  & 2.0\\,--\\,12 keV    & $\\sim\\!15$ \\\\ & 4.5\\,--\\,12 keV  & $\\sim\\!35$ \\\\ \\hline Flux limit (pn) at  $\\sim 90$\\% sky & 0.5\\,--\\,2.0 keV  & $\\sim\\!10$ \\\\ \\ coverage  \\  ($10^{-15} \\rm\\ erg\\ cm^{-2}\\ s^{-1}$ )     & 2.0\\,--\\,12 keV   & $\\sim\\!90$ \\\\ & 4.5\\,--\\,12 keV & $\\sim\\!250$ \\\\ \\hline Astrometric accuracy \\ (1$\\sigma$)& typical    & 1.5\\arcsec \\\\ \\ \\ \\ \\ \\ \\   & best$^{b}$ & 0.35\\arcsec \\\\ \\hline Photometric accuracy & MOS1/2 comparison & $\\le 5$\\% \\\\ & pn/MOS comparison    & $\\le 10$\\% \\\\ \\hline \\end{tabular} \\begin{list}{}{} \\item [$^{\\mathrm{a}}$]overlaps included \\item [$^{\\mathrm{b}}$]limited by systematics \\end{list} \\normalsize  \\end{table}  2XMM is the largest X-ray source catalogue ever produced, containing almost twice as many discrete sources as either the ROSAT survey or ROSAT pointed catalogues. The catalogue complements deeper Chandra and XMM-Newton small area surveys, and probes a large sky area at the flux limit where the bulk of the objects that contribute to the X-ray background lie.  The catalogue has very considerable potential a detailed account of which lies outside the scope of this paper. In particular the catalogue provides a rich resource for generating sizeable, well-defined samples for specific studies, utilising the fact that X-ray selection is a highly efficient (arguably the most efficient) way of selecting certain types of object, notably active galaxies (AGN), clusters of galaxies, interacting compact binaries and active stellar coronae. The large sky area covered by the serendipitous survey, or equivalently the large size of the catalogue, also means that 2XMM is a major resource for exploring the variety of the X-ray source population and identifying rare source types. Although the 2XMM catalogue alone provides a powerful way of studying the X-ray source population, matching the X-ray data with, e.g., optical catalogues can offer an even more effective way to generate considerable samples of particular object types.  Projects that exploit some of these characteristics are already underway.  Finally we note that, since the XMM-Newton spacecraft and instruments remain in good operational health, we can anticipate a substantial growth in the pool of serendipitous X-ray sources detected, increasing at a rate of $\\sim$35\\,000 sources/year. With this backdrop, further XMM-Newton catalogue releases are planned at regular intervals. The first such incremental release is planned for August 2008.   \\label{newapp}  \\begin{description}  \\item [\\bf On-axis position]: the telescope optical axes, defined by the geometry of each of the three X-ray mirror modules, is not coincident with geometrical centres of the EPIC detectors. The target of the observation is preferentially placed close to, but slightly offset from, the optical axis.  \\item [\\bf Point-Spread-Function (PSF)]: the telescope optics spread X-ray photons from a point source into a centrally-peaked distribution which is oversampled by the EPIC cameras. The PSF is energy-dependent and becomes broader with increasing (off-axis) angle from the telescope optical axis but also suffers a distortion which elongates the profile in the azimuthal direction.  \\item [\\bf  Event patterns]: an X-ray photon incident in a given CCD location causes charge deposition in several surrounding CCD pixels, often not symmetrically distributed around the central pixel. Several distinct charge distributions (patterns) are recognised as real events by the on-board processing electronics for the MOS cameras, whilst this processing takes place on the ground for the pn camera.  \\item [\\bf Out-Of-Time (OOT) events]: EPIC camera exposures are composed of many short-duration frames during which the recorded events are rapidly read out and processed by the on-board electronics. The total time between frames (frame-time) depends on the observing mode but is a maximum of 73\\,ms for the pn and 2.6\\,s for the MOS. The cameras are shutterless and record data during the readout (`out-of-time') as well as the processing (`imaging') period, leading to a faint trail of the `out-of-time' events along the readout direction which becomes obvious for bright sources (see Fig.~\\ref{examplesfieldfig}c). The percentage of OOT events is a function of the ratio of the frame readout time to the frame integration time for a given mode. The highest percentage of OOT events at 6.3\\% is for the pn full frame mode, while it never exceeds 0.5\\% for the MOS.  \\item [\\bf Pileup]: for bright sources, pixels in the core of the PSF can receive multiple X-ray photons during an integration frame. The on-board processing electronics cannot recognise them as distinct events within that frame and either treats them as a single event with higher energy or rejects them entirely if the resultant pixel pattern of the combined event lies outside the pre-defined X-ray pattern library. As a result the recorded counts are lower in the core of the source profile, producing a flattening or even depression of the source profile (see Fig.~\\ref{examplesfieldfig}c). In addition, it has an impact on the spectral profile (i.e., a hardening of the spectrum).  \\item [\\bf Optical loading]: The EPIC cameras (more so the MOS detectors) are also sensitive to optical photons so that optically bright objects generate recordable events in EPIC images (Lumb 2000, and references therein). The level of contamination depends on the filters used and the optical brightness of the object. In most observations the filter used is conservatively selected to minimise this effect. Note that in the case of the pn some apparently very soft sources are affected by optical loading.  \\item [\\bf RGA scattered light]: scattering of incident X-rays by the RGAs in the two telescope modules that feed the MOS cameras causes a diffuse bright narrow band in the X-ray images which is detectable for bright X-ray sources (see Fig.~\\ref{examplesfieldfig}c).  \\item [\\bf Good-Time-Interval (GTI)]: data from EPIC camera frames can be accepted or rejected according to the state of various housekeeping and science parameters, e.g., spacecraft attitude stability and particle background level. The `GTI's are the time periods during which the parameter(s) being monitored are within the acceptable thresholds.  \\end{description}  A more detailed description of the instruments can be found in the on-line version of the XMM User Handbook (Ehle et al.\\ 2007) and on the ESAC documentation web-pages for calibration\\footnote{http://xmm2.esac.esa.int/external/xmm\\_sw\\_cal/calib/documenta\\-tion/index.shtml}."
    },
    "0807/0807.3548_arXiv.txt": {
        "abstract": "We introduce two new diagnostics of dark energy (DE). The first, $Om$, is a combination of the Hubble parameter and the cosmological redshift and provides a {\\em null test} of dark energy being a cosmological constant $\\Lambda$. Namely, if the value of $Om(z)$ is the same at different redshifts, then DE $\\equiv \\Lambda$, {\\em exactly}. The slope of $Om(z)$ can differentiate between different models of dark energy even if the value of the matter density is not accurately known. For DE with an unevolving equation of state, a positive slope of $Om(z)$ is suggestive of Phantom ($w < -1$) while a negative slope indicates Quintessence ($w > -1$). The second diagnostic -- {\\em acceleration probe} $\\bar q$ -- is the mean value of the deceleration parameter over a small redshift range. It can be used to determine the cosmological redshift at which the universe began to accelerate, again without reference to the current value of the matter density. We apply the $Om$ and $\\bar q$ diagnostics to the Union data set of type Ia supernovae combined with recent data from the cosmic microwave background (WMAP5) and baryon acoustic oscillations. ",
        "introduction": "\\label{sec:intro} The nature of dark energy (DE) is one of the most intriguing questions facing physics. The fact that DE provides the main contribution to the energy budget of the universe today while remaining subdominant during previous epochs, provides a challenge to model builders attempting to understand the nature of this seemingly all-pervasive ether-like substance.  Theoretical models for DE include the famous cosmological constant, $\\Lambda$, suggested by Einstein in 1917 \\cite{einstein} and shown to be related to the vacuum energy $\\langle T_{ik} \\rangle_{\\rm vac} \\propto \\Lambda g_{ik}$ several decades later \\cite{zel68,wein89}. Indeed, the cosmological constant appears to occupy a privileged position amongst other DE models by virtue of the fact that its equation of state $w=-1$ is Lorentz invariant and so appears the same to any inertial observer. However, within the context of cosmology, an explanation of DE in terms of $\\Lambda$ faces one drawback, namely, in order for the universe to accelerate today the ratio of the energy density in the cosmological constant to that in radiation must have been very small at early times, for instance $\\rho_\\Lambda \\simeq 3\\times 10^{-58}\\rho_{\\rm EW}$ at the time of the electroweak phase transition. Although vacuum energy may conceivably be associated with small numbers such as the neutrino mass ($\\rho_\\Lambda \\sim m_\\nu^4$) or even the fine structure constant \\cite{S99}, a firm theoretical prediction for the value of $\\Lambda$ is currently lacking, allowing room for alternatives including models in which both the DE density and its equation of state (EOS) evolve with time. Alternatives to the cosmological constant include scalar field models called quintessence which have $w > -1$, as well as more exotic `phantom' models with $w < -1$.  Although most recent studies show that a cosmological constant + cold dark matter (LCDM) is in excellent agreement with observational data, dynamical dark energy can explain the data, too \\cite{observations1,snls,HST,essence,union,wmap03,wmap05,metamorphosis,observations2,alam07,arman2,DE_review,ss06}. Indeed, the enormous variety of DE models suggested in the literature (see \\cite{DE_review} for reviews) has been partially responsible for the burgeoning industry of model independent techniques aimed at reconstructing the properties of dark energy directly from observations \\cite{ss06}. It is well known that model independent methods must be wary of several pitfalls which can subvert their efficacy. These relate to priors which are sometimes assumed about fundamental cosmological quantities such as the EOS and the matter density. As first pointed out in \\cite{maor02}, an incorrect prior for the EOS can lead to gross misrepresentations of reality. The same applies to the value of the matter density. Indeed, as we shall demonstrate later in this paper, an incorrect assumption about the value of $\\om$ can lead to dramatically incorrect conclusions being drawn about the nature of dark energy. Clearly the need of the hour, then, is a diagnostic which is able to differentiate LCDM from `something else' with as few priors as possible being set on other cosmological parameters.  In this paper we introduce a new diagnostic, $Om$, which is constructed from the Hubble parameter $H\\equiv \\dot a/a$ determined directly from observational data and provides a {\\em null test} of the LCDM hypothesis. Here $a(t)$ is the scale factor of a Friedmann-Robertson-Walker (FRW) cosmology. We show that $Om$ is able to distinguish dynamical DE from the cosmological constant in a robust manner {\\em both with and without reference} to the value of the matter density, which can be a significant source of uncertainty for cosmological reconstruction. The $Om$ diagnostic is in many respects the logical companion to the statefinder $r = \\atridot/a H^3$ \\cite{statefinder} (otherwise dubbed jerk $j$, see e.g. \\cite{V04,DG08} as well as the earlier paper \\cite{CN98}). We remind the reader that $r = 1$ for LCDM while $ r \\neq 1$ for evolving DE models. Hence $r(z_1) - r(z_2)$ provides a {\\em null test} for the cosmological constant. Similarly, the unevolving nature of $Om(z)$ in LCDM furnishes $Om(z_1) - Om(z_2)$ as a {\\em null test} for the cosmological constant. (For null tests based on gravitational clustering see \\cite{CN07,np08}.) Like the statefinder, $Om$ depends only upon the expansion history of our Universe. However, while the statefinder $r$ involves the third derivative of the expansion factor $a(t)$, $Om$ depends upon its {\\em first} derivative only. Therefore, as we demonstrate in this paper, $Om$ is much easier to reconstruct from observations \\footnote{Note that the $Om$ diagnostic does not use any information about the evolution of inhomogeneities in an FRW background, usually represented by the growth factor $\\delta(z)\\equiv (\\delta\\rho/\\rho)_m$. Therefore our proposal for a null test of LCDM is totally unrelated to earlier null tests \\cite{CN07,np08} (based on the analytical expression for $\\Omega_m$ in terms of $\\delta(z)$ obtained in \\cite{st98}), and to the test of physical nature of dark energy (equivalently, that $G_{eff}$ is a constant) suggested in \\cite{ss06}.} .  It is clear that a determination of $H(z)$ from any single test may suffer from systematic uncertainties. That is why, in order to obtain $Om$ with sufficient accuracy, it is necessary to combine information about $H(z)$ obtained from different independent tests, such as the supernovae luminosity distance, the scale of baryon acoustic oscillations (BAO) in the matter power spectrum as a function of $z$, the acoustic scale in the angular power spectrum of the cosmic microwave background (CMB) temperature fluctuations, etc.   The second diagnostic -- {\\em Acceleration probe} $\\bar q$ -- is constructed out of the Hubble parameter and the lookback time. Like $Om$ it does not depend upon the current value of the matter density. We apply $\\bar q$ to current data and show that it provides an independent test of the present acceleration of the Universe.  The plan of the rest of the paper is as follows. In Sec. 2 the $Om$ diagnostic is introduced for a flat as well as a spatially curved FRW background, and determined using SNLS and Union supernovae data. In Sec. 3 the $\\bar q$ diagnostic is introduced. In Sec. 4 both these diagnostics are determined from a combination of existing data sets including type Ia supernovae, BAO and WMAP5 CMB data. Sec. 5 contains our conclusions. ",
        "conclusions": "In this paper we propose two new diagnostics for determining the properties of dark energy. The first of these, $Om(z)$, is constructed from the Hubble parameter and results in the identity $Om(z) = \\om$ for LCDM. For other DE models $Om(z)$ is a function of the redshift. This allows one to construct a simple {\\em null test} to distinguish the cosmological constant from evolving DE. We find that $Om$ is a robust diagnostic whose value can be determined reasonably well even with current data. Unlike the equation of state $w$, $Om$ relies only on a knowledge of the Hubble parameter and depends neither on $H'(z)$ nor on quantities like the growth factor $\\delta(z)$. Errors in the reconstruction of $Om$ are therefore bound to be smaller than those appearing in the EOS. To diminish systematic uncertainties, different tests which allow for the possibility of reconstructing $H(z)$ from observational data need to be used. $Om$ can be determined using parametric as well as non-parametric reconstruction methods. It can shed light on the nature of dark energy even if the redshift distribution of supernovae is not uniform and the value of the dark matter density is not accurately known. The second diagnostic, {\\em acceleration probe} ${\\bar q}$, is the mean value of the deceleration parameter over a small redshift interval. The {\\em acceleration probe} depends upon the value of the Hubble parameter and the look-back time and can be used to determine the epoch when the universe began to accelerate.  We use the current type Ia SNe data in conjunction with BAO and CMB data to estimate the values of $Om(z)$ and ${\\bar q}$. Our results are consistent with LCDM but do not exclude evolving DE models including phantom and quintessence.  {\\em Note added in proof:} At $z \\ll 1$, \\beq \\frac{Om - \\Omega_{0m}}{1- \\Omega_{0m}} \\simeq 1+w_0~, \\nonumber\\\\ \\eeq which suggests that the $Om$ diagnostic can be used to probe $w_0$ from data available at low $z$, if $\\Omega_{0m}$ is accurately known."
    },
    "0807/0807.4532_arXiv.txt": {
        "abstract": "The  kinematics of  satellite  galaxies reflect  the  masses of  the extended dark matter haloes in which they orbit, and thus shed light on the mass-luminosity relation (MLR) of their corresponding central galaxies.  In  this paper we select  a large sample  of centrals and satellites from the Sloan Digital  Sky Survey (SDSS) and measure the kinematics  (velocity dispersions)  of the  satellite galaxies  as a function of the $r$-band  luminosity of the central galaxies.  Using the analytical framework presented in  Paper I, we use these data to infer {\\it  both} the  mean and  the scatter of  the MLR  of central galaxies, carefully  taking account of selection  effects and biases introduced  by  the   stacking  procedure.   As  expected,  brighter centrals on average reside in  more massive haloes.  In addition, we find  that  the scatter  in  halo masses  for  centrals  of a  given luminosity,  $\\sigma_{\\log  M}$,   also  increases  with  increasing luminosity.  As we demonstrate,  this is consistent with  $\\sigcen$, which reflects  the scatter in the  conditional probability function $\\plcm$,  being  independent of  halo  mass.   Our  analysis of  the satellite  kinematics  yields  $\\sigcen=0.16\\pm0.04$,  in  excellent agreement with constraints from clustering and group catalogues, and with predictions  from a semi-analytical model of  galaxy formation. We  thus  conclude  that  the  amount  of  stochasticity  in  galaxy formation, which is characterized by $\\sigcen$, is well constrained, is independent of  halo mass, and is in  good agreement with current models of galaxy formation. ",
        "introduction": "\\label{sec:intro}  According to  the standard framework of galaxy  formation, dark matter haloes form  gravitational potential wells in  which baryons collapse, dissipate    their    energy    and    form   stars    and    galaxies \\citep{White1978,Blumenthal1984}.   The  complex  process   of  galaxy formation and evolution is believed to  be governed by the mass of the dark  matter halo  in which  it occurs.  To understand  the  halo mass dependence of  this process, it  is important to  statistically relate the observable properties (e.g.  luminosity) of galaxies to the masses of their dark matter haloes.  The dark matter halo is bigger  in extent than the visible matter in a galaxy  due  to  the inability  of  dark  matter  to lose  energy  via dissipative  processes.  Therefore,  a tracer  population, distributed over  large distances  from the  center of  the halo  is  necessary to reliably measure  the mass of  a dark matter halo.  Satellite galaxies satisfy this requirement and their kinematics reflect the depth of the potential  well they orbit.   However, a  reliable measurement  of the kinematics requires a large number of satellites. Whereas a relatively large number  of satellites can  be detected in  cluster-sized haloes, only a handful  of satellites can be detected in  less massive haloes. This leads  to a  large uncertainty on  the kinematic  measurement for less  massive  haloes.  However,  under  the  assumption that  central galaxies of similar properties reside  in similar mass haloes, one can stack  such central  galaxies and  use  their satellites  to obtain  a statistical  measure  of the  kinematics  even  in  low mass  haloes. \\citep{Erickson1987,Zaritsky1993,Zaritsky1994,Zaritsky1997}.  During the  past few  years, data from  large galaxy  redshift surveys such as the Sloan Digital Sky Survey \\citep[SDSS;][]{York2000} and the Two degree Field Galaxy Redshift Survey \\citep[2dFGRS;][]{Colless2001} has become  available. The  kinematics of satellite  galaxies selected from these large datasets have  been used to infer the mass-luminosity relationship      (hereafter     MLR)     of      central     galaxies \\citep{McKay2002,Brainerd2003,    Prada2003,   van    den   Bosch2004, Conroy2005,Conroy2007}  and to  study the  density profiles  of dark matter    haloes    \\citep{Prada2003,Klypin2007}.     These    studies consistently  find that  the  velocity dispersion  of satellites  (and hence  the mass  of the  halo) increases  with the  luminosity  of the central   galaxy.    However,   \\citet[][hereafter   N08]{Norberg2008} demonstrated  a  quantitative disagreement  in  the  results of  these studies and  showed that  this disagreement is  largely due  to subtle differences  in the selection  criteria used  to identify  central and satellite  galaxies.   Therefore,  it  is crucial  to  understand  how selection  effects bias  the  MLR of  central  galaxies inferred  from satellite kinematics.  The MLR of  central galaxies is specified in  terms of the conditional probability  $\\pmlc$, which  describes the  probability for  a central galaxy with  luminosity $\\lc$ to reside in  a halo of mass  $M$. For a completely  deterministic  relation  between  halo  mass  and  central luminosity,   $\\pmlc=\\delta(M-M_0)$,   where   $\\delta$  denotes   the Dirac-delta  function   and  $M_0$  is  a   characteristic  halo  mass corresponding   to  centrals   of  luminosity   $\\lc$.   The  velocity dispersion,   $\\sigma(\\lc)$,  measured   by  stacking   centrals  with luminosity $\\lc$, then  translates into a mass $M_0$  according to the scaling relation $\\sigma^3 \\propto  M$. However, galaxy formation is a stochastic process  and the distribution  $\\pmlc$ is expected  to have non-zero  scatter. If this  scatter is  appreciable then  the stacking procedure  results in combining  the kinematics  of haloes  spanning a wide range  in halo mass.  This complicates the interpretation  of the velocity dispersion.  We addressed this  issue in More, van  den Bosch \\& Cacciato (2008; hereafter Paper I), where we investigated a method to  measure both  the  mean and  the  scatter of  the  MLR of  central galaxies  using  satellite   kinematics.  We  outlined  two  different weighting schemes  to measure  the velocity dispersion  of satellites, {\\it  satellite-weighting} and {\\it  host-weighting}, and  showed that the mean and  the scatter of the MLR can be  inferred by modelling the velocity dispersion measurements in  these two schemes simultaneously. In  this paper,  we  apply  this method  to  the spectroscopic  galaxy catalogue from  SDSS (Data Release 4)  in order to  determine both the mean and the scatter of the MLR of central galaxies.  The aim  of this  paper is twofold.  First, we  carry out a  series of tests on a  realistic mock galaxy catalogue to  validate our method to infer  the  MLR of  central  galaxies  from  a redshift  survey  using satellite   kinematics.     In   particular,   we    show   that   our central-satellite selection  criterion and  the method to  measure the kinematics reliably  recovers the true kinematics present  in the mock catalogue.  We  also show  that the  mean and the  scatter of  the MLR inferred from the kinematics match the corresponding true relations in the mock  catalogue. Second, we  repeat the analysis on  galaxies from the SDSS. In particular, we show that both the mean and the scatter of the  MLR increase  with the  luminosity of  the central  galaxy.  This demonstrates  that the scatter  in halo  masses is  non-negligible and must be  taken into  account when interpreting  measurements involving stacking.  This paper  is organized  as follows. In  Section~\\ref{sec:mockcat} we describe the construction  of the mock catalogue that  is used to test our  method  of  analysis.  In Section~\\ref{sec:selncrit}  we  briefly outline the iterative selection  criterion used to select centrals and satellites.  In  Section~\\ref{sec:veldisp} we  describe  and test  the method used to  measure the kinematics of satellites  as a function of the central luminosity.  The inference of the MLR  from the kinematics of  satellite galaxies requires  the knowledge  of the  number density distribution     of      satellites     within     a      halo.     In Section~\\ref{sec:nsatprof} we show that  this distribution can also be inferred from  the selected satellites.  In Section~\\ref{sec:model} we describe our model to  interpret the measured velocity dispersions and show that this  model is able to recover the true  mean and scatter of the   MLR  of   central   galaxies  from   the   mock  catalogue.   In Section~\\ref{sec:sdssapp}  we apply  our analysis  method to  a volume limited sample  from SDSS  and discuss the  results. We  summarize our results in Section~\\ref{sec:summary}. ",
        "conclusions": "\\label{sec:summary}  The  kinematics  of  satellite  galaxies  have  been  widely  used  to statistically relate the mean halo masses of central galaxies to their luminosities (Zaritsky \\etal 1993; Zaritsky \\& White 1994; Zaritsky \\etal 1997; McKay \\etal 2002; Prada \\etal 2003; Brainerd \\& Specian 2003; Norberg \\etal 2008). These studies use strict criteria to identify central and satellite  galaxies that reside preferentially in isolated environments. Following \\citet{van den Bosch2004}, we applied a relaxed but adaptive selection  criterion to a volume limited sample from  SDSS to  identify centrals  and  their satellites,  not only  in isolated environments  but also in  massive groups and  clusters which allowed us to study the kinematics  of satellites over a wide range of central galaxy luminosity.  We inferred both the mean  and the scatter of  the  mass-luminosity relationship  of  central  galaxies from  the kinematics of satellite galaxies.  Using a  realistic mock catalogue,  we thoroughly tested  the analysis method at  every step. We tested  the performance of  our criterion to identify central and satellite galaxies  and our method to measure the kinematics of  the selected satellites. We showed  that the kinematics recovered from the selected satellites are a fair representation of the true kinematics of  satellite galaxies present in the  mock catalogue. We  presented  an analytical  model  that  properly  accounts for  the selection biases  and showed that  the predictions of  this analytical model  are in  good  agreement  with the  measured  kinematics of  the selected satellites.  In Paper I,  we have shown that the  velocity dispersion of satellites can   be    measured   using   two    different   weighting   schemes: satellite-weighting  and  host-weighting.    We  have  demonstrated  a degeneracy between the  mean and the scatter of  the MLR obtained from either the satellite-weighted or the host-weighted velocity dispersion alone. However, we have also  shown that this degeneracy can be broken by using the  velocity dispersions in the two  schemes simultaneously. In  this  paper  we first  tested  our  method  using a  mock  galaxy catalogue. We fitted the measured satellite-weighted and host-weighted velocity dispersions  simultaneously using a parametric  model for the halo  occupation statistics  of  central and  satellite galaxies,  and demonstrated that we can reliably obtain confidence levels on the true mean and scatter of the mass-luminosity relation of central galaxies.  Next we applied  the above method to a volume  limited sample from the SDSS. The mean of the mass-luminosity relation increases as a function of  the central  host luminosity  indicating that,  as  expected, more massive haloes host brighter centrals. This result is in excellent quantitative agreement with a recent study by \\citet{Cacciato2008}, who use the abundance and clustering properties of galaxies in SDSS to constrain the CLF, and with the SAM of \\citet{Croton2006}. The satellite kinematics obtained in our study are consistent with a model in which $\\plcm$ has a constant scatter, $\\sigcen$, independent of the halo mass $M$. We obtain $\\sigcen=0.16\\pm0.04$ in excellent agreement with other independent measurements suggesting that the amount of stochasticity in galaxy formation is similar in haloes of all masses. This is also suggested by the SDSS group catalog of \\citet{Yang2008} and by the SAM of \\citet{Croton2006}."
    },
    "0807/0807.0127_arXiv.txt": {
        "abstract": "Recent ammonia (1,1) inversion line data on the Galactic star forming region Sgr~B2 show that the column density is consistent with a radial Gaussian density profile with a standard deviation of 2.75~pc.  Deriving a formula for the virial mass of spherical Gaussian clouds, we obtain $M_{\\rm vir} =1.9 \\times 10^6$~M$_\\odot$ for Sgr~B2.  For this matter distribution, a reasonable magnetic field and an impinging flux of cosmic rays of solar neighbourhood intensity, we predict the expected synchrotron emission from the Sgr B2 giant molecular cloud due to secondary electrons and positrons resulting from cosmic ray interactions, including effects of losses due to pion production collisions during diffusive propagation into the cloud complex.  We assemble radio continuum data at frequencies between 330 MHz and 230 GHz. From the spectral energy distribution the emission appears to be thermal at all frequencies.  Before using these data to constrain the predicted synchrotron flux, we first model the spectrum as free-free emission from the known ultra compact H{\\sc ii} regions plus emission from an envelope or wind with a radial density gradient, and obtain an excellent fit.  We thus find the spectrum at all frequencies to be dominated by thermal emission, and this severely constrains the possible synchrotron emission by secondary electrons to quite low flux levels.  The absence of a significant contribution by secondary electrons is almost certainly due to multi-GeV energy cosmic rays being unable to penetrate far into giant molecular clouds. This would also explain why 100~MeV--GeV gamma-rays (from neutral pion decay or bremsstrahlung by secondary electrons) were not observed from Sgr~B2 by EGRET, while TeV energy gamma-rays were observed, being produced by higher energy cosmic rays which more readily penetrate giant molecular clouds. ",
        "introduction": "Molecular clouds have long been studied as laboratories for star formation.  This has led to a wealth of information about the physical characteristics of clouds and their chemical makeup, usually obtained by observing emission/absorption lines of molecules such as CO, OH, etc., which reveal the presence of molecular gas. Most molecular cloud emission is thermal, from the H{\\sc ii} regions and/or dust emission.  Cosmic rays play an important role in molecular cloud evolution by partially ionizing even the cold molecular gas, thereby affecting, through ambipolar diffusion, the dynamics of cloud collapse by coupling the magnetic field to the partially ionized gas.  This in turn could result in amplification of the ambient magnetic field during cloud collapse, and give rise to the correlation found by \\citet{Crutcher1999} between average density of molecular gas of molecular clouds and their line-of-sight magnetic fields.  From VLA observations of Zeeman splitting of the  HI line \\citet{Crutcher1996} found a line of sight magnetic field for Sgr B2 of $B_{LOS}\\approx 0.5$~mG to apply to the outer envelope of the cloud complex, and this would suggest that the magnetic field amplitude could be even higher in the inner parts of the complex.  In fact, both \\citet{Lis1989} and \\citet{Crutcher1996} actually countenance average magnetic field strengths as high as $\\sim$2 mG for the Sgr B2 cloud on the basis of the theoretical prejudice that the cloud be magnetically supported against gravitational collapse; we cannot exclude that such field strengths may actually apply on large scales in the complex.    Recent studies of the ionization rate by \\citet{Tak2006} show that the cosmic ray ionization rate of dense molecular clouds in the Galactic Centre (GC) region may be as much as a factor $\\sim$10 higher than in molecular clouds in the solar neighbourhood.  Other evidence for a higher cosmic ray density in the GC region may come from the observation of 6.4~keV iron line emission.  Assuming that low energy cosmic rays are responsible for heating the molecular gas, \\citet{Yusef-Zadeh2007} estimated the energy densities of cosmic rays in GC molecular clouds to range from 19 eV~cm$^{-3}$ to $6\\times10^4$ eV~cm$^{-3}$, with 51 eV~cm$^{-3}$ for Sgr B2.  This is much higher than the energy density of cosmic rays in the solar neighbourhood which is $\\sim$1 eV~cm$^{-3}$.  Since ionisation is most effective for low-energy cosmic ray nuclei, and because cosmic ray energy spectra typically have an inverse power-law form, it is the cosmic rays with kinetic energies much less than 1~GeV/nucleon that are mainly responsible.  Hence the enhanced ionization in the GC region could be due to an overall enhancement of the cosmic ray flux there or to an additional, steep, low-energy component \\citep{Crockeretal2007}.  High densities of low-energy hadronic cosmic ray nuclei within regions of dense gas will result in enhanced $\\sim$100~MeV gamma-ray emission through enhanced pion production followed by $\\pi^0\\to\\gamma\\gamma$ decay and $\\pi^\\pm \\to \\mu^\\pm \\to e^\\pm$ decay followed by secondary electron bremsstrahlung.  The High Energy Stereoscopic System (HESS) recently completed a survey of the Galactic Centre \\citep{Aharonian2006}, and found a broad scale correlation between the $\\gamma$-ray emission and column density of the molecular gas, but required the GC region cosmic ray flux at $\\sim$10~TeV to be about 10 times higher than that observed in the solar neighbourhood.  Such an enhancement in the GC cosmic ray flux at GeV energies was not inferred by EGRET observations of 100~MeV gamma-rays from the central region of the Galaxy, indeed, it was explicitly noted by \\citet{Mayer-Hasselwander1998} that no localized excess associated with the Sgr B complex was detected by EGRET excluding the possibility of a significantly enhanced CR density in these clouds -- {\\it in the appropriate energy range}, of course.  One possible explanation of this, and the high ionization rate inferred by \\citet{Tak2006}, that has been suggested by \\citet{Crockeretal2007} is that in the Galactic Centre region there may be a steep component $E^{-2.7}$ with a normalization at $\\sim$10~GeV energies comparable to that locally in order to explain the EGRET result, and an even steeper lower energy component to explain the high ionization rate, and finally a flat $E^{-2.2}$ component negligible at GeV energies but giving a tenfold increase at $\\sim$10~TeV to explain the HESS data.  Such a flat component may arise as suggested by \\citet{Cheng2007} if periodic acceleration at Sgr~A* occurs when stars are tidally disrupted at a rate of $10^{-5}$ year$^{-1}$, and diffuse at a distance of $\\sim$500~pc before $pp$ losses steepen the spectrum on a timescale of $\\sim$$10^4$ years.  However, to also explain the broadband radio to gamma ray spectral energy distribution of the Sgr B region, with hadronic models they need a rather strong average magnetic field, viz., 2.2-3.7 mG.  The same interactions of cosmic ray nuclei within regions of dense gas which may lead to enhanced gamma-ray production (at least at TeV energies) should produce copious secondary electrons and positrons which may in turn produce synchrotron emission in GMC magnetic fields which are observed to be higher than elsewhere in the interstellar medium.  This possibility was the motivation for the present work as well as recent observations at 1.4~GHz and 2.4~GHz of the Sgr~B2 GMC \\citep{JonesSgrB2MainCmplx2008}, and of the dense cold starless cores G333.125-0.562 and IRAS 15596-5301 \\citep{Jones_StarlessCores_2008}.  The dense cores were chosen because, unlike Sgr~B2, they are well away from the central region of the Galaxy and would have a cosmic ray environment expected to be similar to that of the Solar region.  They are of much lower mass than the Sgr~B2 GMC, and their magnetic fields are unknown.  The non-detection of these dense cold starless cores in non-thermal emission was used to place upper limits on the magnetic fields of both of $\\sim$0.5~mG.   Here, we investigate whether one should expect to see radio synchrotron emission by secondary $e^\\pm$ from giant molecular cloud (GMC) complexes.  We shall compare our predictions of the expected synchrotron emission with our recent observations \\citep{JonesSgrB2MainCmplx2008} to draw conclusions about the cosmic ray environment around and within the Sgr~B2 GMC.  We chose the Sgr~B2 GMC for this study because of its large mass, its location in the central region of the Galaxy where the cosmic ray density may be higher than that in the solar neighbourhood, and its high magnetic field.  This was in spite of being aware of its complicated nature, and the likely difficulty in disentangling non-thermal from thermal emission in this source -- our search for any comparable molecular cloud with no star formation in the central region of the Galaxy was unsuccessful. ",
        "conclusions": "Radio continuum observations of the Sgr~B2 GMC, including new measurements at 1.4 GHz and 2.4 GHz, are discussed in detail in a separate paper by \\citet{JonesSgrB2MainCmplx2008}.  Since we want to compare the continuum emission in the Sgr~B2 region with our predictions for the region of the dense central region of the giant molecular cloud we use continuum flux estimates for the region 11~pc in size centred midway between Sgr~B2~(M) and Sgr~B2~(N).  Note that other papers may use quite different sizes for the Sgr~B2 region and hence quote very different fluxes.  In this complex region of the Galaxy, when assembling radio spectra from fluxes established by different groups, it is essential to use fluxes obtained over the same solid angle.  No evidence was found for diffuse, non-thermal emission out to $\\sim$11~pc from the centre of the Sgr~B2 GMC with limb-brightening as predicted for synchrotron emission by secondary electrons in the previous section, consistent with the results of \\citet{LangPalmerGoss2008} who find a thermal spectrum for Sgr B2 including its envelope.  Nor was there any suggestion of polarised emission characteristic of synchrotron emission.  However, the major H{\\sc ii} regions Sgr B2(M) and Sgr B2(N) showed up at all frequencies, and there is evidence of a strong non-thermal source to the south, which we have marked as ``Non-Thermal Source'' in Fig.~\\ref{sketch}, which is the subject of a separate paper \\citet{Jones2008b}.  Figure \\ref{flux_thermal} shows the spectral energy distribution for the central region of the Sgr~B2 GMC (excluding the Non-Thermal Source), and therefore includes the combined fluxes from the major H{\\sc ii} regions.   Sgr B2 is by far the most massive molecular cloud in the Galaxy and it is near the Galactic Centre which is almost certainly a region of enhanced cosmic rays.  When selecting Sgr~B2 for our study, we knew that separating out the thermal emission would be extremely difficult.  Our search for any comparable mass molecular cloud with no star formation was unsuccessful.  Before estimating upper limits to any synchrotron emission from secondary electrons in the Sgr~B2 GMC, we shall attempt to fit the observed spectral energy distribution solely by thermal emission processes.  A large number of UCHII regions have been observed at high frequencies, and so our first step will be to extrapolate their spectra to low frequencies.  This will be done under the assumption that each known UCHII region is a homogeneous sphere of ionized interstellar gas, and we shall sum the contributions from all known UCHII regions.  Clearly this will be an approximation as there will also be contributions from as-yet undiscovered UCHII regions, and because many of the  UCHII regions will not be homogeneous, having density gradients, e.g.\\ winds.  We shall find that the UCHII regions account for $\\sim$50\\% of the high frequency flux, and give a negligible fraction of the observed flux at low frequencies under these assumptions.  The second step will be to fit the residual flux as thermal emission.  The shape of the spectral energy distribution between 330~MHz and 1.4~GHz is suggestive of thermal emission from a region, or regions, with a density gradient, and we shall model it as a free-free emission from one, or many identical, single temperature winds.  While this is clearly unrealistic, the data available to us do not justify a more sophisticated treatment.   \\begin{figure} \\includegraphics[width=8cm]{fig9.eps}% \\caption{Observed fluxes summarized by \\citet{JonesSgrB2MainCmplx2008} from the central region of Sgr B2 complex including the major H{\\sc ii} regions but excluding the Southern Non-Thermal Source.  The flux from the known UCHII regions is indicated (chain curve), and the best fitting model of free-free emission from a constant temperature spherical envelope or wind with $n \\propto r^{-2}$ is shown by the dashed curve, and the solid curve gives the sum of the two thermal components.  } \\label{flux_thermal} \\end{figure}   \\subsection{Free-free emission from UCHII regions}  The emission at the higher frequencies (22 GHz and 43 GHz) is clearly thermal.  Regions with very low emission measures and high 22~GHz and 43 GHz fluxes could potentially affect the emission at low frequencies.  In order to investigate how much of the flux at these lower frequencies could be attributed to UCHII regions, we modeled the emission from the $\\sim$60 known individual compact and UCHII regions reported in \\citet{Gaume1995} and \\citet{dePree1998}.  This was achieved by ``bootstrapping'' the flux at the respective frequencies such that \\[ S_\\nu = \\sum_k S_{\\nu_i}^{(k)}\\left({\\nu\\over\\nu_i^{(k)}}\\right)^2 \\left({1- e^{-\\tau^{(k)}_\\nu}\\over 1- e^{-\\tau^{(k)}_{\\nu_i}} }\\right) \\] where $\\nu_i=22$~GHz or 43~GHz, the sum over the $\\sim$60 UCHII regions with the label $(k)$ relates to the $k$th UCHII region, and the frequency dependence of the thermal bremsstrahlung absorption coefficient is taken from \\citet{RybickiLightman1979}. The summed flux from the known UCHII regions has been added to Fig.~\\ref{flux_thermal}, and it can be seen that the total flux from these regions accounts for about 50\\% of the 23 and 230 GHz flux but their contribution below 3~GHz is negligible.  \\subsection{Free-free emission from envelopes or winds with density gradients}  Between 330~MHz and 1.4 GHz the spectrum may be fitted with a single power-law $S_\\nu \\sim \\nu^{0.6}$ characteristic of optically thick emission from a spherical envelope or wind with a density gradient of the form \\[ n_e(r)=n_i(r)= n_0 \\left({r \\over r_0}\\right)^{-2} \\] as described by \\citet{PanagiaFelli1975} who give the expected flux at low frequencies \\begin{eqnarray*} S_\\nu^{\\rm thick}&=&0.611\\left({n_0\\over \\mbox{1 cm$^{-3}$}}\\right)^{4/3} \\left({r_0\\over\\mbox{1~pc}} \\right)^{8/3}\\left({\\nu\\over \\mbox{10~GHz}}\\right)^{0.6}\\\\ &&\\left({T \\over 10^4\\mbox{~K}}\\right)^{0.1} \\left({d\\over\\mbox{1 kpc}}\\right)^{-2}. \\end{eqnarray*} For the case of  optically thin emission we can use \\[ \\int_{r_0}^\\infty 4\\pi r^2 \\,  n_e(r)\\, n_i(r) \\,dr = 4\\pi {n_0}^2  r_0^3 \\] together with the free-free emission coefficient and Gaunt factor $g(T,\\nu)$ from \\citet{RybickiLightman1979} to obtain the flux at high frequencies where it is expected to be optically thin \\begin{eqnarray*} S_\\nu^{\\rm thin}&=&1.6\\times 10^{-5}\\left({n_0\\over\\mbox{1 cm$^{-3}$}}\\right)^{2}\\left({r_0\\over\\mbox{1~pc}} \\right)^{3}\\\\ &&\\left({T \\over 10^4\\mbox{~K}}\\right)^{-0.5} g(T,\\nu)\\left({d \\over \\mbox{1 kpc}}\\right)^{-2}. \\end{eqnarray*} Taking the optical depth to be $\\tau_\\nu= S_\\nu^{\\rm thin}/S_\\nu^{\\rm thick}$ the flux from the envelope or wind is $S_\\nu=S_\\nu^{\\rm thick}(1 - e^{-\\tau_\\nu})$.  We fit this density-gradient model to the observed 330~MHz to 230~GHz fluxes (after subtracting the contributions of known UCHII regions as shown by chain curve in Fig.~\\ref{flux_thermal}) and this is shown by the dashed curve in Fig.~\\ref{flux_thermal}; the solid curve shows the sum of the two thermal components.  For a temperature $T=10^4$~K the best fitting parameters are $n_0=3.47\\times 10^7$ cm$^{-3}$ and $r_0=4.12\\times 10^{-3}$~pc. The high density and small size would indicate that the emission is likely to have come from winds off, or excited by, young stars within the H{\\sc ii} regions.  If the flux is due to $N$ separate identical objects, their wind parameters would be $n_0=3.47\\times 10^7 \\times N$ cm$^{-3}$ and $r_0=4.12\\times 10^{-3} / N$~pc in order to give the same total spectrum.  \\subsection{Synchrotron emission by secondary electrons }  In Fig.~\\ref{flux} we re-plot the spectral energy distribution and show the flux predictions for the synchrotron emission from secondary electrons.  Given that we have found no evidence of synchrotron emission from secondary electrons, and that the observed radio continuum emission is consistent with a thermal origin, we shall require that any flux of synchrotron emission from secondary electrons be well below the observed 330~MHz flux. For this, we shall somewhat arbitrarily adopt an upper limit of $S_\\nu$=1 Jy at 330 MHz for any non-thermal component.  \\begin{figure} \\includegraphics[width=8cm]{fig10a.eps}%  \\includegraphics[width=8cm]{fig10b.eps}% \\caption{Observed fluxes summarized by \\citet{JonesSgrB2MainCmplx2008} from the central region of Sgr B2 complex including the major H{\\sc ii} regions but excluding the Southern Non-Thermal Source plotted together with fluxes predicted for (a) $B_\\perp=0.8$mG, (b) $B_\\perp=0.1$mG, and a cosmic ray spectrum as in the solar neighbourhood.  Predicted fluxes are shown for diffusion as defined by Eq~\\ref{Eq:diffco_Gabici} with $\\chi$ as indicated.  The total estimated thermal flux from the major H{\\sc ii} regions is shown by the upper solid curve.  } \\label{flux} \\end{figure}  Comparing the predicted synchrotron flux in Fig.~\\ref{flux}(a) with our adopted limit, we find that for $B_\\perp$=0.8~mG and $f_{CR}$=1, the diffusion factor must be $\\chi<$0.02.  With a significantly lower magnetic field such as $B_\\perp$=0.1~mG (see Fig.~\\ref{flux}b), even $\\chi\\gg 1$ (unimpeded cosmic ray penetration) is allowed and for this case cosmic ray enhancement in the Galactic Centre region at multi-GeV energies up to a factor $f_{CR}<$2.5 higher than in the Solar neighbourhood is not excluded.  For this magnetic field, and more reasonable diffusion factors, the maximum allowed cosmic ray enhancement is 3 ($\\chi$=1), 3.9 ($\\chi$=0.1), 7.7 ($\\chi$=0.01). We emphasize that these are upper limits for cosmic ray enhancement and that there is no evidence for any cosmic ray enhancement at multi-GeV energies.  In fact for the higher magnetic field, i.e.\\ $B_\\perp$=0.8~mG, which is based on Zeeman splitting observations, the data suggest that there is no enhancement, or that cosmic rays at these low energies are unable to significantly penetrate into the Sgr~B2 GMC.  In conclusion, we have no evidence that synchrotron emission by electrons and positrons produced by cosmic ray interactions has been observed from the Sgr~B2 molecular cloud complex.  The most likely explanation for this is that, for reasonable diffusion models, cosmic rays with multi-GeV energies (that produce secondary electrons with the right energy to radiate at GHz frequencies in $\\sim$0.8~mG fields) cannot penetrate into the dense central regions of Sgr B2 GMC where much of the potential mass of target nuclei is located.  This exclusion is also the likely explanation for non-observation of the Sgr B2 GMC by EGRET because it is again the same multi-GeV energy protons producing pions in $pp$ collisions followed by $\\pi^0\\to \\gamma\\gamma$ that make an important contribution to 100~MeV to multi-GeV gamma-rays.  The observation of the Sgr B2 GMC by HESS \\citep{Aharonian2006} at TeV energies is consistent with more complete penetration of cosmic rays at higher energies into the dense central regions.  In choosing giant molecular clouds in the central region of the Galaxy for future investigation of their synchrotron emission by secondary electrons, one would look for a GMC with with a mass of a few $10^5$~M$_\\odot$, a lower central density than Sgr~B2, e.g.\\ $n_{H_2}$$\\sim$$10^4$~cm$^{-3}$ so that low energy cosmic rays may more easily penetrate it, a magnetic field above 0.1~mG and little star formation. We do not know of any, but such clouds may become apparent with the aid of new infrared surveys.   Finally, we emphasise that as we have no independent knowledge of the diffusion coefficient within the Sgr~B2 GMC, i.e.\\ $\\chi$ is unknown, we are unable to make a definitive statement about the enhancement of the low energy cosmic ray flux in the central region of the Galaxy relative to that in the Solar neighbourhood."
    },
    "0807/0807.2064_arXiv.txt": {
        "abstract": "We present the results of C$^{18}$O observations by the Nobeyama Millimeter Array toward dense clumps with radii of $\\sim 0.3$ pc in six cluster-forming regions including massive (proto)stars. We identified 171 cores, whose radius, line width, and molecular mass range from 0.01 to 0.09 pc, 0.43 to 3.33 km s$^{-1}$, and 0.5 to 54.1 $M_{\\sun}$, respectively. % Many cores with various line widths exist in one clump, and the index of the line width--radius relationship of the cores and the parental clump differs from core to core in the clump. This indicates that the degree of dissipation of the turbulent motion varies for each core in one clump. % Although the mass of the cores increases with the line width, most cores are gravitationally bound by the external pressure. In addition, the line width and the external pressure of the cores tend to decrease with the distance from the center of the clump, and these dependencies may be caused by the inner H$_2$ density structure of the clump that affects the physical properties of the cores. % Moreover, the number density of the cores and the number density of young (proto)stars have a similar relationship to the average H$_2$ density of the clumps. % Thus, our findings suggest that the cluster is formed in the clump through the formation of such multiple cores, whose physical properties would have been strongly related to the H$_2$ density structure and the turbulent motion of the clump. % ",
        "introduction": "Most stars in the galactic disk are formed in dense molecular gas within giant molecular clouds (GMCs) as members of clusters (Lada \\& Lada 2003). Thus, clusters must play a critical role in the origins of some of the most fundamental properties of the galactic stellar population. In addition, such clusters consist of stars of various masses, and the most massive stars exist at the centers of clusters (e.g., Raboud \\& Mermilliod 1998). These features suggest that the physical environment in the cluster-forming region would vary from area to area because the mass of formed stars should be closely related to the physical environment. However, clusters do not always include massive stars. Some clusters contain only intermediate-mass and/or low-mass stars in low-mass star-forming regions (e.g., Tachihara et al. 2002). Based on these features, it appears that although the conditions of massive star formation easily satisfy the physical conditions for cluster formation, the conditions of cluster formation do not necessarily provide the physical conditions for massive star formation.   The clusters appear to be continuously forming in the galactic disk and the direct study of their physical conditions and the processes leading to their formation is basically possible. However, discovering clusters at an early evolutionary stage via optical observations is difficult because such clusters are completely embedded in dense gas and dust. Therefore, observation by infrared instruments is necessary to study cluster formation. Recently, many young embedded clusters within molecular clouds have been discovered by infrared observations (e.g., Bica et al. 2003). On the basis of these observations, to investigate the physical conditions in the dense gas that forms rich clusters, dense molecular gas surveys were carried out on such young clusters that include massive stars, revealing that clusters are formed in gas systems with a size scale of a cluster ($\\sim 0.3$ pc) in GMCs (e.g., Lada 1992; Lada \\& Lada 2003; Saito et al. 2007) and these systems are usually called ``clumps.'' However, a few very compact gas systems having a size scale of $\\sim 0.03$ pc similar to typical (cold) cores in low-mass star-forming regions (e.g., Onishi et al. 2002; Umemoto et al. 2002) have been discovered with very high-angular-resolution observations using an interferometer on young clusters that include massive (proto)stars (e.g., Churchwell et al. 1992; Olmi et al. 1993). These systems have high temperatures ($> 100$ K), a high H$_2$ density ($10^6$--$10^7$ cm$^{-3}$), and a high luminosity ($> 10^4$ $L_{\\sun}$). These results indicate that these systems called ``hot cores,'' would include massive protostar(s) or massive young star(s).  Generally, not a large number of stars form in a core even in cluster-forming regions (e.g., van der Tak et al. 2005). Hence, many cores must form in one clump to create a cluster with a large number of stars. If this multiple-core system is formed by the fragmentation of the inner structure of the clump, a close relationship should exist between the physical parameters of the clump and the characteristics of the formed cluster. Saito et al. (2007) carried out molecular line observations using the Nobeyama 45-m telescope directed toward the clumps in several cluster-forming regions and suggested that a good relationship indeed exists between the average H$_2$ density of the clumps and the stellar number density of formed clusters in these clumps. In addition, they found a good correlation between the line width of the clumps and the maximum mass of formed stars. Saito et al. (2006) investigated the cores in three cluster-forming regions and found two types of cores in a single clump. The first type was a core with a large mass and a large turbulent motion, and the second type was a core with a small mass and a small turbulent motion. They found that massive (proto)stars are associated with the cores with a large mass and large turbulent motion. This suggests a theoretical dependence between the mass of the star formed in a core and the turbulent motion in the core (e.g., McKee \\& Tan 2003). According to the suggestion by McKee \\& Tan (2003), the results of Saito et al. (2006) indicate that massive and low-mass stars can be formed from one clump. Such observations suggest that to understand the mechanism of cluster formation, it is important to elucidate the relationships between the physical properties of cores and physical condition of the clump that includes the cores, and between the mass of the star formed in a core and the physical condition of the core. If the cores in a clump are formed by gravitational instability as the theory suggests, a strong connection must exist between the spatial distribution of the cores in the clump and the inner structure of the clump. However, no detailed relationships between the cores and the clumps have been demonstrated. In addition, although young clusters have a tendency to concentrate massive stars in the center, whether massive cores with large turbulent motion are formed in the center of the clump remains unclear.  Therefore, we must observe several clumps to obtain detailed physical parameters of both clumps and cores and reveal the relationships between the cores and the clumps themselves. To detect the cores in the clump, we must uncover the internal motion and the distribution of the H$_2$ column density in the clump. Thus, we need to carry out molecular line observations suitable for revealing velocity structures of the dense gas in the clump. Such observations are important to find a protocluster that consists of multiple cores --- a ``core cluster.''   We carried out C$^{18}$O observations toward nine IRAS point sources with young clusters containing massive stars. The C$^{18}$O line is a good candidate for tracing a large dynamic range in mass and column density in the dense gas because this line has a thin optical depth and a low critical density. In general, the C$^{18}$O line cannot be used to probe high-density regions ($> 10^{4}$ cm$^{-3}$) in starless cores because CO molecules are depleted onto dust grains in such a cold region with a temperature lower than the CO sublimation temperature of $\\sim 20$ K (e.g., Caselli et al. 1999; Bacmann et al. 2002). However, previous CO observations have revealed that CO lines trace high-density regions ($> 10^{4}$ cm$^{-3}$) well with high temperatures of $\\gtrsim 20$ K (e.g., Momose et al. 1998; Sch\\\"{o}ier et al. 2004). Therefore, if the temperatures of high-density regions ($> 10^{4}$ cm$^{-3}$) are higher than $\\sim 20$ K, C$^{18}$O lines can be used to probe these regions. Typically, the regions where massive stars are formed in GMCs maintain high temperatures of $\\gtrsim 20$ K, and thus C$^{18}$O molecules would not be depleted in such regions. In fact, the structures identified by the interferometric observations toward massive star-forming regions have a high average density of $> 10^{5}$ cm$^{-3}$ (e.g., Hunter et al. 2004; Saito et al. 2006). Therefore, C$^{18}$O observations can be used to reveal the distribution, mass, and kinetic motion of the dense gas in cluster-forming regions if this line is optically thin enough.   The identification of the star-forming and starless cores requires observation with high-angular resolution of several arc-seconds because massive star-forming regions are generally located at a distance of several kpc. However, interferometric observation of C$^{18}$O cores in massive star-forming regions has rarely been performed. To find dense cores in the clumps with cluster formation, we carried out C$^{18}$O molecular line observations with high-angular resolution toward nine cluster-forming regions using the Nobeyama Millimeter Array (NMA). The physical conditions of the clumps in these regions were already derived by Saito et al. (2007). Results of the interferometric observations toward three of these nine regions were already reported in Saito et al. (2006). Here we present the results of the observations toward the other six regions. We present our millimeter interferometric observations in $\\S$ 2. In $\\S$ 3, the observational results are presented for millimeter continuum and C$^{18}$O molecular line. We discuss the implications of our results for the physical condition of the dense cores in $\\S$ 4. In $\\S$ 5, we summarize the main results of our study. ",
        "conclusions": "We made high-angular resolution observations with the Nobeyama Millimeter Array (NMA) of six cluster-forming regions including massive (proto)stars, in the $J\\!=\\!1$--$0$ transition of C$^{18}$O molecular emission line and 98 GHz and 110 GHz continuum emission. We identified 171 C$^{18}$O cores in 12 clumps toward six cluster-forming regions including massive (proto)stars. Based on our observations, we studied the correlation between several physical quantities of the C$^{18}$O cores and clumps, and discussed the formation mechanism of star clusters. The main results of the present study are summarized as follows:  1. The ranges of the radius, line width, and molecular mass of the cores are 0.01--0.09 pc, 0.43--3.33 km s$^{-1}$, and 0.5--54.1 $M_{\\sun}$, respectively.  2. We identified 13 millimeter continuum sources (MCSs). Seven are associated with 2MASS sources and/or centimeter continuum sources (CCSs). Together with the one MCS that shows similar characteristics to the above seven MCSs, these eight MCSs would include massive protostars. The range of the mass of MCSs is 2.7--111 $M_{\\sun}$. In addition, we compared C$^{18}$O cores with MCSs, 2MASS sources, and CCSs. We identified 27 star-forming cores including eight massive-star-forming cores.  3. Some cores have various line widths in a single clump. The range of the index of the relationship between the line width and radius of the core and the parental clump differs from core to core. This implies that the degree of dissipation of the turbulent motion varies within a single clump.  4. For most cores, the virial masses are about twice larger than the LTE masses. This indicates that external pressure is necessary for the cores to be gravitationally bound. Based on the virial analysis, the distribution of the external pressure required to bind the cores is likely to be comparable to that of the pressure in the clump with an H$_2$ density structure of $\\sim r^{-\\gamma}$ ($\\gamma \\sim 1$--2).  5. The mass of the cores increases with the line width. This indicates that the average H$_2$ density of the cores would depend on the kinetic motion of the dense gas in the core because the range of the radius of the core is small.  6. The mass of the star depends on the line width and the mass of the parental core. In addition, we found that the star-formation efficiency (SFE) is roughly constant at $\\sim 50$\\% and that no clear dependence exists between the SFE and the mass of the formed stars.  7. The mass, line width, and external pressure of the core decrease with distance from the center of the clump, and a relationship exists between these decreases and the H$_2$ density structure of the clump. We suggest that the spatial distribution of physical parameters of the cores strongly depends on the H$_2$ density structure in the clump.  8. The number density of the cores and the number density of the young (proto) stars in the clump have a similar relationship to the average H$_2$ density of the clump. In addition, the core-formation rate (CFR) of the clump is estimated to be $\\sim 1$--20\\% per $\\sim 1 \\times 10^{6}$ yr and the CFR gradually increases with the average H$_2$ density of the clump."
    },
    "0807/0807.1499_arXiv.txt": {
        "abstract": "We report on the first results obtained from our development project of focusing gamma-rays ($>$60 keV) by using Laue lenses. The first lens prototype model has been assembled and tested. We describe the technique adopted and the lens focusing capabilities at about 100 keV. ",
        "introduction": "\\label{s:intro}  Experimental hard X--/gamma--ray astronomy is moving from direct sky-viewing telescopes to focusing telescopes. With the advent of focusing telescopes in this energy range, a big leap in sensitivity is expected, by a factor of 100-1000 with respect to the best non-focusing instruments of the current generation, either using coded masks (e.g., INTEGRAL/IBIS, Ref.~\\citenum{Ubertini03}) or not (e.g., BeppoSAX/PDS, Ref.~\\citenum{Frontera97}). A significant increase in angular resolution is also expected (from $\\sim 10$ arcmin of the mask telescopes like INTEGRAL IBIS to less than 1 arcmin).  The next generation of hard X--($<$100keV) and gamma--ray ($>$ 100 keV) focusing telescopes will make use of the Bragg diffraction technique, in the former energy range, from multilayer coatings (ML) in a reflection configuration (supermirrors), while in the latter, from mosaic--like crystals in a transmission configuration (Laue lenses).  The expected astrophysical issues that are expected to be solved with the advent of these telescopes are many and of fundamental importance. For what concerns the higher energy band ($>$100 keV), a summary of the main science goals are discussed in the context of a mission proposal, {\\em Gamma Ray Imager} (GRI), submitted to ESA in response to the first AO of the 'Cosmic Vision 2015--2025' plan \\cite{Knodlseder07}. For the astrophysical importance of the $>$100 keV band, see also Refs.~\\citenum{Frontera05a,Frontera06,Knodlseder06,Knodlseder07}. The GRI proposal was not admitted by ESA for further consideration for a launch in 2017 (the first flight opportunity in the new ESA plan) mainly due to readiness problems of the Laue lenses.  Here we report on the status of our project HAXTEL (= HArd X-ray TELescope) devoted to  developing the technology for building broad energy passband (70/100--600 keV) Laue lenses, highlighting test results of the first lens prototype. ",
        "conclusions": "The results of the Laue lens development activity over the last few years have been reported and discussed (see Refs.~\\citenum{Frontera06,Frontera07} and references therein).  In short, by means of theoretical calculations and the development of a Monte Carlo code, we first established the best geometrical configuration of a Laue lens for space astronomy, the best crystal materials, and the crystal mosaic properties that give the best compromise between reflection efficiency and focusing capabilities. This theoretical study was followed by reflectivity measurements of mosaic crystal samples of Cu(111) \\cite{Pellicciotta06}, given that this material showed good expected performance for our goals and could be developed with the desired properties (e.g., Courtois et al.\\cite{Courtois04}).  Given that a Laue lens focusing telescope for gamma--ray astronomy is required to have a spherical--like shape in addition to a large size (meter scale), while the available mosaic crystal sizes are generally small (cm scale), the only viable solution is to build a lens made of many crystal tiles. This means that, in order to correctly focus the photons reflected from all crystal tiles toward the same point, the direction of the vector perpendicular to the mean lattice plane of each crystal tile has to intersect the lens axis, while its inclination with respect to the focal plane has to be equal to the Bragg angle $\\theta_B$ (see Fig.~1 in Ref.~\\citenum{Pellicciotta06}). The angle $\\theta_B$ also establishes the energy centroid of the reflected photons, and it depends on the distance $r$ of the tile center from the lens axis and on the focal length $f$. For correct focusing, $\\theta_B = 1/2 \\arctan{(r/ f)}$. Once the crystal material is established, the Bragg angle increases with $r/f$, while the energy of the focused photons decreases with $r/f$. More generally, once the focal length is established, the outer  and inner lens radii, $r_{max}$ and $r_{min}$, depend on the nominal energy passband of the lens ($E_{min}$, $E_{max}$) and on the crystal lattice spacing: higher $d_{hkl}$ implies lower radii \\cite{Frontera06}. For a given crystal material, outer lens radius and focal length, the minimum energy that can be focused is established.  For a given inner and outer radius, the lens passband  can be established by  the combination of the different crystal materials used. In addition to Copper, other candidate materials, which exhibit high reflectivity and for which the mosaic technology has been developed, are becoming available for Laue lenses, such as mosaic Ge, Si$_{1-x}$Ge$_x$, with a composition--gradient \\cite{Abrosimov05}, or mosaic Gold \\cite{Barriere08}.  Mosaic spread and crystal thickness are the most crucial parameters for the optimization of the lens performance. The crystal thickness issue was investigated by us in relation to the lens weight (e.g., Ref.~\\citenum{Pisa05b}). Concerning the mosaic spread, a higher spread gives a larger effective area, but also produces a larger defocusing of the reflected photons in the focal plane. We introduced a parameter that characterizes a Laue lens and its sensitivity: the focusing factor $G$ \\begin{equation} G = f_{ph} \\frac{A_{eff}}{A_d} \\end{equation} in which $A_{eff}$ is the effective area of the lens and $A_d$ is the area of the focal spot which contains a fraction $f_{ph}$ of photons reflected by the lens. The higher $G$ is, the higher is the lens sensitivity (minimum detectable source flux $F_{min} \\propto 1/G$).  The spatial distribution of the mosaic crystal tiles influences the shape of the effective area vs. energy. For low focal lengths ($< \\sim10$~m), the best crystal tile disposition is an Archimedes' spiral \\cite{Pisa04}, that provides a smooth behavior of the lens effective area $A_{eff}$ with energy. For high focal lengths, a ring crystal disposition is equally good.  The accuracy of the crystal tile positioning in the lens depends not only on the mosaic spread but also on the focal length. Higher focal lengths require a higher positioning accuracy. At the current stage of development, an accurate positioning of the crystals in the lens is one of the major problems to be faced.  From Monte Carlo (MC) studies, we have derived the properties of various lens configurations, such as their Point Spread Functions (PSF), for both on--axis and off--axis incident photons. We show in Fig.~\\ref{f:psf} the expected  PSF of a Cu (111) gamma--ray Laue lens with a focal length of 40 m and a mosaic spread of 1 arcmin, in which a perfect positioning of the composing crystals is assumed. The  expected  angular resolution of this lens is of the order of or even better than 1 arcmin.   \\begin{figure} \\begin{center} \\includegraphics[angle=0, width=0.4\\textwidth]{on-axis_PSF.eps} \\includegraphics[angle=0, width=0.4\\textwidth]{off-axis_PSF.eps} \\end{center} \\caption{Expected response function PSF of a Laue lens {\\em Left panel:} on-axis PSF.  {\\em Right panel:} off--axis PSF, in which three sources are simulated at 0, 2 and 4 arcmin off--axis. See text for details.} \\label{f:psf} \\end{figure}"
    },
    "0807/0807.4043_arXiv.txt": {
        "abstract": "It has been revealed recently that, in the scale free range, i.e. from the scale of the onset of nonlinear evolution to the scale of dissipation, the velocity and mass density fields of cosmic baryon fluid are extremely well described by the self-similar log-Poisson hierarchy. As a consequence of this evolution, the relations among various physical quantities of cosmic baryon fluid should be scale invariant, if the physical quantities are measured in cells on scales larger than the dissipation scale, regardless the baryon fluid is in virialized dark halo, or in pre-virialized state. We examine this property with the relation between the Compton parameter of the thermal Sunyaev-Zel'dovich effect, $y(r)$, and X-ray luminosity, $L_{\\rm x}(r)$, where $r$ being the scale of regions in which $y$ and $L_{\\rm x}$ are measured. According to the self-similar hierarchical scenario of nonlinear evolution, one should expect that 1.) in the $y(r)$-$L_x(r)$ relation, $y(r)=10^{A(r)}[L_{\\rm x}(r)]^{\\alpha(r)}$, the coefficients $A(r)$ and $\\alpha(r)$ are scale-invariant; 2.) The relation $y(r)=10^{A(r)}[L_{\\rm x}(r)]^{\\alpha(r)}$ given by cells containing collapsed objects is also available for cells without collapsed objects, only if $r$ is larger than the dissipation scale. These two predictions are well established with a scale decomposition analysis of observed data, and a comparison of observed $y(r)$-$L_x(r)$ relation with hydrodynamic simulation samples. The implication of this result on the characteristic scales of non-gravitational heating is also addressed. ",
        "introduction": "Scaling relation of dimensional quantities is very powerful to reveal the dynamical feature of various physical systems. There has been a considerable effort devoting to study the correlations and scaling laws of various observable quantities of galaxy clusters. Since virialized self-gravitational system is characterized by one parameter, mass or virial temperature, one can find a set of scaling relations among mass, size, X-ray luminosity, temperature, and Compton parameter of Sunyaev-Zel'dovich (SZ) effect if the velocity and mass density fields of baryon fluid in clusters are assumed to be similar to the virialized dark matter halos \\citep{1986MNRAS.222..323K}. Observed data of galaxy clusters did yield scaling relations \\citep{1991MNRAS.252..428E,1993ApJ...412..479D,1999ApJ...524...22W, 2000MNRAS.315..356H,2000ApJ...538...65X,2005MNRAS.357..279C}. However, observed scaling relations generally do not support the predictions given by the baryon-dark matter similarity of virialized dark halos \\citep{2000MNRAS.315..356H,2000MNRAS.315..689L}.  Since Newtonian gravity is scale-free, the self gravitational system of collisionless dark matter shows scaling behavior if the power spectrum of initial density perturbations is scale-free. These scaling is regardless of whether the underlying gravitational field is virialized \\citep{1980lssu.book.....P}. Thus, if the velocity and mass density fields of cosmic baryon matter are given by a similar mapping of the fields of dark matter, one may expect the scaling relations of clusters. However, the similar mapping assumption is correct only in linear regime \\citep{1992A&A...266....1B}, but is baseless in nonlinear regime \\citep{1989RvMP...61..185S}. The nonlinear evolution of cosmic baryon fluid leads to statistically decouple of the fluid from dark matter. The statistical properties of the velocity and mass density fields of baryon fluid do show deviation from the underlying dark matter field \\citep{2004ApJS..154..475P, 2005ApJ...623..601H,2005ApJ...625..599K}.  Nevertheless, it has been pointed out by \\cite{1989RvMP...61..185S}: the dynamics of cosmic baryon fluid in the expanding universe is scale-free, i.e. no preferred special scales can be identified in the range from the onset of nonlinear evolution down to the length scale of dissipation. It likes fully developed turbulence in inertial range. This idea recently received substantial developments. With the hydrodynamic simulation sample of the concordance $\\Lambda$CDM model, the velocity field of cosmic baryon fluid is found to be extremely well described by She-Leveque's (SL) scaling formula \\citep{1994PhRvL..72..336S} in the ``inertial range'' \\citep{2006PhRvL..96e1302H}. The SL formula is considered to be the basic statistical features of the scale-free evolution of fully developed turbulence. Moreover, the SL formula comes from self-similar log-Poisson hierarchy, which is related to the hidden symmetry of the Navier-Stokes equations \\citep{1994PhRvL..73..959D, 1995PhRvL..74..262S}. Very recently, it has been shown that the clustering of the mass density field of the cosmic baryon fluid can indeed be well described by a log-Poisson hierarchical cascade \\citep{2008ApJ...672...11L}. All the scaling relations and non-Gaussian features predicted from the log-Poisson hierarchy are in very good agreement with the hydrodynamic simulation samples.  These results indicate that, in the scale-free range, the nonlinear evolution of cosmic baryon fluid reaches a statistically quasi-steady state similar to a fully developed turbulence. For turbulence of incompressible fluid, the fluid undergoes a self-similar hierarchical evolution from largest to the smallest eddies and finally dissipates into thermal motion. For cosmic baryon fluid, the clustering on different scales can also be described by a self-similar hierarchy, and the fluid finally falls and dissipates into thermal motion.  This scenario motives us to investigate the scaling relations of clusters from the self-similar hierarchy of cosmic baryon fluid. If the observed scaling relations come from the self-similar hierarchy, one can expect that 1.) the relations of dimensional quantities should be scale-free, i.e. all the scale-dependent coefficients of the scaling relations are scale-invariant; 2.) the relations should be held only if the scales of considered regions are larger than Jeans length, regardless of whether the underlying gravitational field is virialized, i.e. the relations given by cells containing collapsed objects is also available for cells without collapsed objects, only if the scale of cells is larger than the dissipation scale.  Other relevant motivation comes from the non-gravitational heating of baryon gas of clusters. In order to solve the deviation from the similarity of virialized dark halos, various models of non-gravitational heating and cooling of baryonic gas have been proposed \\citep[e.g.][]{1999A&A...347....1V,2001ApJ...546...63T,2002ApJ...576..601V, 2003ApJ...588..704Z,2003ApJ...584...34X,2007ApJ...668....1N}. Since these cooling and heating may introduce characteristic scales, the self-similar hierarchy will no longer work on these characteristic scales. Therefore, it would be worth to detect the scale on which the above-mentioned two predictions to be broken.  We study these properties with the relation between the Compton parameter $y$ of SZ effect and X-ray luminosity $L_{\\rm x}$. The thermal SZ effect is due to the inverse Compton scattering of cosmic microwave background (CMB) photons by hot electrons of baryon fluid. The Compton parameter $y$ depends on the pressure of electron gas \\citep{1969Ap&SS...4..301Z,1980ARA&A..18..537S}. There are many works on the $y$-$L_x$ relation \\citep[e.g.][]{2004MNRAS.348.1401D, 2007ApJ...668..772M,2007arXiv0708.0815B}. We will, however, focus on the above-mentioned two points, which have not yet been addressed right now.  The outline of this paper is as follows. \\S 2 presents the scaling relations $y=10^{A(r)}L_x^{\\alpha(r)}$ with observed samples, and shows that $A(r)$ and $\\alpha(r)$ are scale-invariant. \\S 3 describes the hydrodynamic cosmological simulation samples. The comparison of the scaling relations of simulation samples with observed results is presented in \\S 4. The conclusions and discussion are given in \\S 5. ",
        "conclusions": "The scaling studied in this paper comes from the scale-free dynamics of cosmic baryon fluid in the expanding universe. That is, there is no preferred special scales can be identified in the range from the onset of nonlinear evolution down to the characteristic length of dissipation. The clustering evolution is described by self-similar hierarchy of baryon fluid. This self-similar hierarchy is different from the gravitational self-similarity of virialzed system, which is characterized by the mass of the system. On the other hand, the scaling from the scale-free dynamics is characterized by the scale range, on which physical quantities of the mass and velocity field of cosmic baryon fluid are measured. That is, the dimensional quantities measured in cells with size larger than the dissipation characteristic length should satisfy the same scaling relations, regardless whether the cells contains massive collapsed objects. The scaling will be broken if the scale is less than the characteristic length of dissipation.  We demonstrated this scaling with the $y(r)$-$L_{\\rm x}(r)$ relation of baryon field. The observed $y(r)$-$L_{\\rm x}(r)$ relation can be well reproduced with the hydrodynamical simulation. The important point is that the fitting of $y(r)$-$L_{\\rm x}(r)$ scaling is not based on the identified clusters in simulation samples, but using all cells on a given scale $r$, regardless whether the cell contains rich clusters with strong X-ray emission. In other words, the scaling relations $y(r) =10^{A(r)}[L_{\\rm x}(r)]^{\\alpha(r)}$ can be used to describe cells containing strong X-ray emission, $L_{\\rm x}>10^{43}$ erg s$^{-1}$, as well as very weak X-ray emission regions $L_{\\rm x}\\sim10^{40}$ erg s$^{-1}$. This result supports the self-similar hierarchical scenario of the clustering of baryon fluid.  The observed $y(r)$-$L_{\\rm x}(r)$ relations starts to show a deviation from simulation results on scales smaller than 0.78 h$^{-1}$ Mpc, which would be the characteristic scale of the dissipation. We calculated the Jeans lengths for each objects used in our statistics. All of the Jeans lengths are found to be less than 0.8 h$^{-1}$ Mpc. It is very well consistent with our result. Moreover, this result is also consistent with previous studies on the non-gravitational heating of clusters. The non-gravitational heating generally is considered to be due to the injection of hot gas and energy from SN and AGN. However, the significant effect of injecting hot gas and energy by SN is mostly arising from dwarf galaxies, i.e. on scales much smaller than clusters. The Ly$\\alpha$ observations of protoclusters shows that the feedback of AGN is not strong enough to heat the gas of clusters within comoving size 0.5 h$^{-1}$ Mpc \\citep{2003ApJ...584...45A}. Therefore, the effect of the SN and AGN heating would be dramatic only on scales smaller than about 0.5 h$^{-1}$ Mpc, but probably cannot heat gas within cells on scales $\\geq 1$ h$^{-1}$ Mpc to amount of the order of 1 keV per nucleon \\citep{1999ApJ...510L...1P,1999ApJ...524...22W}. Actually we find that scaling relations on scales $>0.78$h$^{-1}$ Mpc are still held after subtracting the contributions from the central part ($<0.2$ h$^{-1}$ Mpc) of clusters to both $L_{\\rm x}$ and $y$, as done in \\cite{1998ApJ...504...27M,2007ApJ...668..772M}.  As an application, the scaling relations of $y(r)$-$L_{\\rm x}(r)$ would be useful for estimating the contamination of SZ effect on CMB, which is important, especially, on small scales \\citep{2006MNRAS.369..645C}. Recent simulation has shown that the Planck project would be capable of probing $y$ on the order of $y=10^{-7}-10^{-8}$ \\citep{2005MNRAS.363...29D}. It might give a direct test on the universal scaling relations of $y(r)$-$L_{\\rm x}(r)$ given by the self-similar hierarchical evolution.  Our simulation of the baryon fluid is within the Eulerian framework. One can also study the nature of intermittency of fluid with Lagrangian point of view. In this approach, the hierarchical clustering can be tracked with Lagrangian trajectories. It has been found that the intermittent scaling is related to the long time correlations in the particle acceleration, namely, the random forces driving the particle motion is long range correlated \\citep{2002PhRvL..89y4502M}. It would be interesting to investigate the scaling in the Lagrangian scheme.  \\noindent{\\bf Acknowledgments.}  T.-J.Z. is supported by the Fellowship of the World Laboratory. H.-Y.W., T.-J.Z. and L.-L.F. acknowledge support from the National Science Foundation of China (grants 10473002, 10573036 and 10545002). T.-J.Z thanks Xiang-Ping Wu for his valuable discussion. This work was also partially supported by US NSF AST 05-07340."
    },
    "0807/0807.3663_arXiv.txt": {
        "abstract": "We account for all the image distortions relevant to weak gravitational lensing to second order. Besides the familiar shear, convergence, rotation and flexions, we find a new image distortion with two distinct descriptions, the twist and the turn. Like rotation, this distortion is not activated gravitationally to first order, but will be activated by systematic effects. We examine the rotational properties of twist and turn, and their effect on images in real and shapelet space. We construct estimators for the new distortion, taking into account the centroid shift which it generates. We then use these estimators to make first constraints on twist using the STAGES HST survey; we find that the mean twist estimator is consistent with zero. We measure correlation functions for our twist estimator on the survey, again finding no evidence of systematic effects. ",
        "introduction": "Weak gravitational flexion is a relatively new addition to the panoply of gravitational lensing effects, but has considerable potential for measuring substructure in the density distribution of matter in the Universe \\citep[see e.g.][]{2002ApJ...564...65G, 2005ApJ...619..741G, 2005NewAR..49...83I, 2006MNRAS.365..414B, 2008ApJ...680....1O, 2008A&A...485..363S}.  Flexion is proportional to third angular derivatives of the projected gravitational potential along the line of sight. As such, it is at the next order of differentiation compared to shear and convergence, which are the more studied weak lensing measures \\citep[see][ for an extensive review]{2001PhR...340..291B}. Since, as we shall see, there are two independent combinations of third derivatives, there are two different flexion effects: the 1-flexion, which is a vector distortion leading to objects being skewed; and the 3-flexion, which is a spin three distortion changing circular objects into trefoils.  Up until now, these have been the only known image distortions at this order. However, in this paper we will show that there is a further neglected image distortion at the flexion level, with two alternative descriptions which we call twist and turn for reasons which will become obvious. This distortion is not activated by gravity under the most straightforward approximations; but it will be activated by systematic effects. The latter are of great concern to weak lensing, so finding a further signature of systematics is potentially very valuable to upcoming lensing surveys.  In this paper we show how twist or turn arises, and account for why it has not been noticed before. We show how it affects images in real and shapelet space, and give details of how it can be measured with fairly straightforward estimators. We then measure twist for the first time using the Space Telescope A901/902 Galaxy Evolution Survey \\citep[STAGES,][]{2007AAS...21113220G}, a large mosaic observed with the Hubble Space Telescope (HST). We show that the twist is consistent with zero for STAGES on all scales, both in terms of its mean values and its correlation functions, incrementally adding confidence in the management of systematics for this survey.  The paper is organised as follows. In Section~\\ref{sect_lensing_first_order}, we recount the theory of image distortions in weak lensing at the more studied first order. We note that there is already a non-gravitational mode at this order; image rotation. We write the distortions in terms of Pauli matrices, which will give us the necessary clues for how to treat higher order distortions later.  In Section~\\ref{sect_lensing_second_order}, we extend the account to second order. We find that there are combinations of Pauli matrices orthogonal to those describing the conventional flexion degrees of freedom; these orthogonal combinations give twist and turn distortions. We are therefore able to write down for the first time the complete weak image distortion to second order, and show how twist and turn are related to one another observationally.  Section~\\ref{sect_twist_turn} describes the behaviour of twist/turn. The rotational properties of the distortion is worked out, and we find that is is a vector quantity. We show the impact of twist and turn on simple images; we find that they do not affect the shape of circularly symmetric images, but only images with non-zero ellipticity. We show explicitly the nature of twist and turn in shapelet space, proving that they have no impact on circularly symmetric sources, and derive how they move power between shapelet coefficients.  In Section~\\ref{sect_estimators} we go about finding practical estimators for measuring twist and turn. We derive simple estimators in shapelet space. Noting that like flexion, twist and turn affect the centroids of objects, we correct the estimators by constructing slightly more complicated expressions which take this shift into account. However, we will show that our estimators are not perfect; they would respond to flexion if it is present, and should therefore be treated as estimators of any second-order systematics, or of (real or systematic) twist on scales where flexion is negligible.  In Section~\\ref{sect_measurement} we use these estimators to constrain twist observationally for the first time, using the STAGES HST survey. We find that our twist estimator has a larger variance than flexion estimators, and that its mean value is consistent with zero in the STAGES data. We measure correlation functions for twist estimators, again finding that they are consistent with zero systematic in STAGES. We summarise our results and conclude in Section~\\ref{sect_summary}. ",
        "conclusions": "\\label{sect_summary}  In this paper, we have written down for the first time the full weak image distortion relevant to weak lensing, to second order. This involved the discovery of a new image distortion, the twist, which can be written in an alternative form as the turn.  We reviewed weak lensing distortions to first order, recouching the lens mapping in terms of Pauli matrices. We noted the existence of a non-gravitational mode, the rotation. This sets a precedent which we also see at second order.  We then extended the formalism to second order; at this point it became clear that the gradient of the rotation gives a new mode, which we call the turn. A further orthogonal mode in $D$ was found by seeking a combination of Pauli matrices orthogonal to all known modes; this new mode was called the twist. We showed that twist and turn can be written in terms of each other regarding their observational consequences, and were then able to write down the full image distortion mapping to second order.  We explored the properties of twist/turn, finding that it is a vector quantity. Its visual effect was shown, as was its impact in shapelet space. With the ladder operator formalism we showed that twist/turn has no effect on circularly symmetric objects, but only objects with non-zero ellipticity. We saw how twist/turn moves power from spin-2 modes to modes with spin-1 and spin-3.  Using our ladder operator forms for twist and turn, we found simple estimators for the distortions; however, twist/turn causes a centroid shift which needs to be taken into account. This leads to the inclusion of a common factor $B$ in the estimators for twist and turn. In addition, any flexion present would contribute to our simple estimator, which should therefore be used either as a means to detect any second-order systematic, or to measure real or systematic twist on scales where flexion is negligible.  We used these estimators to constrain twist for the first time in the HST STAGES survey. We noted that our estimator has a larger intrinsic noise scatter compared to 1- or 3-flexion, but its mean value across the survey is already at an interesting level for checking large flexion systematics. We found that mean twist is consistent with zero in this survey. We also measured twist correlation functions, and found that they too were consistent with zero.  The two quantities introduced in this study complete the set of distortions to second order. They will be of use in testing for systematic effects, and have a certain elegance of their own."
    },
    "0807/0807.4934_arXiv.txt": {
        "abstract": "We  use galaxy  groups selected from the  Sloan Digital  Sky Survey (SDSS) together with mass models  for individual groups to study the galaxy-galaxy  lensing signals expected  from galaxies  of different luminosities  and   morphological  types.   We   compare  our  model predictions with the observational results obtained from the SDSS by Mandelbaum et  al.  (2006)  for the same  samples of  galaxies.  The observational results  are well  reproduced in a  $\\Lambda$CDM model based on the WMAP 3-year  data, but a $\\Lambda$CDM model with higher $\\sigma_8$,  such  as  the  one  based  on  the  WMAP  1-year  data, significantly  over-predicts the  galaxy-galaxy lensing  signal.  We model,  separately, the contributions  to the  galaxy-galaxy lensing signals   from  different   galaxies:   central  versus   satellite, early-type  versus late-type,  and galaxies  in haloes  of different masses.   We also  examine how  the predicted  galaxy-galaxy lensing signal depends  on the shape,  density profile, and the  location of the central galaxy with respect to its host halo. ",
        "introduction": "According  to the  current paradigm  of structure  formation, galaxies form and reside inside extended cold dark haloes.  While the formation and evolution  of dark  matter haloes in  the cosmic density  field is mainly  determined  by  gravitational  processes,  the  formation  and evolution  of  galaxies involves  much  more  complicated, and  poorly understood processes,  such as radiative cooling,  star formation, and all  kinds  of feedback.   One  important  step  in understanding  how galaxies form and  evolve in the cosmic density  field is therefore to understand how  the galaxies  of different physical  properties occupy dark matter haloes of different masses.  Theoretically, the connection between galaxies and dark matter haloes can be studied using numerical simulations (e.g.,  Katz, Weinberg  \\& Hernquist 1996;  Pearce et  al. 2000; Springel 2005; Springel  et al.  2005) or semi-analytical models (e.g.  White \\&  Frenk 1991; Kauffmann et al.   1993, 2004; Somerville \\& Primack 1999;  Cole et al.  2000;  van den Bosch 2002; Kang  et al. 2005; Croton et al.  2006).  These approaches try to model the process of  galaxy  formation  from  first  principles.   However,  since  our understanding of  the relevant processes is still  poor, the predicted connection between  the properties of galaxies and  dark matter haloes needs  to be  tested against  observations.  More  recently,  the halo occupation model  has opened another  avenue to probe  the galaxy-dark matter halo  connection (e.g.  Jing,  Mo \\& B\\\"orner 1998;  Peacock \\& Smith 2000; Berlind  \\& Weinberg 2002; Cooray \\&  Sheth 2002; Scranton 2003; Yang, Mo \\& van den Bosch  2003; van den Bosch, Yang \\& Mo 2003; Yan, Madgwick \\&  White 2003; Tinker et al.  2005;  Zheng et al. 2005; Cooray 2006; Vale \\& Ostriker 2006; van den Bosch et al.  2007).  This technique uses the observed  galaxy luminosity function and clustering properties  to  constrain the  average  number  of  galaxies of  given properties that occupy a dark matter halo of given mass.  Although the method has  the advantage that  it can yield  much better fits  to the data  than the  semi-analytical models  or numerical  simulations, one typically  needs  to  assume  a  somewhat ad-hoc  functional  form  to describe the halo occupation model.  A more  direct way  of studying the  galaxy-halo connection is  to use galaxy groups\\footnote{In this paper, we refer to a system of galaxies as a group regardless of its richness, including isolated galaxies (i.e., groups with a single member) and rich clusters of galaxies.}, provided that they are defined as sets of galaxies that reside in  the same dark matter  halo.  Recently, Yang et al. (2005; 2007) have developed  a halo-based group finder that is optimized for grouping galaxies that reside in the same dark matter halo. Using mock galaxy  redshift surveys constructed  from the conditional luminosity function  model  (e.g.   Yang et al. 2003) and a semi-analytical model (Kang et al. 2005), it is found that this group finder is  very successful in associating galaxies with their common dark matter haloes (see Yang et al. 2007; hereafter Y07). The group finder  also performs reliably for poor systems, including isolated galaxies in small mass haloes, making it ideally suited for the study of the relationship between galaxies and dark matter haloes  over a wide range  of  halo masses. However, in order to interpret  the properties of the galaxy systems in terms of dark matter haloes, one needs to  know the halo mass associated with each of the groups. One approach commonly adopted is to  use some halo mass indicator (such as the total stellar mass or luminosity contained in member galaxies) to rank the groups.   With the assumption that the  corresponding halo masses have the same ranking and  that the mass function of the haloes associated  with groups is the same as that given  by  a model  of structure formation, one can assign  a halo  mass  to each of  the observed groups. This approach was adopted by  Y07 for the  group catalogue used in this paper. There are three potential problems with this approach. First, the approach is model-dependent, in  the sense that the assumption  of a different model of  structure formation will lead to  a different halo mass function,  and hence  assign different masses to the groups. Second, even if the assumed model of structure formation is correct, it  is still  not  guaranteed that the  mass assignment based on the ranking  of group stellar mass (or luminosity) is valid. Finally, even if  all groups are assigned with accurate halo masses, the question how dark  matter is distributed within the galaxy groups remains open. Clearly,  it is important to have independent mass measurements of the haloes associated with galaxy groups to test the validity of the mass estimates based on the stellar mass (luminosity) ranking.  Gravitational   lensing   observations,   which  measure   the image distortions of  background galaxies caused by  the gravitational field of the matter distribution in the foreground, provide a promising tool to  probe  the  dark  matter distribution directly. In  particular, galaxy-galaxy  weak lensing,  which focuses  on the image distortions around lensing galaxies, can be used to probe the distribution of dark matter  around   galaxies, hence  their  dark matter  haloes.   The galaxy-galaxy  lensing signal  produced  by individual  galaxies  is usually very weak, and so one has to stack the  signal from many lens galaxies  to have a statistical measurement.  The  first  attempt to detect such galaxy-galaxy lensing signal  was reported by Tyson et al. (1984). More recently, with the  advent of  wide and  deep surveys, galaxy-galaxy lensing can be studied  for lens galaxies  of different luminosities, stellar masses,  colors and  morphological types  (e.g. Brainerd et al. 1996; Hudson et al. 1998; McKay et al. 2001; Hoekstra et  al. 2003; Hoekstra 2004;  Sheldon et  al.  2004; Mandelbaum et  al. 2005, 2006; Sheldon et  al.  2007a;  Johnston et al. 2007;  Sheldon et al. 2007b; Mandelbaum et  al. 2008).  Given that galaxies reside in dark matter haloes,  these  results provide  important constraints on the mass distribution associated with galaxies in a statistical way.  In this paper, we use the galaxy groups of Y07 selected from the Sloan Digital Sky  Survey (SDSS), together  with mass models  for individual groups, to predict the galaxy-galaxy lensing signal expected from SDSS galaxies.   We compare  our model  predictions with  the observational results obtained by  Mandelbaum et al. (2006) for  the same galaxies. Our goal  is threefold. First, we  want to test whether  the method of halo-mass assignment to  groups adopted by Y07 is  reliable. Since the method provides a  potentially powerful way to obtain  the halo masses associated  with the  galaxy  groups, the  test  results have  general implications for  the study of  the relationship between  galaxies and dark  matter  haloes.   Second,  we  want to  examine  in  detail  the contributions  to  the  galaxy-galaxy  lensing signal  from  different systems, such as central  versus satellite galaxies, early-type versus late-type galaxies, and groups  of different masses. Such analysis can help us interpreting the observational results. Finally, we would like to  study how the  predicted galaxy-galaxy  lensing signal  depends on model  assumptions, such  as the  cosmological model  and  the density profiles of dark matter haloes. In  a companion paper (Cacciato et al. 2008,  hereafter C08), we  use the  relationship between  galaxies and dark matter  haloes obtained from the  conditional luminosity function (CLF)  modeling (Yang  et al.   2003; van  den Bosch  et al.  2007) to predict  the  galaxy-galaxy cross  correlation  and  to calculate  the lensing signal, while here we  directly use the observed galaxy groups and their galaxy memberships.  This paper  is organized as follows. In  Section \\ref{sec_analysis} we define  the statistical measure  that characterizes  the galaxy-galaxy lensing effect expected from the mass distribution associated with the galaxy groups.   We provide  a brief description of the  galaxy group catalogue  and the  models of  the mass distribution  associated with galaxy  groups in Section \\ref{sec_SDSS}. We  present our  results in Section \\ref{sec_results}      and     conclude     in     Section \\ref{sec_conclusions}.   Unless   specified  otherwise,  we   adopt a $\\Lambda$CDM cosmology  with parameters given by the  WMAP 3-year data (Spergel  et al.  2007,  hereafter WMAP3  cosmology) in  our analysis: $\\Omega_{\\rm m}=0.238$,  $\\Omega_{\\Lambda}=0.762$, and $h\\equiv H_0/(100 \\rm {km\\,s^{-1} \\,Mpc^{-1}})=0.73$, $\\sigma_8=0.75$ . ",
        "conclusions": "\\label{sec_conclusions}  In this paper  we model the galaxy-galaxy lensing  signal expected for SDSS  galaxies, using  the galaxy  groups  selected from  the SDSS  to represent the dark matter haloes  within which the galaxies reside. We use the properties of the dark halo population, such as mass function, density profiles  and shapes,  expected from the current $\\Lambda$CDM cosmogony to model the dark  matter distribution in each of the groups identified  in the  SDSS volume. The use  of the  real  galaxy groups allows us to predict  the galaxy-galaxy lensing signals separately for galaxies  of  different luminosity,  morphological  types,  and  in different  environments (e.g. central  versus satellite  galaxies). We check  the robustness of  our  model  predictions  by  changing  the assumptions about the dark  matter distribution in  individual groups (such as the shape, density  profile, and center offset of dark matter haloes), as  well as  the  cosmological  model  used to predict  the properties of the halo population.  We compare  our model predictions with the observational data of  Mandelbaum et al. (2006)  for similar samples of lens galaxies. Although there is some discrepancy for lens galaxies in the low-luminosity bins, the overall observational results can be well understood in the current $\\Lambda$CDM cosmogony. In particular, the observed results can be well reproduced in a $\\Lambda$CDM model with parameters based on the WMAP3 data, but a $\\Lambda$CDM model with a significantly higher $\\sigma_8$, such as the one  based on the WMAP1 data, significantly over-predicts the galaxy-galaxy lensing signal. Our results  also suggest that, once a correct model of structure formation is adopted,  the halo masses assigned to galaxy groups based on ranking their  stellar masses with the  halo mass function, are statistically  reliable. The results obtained imply that galaxy-galaxy  lensing  is  a powerful tool  to constrain both the mass  distribution associated with galaxies  and cosmological models. In the  future, when deep imaging surveys provide more sources  with high  image quality in the SDSS sky coverage, the galaxy-galaxy  lensing signals produced by the SDSS galaxies  can be estimated to much higher accuracy. The same analysis as presented here is expected to provide stringent constraints on the properties of the dark  matter haloes associated with galaxies  and galaxy systems, as well as on cosmological parameters."
    },
    "0807/0807.2723_arXiv.txt": {
        "abstract": "The neutrino-nucleus reaction cross sections of $^4$He and $^{12}$C are evaluated using new shell model Hamiltonians. Branching ratios of various decay channels are calculated to evaluate the yields of Li, Be, and B produced through the $\\nu$-process in supernova explosions. The new cross sections enhance the yields of $^7$Li and $^{11}$B produced during the supernova explosion of a 16.2 $M_\\odot$ star model compared to the case using the conventional cross sections by about $10\\%$. On the other hand, the yield of $^{10}$B decreases by a factor of two. The yields of $^6$Li, $^9$Be, and the radioactive nucleus $^{10}$Be are found at a level of $\\sim 10^{-11} M_\\odot$. The temperature of $\\nu_{\\mu,\\tau}$- and $\\bar{\\nu}_{\\mu,\\tau}$-neutrinos inferred from the supernova contribution of $^{11}$B in Galactic chemical evolution models is constrained to the $4.3-6.5$ MeV range. The increase in the $^7$Li and $^{11}$B yields due to neutrino oscillations is demonstrated with the new cross sections. ",
        "introduction": "Supernova (SN) explosions constitute one of several production sites of the relatively rare light elements Li, Be, and B. In SN environments these elements are produced through neutrino-nucleus reactions \\citep[the $\\nu$-process;][]{de78,wh90}. Neutrinos of all flavors are emitted in large numbers from a proto-neutron star, created during core-collapse of massive stars and the subsequent supernova explosion. Among the light elements, $^7$Li and $^{11}$B are abundantly produced through the $\\nu$-process \\citep{wh90,yt04,yk05,hk05}. Production of these light element isotopes in core-collapse supernovae (ccSNe) can contribute significantly to the increase in their abundances during Galactic chemical evolution \\citep[GCE;][]{fo00,rl00,rs00}.   Cross sections for neutrino-nucleus interactions are some of the most important data required to reliably estimate the $^7$Li and $^{11}$B yields in supernovae. The $\\nu$-process cross sections have been evaluated for a wide range of nuclear species in \\citet{wh90}. The data are tabulated in Hoffman \\& Woosley (1992, hereafter referred to as HW92) \\footnote{See http://www-phys.llnl.gov/Research/RRSN/nu\\_csbr/neu\\_rate.html}. Since the evaluation by HW92, further development of shell model calculations now enable us to more accurately evaluate these essential cross sections.   The $\\nu$-process cross sections are often presented as a function of neutrino temperature, based on averaging energy dependent cross sections over a Fermi-Dirac distribution of given temperature and chemical potential (for simplicity often assumed to be zero). However, it is more appropriate to consider the energy dependence as the primary information, as studies of SN neutrino transport show that their spectra do not exactly follow Fermi-Dirac distributions with zero-chemical potential \\citep[e.g.,][]{kr03}. Furthermore, when considering neutrino oscillations in SNe, the spectra are non-thermal after the neutrino flavors change, even if the Fermi-Dirac distribution approximates the spectra at the neutrino sphere reasonably well \\citep[e.g.,][]{ds00,tw01}.      The main purpose of this study is the re-evaluation of neutrino-nucleus reaction cross section for $^{12}$C and $^4$He using new shell-model Hamiltonians. We evaluate the branching ratios of many decay channels for light element species. Then, we evaluate the yields of the light elements, $^6$Li, $^7$Li, $^9$Be, $^{10}$Be, $^{10}$B, and $^{11}$B and discuss their production processes. We re-estimate the allowed range of the neutrino temperatures derived from constraints on the SN contribution of $^{11}$B in GCE models \\citep[following][]{yk05}. We also investigate the dependence of the neutrino oscillation parameters, i.e., mass hierarchy and the mixing angle $\\theta_{13}$, on the $^7$Li and $^{11}$B yields using the new cross sections.   The paper is organized as follows. In \\S 2 new cross sections for neutrino- $^{12}$C reactions are derived using the SFO and PSDMK2 Hamiltonians. New cross sections for neutrino-$^4$He reactions are evaluated using the WBP and SPSDMK Hamiltonians, as also shown in this section. The temperature dependence of the cross sections is discussed. The supernova explosion model and the supernova neutrino models employed are introduced and explained in detail in \\S 3. The nuclear reaction network used in this study is presented briefly. Light-element production mechanisms are discussed in \\S 4. The yields obtained using the new cross sections and the differences from those obtained with old cross sections are shown. The dependence of the light element yields on neutrino chemical potential is also discussed. The dependence of the yields of $^7$Li and $^{11}$B on the neutrino oscillation parameters, mass hierarchy and the mixing angle $\\theta_{13}$ is shown in \\S 5. The dependence of neutrino oscillation parameters on the $^7$Li/$^{11}$B ratio, the elemental abundance ratios of the light elements, is considered, and the possibility of constraining mass hierarchies and the mixing angle $\\theta_{13}$ is evaluated. Other effects on flavor-exchange of neutrinos in supernovae are discussed in \\S 6, and our conclusions are finally presented in \\S 7.   \\begin{figure*}[p] \\epsscale{1.05} \\plottwo{f1a.eps}{f1b.eps} \\plottwo{f1c.eps}{f1d.eps} \\caption{ Cross sections for $^4$He as a function of neutrino temperature $T_\\nu$. The neutrino energy spectrum is assumed to follow a Fermi-Dirac distribution with zero chemical potential. ($a$): neutral-current reactions with the WBP Hamiltonian, ($b$): charged-current reactions with the WBP Hamiltonians, ($c$): neutral-current reactions with the SPSDMK Hamiltonian, ($d$): charged-current reactions with the SPSDMK Hamiltonian. } \\label{crosshe4} \\vspace{10cm} \\end{figure*}    \\begin{figure*} \\epsscale{1.0} \\plottwo{f2a.eps}{f2b.eps} \\plottwo{f2c.eps}{f2d.eps} \\plottwo{f2e.eps}{f2f.eps} \\caption{ Averaged cross sections of $^{12}$C as a function of neutrino temperature, $T_\\nu$, for the SFO Hamiltonian. The neutrino energy spectrum is assumed to follow a Fermi-Dirac distribution with zero chemical potential. Top, meddle, and bottom panels correspond to neutral-current reactions, charged-current reactions for $\\nu_e$, and charged-current reactions for $\\bar{\\nu}_e$, respectively. The line with labeled $^{12}$C corresponds to the total decomposition rate of $^{12}$C, $\\sigma_{^{12}{\\rm C},\\nu}$ (see eq. [1]). } \\label{crosssc12sfo} \\end{figure*}   \\begin{figure*} \\epsscale{1.0} \\plottwo{f3a.eps}{f3b.eps} \\plottwo{f3c.eps}{f3d.eps} \\plottwo{f3e.eps}{f3f.eps} \\caption{ As in Fig.~\\ref{crosssc12sfo}, but for the PSDMK2 Hamiltonian. } \\label{crosssc12mk} \\end{figure*}   \\begin{deluxetable*}{ccccccccc} \\tabletypesize{\\tiny} \\tablecaption{ Neutrino-induced reaction cross sections of $^4$He in units of $10^{-42}$ cm$^2$ with the use of the WBP Hamiltonian. } \\tablehead{ \\colhead{$E_\\nu$} & & & & & & & & \\\\ \\colhead{(MeV)} & \\colhead{($\\nu,\\nu'p)^3$H} & \\colhead{($\\nu,\\nu'n)^3$He} & \\colhead{($\\nu,\\nu'd)^2$H} & \\colhead{($\\nu,\\nu'nnp)^1$H} & \\colhead{($\\nu_e,e^-p)^3$He} & \\colhead{($\\bar{\\nu}_e,e^+n)^3$H} & \\colhead{($\\nu_e,e^-pp)^2$H} & \\colhead{($\\bar{\\nu}_e,e^+nn)^2$H} } \\startdata 10.0 & 0.000E+00 & 0.000E+00 & 0.000E+00 & 0.000E+00 & 0.000E+00 & 0.000E+00 & 0.000E+00 & 0.000E+00 \\\\ 20.0 & 0.000E+00 & 0.000E+00 & 0.000E+00 & 0.000E+00 & 0.000E+00 & 0.000E+00 & 0.000E+00 & 0.000E+00 \\\\ 30.0 & 4.018E-02 & 3.829E-02 & 2.168E-11 & 3.538E-08 & 1.604E-01 & 1.264E-01 & 0.000E+00 & 0.000E+00 \\\\ 40.0 & 4.609E-01 & 4.425E-01 & 3.169E-04 & 9.746E-03 & 2.094E+00 & 1.556E+00 & 1.054E-04 & 1.587E-04 \\\\ 50.0 & 1.802E+00 & 1.738E+00 & 7.218E-02 & 1.730E-01 & 8.957E+00 & 5.992E+00 & 3.140E-01 & 2.211E-01 \\\\ 60.0 & 4.777E+00 & 4.620E+00 & 3.381E-01 & 7.782E-01 & 2.564E+01 & 1.529E+01 & 1.670E+00 & 1.053E+00 \\\\ 70.0 & 1.017E+01 & 9.856E+00 & 8.064E-01 & 1.991E+00 & 5.842E+01 & 3.108E+01 & 4.243E+00 & 2.409E+00 \\\\ 80.0 & 1.874E+01 & 1.818E+01 & 1.485E+00 & 4.021E+00 & 1.145E+02 & 5.453E+01 & 8.148E+00 & 4.167E+00 \\enddata \\tablecomments{Table~\\ref{tab4hewbp} is published in its entirety in the electronic edition of the {\\it Astrophysical Journal}. A portion is shown here for guidance regarding its form and content.} \\label{tab4hewbp} \\end{deluxetable*}   \\begin{deluxetable*}{ccccccccc} \\tabletypesize{\\tiny} \\tablecaption{ Neutrino-induced reaction cross sections of $^4$He in units of $10^{-42}$ cm$^2$ with the use of the SPSDMK Hamiltonian. } \\tablehead{ \\colhead{$E_\\nu$} & & & & & & & & \\\\ \\colhead{(MeV)} & \\colhead{($\\nu,\\nu'p)^3$H} & \\colhead{($\\nu,\\nu'n)^3$He} & \\colhead{($\\nu,\\nu'd)^2$H} & \\colhead{($\\nu,\\nu'nnp)^1$H} & \\colhead{($\\nu_e,e^-p)^3$He} & \\colhead{($\\bar{\\nu}_e,e^+n)^3$H} & \\colhead{($\\nu_e,e^-pp)^2$H} & \\colhead{($\\bar{\\nu}_e,e^+nn)^2$H} } \\startdata 10.0 & 0.000E+00 & 0.000E+00 & 0.000E+00 & 0.000E+00 & 0.000E+00 & 0.000E+00 & 0.000E+00 & 0.000E+00 \\\\ 20.0 & 0.000E+00 & 0.000E+00 & 0.000E+00 & 0.000E+00 & 0.000E+00 & 0.000E+00 & 0.000E+00 & 0.000E+00 \\\\ 30.0 & 6.992E-02 & 6.399E-02 & 0.000E+00 & 2.706E-04 & 2.045E-01 & 1.694E-01 & 0.000E+00 & 0.000E+00 \\\\ 40.0 & 7.360E-01 & 6.805E-01 & 3.413E-07 & 5.913E-03 & 2.709E+00 & 1.968E+00 & 0.000E+00 & 0.000E+00 \\\\ 50.0 & 2.879E+00 & 2.675E+00 & 6.340E-03 & 4.657E-02 & 1.211E+01 & 7.761E+00 & 2.338E-02 & 1.717E-02 \\\\ 60.0 & 7.633E+00 & 7.112E+00 & 5.841E-02 & 2.429E-01 & 3.557E+01 & 2.025E+01 & 2.516E-01 & 1.628E-01 \\\\ 70.0 & 1.616E+01 & 1.509E+01 & 1.800E-01 & 8.009E-01 & 8.209E+01 & 4.173E+01 & 8.673E-01 & 4.997E-01 \\\\ 80.0 & 2.944E+01 & 2.753E+01 & 3.897E-01 & 2.037E+00 & 1.614E+02 & 7.367E+01 & 2.062E+00 & 1.063E+00 \\enddata \\tablecomments{Table~\\ref{tab4hemk} is published in its entirety in the electronic edition of the {\\it Astrophysical Journal}. A portion is shown here for guidance regarding its form and content.} \\label{tab4hemk} \\end{deluxetable*}   \\begin{deluxetable*}{ccccccccccc} \\tabletypesize{\\tiny} \\tablecaption{ Neutrino-induced neutral-current reaction cross sections of $^{12}$C in units of $10^{-42}$ cm$^2$ based on the SFO Hamiltonian. } \\tablehead{ \\colhead{$E_\\nu$} &  & & & & & & & & &  \\\\ \\colhead{(MeV)} & \\colhead{$^{12}$C\\tablenotemark{a}} & \\colhead{$n$} & \\colhead{$p$} & \\colhead{$d$} & \\colhead{$t$} & \\colhead{$^3$He} & \\colhead{$^4$He} & \\colhead{$^6$He} & \\colhead{$^6$Li} & \\colhead{$^7$Li} } \\startdata 10.0 & 0.000E+00 & 0.000E+00 & 0.000E+00 & 0.000E+00 & 0.000E+00 & 0.000E+00 & 0.000E+00 & 0.000E+00 & 0.000E+00 & 0.000E+00 \\\\ 20.0 & 1.484E-03 & 1.482E-05 & 1.330E-03 & 0.000E+00 & 0.000E+00 & 0.000E+00 & 4.201E-04 & 0.000E+00 & 0.000E+00 & 0.000E+00 \\\\ 30.0 & 4.062E-01 & 8.046E-02 & 3.147E-01 & 1.329E-04 & 9.846E-05 & 1.462E-04 & 3.296E-02 & 0.000E+00 & 1.644E-13 & 5.526E-04 \\\\ 40.0 & 3.206E+00 & 7.104E-01 & 2.308E+00 & 1.259E-02 & 1.136E-02 & 1.458E-02 & 5.987E-01 & 2.569E-07 & 1.256E-04 & 5.501E-02 \\\\ 50.0 & 1.123E+01 & 2.576E+00 & 7.890E+00 & 6.209E-02 & 6.915E-02 & 8.576E-02 & 2.637E+00 & 9.178E-05 & 3.726E-03 & 2.645E-01 \\\\ 60.0 & 2.705E+01 & 6.346E+00 & 1.877E+01 & 1.765E-01 & 2.110E-01 & 2.574E-01 & 7.050E+00 & 7.484E-04 & 1.801E-02 & 7.078E-01 \\\\ 70.0 & 5.199E+01 & 1.240E+01 & 3.578E+01 & 3.657E-01 & 4.583E-01 & 5.560E-01 & 1.439E+01 & 2.094E-03 & 4.453E-02 & 1.442E+00 \\\\ 80.0 & 8.574E+01 & 2.068E+01 & 5.862E+01 & 6.252E-01 & 8.167E-01 & 9.890E-01 & 2.471E+01 & 3.979E-03 & 8.175E-02 & 2.486E+00 \\\\ \\hline \\hline \\colhead{$E_\\nu$} &  & & & & & & & & & \\\\ \\colhead{(MeV)} & \\colhead{$^8$Li} & \\colhead{$^9$Li} & \\colhead{$^7$Be} & \\colhead{$^9$Be} & \\colhead{$^{10}$Be} & \\colhead{$^8$B} & \\colhead{$^{10}$B} & \\colhead{$^{11}$B} & \\colhead{$^{10}$C} & \\colhead{$^{11}$C} \\\\ \\hline 10.0 & 0.000E+00 & 0.000E+00 & 0.000E+00 & 0.000E+00 & 0.000E+00 & 0.000E+00 & 0.000E+00 & 0.000E+00 & 0.000E+00 & 0.000E+00 \\\\ 20.0 & 0.000E+00 & 0.000E+00 & 0.000E+00 & 0.000E+00 & 0.000E+00 & 0.000E+00 & 0.000E+00 & 1.330E-03 & 0.000E+00 & 1.482E-05 \\\\ 30.0 & 0.000E+00 & 0.000E+00 & 4.367E-05 & 1.459E-04 & 7.516E-09 & 0.000E+00 & 1.331E-04 & 3.141E-01 & 0.000E+00 & 8.042E-02 \\\\ 40.0 & 0.000E+00 & 0.000E+00 & 9.063E-03 & 9.957E-03 & 5.222E-04 & 0.000E+00 & 1.386E-02 & 2.239E+00 & 5.356E-06 & 6.950E-01 \\\\ 50.0 & 1.277E-06 & 1.657E-12 & 5.537E-02 & 4.668E-02 & 5.365E-03 & 1.759E-08 & 7.018E-02 & 7.520E+00 & 1.747E-04 & 2.455E+00 \\\\ 60.0 & 6.709E-05 & 2.977E-08 & 1.657E-01 & 1.272E-01 & 1.904E-02 & 6.102E-06 & 1.920E-01 & 1.770E+01 & 9.144E-04 & 5.937E+00 \\\\ 70.0 & 2.643E-04 & 1.879E-07 & 3.596E-01 & 2.610E-01 & 4.343E-02 & 2.902E-05 & 3.891E-01 & 3.352E+01 & 2.295E-03 & 1.147E+01 \\\\ 80.0 & 5.593E-04 & 4.514E-07 & 6.455E-01 & 4.499E-01 & 7.859E-02 & 6.482E-05 & 6.616E-01 & 5.466E+01 & 4.216E-03 & 1.899E+01 \\enddata \\tablecomments{Table~\\ref{tab12c_ncsfo} is published in its entirety in the electronic edition of the {\\it Astrophysical Journal}. A portion is shown here for guidance regarding its form and content.} \\tablenotetext{a}{Decomposition rate. See \\S 2 and Eq. \\ref{csc12}.} \\label{tab12c_ncsfo} \\end{deluxetable*}   \\begin{deluxetable*}{cccccccccc} \\tabletypesize{\\tiny} \\tablecaption{ Neutrino-induced charged-current reaction cross sections for $\\nu_e$ of $^{12}$C in units of $10^{-42}$ cm$^2$ based on the SFO Hamiltonian. } \\tablehead{ \\colhead{$E_\\nu$} & & & & & & & & & \\\\ \\colhead{(MeV)} & \\colhead{$^{12}$C\\tablenotemark{a}} & \\colhead{$n$} & \\colhead{$p$} & \\colhead{$d$} & \\colhead{$t$} & \\colhead{$^3$He} & \\colhead{$^4$He} & \\colhead{$^6$Li} & \\colhead{$^7$Li} } \\startdata 10.0 & 0.000E+00 & 0.000E+00 & 0.000E+00 & 0.000E+00 & 0.000E+00 & 0.000E+00 & 0.000E+00 & 0.000E+00 & 0.000E+00 \\\\ 20.0 & 2.796E-01 & 0.000E+00 & 4.275E-05 & 0.000E+00 & 0.000E+00 & 0.000E+00 & 0.000E+00 & 0.000E+00 & 0.000E+00 \\\\ 30.0 & 5.516E+00 & 0.000E+00 & 5.781E-01 & 4.584E-07 & 0.000E+00 & 1.691E-05 & 9.536E-04 & 0.000E+00 & 0.000E+00 \\\\ 40.0 & 2.150E+01 & 1.330E-03 & 6.374E+00 & 2.059E-02 & 0.000E+00 & 6.608E-02 & 3.252E-01 & 1.396E-04 & 0.000E+00 \\\\ 50.0 & 5.655E+01 & 2.281E-02 & 2.658E+01 & 1.258E-01 & 2.127E-05 & 4.985E-01 & 2.159E+00 & 5.933E-03 & 6.552E-08 \\\\ 60.0 & 1.204E+02 & 1.056E-01 & 7.247E+01 & 4.043E-01 & 5.022E-04 & 1.685E+00 & 7.021E+00 & 3.680E-02 & 1.010E-04 \\\\ 70.0 & 2.209E+02 & 2.807E-01 & 1.537E+02 & 9.218E-01 & 1.981E-03 & 3.980E+00 & 1.630E+01 & 1.044E-01 & 6.282E-04 \\\\ 80.0 & 3.611E+02 & 5.595E-01 & 2.751E+02 & 1.703E+00 & 4.387E-03 & 7.625E+00 & 3.092E+01 & 2.108E-01 & 1.567E-03 \\\\ \\hline \\hline \\colhead{$E_\\nu$} & & & & & & & & & \\\\ \\colhead{(MeV)} & \\colhead{$^8$Li} & \\colhead{$^7$Be} & \\colhead{$^9$Be} & \\colhead{$^8$B} & \\colhead{$^{10}$B} & \\colhead{$^9$C} & \\colhead{$^{10}$C} & \\colhead{$^{11}$C} & \\colhead{$^{12}$N} \\\\ \\hline 10.0 & 0.000E+00 & 0.000E+00 & 0.000E+00 & 0.000E+00 & 0.000E+00 & 0.000E+00 & 0.000E+00 & 0.000E+00 & 0.000E+00 \\\\ 20.0 & 0.000E+00 & 0.000E+00 & 0.000E+00 & 0.000E+00 & 0.000E+00 & 0.000E+00 & 0.000E+00 & 4.275E-05 & 2.795E-01 \\\\ 30.0 & 0.000E+00 & 1.767E-04 & 0.000E+00 & 7.430E-04 & 5.340E-07 & 0.000E+00 & 4.584E-07 & 5.779E-01 & 4.937E+00 \\\\ 40.0 & 0.000E+00 & 1.345E-01 & 6.658E-07 & 5.803E-02 & 1.235E-02 & 0.000E+00 & 2.174E-02 & 6.147E+00 & 1.506E+01 \\\\ 50.0 & 0.000E+00 & 7.781E-01 & 2.669E-04 & 3.465E-01 & 1.246E-01 & 2.064E-05 & 1.328E-01 & 2.499E+01 & 2.966E+01 \\\\ 60.0 & 2.761E-10 & 2.296E+00 & 2.518E-03 & 1.102E+00 & 4.533E-01 & 3.532E-04 & 3.996E-01 & 6.718E+01 & 4.714E+01 \\\\ 70.0 & 3.717E-09 & 5.054E+00 & 7.803E-03 & 2.556E+00 & 1.099E+00 & 1.161E-03 & 8.766E-01 & 1.413E+02 & 6.558E+01 \\\\ 80.0 & 1.062E-08 & 9.328E+00 & 1.606E-02 & 4.855E+00 & 2.136E+00 & 2.402E-03 & 1.597E+00 & 2.515E+02 & 8.314E+01 \\enddata \\tablecomments{Table~\\ref{tab12c_ccesfo} is published in its entirety in the electronic edition of the {\\it Astrophysical Journal}. A portion is shown here for guidance regarding its form and content.} \\tablenotetext{a}{Decomposition rate. See \\S 2 and Eq. \\ref{csc12}.} \\label{tab12c_ccesfo} \\end{deluxetable*}   \\begin{deluxetable*}{ccccccccccc} \\tabletypesize{\\tiny} \\tablecaption{ Neutrino-induced charged-current reaction cross sections for $\\bar{\\nu}_e$ of $^{12}$C in units of $10^{-42}$ cm$^2$ based the SFO Hamiltonian. } \\tablehead{ \\colhead{$E_\\nu$} & & & & & & & & & & \\\\ \\colhead{(MeV)} & \\colhead{$^{12}$C\\tablenotemark{a}} & \\colhead{$n$} & \\colhead{$p$} & \\colhead{$d$} & \\colhead{$t$} & \\colhead{$^3$He} & \\colhead{$^4$He} & \\colhead{$^6$He} & \\colhead{$^6$Li} & \\colhead{$^7$Li} } \\startdata 10.0 & 0.000E+00 & 0.000E+00 & 0.000E+00 & 0.000E+00 & 0.000E+00 & 0.000E+00 & 0.000E+00 & 0.000E+00 & 0.000E+00 & 0.000E+00 \\\\ 20.0 & 7.004E-01 & 5.771E-03 & 0.000E+00 & 0.000E+00 & 0.000E+00 & 0.000E+00 & 0.000E+00 & 0.000E+00 & 0.000E+00 & 0.000E+00 \\\\ 30.0 & 5.252E+00 & 8.317E-01 & 1.149E-05 & 7.809E-04 & 1.223E-04 & 0.000E+00 & 3.422E-03 & 0.000E+00 & 0.000E+00 & 8.886E-04 \\\\ 40.0 & 1.592E+01 & 5.370E+00 & 5.521E-03 & 2.646E-02 & 1.832E-02 & 0.000E+00 & 1.375E-01 & 6.683E-05 & 3.032E-06 & 5.223E-02 \\\\ 50.0 & 3.482E+01 & 1.664E+01 & 4.072E-02 & 1.090E-01 & 1.185E-01 & 5.077E-05 & 6.556E-01 & 2.976E-03 & 7.866E-04 & 2.640E-01 \\\\ 60.0 & 6.290E+01 & 3.647E+01 & 1.325E-01 & 2.743E-01 & 3.579E-01 & 4.628E-04 & 1.759E+00 & 1.404E-02 & 5.310E-03 & 7.095E-01 \\\\ 70.0 & 9.941E+01 & 6.470E+01 & 2.854E-01 & 5.194E-01 & 7.521E-01 & 1.296E-03 & 3.494E+00 & 3.240E-02 & 1.341E-02 & 1.419E+00 \\\\ 80.0 & 1.422E+02 & 9.948E+01 & 4.916E-01 & 8.266E-01 & 1.290E+00 & 2.367E-03 & 5.803E+00 & 5.553E-02 & 2.362E-02 & 2.379E+00 \\\\ \\hline \\hline \\colhead{$E_\\nu$} & & & & & & & & & & \\\\ \\colhead{(MeV)} & \\colhead{$^8$Li} & \\colhead{$^9$Li} & \\colhead{$^7$Be} & \\colhead{$^9$Be} & \\colhead{$^{10}$Be} & \\colhead{$^{11}$Be} & \\colhead{$^{10}$B} & \\colhead{$^{11}$B} & \\colhead{$^{12}$B} & \\\\ \\hline 10.0 & 0.000E+00 & 0.000E+00 & 0.000E+00 & 0.000E+00 & 0.000E+00 & 0.000E+00 & 0.000E+00 & 0.000E+00 & 0.000E+00 \\\\ 20.0 & 0.000E+00 & 0.000E+00 & 0.000E+00 & 0.000E+00 & 0.000E+00 & 0.000E+00 & 0.000E+00 & 5.771E-03 & 6.946E-01 \\\\ 30.0 & 2.533E-03 & 0.000E+00 & 0.000E+00 & 1.223E-04 & 7.809E-04 & 1.149E-05 & 1.220E-16 & 8.309E-01 & 4.417E+00 \\\\ 40.0 & 7.545E-02 & 0.000E+00 & 0.000E+00 & 1.361E-02 & 2.765E-02 & 4.090E-03 & 8.826E-04 & 5.310E+00 & 1.044E+01 \\\\ 50.0 & 2.821E-01 & 5.011E-05 & 5.056E-12 & 7.359E-02 & 1.164E-01 & 2.238E-02 & 1.154E-02 & 1.627E+01 & 1.772E+01 \\\\ 60.0 & 6.551E-01 & 3.794E-04 & 3.581E-07 & 2.125E-01 & 2.872E-01 & 6.241E-02 & 4.018E-02 & 3.537E+01 & 2.536E+01 \\\\ 70.0 & 1.206E+00 & 9.657E-04 & 1.865E-06 & 4.418E-01 & 5.389E-01 & 1.282E-01 & 8.540E-02 & 6.241E+01 & 3.273E+01 \\\\ 80.0 & 1.922E+00 & 1.693E-03 & 4.074E-06 & 7.559E-01 & 8.586E-01 & 2.189E-01 & 1.438E-01 & 9.560E+01 & 3.950E+01 \\enddata \\tablecomments{Table~\\ref{tab12c_ccbsfo} is published in its entirety in the electronic edition of the {\\it Astrophysical Journal}. A portion is shown here for guidance regarding its form and content.} \\tablenotetext{a}{Decomposition rate. See \\S 2 and Eq. \\ref{csc12}.} \\label{tab12c_ccbsfo} \\end{deluxetable*}   \\begin{deluxetable*}{ccccccccccc} \\tabletypesize{\\tiny} \\tablecaption{ Neutrino-induced neutral-current reaction cross sections of $^{12}$C in units of $10^{-42}$ cm$^2$ based on the PSDMK2 Hamiltonian. } \\tablehead{ \\colhead{$E_\\nu$} & & & & & & & & & & \\\\ \\colhead{(MeV)} & \\colhead{$^{12}$C\\tablenotemark{a}} & \\colhead{$n$} & \\colhead{$p$} & \\colhead{$d$} & \\colhead{$t$} & \\colhead{$^3$He} & \\colhead{$^4$He} & \\colhead{$^6$He} & \\colhead{$^6$Li} & \\colhead{$^7$Li} } \\startdata 10.0 & 0.000E+00 & 0.000E+00 & 0.000E+00 & 0.000E+00 & 0.000E+00 & 0.000E+00 & 0.000E+00 & 0.000E+00 & 0.000E+00 & 0.000E+00 \\\\ 20.0 & 3.289E-04 & 2.518E-07 & 2.648E-04 & 0.000E+00 & 0.000E+00 & 0.000E+00 & 1.914E-04 & 0.000E+00 & 0.000E+00 & 0.000E+00 \\\\ 30.0 & 2.976E-01 & 4.222E-02 & 2.504E-01 & 4.759E-06 & 3.311E-05 & 1.551E-04 & 1.481E-02 & 0.000E+00 & 1.826E-24 & 3.170E-04 \\\\ 40.0 & 2.730E+00 & 5.182E-01 & 2.114E+00 & 7.058E-03 & 1.314E-02 & 1.741E-02 & 3.537E-01 & 1.316E-07 & 3.610E-05 & 4.747E-02 \\\\ 50.0 & 1.016E+01 & 2.077E+00 & 7.658E+00 & 5.211E-02 & 9.507E-02 & 1.138E-01 & 1.778E+00 & 9.817E-05 & 2.592E-03 & 2.633E-01 \\\\ 60.0 & 2.526E+01 & 5.382E+00 & 1.881E+01 & 1.827E-01 & 3.082E-01 & 3.575E-01 & 5.089E+00 & 1.037E-03 & 1.452E-02 & 7.470E-01 \\\\ 70.0 & 4.950E+01 & 1.081E+01 & 3.657E+01 & 4.194E-01 & 6.883E-01 & 7.909E-01 & 1.079E+01 & 3.135E-03 & 3.761E-02 & 1.571E+00 \\\\ 80.0 & 8.259E+01 & 1.833E+01 & 6.067E+01 & 7.514E-01 & 1.242E+00 & 1.423E+00 & 1.894E+01 & 6.070E-03 & 6.935E-02 & 2.765E+00 \\\\ \\hline \\hline \\colhead{$E_\\nu$} & & & & & & & & & & \\\\ \\colhead{(MeV)} & \\colhead{$^8$Li} & \\colhead{$^9$Li} & \\colhead{$^7$Be} & \\colhead{$^9$Be} & \\colhead{$^{10}$Be} & \\colhead{$^8$B} & \\colhead{$^{10}$B} & \\colhead{$^{11}$B} & \\colhead{$^{10}$C} & \\colhead{$^{11}$C} \\\\ \\hline 10.0 & 0.000E+00 & 0.000E+00 & 0.000E+00 & 0.000E+00 & 0.000E+00 & 0.000E+00 & 0.000E+00 & 0.000E+00 & 0.000E+00 & 0.000E+00 \\\\ 20.0 & 0.000E+00 & 0.000E+00 & 0.000E+00 & 0.000E+00 & 0.000E+00 & 0.000E+00 & 0.000E+00 & 2.648E-04 & 0.000E+00 & 2.518E-07 \\\\ 30.0 & 0.000E+00 & 0.000E+00 & 3.331E-06 & 1.525E-04 & 3.542E-09 & 0.000E+00 & 4.759E-06 & 2.501E-01 & 0.000E+00 & 4.221E-02 \\\\ 40.0 & 0.000E+00 & 0.000E+00 & 8.400E-03 & 1.041E-02 & 4.961E-04 & 0.000E+00 & 6.689E-03 & 2.050E+00 & 4.332E-06 & 5.005E-01 \\\\ 50.0 & 1.008E-05 & 2.562E-13 & 5.944E-02 & 6.500E-02 & 6.162E-03 & 1.030E-08 & 4.242E-02 & 7.248E+00 & 2.332E-04 & 1.929E+00 \\\\ 60.0 & 3.630E-04 & 6.660E-08 & 1.857E-01 & 2.028E-01 & 2.324E-02 & 1.494E-05 & 1.278E-01 & 1.753E+01 & 1.294E-03 & 4.860E+00 \\\\ 70.0 & 1.361E-03 & 5.621E-07 & 4.095E-01 & 4.464E-01 & 5.411E-02 & 7.867E-05 & 2.698E-01 & 3.374E+01 & 3.327E-03 & 9.602E+00 \\\\ 80.0 & 2.812E-03 & 1.436E-06 & 7.399E-01 & 7.994E-01 & 9.837E-02 & 1.795E-04 & 4.658E-01 & 5.561E+01 & 6.111E-03 & 1.612E+01 \\enddata \\tablecomments{Table~\\ref{tab12c_ncmk2} is published in its entirety in the electronic edition of the {\\it Astrophysical Journal}. A portion is shown here for guidance regarding its form and content.} \\tablenotetext{a}{Decomposition rate. See \\S 2 and Eq. \\ref{csc12}.} \\label{tab12c_ncmk2} \\end{deluxetable*}   \\begin{deluxetable*}{cccccccccc} \\tabletypesize{\\tiny} \\tablecaption{ Neutrino-induced charged-current reaction cross sections for $\\nu_e$ of $^{12}$C in units of $10^{-42}$ cm$^2$ based on the PSDMK2 Hamiltonian. } \\tablehead{ \\colhead{$E_\\nu$} & & & & & & & & & \\\\ \\colhead{(MeV)} & \\colhead{$^{12}$C\\tablenotemark{a}} & \\colhead{$n$} & \\colhead{$p$} & \\colhead{$d$} & \\colhead{$t$} & \\colhead{$^3$He} & \\colhead{$^4$He} & \\colhead{$^6$Li} & \\colhead{$^7$Li} } \\startdata 10.0 & 0.000E+00 & 0.000E+00 & 0.000E+00 & 0.000E+00 & 0.000E+00 & 0.000E+00 & 0.000E+00 & 0.000E+00 & 0.000E+00 \\\\ 20.0 & 2.363E-01 & 0.000E+00 & 2.663E-06 & 0.000E+00 & 0.000E+00 & 0.000E+00 & 0.000E+00 & 0.000E+00 & 0.000E+00 \\\\ 30.0 & 4.557E+00 & 0.000E+00 & 3.901E-01 & 1.076E-06 & 0.000E+00 & 6.638E-05 & 1.909E-03 & 0.000E+00 & 0.000E+00 \\\\ 40.0 & 1.802E+01 & 1.183E-03 & 5.282E+00 & 1.233E-02 & 0.000E+00 & 5.975E-02 & 3.247E-01 & 3.825E-05 & 0.000E+00 \\\\ 50.0 & 4.884E+01 & 2.703E-02 & 2.365E+01 & 1.266E-01 & 3.300E-05 & 5.269E-01 & 2.478E+00 & 3.558E-03 & 1.215E-07 \\\\ 60.0 & 1.071E+02 & 1.370E-01 & 6.690E+01 & 5.133E-01 & 1.394E-03 & 1.897E+00 & 8.601E+00 & 2.519E-02 & 1.707E-04 \\\\ 70.0 & 2.011E+02 & 3.824E-01 & 1.450E+02 & 1.313E+00 & 6.217E-03 & 4.621E+00 & 2.065E+01 & 7.503E-02 & 1.126E-03 \\\\ 80.0 & 3.346E+02 & 7.755E-01 & 2.632E+02 & 2.567E+00 & 1.418E-02 & 8.990E+00 & 3.981E+01 & 1.523E-01 & 2.829E-03 \\\\ \\hline \\hline \\colhead{$E_\\nu$} & & & & & & & & & \\\\ \\colhead{(MeV)} & \\colhead{$^8$Li} & \\colhead{$^7$Be} & \\colhead{$^9$Be} & \\colhead{$^8$B} & \\colhead{$^{10}$B} & \\colhead{$^9$C} & \\colhead{$^{10}$C} & \\colhead{$^{11}$C} & \\colhead{$^{12}$N} \\\\ \\hline 10.0 & 0.000E+00 & 0.000E+00 & 0.000E+00 & 0.000E+00 & 0.000E+00 & 0.000E+00 & 0.000E+00 & 0.000E+00 & 0.000E+00 \\\\ 20.0 & 0.000E+00 & 0.000E+00 & 0.000E+00 & 0.000E+00 & 0.000E+00 & 0.000E+00 & 0.000E+00 & 2.663E-06 & 2.363E-01 \\\\ 30.0 & 0.000E+00 & 7.641E-04 & 0.000E+00 & 1.012E-03 & 6.827E-06 & 0.000E+00 & 1.076E-06 & 3.892E-01 & 4.166E+00 \\\\ 40.0 & 0.000E+00 & 1.398E-01 & 1.573E-05 & 6.329E-02 & 1.358E-02 & 0.000E+00 & 1.245E-02 & 5.052E+00 & 1.268E+01 \\\\ 50.0 & 0.000E+00 & 9.564E-01 & 1.355E-03 & 3.774E-01 & 1.405E-01 & 3.162E-05 & 1.102E-01 & 2.176E+01 & 2.493E+01 \\\\ 60.0 & 4.085E-09 & 3.061E+00 & 9.674E-03 & 1.201E+00 & 5.190E-01 & 8.437E-04 & 3.828E-01 & 6.016E+01 & 3.953E+01 \\\\ 70.0 & 9.469E-08 & 7.057E+00 & 2.881E-02 & 2.776E+00 & 1.262E+00 & 3.024E-03 & 9.070E-01 & 1.287E+02 & 5.488E+01 \\\\ 80.0 & 2.948E-07 & 1.339E+01 & 5.859E-02 & 5.242E+00 & 2.442E+00 & 6.394E-03 & 1.733E+00 & 2.315E+02 & 6.945E+01 \\enddata \\tablecomments{Table~\\ref{tab12c_ccemk2} is published in its entirety in the electronic edition of the {\\it Astrophysical Journal}. A portion is shown here for guidance regarding its form and content.} \\tablenotetext{a}{Decomposition rate. See \\S 2 and Eq. \\ref{csc12}.} \\label{tab12c_ccemk2} \\end{deluxetable*}   \\begin{deluxetable*}{ccccccccccc} \\tabletypesize{\\tiny} \\tablecaption{ Neutrino-induced charged-current reaction cross sections for $\\bar{\\nu}_e$ of $^{12}$C in units of $10^{-42}$ cm$^2$ based on the PSDMK2 Hamiltonian. } \\tablehead{ \\colhead{$E_\\nu$} & & & & & & & & & & \\\\ \\colhead{(MeV)} & \\colhead{$^{12}$C\\tablenotemark{a}} & \\colhead{$n$} & \\colhead{$p$} & \\colhead{$d$} & \\colhead{$t$} & \\colhead{$^3$He} & \\colhead{$^4$He} & \\colhead{$^6$He} & \\colhead{$^6$Li} & \\colhead{$^7$Li} } \\startdata 10.0 & 0.000E+00 & 0.000E+00 & 0.000E+00 & 0.000E+00 & 0.000E+00 & 0.000E+00 & 0.000E+00 & 0.000E+00 & 0.000E+00 & 0.000E+00 \\\\ 20.0 & 5.917E-01 & 2.368E-03 & 0.000E+00 & 0.000E+00 & 0.000E+00 & 0.000E+00 & 0.000E+00 & 0.000E+00 & 0.000E+00 & 0.000E+00 \\\\ 30.0 & 4.382E+00 & 6.548E-01 & 2.563E-05 & 5.835E-05 & 6.821E-05 & 0.000E+00 & 1.867E-03 & 0.000E+00 & 0.000E+00 & 1.566E-04 \\\\ 40.0 & 1.363E+01 & 4.810E+00 & 1.195E-02 & 1.714E-02 & 3.012E-02 & 0.000E+00 & 1.608E-01 & 1.324E-04 & 1.989E-06 & 6.479E-02 \\\\ 50.0 & 3.079E+01 & 1.571E+01 & 6.777E-02 & 1.026E-01 & 1.932E-01 & 1.069E-04 & 8.939E-01 & 5.357E-03 & 5.897E-04 & 3.480E-01 \\\\ 60.0 & 5.724E+01 & 3.543E+01 & 2.043E-01 & 3.160E-01 & 5.845E-01 & 1.392E-03 & 2.582E+00 & 2.564E-02 & 4.866E-03 & 9.508E-01 \\\\ 70.0 & 9.248E+01 & 6.393E+01 & 4.246E-01 & 6.618E-01 & 1.227E+00 & 4.247E-03 & 5.325E+00 & 6.002E-02 & 1.306E-02 & 1.914E+00 \\\\ 80.0 & 1.344E+02 & 9.928E+01 & 7.099E-01 & 1.109E+00 & 2.095E+00 & 7.954E-03 & 8.998E+00 & 1.028E-01 & 2.338E-02 & 3.218E+00 \\\\ \\hline \\hline \\colhead{$E_\\nu$} & & & & & & & & & & \\\\ \\colhead{(MeV)} & \\colhead{$^8$Li} & \\colhead{$^9$Li} & \\colhead{$^7$Be} & \\colhead{$^9$Be} & \\colhead{$^{10}$Be} & \\colhead{$^{11}$Be} & \\colhead{$^{10}$B} & \\colhead{$^{11}$B} & \\colhead{$^{12}$B} & \\\\ \\hline 10.0 & 0.000E+00 & 0.000E+00 & 0.000E+00 & 0.000E+00 & 0.000E+00 & 0.000E+00 & 0.000E+00 & 0.000E+00 & 0.000E+00 \\\\ 20.0 & 0.000E+00 & 0.000E+00 & 0.000E+00 & 0.000E+00 & 0.000E+00 & 0.000E+00 & 0.000E+00 & 2.368E-03 & 5.893E-01 \\\\ 30.0 & 1.687E-03 & 0.000E+00 & 0.000E+00 & 5.622E-05 & 5.844E-05 & 2.555E-05 & 1.521E-08 & 6.546E-01 & 3.726E+00 \\\\ 40.0 & 6.814E-02 & 0.000E+00 & 0.000E+00 & 1.674E-02 & 1.892E-02 & 9.553E-03 & 9.832E-04 & 4.727E+00 & 8.713E+00 \\\\ 50.0 & 2.989E-01 & 1.056E-04 & 5.199E-12 & 9.082E-02 & 1.084E-01 & 3.839E-02 & 1.462E-02 & 1.516E+01 & 1.461E+01 \\\\ 60.0 & 7.605E-01 & 1.117E-03 & 1.744E-06 & 2.612E-01 & 3.058E-01 & 8.812E-02 & 5.364E-02 & 3.373E+01 & 2.064E+01 \\\\ 70.0 & 1.478E+00 & 3.086E-03 & 1.104E-05 & 5.345E-01 & 6.164E-01 & 1.578E-01 & 1.161E-01 & 6.031E+01 & 2.635E+01 \\\\ 80.0 & 2.424E+00 & 5.555E-03 & 2.528E-05 & 8.968E-01 & 1.027E+00 & 2.440E-01 & 1.955E-01 & 9.311E+01 & 3.154E+01 \\\\ \\enddata \\tablecomments{Table~\\ref{tab12c_ccbmk2} is published in its entirety in the electronic edition of the {\\it Astrophysical Journal}. A portion is shown here for guidance regarding its form and content.} \\tablenotetext{a}{Decomposition rate. See \\S 2 and Eq. \\ref{csc12}.} \\label{tab12c_ccbmk2} \\end{deluxetable*} ",
        "conclusions": "We evaluated the neutrino-nucleus reaction cross sections for $^4$He and $^{12}$C using new shell-model Hamiltonians. These cross sections are important for the yields of the light elements Li, Be, and B produced through the $\\nu$-process in SNe. We investigated the nucleosynthesis of the light elements in a SN model corresponding to SN 1987A using these new cross sections. We investigated the dependence of the light element yields on cross sections, neutrino energy spectra, and neutrino oscillations. We obtained the following results.   1. The neutrino-nucleus reaction cross sections are evaluated using WBP and SPSDMK Hamiltonians for $^4$He and SFO and PSDMK2 Hamiltonians for $^{12}$C. Main production channels of the $^{12}$C reaction are the nuclei $n$, $p$, $^4$He, $^{11}$B, and $^{11}$C. Production of $^6$Li, $^9$Be, $^{10}$Be, and $^{10}$B is also important.  2. For a given neutrino temperature set, the yields of $^7$Li and $^{11}$B with the cross sections of WBP+SFO model are larger than those of the cross sections in \\citet{wh90}. This is mainly due to larger cross sections of $^4$He($\\nu,\\nu'p)^3$H and $^4$He($\\nu,\\nu'n)^3$He.  3. Larger yields of $^6$Li and $^9$Be result from the new cross sections, with the exception of the yield of $^{10}$B, which is reduced. These changes reflect the difference of the cross sections of the $\\nu$-process for $^{12}$C and deuteron production branches of $^4$He. Radioactive $^{10}$Be is produced with abundance levels similar to $^6$Li. The channel for producing $^6$Li and $^{10}$Be have been evaluated in the $\\nu$-process of $^{12}$C.  4. The larger cross sections slightly decrease the range of acceptable neutrino temperature, as constrained by the  $^{11}$B abundance evolution during GCE. The range of the neutrino temperature is $4.3 {\\rm MeV} \\la T_{\\nu_{\\mu,\\tau}} \\la 6.5 {\\rm MeV}$ for WBP+SFO model and $4.0 {\\rm MeV} \\la T_{\\nu_{\\mu,\\tau}} \\la 6.0 {\\rm MeV}$ for the SPSDMK+PSDMK2 model.  5. The dependence of the $^7$Li and $^{11}$B yields on neutrino oscillation parameters, such as mass hierarchy and the mixing angle $\\theta_{13}$, is not changed by the new cross sections. Yield enhancements are smaller.  6. The $^7$Li/$^{11}$B abundance ratio depends on mass hierarchy and the mixing angle $\\theta_{13}$, even when considering uncertainties of neutrino temperatures and the total neutrino energy. For a normal mass hierarchy and $\\sin^22\\theta_{13} \\ga 2 \\times 10^{-3}$, i.e., adiabatic H resonance, the $^7$Li/$^{11}$B ratio is larger than 0.78. In the case of an inverted mass hierarchy or the case without neutrino oscillations, the $^7$Li/$^{11}$B ratio is smaller than 0.61. Smaller uncertainty of neutrino temperatures extends the difference in the $^7$Li/$^{11}$B ratio.  7. The Be/Li abundance ratio is very small. A decrease in the Be/Li ratio is seen for a normal mass hierarchy and $\\sin^22\\theta_{13} \\ga 2 \\times 10^{-3}$."
    },
    "0807/0807.0660_arXiv.txt": {
        "abstract": "We present multi-band photometric and optical spectroscopic observations of SN2007ax, the faintest and reddest Type Ia supernova (SN\\,Ia) yet observed. With M$_{\\rm B}$ = $-$15.9 and ($B-V$)$_{\\rm max}$ = 1.2, this SN is over half a magnitude fainter at maximum light than any other SN\\,Ia. Similar to subluminous SN2005ke, SN2007ax also appears to show excess in UV emission at late time. Traditionally, $\\Delta$m$_{15}$($B$) has been used to parameterize the decline rate for SNe\\,Ia. However, the B-band transition from fast to slow decline occurs sooner than 15 days for faint SNe\\,Ia. Therefore we suggest that a more physically motivated parameter, the time of intersection of the two slopes, be used instead. Only by explaining the faintest (and the brightest) supernovae, we can thoroughly understand the physics of thermonuclear explosions. We suggest that future surveys should carefully design their cadence, depth, pointings and follow-up to find an unbiased sample of extremely faint members of this subclass of faint SNe\\,Ia. ",
        "introduction": "Inspired by the application as a standard cosmological candle, the progress in understanding Type Ia supernovae (SNe\\,Ia) has grown in leaps and bounds. However, the understanding of their weakest subluminous cousins has been purposefully overlooked as their atypical light curve and atypical spectra make them contaminants for cosmological studies. We suggest here some characteristics that make the physics of the explosions of faint SNe\\,Ia intriguing in their own right.  In this paper, we present SN2007ax which, with a peak absolute magnitude of M$_{\\rm B}$ = $-$15.9 and ($B-V$)$_{\\rm max}$ = 1.2, is the faintest and reddest Type Ia supernova yet discovered. Although the class of SNe\\,Ia is remarkably homogenous, subluminous SNe\\,Ia show atypical spectral and light curve features (\\citealt{g+04}, \\citealt{t+07}). Photometrically, not only do they fade much faster than predicted by the Phillips relation, they are also very red at maximum and (at least SN2005ke and SN2007ax) appear to show UV excess at late-time. Spectroscopically, they have broad \\ion{Ti}{2} features and moderate expansion velocities.  SN2007ax was discovered in NGC\\,2577 on UT 2007 Mar 21.978 by \\citet{a07} at an unfiltered magnitude of 17.2. Upper limits of $>$18.5 mag on Mar 17.636 and $>$19.0 mag on Mar 9.959 were also reported. Spectra obtained on Mar 26 by \\citet{bm+07} and \\citet{mf07} showed that it was a SN\\,Ia near maximum light similar to SN1991bg.  In this paper, we present multi-epoch, multi-band imaging and spectroscopic follow up of SN2007ax including optical, ultraviolet, and near-infrared. We summarize our observations in \\S\\,2, present our analysis and comparison with other faint SNe\\,Ia in \\S\\,3 and discuss possible scenarios for faint thermonuclear explosions in \\S\\,4. We conclude with how future surveys can systematically design their cadence, limiting magnitude and pointings to search for more members belonging to this subclass of faint SNe\\,Ia. % ",
        "conclusions": "To summarize, the primary observational characteristics of subluminous SNe\\,Ia (of which SN2007ax is an extreme case) are small t$_{\\rm b}$ in the optical $B$-band light curve, extremely red $B-V$ color at maximum, possible excess in UV emission at late-time, presence of intermediate mass elements in spectra, medium ejecta velocities, low $^{56}$Ni mass in ejecta and short rise-times.  Several theoretical models have been proposed to explain faint SNe\\,Ia -- complete detonation of a sub-Chandrashekhar mass white dwarf, a delayed detonation model, a failed neutron star model and a small scale deflagration model. The detonation of a sub-Chandrashekhar C-O white dwarf (e.g. \\citealt{l90}, \\citealt{ww94}) produces more $^{56}$Ni than observed and is more blue at maximum than observed \\citep{hk96}. If we consider detonation of a sub-Chandrashekhar O-Ne-Mg white dwarf \\citep{i+91}, the total nuclear energy is smaller and the predicted ejecta velocities are lower than observed \\citep{f+92}. \\citet{m+07} use detailed spectral modeling to show a common explosion mechanism for all SNe\\,Ia, likely delayed detonation. The failed neutron star model \\citep{ni85} suggests that if the accretion rate of carbon and oxygen from a companion onto a white dwarf is high enough, it may prematurely ignite CO on the white dwarf surface. Thus, instead of a neutron star, we may see a faint SNe\\,Ia. Small-scale deflagration models suggest that either the burning is restricted to the outer layers or that it occurs slowly.  Another intriguing theoretical possibility recently proposed by \\citet{b+07} is faint thermonuclear supernovae from ultracompact double degenerate AM CVn systems. This supernova is tantalizingly at the brightest end of their predictions (M$_{V}$=$-$14 to $-$16, timescale = 2--6 days, M$_{ej} <$ 0.1 M$\\odot$). However, the decay time predicted by these models is much shorter and the $^{56}$Ni mass less than that observed in SN2007ax. Also, the spectrum does not show any feature which suggests being powered by different radioactive material ($^{48}$Cr, $^{44}$Ti, $^{52}$Fe) produced by some of these models.  None of the above models convincingly explain all the observed characteristics of subluminous SNe\\,Ia. SN2007ax compels the question of what is the (and whether there is) lower limit of $^{56}$Ni mass in a thermonuclear explosion. % Only if we can explain the extremely faint (and the extremely bright) supernovae will we thoroughly understand the limitations in physical processes involved in the thermonuclear explosion, in particular, the $^{56}$Ni mass production.  Future supernova surveys which have a shorter cadence and a  deeper limiting magnitude will provide invaluable clues to understanding the nature of subluminous SNe\\,Ia. Follow-up of these supernovae with well-sampled UV light curves and well-calibrated multi-epoch UV spectra would also be important to understand the apparent excess at late-time.  We suggest how a near-future survey, for example, the Palomar Transient Factory\\footnote{The Palomar Transient Factory is a dedicated time-domain astronomy project to come online on Palomar 48-inch in Nov 08.}, can systematically search for faint SNe\\,Ia. The parameters of the survey design are sky coverage, cadence, depth, filter and choice of pointings. \\citet{h01} shows that faint SNe\\,Ia occur preferentially in early-type galaxies and \\citet{t+07} suggest that they occur in lower metallicity, old stellar mass populations. Since they decline by a magnitude in five days, the cadence of the search should be faster than five days so that the detection sample is complete. Since faint SNe\\,Ia are extremely red at maximum, we should choose a red filter for the search. To maximize sky coverage, searching with a single red filter should suffice (with multi-band follow-up). Since the local universe is clumpy (e.g. : $\\approx$25\\% of the total light at the distance of Virgo is in the Virgo supercluster), the sky coverage must include concentrations in stellar mass, such as the Virgo, Perseus and Coma galaxy clusters. The rate of normal SNe\\,Ia is 3 per 10$^{11}$L$_{\\odot}$ per century \\citep{sb05}. \\citet{l+01} estimate a rate for subluminous SNe\\,Ia to be 16\\% of normal SNe\\,Ia rate based on LOSS and BAOSS surveys.  To a depth of absolute magnitude of $-$15.5, and with a limiting magnitude of 20.5, the survey volume would be 1.5$\\times$ 10$^{7}$ Mpc$^3$. Using the 2MASS K-band luminosity function of 5.1$\\times $10$^{8} $L$_{\\odot}$Mpc$^{-3}$ (\\citealt{kk05},~\\citealt{k+01}), we expect a rate of the faintest subluminous supernovae to be $\\approx$370 all sky per year. The Palomar Transient Factory plans a 5-day cadence 2700 sq. deg. experiment which would give $\\approx$24 faint SNe\\,Ia per year."
    },
    "0807/0807.0456_arXiv.txt": {
        "abstract": "The possibility of light sterile neutrinos allows for the resonant production of lepton number in the early universe through matter-affected neutrino mixing.  For a given a mixing of the active and sterile neutrino states it has been found that the lepton number generation process is chaotic and strongly oscillatory.  We undertake a new study of this process' sensitivity to initial conditions through the quantum rate equations.  We confirm the chaoticity of the process in this solution, and moreover find that the resultant lepton number and the sign of the asymmetry produces a fractal in the parameter space of mass, mixing angle and initial baryon number.  This has implications for future searches for sterile neutrinos, where arbitrary high sensitivity could not be determinate in forecasting the lepton number of the universe. ",
        "introduction": "Neutrino physics is entering an era of increasing precision as initial evidence for neutrino oscillations is leading to confirmation and refinement of the neutrino mass and mixing parameters.  The solar neutrino deficit is consistent with a matter-affected Mikheyev-Smirnov-Wolfenstein (MSW) transformation of electron neutrinos, $\\nu_e$~\\cite{Ahmed:2003kj}, and was confirmed in the CPT conjugate channel by reactor electron antineutrino disappearance at KamLAND~\\cite{Araki:2004mb}.  Atmospheric neutrino disappearance due to apparent neutrino oscillations between what is predominantly $\\nu_\\mu$ and $\\nu_\\tau$ and their CPT conjugates has been confirmed and is being determined precisely by long-baseline neutrino experiments ~\\cite{Michael:2006rx,Ahn:2006zza}.  There were also indications of neutrino oscillations governed by a large mass scale from the short-baseline Liquid Scintillator Neutrino Detector (LSND)~\\cite{Athanassopoulos:1997pv} that would require a third mass scale difference and therefore a fourth neutrino or more~\\cite{Sorel:2003hf}.  This neutrino could not participate in weak interactions due to the $Z^0$ width~\\cite{Yao:2006px}, and therefore is dubbed ``sterile.''  Sterile neutrinos are prevalent in most models of neutrino mass generation mechanisms~\\cite{Mohapatra:2006gs}, and could be light enough to participate in neutrino oscillations.  The four neutrino oscillation hypothesis for the LSND result is in conflict with new results from the from the short-baseline MiniBOONE experiment\\cite{collaboration-2007-98}.  The parameter space of interest for the LSND and MiniBOONE results are a small fraction of that of interest for lepton number generation, and do not exclude the possibility of this mechanism. Moreover, such light sterile neutrinos in the parameter space of interest for lepton number generation, and outside of the MiniBooNE bounds are favored to produce the electron fraction required by successful $r$-process nucleosynthesis in Type-II supernovae \\cite{McLaughlin:1999pd,Fetter:2002xx}.  See-saw type mass generation mechanisms could readily include such light sterile neutrinos \\cite{deGouvea:2005er}.  The effects of sterile neutrinos in the early universe have been under study for some time.  Early work studied the effect of the partial or complete thermalization of the sterile neutrino through oscillations, which would alter the energy density and expansion rate of the universe prior to primordial nucleosynthesis, and therefore the production of the light elements~\\cite{Langacker:1989sv,Barbieri:1989ti,Kainulainen:1990ds,Shi:1993hm}. The primordial helium yield is very sensitive to the expansion rate, and has been used to place bounds on four neutrino models with appreciable mixing between active and sterile neutrinos of the kind required for the LSND indications~\\cite{DiBari:2001ua,Abazajian:2002bj}.  It was also discovered that an MSW resonance can cause the generation of a large neutrino asymmetry that could modify primordial nucleosynthesis even with very small mixing angles between the active and sterile neutrinos~\\cite{Foot:1995qk}.  This was confirmed in Ref.~\\cite{Shi:1996ic}, and further found rapid oscillations in the asymmetry generation, extreme sensitivity to initial conditions, and chaoticity in certain regions of parameter space of mass and mixing angle.  A key question regarding the chaoticity of the lepton generation mechanism is the determinacy of the sign of the asymmetry given a mixing angle and mass-splitting for the active-sterile neutrino pair. That is, if a mass and mixing angle were experimentally measured for the pair, could the sign and magnitude of the cosmological neutrino asymmetry be predicted?  One way of assessing the predictability of the final lepton number given a potentially experimentally determined mass and mixing angle is the level of information loss in the lepton number generation process.  The level of information loss in the chaotic process was explored in Ref.~\\cite{Braad:2000zw}, which found appreciable loss of initial conditions due to chaoticity in small regions of parameter space.  The sensitivity of the sign and amplitude of the asymmetry to the mixing parameter space was found to be ``complex'' in Ref.~\\cite{enqvist-1999-464}.  All of the work discussed thus far was done through the quantum rate equations (QRE's), where the neutrino system is modeled to follow the average neutrino momentum $\\langle p \\rangle \\approx 3.15 T$.  The strong momentum dependence of the resonant process of lepton number generation suggests that an accurate solution should take into account the full spectrum in the solution of the equations through the so-called Quantum Kinetic Equations (QKE's).  The narrowness of the resonance requires exceptionally fine momentum binning and therefore following the exponential growth in a very large number of momenta modes.  Several authors found limitations to solving the equations numerically \\cite{bari-foot,Volkas:2000ei}.  For example, Ref.~\\cite{bari-foot} found oscillatory behavior in a range of parameter space with the QKE's, but could not conclude that it was physical and not numerical.  The QKE's were solved through the oscillatory phase by Ref.~\\cite{kainulainen-2002-0202}, who used a new binning procedure following the resonance to simulaneously have fine binning yet a manageable number of momenta modes. They found that the chaotic nature of the generation process remained.   In this work, we explore the nature of the lepton number generation process and its chaoticity.  We confirm that the resonant generation is highly oscillatory in regions of the parameter space and extremely sensitive to the initial conditions.  Moreover, we explore the dependence of the final sign of the asymmetry on the oscillation parameter space and find that the extreme sensitivity to the parameter space gives rise to a fractal structure in the final sign dependence on mass-splitting and mixing angle.  This has significant implications for the determinacy of the cosmological lepton number in the case of discovering a small sterile neutrino mixing with the active neutrino sector.  If the parameters are found to lie in a region of fractal nature, even arbitrarily high precision in the determination of the parameters will not allow for a prediction of the cosmological lepton number.    We will briefly introduce our system in Sec. \\ref{sec:QKEs}, including the effects of our constant momentum approximation. We go on in Sec. \\ref{sec:results} to present the solution under the above approximation, where we show that the final sign of the asymmetry is not just chaotic, but exhibits a self-similarity like that of a fractal at very small scales.  We compare our results with Ref. \\cite{enqvist-1999-464}, which shows that the resolution and numerical error are responsible for the seeming ``white noise'' structures in the parameter space.  We conclude in Sec. \\ref{sec:Conclusion}. ",
        "conclusions": "}  We have explored the generation of cosmological lepton number in the case of the presence of a light sterile neutrino state with small mixing with an active neutrino.  Consistent with previous studies of lepton number generation using the momentum averaged quantum rate equations, we find that the process is chaotic and highly sensitive to the mixing parameters coupling the active and sterile neutrino states.  Our main results are the fractal nature of the dependence on the sign of the lepton number to the mixing parameters as well as the initial baryon number.  This leaves significant implications for the predictability of the lepton number of the universe, even in the case of arbitrarily precise determinations of the mixing parameters between the active and sterile states in certain regions of the parameter space.  Future searches for subdominant mixings with the active neutrinos and unitarity in the active neutrino mixing matrix therefore could be indeterminant in predicting the sign and magnitude of the cosmological lepton number.  The fractal or correlation dimension depends on the region in parameter space that is examined, which is expected.  The banded structure of the sign of the parameter space undergoes changes in its structure throughout the space, with some definitive lines in the pattern change.  The fractal dimension therefore shifts in the parameter space.  This is likely due to the fact that the frequency of the rapid oscillations are deterimined by the oscillation dynamics, and the point at which they stop or freeze-out is determined by the expansion dynamics, which depend on the total radiation energy density of the universe independent of the neutrino physics.  Our work has focused on the single-momentum quantum rate equations for the solution of this system, and we acknowledge that it has been shown that the behavior changes in the full momentum-dependent quantum kinetic description.  Studies have found that the oscillatory behavior persists and may be chaotic in the full quantum kinetic equation solution \\cite{bari-foot,kainulainen-2002-0202}.  Since there are numerical difficulties in quantifying the nature of the solution in the full momentum case, which we find in our own work, the resolution of this question may require a completely new formulation of the quantum kinetic equations and their solution in order to make them numerically stable to provide a convincing solution of the evolution. The work presented here highlights the new type of chaotic, fractal physics that can arise in this system, and its consequences in determining the lepton number of the universe given the existence of a light sterile neutrino.  If a light sterile neutrino is discovered within this parameter space, a new approaches to quantifying the lepton number of the universe will be strongly motivated. Though if the sign of the lepton produces such fractal structure, the sign of the lepton asymmetry may be indeterminate due to arbitrarily fine fractal structures."
    },
    "0807/0807.2453_arXiv.txt": {
        "abstract": "This work presents a detailed analysis of the overall flow structure and unique features of the inner region of the tilted disk simulations described in Fragile et al. (2007). The primary new feature identified in the main disk body is a latitude-dependent radial epicyclic motion driven by pressure gradients attributable to the gravitomagnetic warping of the disk. The induced motion of the gas is coherent over the scale of the entire disk and is fast enough that it could be observable in features such as relativistic iron lines. The eccentricity of the associated fluid element trajectories increases with decreasing radius, leading to a crowding of orbit trajectories near their apocenters. This results in a local density enhancement akin to a compression. These compressions are sufficiently strong to produce a pair of weak shocks in the vicinity of the black hole. These shocks are roughly aligned with the line-of-nodes between the black-hole symmetry plane and disk midplane, with one shock above the line-of-nodes on one side of the black hole and the other below on the opposite side. These shocks enhance angular momentum transport and energy dissipation near the hole, forcing some material to plunge toward the black hole from well outside the innermost stable circular orbit. A new, extended simulation, which was evolved for more than a full disk precession period, allows us to confirm that these shocks and the previously identified ``plunging streams'' precess with the disk in such a way as to remain aligned relative to the line-of-nodes, as expected based on our physical understanding of these phenomena. Such a precessing structure would likely present a strong quasi-periodic signal. ",
        "introduction": "\\label{sec:intro}  The purpose of this paper is to present a more detailed analysis of the numerical simulation results described in \\citet{fra07b} (hereafter Paper I). That paper described a global numerical simulation of an accretion disk subject to the magneto-rotational instability (MRI) that was misaligned (tilted) with respect to the rotation axis of a modestly fast rotating ($a/M=0.9$) Kerr black hole. It was the first such numerical simulation of a tilted disk to fully incorporate the effects of the black hole spacetime and not rely on an {\\em ad hoc} prescription of angular momentum transport.  A clear identification of a tilted black-hole accretion disk in Nature has yet to be made. Nevertheless, there is reason to believe they may be quite common. In any accretion disk system, the orientation of the disk is set by the net angular momentum of the gas reservoir on large scales. The orientation of the black hole, on the other hand, can either be set by its formation or its evolution. For stellar-mass black holes in low-mass binary systems, only the formation is likely to matter (any evolution subsequent to the formation is unlikely to change the orientation of the black hole significantly). Since the formation of the black hole is largely independent of the angular momentum of the gas reservoir, alignment should generally not be expected in this case. For high mass binaries and active galactic nuclei (AGN), on the other hand, a large episode of misaligned accretion could reorient the spin of the black hole \\citep{nat99}.  Therefore a more detailed understanding of the evolutionary history of the system, including mergers in the case of AGN, would be needed to know if a tilted configuration is expected.  A tilted disk is subject to differential warping due to the effect of Lense-Thirring precession. In disks with a large ``viscous'' stress (parametrized by the \\citet{sha73} $\\alpha$ parameter) the warping results in the Bardeen-Petterson configuration \\citep{bar75,kum85}, characterized by an alignment of the disk with the equatorial plane of the black hole inside some characteristic warp radius. The Bardeen-Petterson effect has been invoked by a number of authors to explain peculiar observations, such as misaligned jets in AGN \\citep{kon05,cap06,cap07} and X-ray binaries \\citep{fra01a,mac02}.  However, we did not see evidence for the Bardeen-Petterson effect in our simulation in Paper I. This was not surprising since our simulation was carried out in the low stress regime, with $\\alpha < H/r$, where $H$ is the half-height of the disk. In this limit warps produced in the disk propagate as waves \\citep{pap95a}, rather than diffusively as in the Bardeen- Petterson case. Instead of a smooth transition between an untilted disk at small radii and tilted disk at large radii, as with the Bardeen-Petterson effect, we found the tilt to be a long-wavelength oscillatory function of radius.  The tilted simulation in Paper I also showed other dramatic differences from comparable simulations of untilted disks. Accretion onto the hole occurred predominantly through two opposing plunging streams that started from high latitudes with respect to both the black-hole and disk midplanes. These plunging streams also started from a larger radius than the innermost stable circular orbit (ISCO), which is often assumed to represent the inner edge of untilted disks.  We interpreted this as being due to the fact that the tilted disk encounters a generalized ISCO surface at a larger cylindrical radius than an untilted disk \\citep{fra07a}. In this regard the tilted black hole effectively acts like an untilted black hole of lesser spin.  Because of the fast sound-crossing time in the disk, the torque of the black hole acted globally rather than differentially. Instead of strongly warping the disk, the torque caused the disk to experience global (solid-body) precession. The precession had a frequency of $\\nu_\\mathrm{prec} = 3 (M_\\odot/M)$ Hz, a value consistent with many observed low-frequency quasi-periodic oscillations (QPOs). However, this value is strongly dependent on the size of the disk ($\\nu_\\mathrm{prec} \\propto r_o^{-5/2}$, where $r_o$ is the outer radius), so this frequency may be expected to be vary over a large range. In the limit of a very large disk or in a case where the disk is strongly coupled to a gas reservoir at large radii, this precession frequency may be expected to drop to zero.  In the present work we expound on some of the features of the tilted-disk simulation which were not described in Paper I. This paper is organized as follows: In \\S \\ref{sec:methods} we briefly review the details of the simulations presented in Paper I, in particular focusing on models 90h (untilted) and 915h (tilted), which are identical other than the initial tilt of the black hole relative to the disk ($\\beta_0=0$ and $15^\\circ$, respectively), and model 915m, which is a lower resolution version of 915h. In \\S \\ref{sec:bending}, we describe the large-scale epicyclic motion that appears in the tilted simulation but is absent in the comparable untilted simulation. This epicyclic motion is driven by radial pressure gradients associated with the stationary bending wave created by the warping action of the gravitomagnetic torque of the black hole (e.g. \\citealt{nel99}). Next, in \\S \\ref{sec:shock}, we describe a pair of standing shocks that again have no counterparts in untilted simulations. The shocks form roughly along the line-of-nodes between the disk midplane and black-hole symmetry plane, one shock on each side of the black hole. We verify that this is not a chance alignment by following the long-term evolution of model 915m and confirming that the shocks remain aligned with the line-of-nodes as it precesses with the disk. Finally, in \\S \\ref{sec:discussion}, we consider some of the implications of our discoveries, particularly in the context of relativistic iron lines. ",
        "conclusions": "\\label{sec:discussion}  In this paper we have compared two simulations that are identical to one another in every respect except one: the initial tilt between the black-hole and disk angular momenta. Despite the apparent similarity, we observed remarkable differences in the evolution of these two simulations.  The primary new feature that we describe in the main disk body is a strong epicyclic driving attributable to the gravitomagnetic torque of the misaligned (tilted) black hole. The induced motion of the gas is coherent over the scale of the entire disk. An interesting point about this epicyclic motion that has not been made before is that it could be detectable, for instance in the profile of relativistically broadened iron lines. We have previously pointed out the importance of the iron lines in directly probing tilted accretion disks \\citep{fra05c}, but at the time we were not in a position to recognize the importance of the epicyclic motion. From Figure \\ref{fig:rhoVph1} we get that the velocities associated with the epicyclic motion represent a significant fraction ($\\lesssim40\\%$) of the orbital velocity of the gas. The corresponding shift in a reflection feature such as the iron line should be of a similar magnitude. The interesting thing is to note that an observer viewing the disk from the vantage point of Figure \\ref{fig:rhoVph1}{\\em a} and seeing the ``top'' of the disk would see a smaller than expected red shift (the gas going away is not moving as fast as expected) and a larger than expected blue shift (the approaching gas is moving faster than expected). An observer viewing the same disk from the same vantage point but seeing the ``bottom'' of the disk would see exactly the opposite shift. Therefore, depending on the viewing angle, the entire line profile of a misaligned disk could be shifted toward the blue {\\em or} the red relative to an aligned disk. On the redshifted side this effect might be confused with gravitational redshifting, making it appear that the line is coming from deeper in the potential well than is actually the case.  If the disk actually precesses then there is a clear and simple way to disentangle this effect because the epicyclic motion is phased with the orientation of the disk, as discussed in \\S \\ref{sec:bending} above. This means that the line shift we are describing would reverse itself twice per precession period, appearing for half a precession period as an overall blueshift and for the other half as an overall redshift. If the integration time of the detector is longer than the precession period, the net effect would generally be to smear or broaden the line. This may provide an alternative explanation for an exceptionally broad line profile such as GX 339-4 \\citep{mil04}, without requiring a high black-hole spin or small disk radius. On the other hand, if the integration time is shorter than the precession period, then one would expect the iron line to vary in phase with changes in the X-ray flux (i.e. in conjunction with a low frequency QPO that would correspond to the precession frequency of the disk). Such behavior has been identified in at least one source, GRS 1915+105 \\citep{mil05}.  The second new feature of our tilted disk simulation that we described is a pair of standing shocks, roughly aligned with the line-of-nodes between the disk and black hole symmetry planes. This asymmetric shock provides additional angular momentum transport and energy dissipation in the tilted disk, relative to the untilted one. This enhanced dissipation may help compensate for the loss of radiative efficiency (relative to an untilted disk) due to the plunging region starting further out than the equatorial ISCO radius of the spinning black hole. In collisionless accretion flows, such as those thought to be relevant for Sgr~A* \\citep[e.g.][]{qua03}, it is conceivable that such shocks could also be sites of particle acceleration.  The precession of these shocks might then result in periodic variations of nonthermal radiation.  In considering the magnetic fields in the inner part of the tilted disk, as noted in Paper I, the magnetic fields remain largely subthermal. They do not play an important role in the physics of the epicyclic motion, plunging streams, or standing shocks, although we do find some enhancement of $\\beta_\\mathrm{mag}$ associated with the shocks.  We finish with a figure (\\ref{fig:schematic}) which provides a visual summary of our findings. We note the overall consistency of our interpretation of the results: 1) The observed epicyclic motion is in agreement with expectations for warped disks in the wave-like propagation limit; and 2) The shocks are located near the apocenters of the epicyclic motion, as one expects when the eccentricity increases with decreasing radius.  \\begin{figure} \\plotone{f14.eps} \\caption{Schematic diagram of the inner region of the tilted accretion disk, showing the pattern of epicyclic motion, the standing shock, and the plunging streams. \\label{fig:schematic}} \\end{figure}"
    },
    "0807/0807.3043_arXiv.txt": {
        "abstract": "Gravitational waves carry away both energy and angular momentum as binary black holes inspiral and merge.  The relative efficiency with which they are radiated determines whether the final black hole of mass $M_f$ and spin $S_f$ saturates the Kerr limit ($\\chi_f \\equiv S_f/M_f^2 \\leq 1$). Extrapolating from the test-particle limit, we propose expressions for $S_f$ and $M_f$ for mergers with initial spins aligned or anti-aligned with the orbital angular momentum.  We predict the the final spin at plunge for equal-mass non-spinning binaries to better than 1\\%, and that equal-mass maximally spinning aligned mergers lead to nearly maximally spinning final black holes ($\\chi_f \\simeq 0.9988$).  We also find black holes can always be spun up by aligned mergers provided the mass ratio is small enough. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0807/0807.3872_arXiv.txt": {
        "abstract": "Keeping the exact general relativistic treatment of light bending as a reference, we compare the accuracy of commonly used approximate lens equations. We conclude that the best approximate lens equation is the Ohanian lens equation, for which we present a new expression in terms of distances between observer, lens and source planes. We also examine a realistic gravitational lensing case, showing that the precision of the Ohanian lens equation might be required for a reliable treatment of gravitational lensing and a correct extraction of the full information about gravitational physics. ",
        "introduction": "Gravitational lensing is a well-established research subject treating the bending of light trajectories by gravitational fields. Its methodology is traditionally developed within the weak field approximation of General Relativity, which describes photon trajectories through the geodesics equation. In this context, the true position of a source in the sky and its apparent position after deflection of the light by a massive body are related by the so-called lens equation \\cite{Liebes,Refsdal} (for textbook reviews, see \\cite{SEF,PLW,MolRou}). This relation and its mathematical properties have been extensively studied for generic lens models and represents the basis of the whole gravitational lensing theory.  The lens equation is typically established in the small angle approximation, in conjunction with the weak deflection hypothesis. In more recent years, there has been a renewed interest in astrophysical situations in which the deflection angle is not small. For example, if the bending of light emitted by sources near a black hole is considered, the deflection angle may reach arbitrary large values \\cite{Dar}. For such gravitational lenses the old small angles lens equation must be obviously revised. Several proposals for generalized lens equations have then appeared in the literature and have been applied to specific cases \\cite{FriNew,FKN,Oha,VNC,VirEll,DabSch,BozSer}. Such lens equations principally differ among each other for what concerns the variables in which they are expressed. However they also lie at different approximation levels below the full general relativistic description of the photon motion.  In this work, we review all these lens equations for spherically symmetric bodies appeared in the literature and introduce a new lens equation that fills a gap in the present taxonomy. We also present a detailed discussion of the order of magnitude of the errors committed in the use of different lens equations, giving a complete interpretation for their origin. It is worth mentioning that a critical review of the approximations leading to the weak deflection lens equation has recently appeared \\cite{KliFri}. The present work shares the same spirit by testing different approximations leading to large deflection lens equations, instead.  The plan of the paper is as follows. In Section II we establish the notations by describing the basic geometric configuration for gravitational lensing. In Section III, we review the exact lens equation by Frittelli, Kling and Newman, derived in a fully general relativistic context. In Section IV, we discuss the asymptotic approximation and all lens equations making use of this approximation. We also introduce our new proposal for an improved lens equation. In Section V we present a numerical example about a realistic gravitational lensing situation, in order to compare the different approximations in the lens equations previously discussed. Section VI contains a discussion about the precision needed in gravitational lensing observations and the conclusions of the work. ",
        "conclusions": "Of course, an approximation can be valid or not depending on the precision of the observational instrumentation, the level of environmental noise and the accuracy needed for the extraction of the interesting information.  The best resolution we can expect to reach in a reasonably near future is that of the MAXIM project (http://maxim.gsfc.nasa.gov), which amounts to a fraction of $\\mu$as. Then, in the example considered in this section, the weak deflection lens equation would not be able to ensure such a precision. The VE and the small angles lens equations would be marginally adequate, whereas the Ohanian lens equation would be largely sufficient, without need to resort to the exact lens equation. So, in order to exploit the resolution of future instruments, it might be necessary to use the Ohanian lens equation or one of its variants.  Environmental noise can be provided by scattering and absorption processes affecting the photons involved in gravitational lensing. Such processes should be very effective in the environment of the supermassive black hole at the Galactic center. However, they act in a statistical way and a good average process should be able to restore the original information regarding gravitational lensing. More subtle are the effects that introduce unwanted systematics, such as gravitational lensing by secondary objects. However, the gravitational lensing optical depth towards the center of the bulge is of the order $10^{-6}$ \\cite{MolRou}, which indicates that lensing or microlensing effects by stars or other compact objects on the line of sight is absolutely negligible (see also \\cite{CGM}). In the absence of relevant systematics, the position of the centroids of the images can be used for a clean reconstruction of the gravitational lensing event at the angular resolution level reached by the instruments.  Finally, we note that there are many theoretical reasons to require a reconstruction of a gravitational lensing event as precise as possible. As anticipated, the study of higher order terms in the deflection angle requires a precision better than $(r_g/D_{LS})^{1/2}$ at second PPN order and $(r_g/D_{LS})$ at third PPN order. Therefore, if one wants to compare the results of PPN formalism with different black hole metrics, it is mandatory to use a lens equation that guarantees the necessary accuracy. In this respect the VE lens equation does not represent an improvement with respect to the small angles approximation, as it is affected by an error of the same order of magnitude.  In this work we have compared the level of accuracy of several lens equations, keeping the exact general relativistic lens equation introduced by Frittelli, Kling and Newman \\cite{FriNew,FKN} as the reference equation. We have shown that the Ohanian lens equation \\cite{Oha} and its close relatives \\cite{BozMan,BozSer} are the best approximations of the exact lens equation, in that they adopt the asymptotic approximation only. Other lens equations are not expressed in terms of relative positions of source, lens and observer \\cite{VNC} or introduce additional approximations \\cite{VirEll,DabSch}.  Furthermore, we have presented a new formulation of the Ohanian lens equation in terms of the distances between the observer, lens and source planes, which fills a gap in the lens equation zoo. As shown by a numerical example describing a realistic gravitational lensing event by the black hole at the center of our Galaxy, such an equation represents a noteworthy improvement with respect to previous commonly used lens equations expressed in terms of the same quantities."
    },
    "0807/0807.3796_arXiv.txt": {
        "abstract": "All far ultraviolet observations of HD209458 tend to support a scenario in which the inflated hydrogen atmosphere of its planetary companion strongly absorbs the stellar \\lya\\ flux during transit. However, it was not clear how the transit absorption depends on the selected wavelength range in the stellar line profile, nor how the atomic hydrogen cloud was distributed spatially around HD209458b. Here we report a sensitivity study of observed time and spectral variations of the stellar flux. In particular, the sensitivity of the absorption depth during transit to the assumed spectral range in the stellar line profile is shown to be very weak, leading to a transit depth in the range $(8.4-8.9)\\%\\pm 2.0\\%$ for all possible wavelength ranges, and thereby confirming our initially-reported absorption rate. Taking the ratio of the line profile during transit to the unperturbed line profile, we also show that the spectral signature of the absorption by the exoplanetary hydrogen nebula is symmetric and typical of a Lorentzian, optically thick medium. Our results question the adequacy of models that require a huge absorption and/or a strong asymmetry between the blue and red side of the absorption line during transit as no such features could be detected in the HST FUV absorption profile. Finally, we show that standard atmospheric models of HD209458b provide a good fit to the observed absorption profile during transit. Other hybrid models assuming a standard model with a thin layer of superthermal hydrogen on top remain possible. ",
        "introduction": "Just like our Sun, the \\lya\\ line profile of HD209458 shows a double-horn emission with the distinguishing feature of intervening interstellar gas strongly absorbing the line core. From the signal left at the Earth's orbit, \\cite{vid03} (hereafter \\vma ) initially claimed that HST/STIS medium resolution observations of HD209458 are showing a strong absorption of 15\\%  during transit and that a large asymmetry appears on the blue side of the \\lya\\ absorption profile. Those authors used the 15\\% drop-off of the stellar signal during transit as an indication of the presence of a hydrogen cloud that extends beyond the Roche lobe of the planet. The next step for \\vma\\ was to use the asymmetry claimed in the blue side of the line during transit as an indication that the evaporated gas is accelerated by the radiation pressure from the star, an acceleration that results into the build-up of a cometary shape of the nebulosity that escapes from the planet away from the star. \\\\  Later on, \\cite{hol08} rejected the acceleration process by radiation pressure, arguing that neutrals will not have the time to be accelerated by the radiation pressure before they get ionized. Furthermore, they proposed that observations thus far reported by \\vma\\ on HD209458b could be explained by the presence of a population of energetic neutral atoms produced by charge exchange between the neutral atmosphere and the stellar wind ions beyond an obstacle region corresponding to the planetary magnetopause assumed around $\\sim 4.2 $ R$_p$ from the center of the planet. In the meantime, we revisited the HST/STIS medium resolution archive data to report that exoplanet HD209458b absorbs $\\sim 8.9\\% \\pm 2.1\\%$ of the parent star \\lya\\ flux during its transit, much less than previously reported by \\vma\\ (\\cite{ben07a}, hereafter {\\bjl}).  We also reported that no sign of asymmetry could be found in the stellar line profile during transit, a finding that directly contradicts the claims initially reported in {\\vma}. Soon thereafter, the absorption drop-off reported by {\\bjl} for a specific wavelength range was used out of the spectral information provided in Fig.3 of \\bjl\\ to claim that evaporation was strengthened with a factor $\\sim 2$ asymmetry in the absorption between the blue ($9.8\\%\\pm 1.8\\%$) and red ($\\sim 5.2\\%\\pm 1.0\\%$) sides of the \\lya\\ line during transit (\\cite{vid08} hereafter {\\vml}). In any case, it appears that our initial selection of the wavelength window to estimate the transit absorption created some confusion as to the reading of the data analysis presented in {\\bjl}, and this confusion adversely affects efforts to discriminate between the processes thus far proposed to explain the hydrogen distribution around HD209458b ({\\vma}, \\cite{mun07,sch07,hol08}). In the following, we revisit available \\hbox{\\lya\\ } HST/STIS medium resolution archive data to conduct a sensitivity study that should dissipate the existing confusion, an important step in formulating clear observational constraints on the steady-state properties of the hydrogen nebula around HD209458b. \\\\  \\begin{table*} \\begin{center} \\scalebox{1.}{ \\begin{tabular}{lcccccc}  Dataset name &\tProgram ID & Date Obs.& Time Obs. & Start time - TCT (s) & Duration (s) & Earth's Velocity \\\\ O4ZEA4010 & 7508 & 2001-08-28 &\t06:50:07 & -7980.28  & 2600.\t& 2.05 \\\\ O4ZEA4020 & 7508 & 2001-08-28 &\t08:26:25& -2202.29  &\t2600  & 2.05\\\\ O6E201010 & 9064 & 2001-09-07 &20:35:27\t& -8075.11\t& 1780.\t& -2.59\\\\ O6E201020 & 9064 & 2001-09-07 &22:12:36\t& -2246.12\t& 2100.\t&  -2.62\\\\ O6E201030 & 9064 & 2001-09-07 &23:48:54\t& 3531.89\t  & 2100. & -2.65\\\\ O6E202010 & 9064 & 2001-09-14 &\t21:11:51& -10167.67\t& 1780.\t& -5.68\\\\ O6E202020 & 9064 & 2001-09-14 & 22:42:08 & -4750.68\t& 2100.\t&  -5.71\\\\ O6E202030 & 9064 & 2001-09-15 &00:18:24\t& 1025.31\t  & 2100. &   -5.74\\\\ O6E203010$^*$ & 9064 & 2001-10-20 & 02:50:03\t& -11258.45\t& 1780.\t&  -19.30\\\\ O6E203020$^*$ & 9064 &2001-10-20 & 04:19:45\t& -5876.47\t& 2100.\t& -19.32\\\\ O6E203030$^*$ & 9064 & 2001-10-20 & 05:55:59\t& -102.45\t  & 2100.\t& -19.34\\\\  \\end{tabular} } \\end{center} \\caption{HST/STIS data set on HD209458 used in this study. All observations were obtained with the G140M grating and the 52\"x0.1\" long slit. Transit central time (TCT) is defined by 2,452,826.628521 HJD (\\cite{knu07}). Earth's velocity is projected along line of sight in the heliocentric reference frame. ($^*$) indicate data set that shows most time variability. \\label{tbl1}} \\end{table*}  Data analysis is carefully described in section 2 so that light curves may be extracted (in selected wavelength windows within the stellar line) and studied in section 3.  In section 4, time variability of the HD209458 \\lya\\ emission is examined and further evidence is provided on the variable behaviour of the HD209458 emission line, a property that may not be well-described by proxies derived for the solar \\lya\\ line. The absorption spectral profile is then derived as the ratio between the stellar flux during transit to the unperturbed spectrum of the star with no transit effect.  Finally, the adequacy of models thus far proposed to explain the hydrogen distribution around hot giant exoplanets is discussed in the light of the constraints derived here, with new solutions proposed for the H nebula distribution. ",
        "conclusions": "HST archive observations of the \\lya\\ emission of HD209458 are revisited to test the absorption rate sensitivity to the spectral range considered inside the \\lya\\ emission line and to derive the absorption line profile during transit. Our first conclusion is that the absorption depth during transit is weakly dependent on the assumed spectral range assumed, with rates in the range of $(8.4-8.9)\\%\\pm 2.\\%$ for all possible wavelength windows as reported in our initial study (\\bjl). Self-consistent derivation of the absorption line profile during transit shows a symmetrical Lorentzian shape, typical of optically thick conditions of the hydrogen nebula around HD209458b. Whether using direct light curves from the time series or inspecting the absorption line profile during transit, neither method shows any indication of preferred absorption on the blue side of the stellar line in the spectral domain $\\sim\\pm 200 {\\rm km\\, s^{-1}}$ from line center.  Finally, we stress that because the planet is not resolved spatially, several scenarios of the gas distribution may lead to the observed absorption line profile. In that context, two scenarios that fit rather well to the absorption profile are suggested. A first model assumes an optically thick atmosphere of HD209458b obtained from the standard model of \\cite{mun07}. A second class of models consists of a standard model that may have a thin layer of superthermal neutrals (ENAs, turbulence energized particles, etc.) on top of it, much like what we observe for the upper atmosphere of Jupiter (\\cite{ben93,eme96}). In all cases, the natural broadening that was missed in most past studies is required in the modelling of the \\lya\\ absorption profile. Moreover, for Enas production, a self-consistent treatment of the charge exchange process in conjunction with the size of the magnetospheric cavity is much needed to test the validity of such process. In any case, high resolution observations with a {\\it S/N} better than $(0.1-1.0)\\% $ are required to efficiently constrain the hydrogen distribution and its time and phase variations along the planetary orbit."
    },
    "0807/0807.2790_arXiv.txt": {
        "abstract": "We compute one-loop corrections to the power spectrum of the curvature perturbation in single-field slow-roll inflation arising from gravitons and inflaton interactions. The quantum corrections due to gravitons to the power spectrum of the inflaton field are computed around the time of horizon crossing and their effect on the curvature perturbation is obtained on superhorizon scales through the $\\delta N$ formalism. We point out that one-loop corrections from the tensor modes are of the same magnitude as those coming from scalar self-interactions, therefore they cannot be neglected in a self-consistent calculation. ",
        "introduction": "Inflation has become the standard paradigm to understand the initial conditions for structure formation and Cosmic Microwave Background (CMB) anisotropies. In the inflationary picture, primordial density and gravity-wave fluctuations are generated from quantum fluctuations and stretched to superhorizon scales during an early period of accelerated expansion of the universe~\\cite{lrreview}. Such fluctuations can trigger the formation of the CMB anisotropies taking place in a later radiation- or matter-dominated epoch. CMB observations~\\cite{smoot92,bennett96,gorski96,wmap3,wmap5} show that the cosmological perturbations are very small, of order $10^{-5}$ compared to the homogeneous background. This has justified the use of linear perturbation theory for a comparison with observations. There are however two good reasons to go beyond linear order. One is the possibility of observing some amount of non-Gaussianity in the CMB anisotropies, revealing fundamental details about the mechanism generating the primordial density perturbations. This issue is being deeply investigated at present, especially because of the high sensitivity to non-Gaussianity of the Planck satellite~\\cite{planck} and its successors. Such non-Gaussianities are sourced by self-interactions in the early universe and become visible at the level of second- or higher-order perturbation theory~\\cite{review}. Another reason to go beyond linear order is to look at the loop corrections the self-interactions of any scalar field during an inflationay phase (and more interestingly of the curvature perturbation $\\zeta$) can generate in the cosmological correlators, including the observationally interesting cases of the two- and the three-point correlation functions (or, in Fourier space, the power spectrum and the bispectrum, respectively).  Loop corrections may lead to interesting effects which scale like the power of the number of e-folds between horizon exit of a given mode $k$ and the end of inflation~\\cite{Mukhanov:1996ak,Abramo:1997hu,boy1,boy2,olandesi}. Recently there has been some renewed interest in loop corrections to the correlators of cosmological perturbations generated during an early epoch of inflation, stimulated by two papers of Weinberg~\\cite{weinberg1,weinberg2}. The reason is that one-loop corrections to the power spectrum of the curvature perturbation $\\zeta$ seem to show some infra-red divergences which scale like $\\ln(kL)$, where $L^{-1}$ is some infra-red comoving momentum cut-off. \\cite{sloth,sloth2,seery1,seery2}. \\footnote{By choosing the smallest possible value for $L^{-1}$, $L^{-1}=a_i H$, where $a_i$ is the scale factor at the beginning of inflation, the logarithm becomes $\\ln(kL)=\\ln(a_k/a_i)$, which is proportional to the number of e-folds between the beginning of inflation and the time the mode $k$ leaves the horizon (when $k=a_k H$). This would have a dramatic impact for the control of cosmological perturbation theory and in particular for the comparison of inflationary predictions for the primordial density perturbations with observations.}. However it has been discussed in Refs.~\\cite{box4,box5} (see also~\\cite{Enqvistetal}) that such potentially large corrections do not appear in quantities that are directly observable. Indeed they concern variations of background quantities defined within a region of comoving size $\\ell$ not much larger than the present horizon $H_0^{-1}$ when one considers how they change on scales $L$ much larger than the size $\\ell$. Such variations will be then associated to infrared corrections which however are not related to the local observable quantities. This means that loop corrections are under control because it is justified to choose for the infrared momentum cut-off the scale $\\ell$ not much larger than the present horizon $H_0^{-1}$. \\footnote{In some special cases a non-linear resummation of the infra-red divergences discussed here is doable, see~\\cite{RiottoSloth}.}   Independently of this issue, up to now the computations of one-loop corrections to the power spectrum of curvature perturbations have neglected any contribution arising from the graviton excitations (or tensor perturbation modes) which are unavoidably generated during a quasi de-Sitter phase. In particular, when going beyond linear order, a mixing of scalar and tensor perturbations is inevitable as well. In this paper we compute for the first time the graviton effects at one-loop level to the power spectrum of the curvature perturbation $\\zeta$ in single-field models of slow-roll inflation. Our results show that such corrections are of the same order as those obtained by considering only the loops from scalar self-interactions (see for example~\\cite{seery1,seery2}). This further strengthens the point that a full self-consistent computation of one- (or higher-order) loop corrections to the correlation functions of primordial curvature perturbations must necessarily account for the graviton contribution inside the loops. Loop corrections in general appear to be small, neverthless an evaluation as precise as possible is required. However notice that, interestingly enough, we find also for the graviton-loop contributions terms which scale like $\\ln(k \\ell)$.  Most of the computation consists in deriving the quantum corrections due to gravitons to the two-point function of the inflaton field $\\delta \\phi$ around the time of horizon crossing. This is achieved using the in-in, also dubbed Schwinger-Keldysh, formalism~\\cite{in-in1,in-in2,in-in3}. Their contribution to the power spectrum of the curvature perturbation $\\zeta$ on superhorizon scales is then obtained by exploiting the $\\delta N$ formalism~\\cite{deltaN1,deltaN2,deltaN3,deltaN4}. The $\\de N$ formalism allows to obtain an expression for the curvature perturbation $\\ze$ (in the uniform density gauge) in terms of the scalar field(s) fluctuations \\be \\label{DN} \\ze(\\vec{x},t)=\\sum_{n}{\\frac{N^{(n)}(t,t_*)}{n!}\\left(\\df(\\vec{x},t_*)\\right)^{n}}\\, , \\ee where $N^{(n)}(t,t_*)=(\\partial /\\partial \\phi_*)^n N(t,t_*)$ represents the derivative of the number of e-folds (between the final time $t$ and some initial time $t_*$) with respect to the initial value of the scalar field (in the spatially flat gauge). On superhorizon scales during inflation, Eq.~(\\ref{DN}) represents the perturbation in the number of e-folds of expansion between widely separated Hubble regions due to the initial fluctuations of the scalar field $\\phi_*$. In fact the time $t_*$ is left to be freely chosen, as long as it is set after horizon crossing, since at that time the perturbations will have settled to their classical value. One possibility is to take $t$ to fall a few e-foldings after horizon crossing like for example in \\cite{byrnes,seery1,seery2}; we will adopt this choice, motivated by the possibility of simplifying our analysis and calculations (we will provide details about this matter in section 3).\\\\  The reader might be worried about the validity of the $\\delta N$ expansion~(\\ref{DN}) when tensor perturbation modes $\\gamma_{ij}$ are accounted for, since they introduce an additional degree of freedom. In fact the $\\delta N$ formalism holds also in this case. The reason can be simply understood by considering that the spatial perturbed metric can be written as~\\cite{SalopekBond} $h_{ij}\\propto a^2(t) e^{2 \\zeta} (e^{\\gamma})_{ij}$ (where $a(t)$ is the unperturbed scale factor, see also Eq.~(\\ref{smetric})) and, on superhorizon scales, the tensor modes can always be reabsorbed into a rescaling of the coordinates.~\\footnote{ A rigorous proof that the presence of tensor modes does not alter Eq.~(\\ref{DN}) can be found, e.g., in~\\cite{deltaN3}, and amounts to consider that $(e^{\\gamma})_{ij}$ has unit determinant so that it does not modify the local volume defined in terms of the local scale factor $a(t,{\\bf x})=a(t) e^\\zeta$.}     One-loop corrections to the power spectrum of $\\ze$ are provided by the action of $\\delta \\phi$ computed up to $4$th order. Both corrections due to scalar and graviton loops will be taken into account. The tensor fluctuations will enter the calculations up to second order (higher order tensor terms are unnecessary to our purpose). A possible classification for theories of inflation is the one that distinguishes between a canonical and non-canonical form for the kinetic term in the Lagrangian~\\cite{non canonical.1,non canonical.2}. The latter are particularly interesting for the study of non-Gaussianity \\cite{gauss1,gauss2,gauss3}. We derive the solutions to the constraint equations for a general non-canonical Lagrangian but our loop calculations will be focused on the canonical case (we will cover the study of the tensor corrections in the non canonical case both for the power spectrum and the trispectrum of the curvature perturbations in a companion work \\cite{noi}).  The paper is organized as follows: in section 2 we calculate the complete action to fourth order in the spatially flat gauge; in section 3 we study the graviton one-loop corrections to the power spectrum of the inflaton field perturbations; in section 4 we collect our results together with the one-loop corrections from scalar perturbation modes which have been already computed in Ref.~\\cite{seery1,seery2}, to be combined in a final formula for the corrected power spectrum of the curvature perturbation; in section 4 we also draw our conclusions. Appendices are organized as follows: in Appendix A we face the problem of dealing with a lagrangian with time derivative interactions; Appendix B contains a detailed study of the slow-roll order of the different terms of the $4$th order action together with the calculation of some of the diagrams with graviton loops; Appendix C collects the complete expressions for the one-loop two-vertex diagrams. ",
        "conclusions": "Let us now collect our results in the final formula for the power spectrum of the curvature perturbation $\\zeta$ computed up to one-loop level. This can be derived from Eq.~(\\ref{ZZ}), which follows from the $\\delta N$ formula, and the final expression reads~\\cite{byrnes,seery2} \\bea\\label{QQQ} \\langle \\zeta_{\\vec{k_{1}}}(t)\\zeta_{\\vec{k_{2}}}(t) \\rangle&=&(2 \\pi)^{3}\\delta^{(3)}(\\vec{k_{1}}+\\vec{k_{2}})\\Big[\\left( N^{(1)} \\right)^{2}\\Big(P_{{\\rm tree}}(k_1) +P_{{\\rm one-loop}}(k_1)\\Big) \\nonumber\\\\ &+&N^{(1)}N^{(2)}\\int \\frac{d^{3}q}{(2 \\pi)^{3}} B(k_1,q,|\\vec{k}_1-\\vec{q}|)\\nonumber\\\\ &+&\\frac{1}{2}\\left( N^{(2)} \\right)^{2}\\int \\frac{d^{3}q}{(2 \\pi)^{3}} P_{{\\rm tree}}(q)P_{{\\rm tree}}(|\\vec{k}_1-\\vec{q}|) \\nonumber\\\\ &+&N^{(1)}N^{(3)}P_{{\\rm tree}}(k)\\int \\frac{d^{3}q}{(2 \\pi)^{3}} P_{{\\rm tree}}(q)\\Big]\\, , \\eea where $P_{{\\rm tree}}(k)$ is the tree level power spectrum~(\\ref{tree}) and $B(k_{1},k_{2},k_{3})$ is the bispectrum of the scalar field defined by \\bea \\langle {\\df}_{\\vec{k_{1}}}{\\df}_{\\vec{k_{2}}}{\\df}_{\\vec{k_{3}}} \\rangle \\equiv (2 \\pi)^{3}\\delta^{(3)}(\\vec{k_{1}}+\\vec{k_{2}}+\\vec{k_{3}})B(k_{1},k_{2},k_{3})\\, . \\eea In previous works the computations of one-loop corrections accounted just for the contributions coming from the inflaton field peturbations (see, for example, Refs.~\\cite{sloth,sloth2,seery1,seery2}. The focus of this paper is on including also the one-loop terms arising from interactions between the tensor (graviton) modes and the scalar field in the case of single-field models of slow-roll inflation. At this order the graviton contribution enter the final formula as a correction to the power spectrum of the inflaton field around the time of horizon crossing \\begin{equation} P_{{\\rm one-loop}}(k)=P_{{\\rm scalar}}(k)+P_{{\\rm tensor}}(k)\\, , \\end{equation} and as such it is due to genuine quantum effects, while it does not affect the remaining terms of Eq.~(\\ref{QQQ}), which, according to Ref.~\\cite{seery2}, can be considered as classical contributions arising after the perturbation modes leave the horizon. \\footnote{Notice however that the computation of $\\langle\\zeta_{\\vec{k_{1}}}(t)\\zeta_{\\vec{k_{2}}}(t)\\rangle$ could be carried out without using the $\\delta N$ formalism, working in the comoving gauge and making use only of the Schwinger-Keldysh path integral formula as in \\cite{maldacena}. In this respect, the two types of terms, ``classical' and  ``genuine quantum effects'', are not really supposed to be entirely different.}  Summing the main results of the previous section, Eqs. (\\ref{FR1}), (\\ref{appC1}) and (\\ref{appC2}), we finally find the one-loop graviton correction to the inflaton power spectrum  \\be\\label{FR} \\langle\\df_{\\vec{k_{1}}}(\\eta^{*})\\df_{\\vec{k_{2}}}(\\eta^{*})\\rangle_{1L}=\\pi\\delta^{(3)}(\\vec{k_{1}}+\\vec{k_{2}})\\frac{H_{*}^{4}}{k^{3}} \\left[f_{1}\\ln(k)+f_{2}\\ln(k \\ell)+f_{3}\\right]\\, , \\ee where \\bea f_{1}=-\\frac{4}{15}\\left(5+5x^{*2}+2x^{*4}\\right)\\, , \\eea \\bea f_{2} =a_{2}+c_{2}\\, , \\eea and $f_{3}$ is given by a left-over scheme-dependent renormalization constant plus contributions of order ${\\cal O}(1)$ (see Appendix C for the complete expressions of $a_{2}$, $c_{2}$ and $f_3$). If we calculate the two point funtion of $\\df$ a few e-foldings after horizon crossing, i.e. $x^{*}$ ranges for example between $10^{-1}$ and $10^{-2}$, $f_{1}$ reduces to a negative constant of order ${\\cal O}(1)$ and $f_{2}\\sim 4\\left(1-{\\rm Ci}(2x^{*})\\right) \\sim {\\cal O}(10)$. In the limit where $x^{*} \\rightarrow 1$ both $f_{1}$ and $f_2$ turn out to be of order unity.  In order to understand which is the dominant contribution in Eq.~(\\ref{QQQ}) and how big it is, one needs to (i) know the slow-roll order of the coefficients $N^{(i)}$: $N^{(1)} \\sim {\\epsilon}^{-1/2}$, $N^{(2)} \\sim \\epsilon^{0}$, $N^{(3)} \\sim \\epsilon^{1/2}$; (ii) compute the integrals involving the power spetrum $P(q)$. This is discussed in details in Ref.~\\cite{seery2} (for the case of scalar perturbations only), see in particular Sec IV of~\\cite{seery2}. It turns out that the crucial quantity is represented by the number of e-foldings of inflation between the times of horizon exit of the mode $\\ell^{-1}$, which corresponds to the infrared cutoff, and the time of horizon exit of the mode $k$ we want to observe. However, to deal with observable quantitites one has to choose $\\ell$ not much bigger than the present cosmological horizon $H_0^{-1}$~\\cite{box4,box5}.\\\\ The relevant point about Eq. (\\ref{FR}) is that it gives in Eq.~(\\ref{QQQ}) a contribution which is of the same order of magnitude as those coming from loops which accounts for scalar perturbations only. Since in terms of the slow-roll parameters $\\left( N^{(1)} \\right)^2 \\sim {\\epsilon_*}^{-1}$ the magnitude of the one-loop graviton correction turns out to be \\begin{equation} \\label{fin} \\Delta P_{\\zeta}^{1\\rm{loop}}(k) \\sim \\frac{2 \\pi^2}{k^3} \\alpha(k) \\frac{1}{\\epsilon_*}{\\cal P}^2_*(k)\\, , \\end{equation} where we have used Eq.~(\\ref{tree}) for the power spectrum of the inflaton field. In Eq.~(\\ref{fin}) $\\alpha(k)$ includes the various coefficients of Eq.~(\\ref{FR}), and it is ${\\cal O}(1)$. Eq.~(\\ref{fin}) allows a more direct comparison with the results of Ref.~\\cite{seery2}, showing that the graviton contributions to the one-loop corrections are comparable to the ones computed only from scalar interactions. Notice that also for the tensor contributions we find terms of the form $\\ln(k\\ell)$.  One-loop corrections to the power spectrum of density perturbations from inflation are small, still a precise and self-consistent computation requires to account also for the loops which are switched on by the tensor (graviton) modes. First, going beyond linear order, the tensor perturbation modes produced during inflation unavoidably mix with scalar modes. This fact alone would require to include the tensor modes for a self-consistent computation. Most importantly, despite a naive expectation suggested by the fact that the power spectrum of the tensor modes is suppressed on large scales with respect to that of the curvature (scalar) perturbations, our results show explicitly that their inclusion is necessary since their contribution is not at all negligible with respect to the loop corrections arising from interactions involving the inflaton field only."
    },
    "0807/0807.0040_arXiv.txt": {
        "abstract": "We use the bulge S\\'ersic index $n$ and bulge-to-total ratio ($B/T$) to explore the fundamental question of how bulges form. We perform  $2D$ bulge-disk-bar decomposition on $H$-band images of 143 bright, high mass ($M_\\star \\geq 1.0 \\times 10^{10} M_\\odot$) low-to-moderately inclined ($i<70^\\circ$) spirals. Our results are: (1)~Our $H$-band bar  fraction ($\\sim58\\%$) is consistent with that from ellipse fits. (2)~70\\% of the stellar mass is in disks, 10\\% in bars, and 20\\% in bulges. (3)~{\\it A  large fraction ($\\sim$~69\\%) of bright  spirals have $B/T \\leq$~0.2, and $\\sim$~76\\% have low $n \\leq 2$ bulges. These bulges exist in barred and unbarred galaxies across a wide range of Hubble types.} (4)~About 65\\% (68\\%) of bright spirals with $n\\leq2$ ($B/T \\le 0.2$) bulges host bars, suggesting a possible link between bars and bulges. (5)~We compare the results with  predictions from a set of $\\Lambda$CDM  models. In the models,  a high mass spiral can have a bulge with a present-day low  $B/T \\leq$~0.2  only if it did not undergo a major merger since $z\\le 2$. The predicted fraction ($\\sim$~1.6\\%) of high mass spirals, which have undergone a major merger since $z\\le 4$  and host a bulge with a present-day low  $B/T \\le 0.2$, is  {\\it a factor of over thirty smaller} than the observed fraction  ($\\sim66\\%$) of high mass spirals with   $B/T \\le 0.2$. Thus, {\\it contrary to common perception, bulges built  via major mergers  since $z\\le 4$ seriously fail to account for the bulges present in $\\sim66\\%$ of  high  mass spirals}. Most of these present-day low $B/T \\leq 0.2$ bulges are likely to have been built by a combination of minor mergers and/or secular processes since  $z\\le4$. ",
        "introduction": "The formation of galaxies is a classic problem in astrophysics. Contemporary galaxy formation models combine the well-established $\\Lambda$ Cold Dark Matter ($\\Lambda$CDM) cosmology, which describes behavior of dark matter on very large scales, with baryonic physics to model galaxy formation. In the early Universe, pockets of dark matter decoupled from the Hubble flow, collapsed into virialized halos, and then clustered hierarchically into larger structures. Meanwhile, gas aggregated in the interiors of the halos to form rotating disks, which are the building blocks of galaxies (Steinmetz \\& Navarro 2002; Cole \\etal 2000). Such disks are typically destroyed during major mergers of galaxies with mass ratio $M_1/M_2 > 1/4$ (e.g. Steinmetz \\& Navarro 2002; Naab \\& Burkert 2003; Burkert \\& Naab 2004; Mihos \\& Hernquist 1996). When the mass ratio is close to unity, the remnant is a spheroid with properties close to that of a classical bulge, namely a steep de Vaucouleurs $r^{1/4}$ surface brightness profile and a high ratio of ordered-to-random motion ($v/\\sigma$). We shall return to this point in $\\S$~\\ref{sresult}. Within this hierarchical framework, the disk of  spiral  galaxies forms when gas of higher specific angular momentum subsequently accretes around the bulge (Steinmetz \\& Navarro 2002; Burkert \\& Naab 2004).  $\\Lambda$CDM-based simulations of galaxy formation face several challenges. One issue is the angular momentum problem; simulated galaxy disks have smaller scalelengths and, therefore, less specific angular momentum than their counterparts in nature (Navarro \\& Steinmetz 2000; Burkert \\& D'Onghia 2004; D'Onghia \\etal  2006). A second issue is the problem of bulgeless or low bulge-to-total mass ratio ($B/T$)  spirals.  Within the $\\Lambda$CDM paradigm, galaxies that had a past major merger at a time when its mass was a fairly large fraction of its present-day mass are expected to have a significant bulge  with large $B/T$ and high S\\'ersic index.  Depending on the merger history and hence the fraction of spiral galaxies that  fulfill this criterion (see $\\S$~\\ref{smodel1}) we can end up with a small or  large fraction of present-day  galaxies with low $B/T$.  There is rising evidence that low $B/T$ and bulgeless galaxies are quite common in the local Universe, especially in low mass or late-type galaxies.  Late-type Sd galaxies often harbor no bulge (B{\\\"o}ker \\etal  2002). Kautsch \\etal (2006) and Barazza, Jogee \\& Marinova (2007, 2008) also find from the analysis of several thousand late-type SDSS galaxies that 15-20\\% of such disk galaxies out to $z \\sim 0.03$ appear bulgeless. Of the 19 local galaxies ($D<8$~Mpc) with circular velocity $V_c>150$~km~s$^{-1}$, 11 (58\\%) have pseudobulges instead of merger-built classical bulges (Kormendy \\& Fisher 2008).  Theoretical work by Koda \\etal (2007) conclude the survival of disk-dominated systems in a $\\Lambda$CDM universe is compatible with observational constraints provided classical bulges form only in mergers where $M_1/M_2>0.3$ and the primary halo has virial velocity $V_{vir}>55$~km~s$^{-1}$.  Evidence also suggests that bulges with  low $B/T$ and low S\\'ersic index $n$ may be common even in high mass and/or early-type spirals. Balcells \\etal (2003) report that early-type disk galaxies tend to have $n <$~3 and often from 1 to 2. Laurikainen \\etal (2007) find barred and unbarred early-type disk galaxies to have mean $B/T$ between 0.25 and 0.35, while later Hubble types have $B/T<0.2$; they also find mean bulge S\\'ersic indices to be $\\sim2.5$ or less across the Hubble sequence. Graham \\& Worley (2008) report low $B/D$ ratios across the Hubble sequence based on bulge-disk decomposition of $K$-band images of local spiral galaxies. They suggest that these low values  are problematic for $\\Lambda$CDM simulations, but no quantitative assessment of the extent of the problems is presented.  These emerging statistics on the  fraction of bulgeless ($B/T\\sim0$) galaxies, and  galaxies with low $B/T$ and low $n$ bulges provide important first constraints. More work is needed to fully explore the the distribution of bulge properties in both high and low mass galaxies. In particular, we need to explore how the observed distributions of bulge $B/T$  and $n$ compare with the predictions from $\\Lambda$CDM-based simulations of galaxy evolution. To the best of our knowledge, few such quantitative comparisons have  been attempted, so that  it remains unclear how serious the problem of low $B/T$ galaxies is. This study is an attempt to derive robust observational constraints on bulge properties in high mass spirals and to  attempt such a comparison with models.  Completely resolving the issue of low $B/T$ systems will require understanding the different types of bulges and their formation pathways.  Bulges are commonly divided in  several groups:  classical bulges, boxy/peanut bulges, and `pseudobulges' or  disky bulges.  Classical bulges are believed to be built by major mergers ($M_1/M_2 \\ge 1/4$)  and the associated violent relaxation of stars. They are associated with modest-to-high bulge  S\\'ersic indices, in the range $2<n<6$ (Hopkins \\etal 2008; Springel \\etal 2005; Robertson \\etal 2006; $\\S$~\\ref{smodel1}). Boxy/peanut bulges are believed to be the result of vertical resonances and buckling instabilities in bars, which are viewed at high inclination (Combes \\& Sanders 1981; Combes \\etal 1990; Pfenniger  \\& Norman 1990;  Bureau \\& Athanassoula 2005; Athanassoula 2005; Martinez-Valpuesta \\etal 2006). Pseudobulges or  disky bulges are believed to form as a result of gas inflow into the central kiloparsec and  subsequent star formation building a  compact disky, high $v/\\sigma$  stellar component (Kormendy 1993; Jogee 1999; Kormendy \\& Kennicutt 2004, hereafter KK04;  Jogee, Scoville, \\& Kenney 2005;  Athanassoula 2005; Kormendy  \\& Fisher 2005). Pseudobulges  tend to have a bulge $n<2.5$ (Kormendy  \\& Fisher 2005; Fisher \\& Drory 2008).  One possibility for the formation of  disky bulges or pseudobulges  is the idea of secular evolution (Kormendy 1993; KK04; Jogee, Scoville, \\& Kenney 2005),  where a stellar bar or globally oval structure in a  {\\it non-interacting} galaxy drives the gas inflow into the inner kpc via  shocks and gravitational torque. Another idea for building disky bulges is that the gas inflow into the inner kiloparsec is driven by  {\\it external non-secular processes}, such as tidal interaction and minor mergers. The gas inflow in such cases can be caused by a  tidally induced non-axisymmetric feature, such as a bar (e.g.,  Quinn et al. 1993; Hernquist \\& Mihos 1995), and by tidal torques from  the companion.  The  subsequent central star formation can still form a compact  high $v/\\sigma$  stellar component, $aka$ a pseudobulge.  Throughout this paper,  we avoid making any {\\it a priori} assumptions about the origin of different types of bulges by simply referring to them according to their bulge S\\'ersic index $n$ or  bulge-to-total mass ratio ($B/T$).  We consider bulges of high  ($n\\ge$~4), intermediate ($2<n<4$) and low ($n \\leq 2$) index, as well as those of low or high $B/T$.  The structural properties of galaxy components, such as  bulges, disks, and bars can be derived through the decomposition of the $2D$ light distribution, taking into account the PSF. Many early studies have performed  only   two component $2D$  bulge-disk decomposition (e.g., Allen \\etal 2006; Byun \\& Freeman 1995; de Jong 1996; Simard 1998; Wadadekar \\etal 1999), ignoring the contribution of the bar, even in strongly barred galaxies. However, recent work has shown that it is important to include the bar in $2D$ decomposition of barred galaxies, else the $B/T$ ratio can be artificially inflated, and bulge properties skewed (e.g., Laurikainen \\etal  2005, 2007). Furthermore, since most ($\\ge60\\%$) bright spiral galaxies are barred in the NIR (Eskridge  \\etal  2000; Laurikainen \\etal 2004; Marinova \\& Jogee 2007, hereafter MJ07;  Menendez-Delmestre \\etal 2007), the inclusion of the bar is quite important.  This has led to several recent studies, where $2D$ bulge-disk-bar decomposition are being performed  (e.g.   Laurikainen \\etal  2007; Reese \\etal 2007; Gadotti \\& Kauffmann 2007).  Another advantage of bulge-disk-bar decomposition over bulge-disk decomposition is that the former allows us  to constrain the properties of the bar itself. Bars provide the most important internal mechanism for redistributing angular momentum in baryonic and dark matter components (e.g.  Weinberg 1985; Debattista \\& Sellwood 1998, 2000; Athanassoula 2002; Berentzen, Shlosman, \\& Jogee 2006). They efficiently drive gas inflows into the central kpc, feed central starbursts  (Elmegreen 1994; Knapen \\etal  1995; Hunt \\& Malakan 1999; Jogee \\etal 1999;  Jogee, Scoville, \\& Kenney 2005; Jogee 2006) and lead to the formation of disky or pseudobulges (see above). Furthermore, the prominence of strong bars out to $z\\sim 1$  over the last 8 Gyr  (Jogee \\etal 2004; Sheth \\etal 2008) suggest that bars have been present over cosmological times and can shape the  dynamical  and secular evolution of disks. Thus, quantifying bar properties, such  as the fractional light and mass ratio (Bar/$T$), can yield insight into these processes.  In this paper, we constrain the properties of bulges and bars along the Hubble sequence, and compare our results with $\\Lambda$CDM-based simulations of galaxy evolution. In $\\S$~\\ref{sdata}, we define our complete sample of 143 bright ($M_B\\leq-19.3$) low-to-moderately inclined ($i \\le 70^\\circ$) spirals from the Ohio State University Bright Spiral Galaxy Survey (OSUBSGS; Eskridge \\etal 2002), which is widely used as the local reference sample for  bright spirals by numerous studies  (e.g.,  Eskridge \\etal 2000; Block \\etal 2002;  Buta \\etal 2005; MJ07 ; Laurikainen \\etal 2004, 2007). In $\\S$~\\ref{sanaly}, we  perform $2D$ bulge-disk and bulge-disk-bar decompositions of $H$-band images using GALFIT (Peng \\etal 2002), and derive  fractional light ratios ($B/T$, Bar/$T$, Disk/$T$), as well as  S\\'ersic indices and half light radii or scale lengths. Tests to verify the  robustness  of our decompositions  are presented in $\\S$~\\ref{stests}. In $\\S$~\\ref{sresult}, we present our results. Specifically, the total stellar mass present in bulges, disks, and bars is calculated  ($\\S$~\\ref{smass}). In  $\\S$~\\ref{sbulgenbt}, the distribution of bulge S\\'ersic  index $n$ and $B/T$ as a function of galaxy Hubble type and stellar mass is presented, and the surprising prevalence of bulges with low S\\'ersic index $n$ and low $B/T$ is established. A comparison with other works is presented in $\\S$~\\ref{scomp1}. We  examine how Bar/$T$ and bar fraction (the fraction barred disks) change as a function of host galaxy properties in $\\S$~\\ref{sbar2}. In  $\\S$~\\ref{smodel1}, we compare our  observed  distribution of bulge $B/T$ and $n$ in  high mass ($M_\\star \\geq  1.0\\times 10^{10} M_\\odot$) spirals with predictions from $\\Lambda$CDM  cosmological semi-analytical models. $\\S$~\\ref{ssumm} summarizes our results. ",
        "conclusions": "\\label{sresult}  \\subsection{Impact of Bars in $2D$ Decomposition }\\label{sbar1} From the Stage 2 bulge-disk decomposition  and Stage 3 bulge-disk-bar decompositions, which we performed on all objects ($\\S$~\\ref{sdecom2}) we saw firsthand the effects of adding a bar to the fit of a barred galaxy. We summarize below some of these effects in order to underscore the importance of including a bar component in the $2D$ luminosity decomposition of barred galaxies  \\begin{enumerate} \\item During the Stage 2 bulge-disk decomposition of a barred galaxy, the luminosity which comes from the galaxy's disk, bulge, {\\it and bar} gets distributed only between two model components: the model bulge and disk. Since the disk $b/a$ and $PA$ are measured independently and held constant during the fits, the Stage 2 model tends to distort the bulge in order to fit the bar. Thus, the bulge in the Stage 2 bulge-disk decomposition  of a barred galaxy can be artificially long or too bright and extended. When a model bar component is added in the Stage 3  bulge-disk-bar decomposition of a barred galaxy, it forces a reshuffling of the luminosity between the three components. Generally, the bulge declines in luminosity, whereas light can be either taken from, or added back, to the disk. \\item We find that the inclusion of a bar component in the  Stage 3 bulge-disk-bar decomposition  of a barred galaxy   reduces the bulge fractional luminosity  $B/T$, compared with the Stage 2 bulge-disk decomposition. For our 75 barred galaxies, the reductions correspond to factors of less than two, 2 to 4, and above 4, in   36\\%,  25\\%, and 39\\% of barred galaxies, respectively. The larger changes in $B/T$  occur in very strongly barred galaxies, where a prominent bar cause the Stage 2 bulge-disk decomposition to overestimate the  bulge. For instance, $B/T$ declines in both of NGC 4643 (Figure~\\ref{ngc4643} and Table~\\ref{tngc4643}) and NGC 4548 (Figure~\\ref{ngc4548} and Table~\\ref{tngc4548}). In the latter case, $B/T$ is reduced by a factor of $\\sim5$ between Stage 2 and Stage 3. These examples underscore the importance of including bars in $2D$ luminosity decomposition of very strongly barred galaxies. \\item The scalelength of the disk is generally unchanged by including the bar. NGC 4548 (Figure~\\ref{ngc4548} and Table~\\ref{tngc4548}) is a good example.  Sometimes, however, the disk from the Stage 2 bulge-disk decomposition of a barred galaxy is erroneous due to a poor fit.  The disk parameters from the Stage 3 bulge-disk-bar decomposition are quite different in such cases. NGC 4643 (Figure~\\ref{ngc4643} and Table~\\ref{tngc4643}) illustrates this behavior. \\end{enumerate}  We find that for  our sample S1 of 143  bright ($M_B\\leq-19.3$) low-to-moderately inclined  ($i \\le 70^\\circ$) spirals  (Figure~\\ref{2posu-hist-hubb})  in the OSUBSGS survey, 75 of 143  or $\\sim52\\%$ are better fit with a Stage 3 bulge-disk-bar decomposition than a Stage 2 bulge-disk decomposition. There are also eight galaxies with pure bar-disk fits.  As stated in $\\S$~\\ref{schoice23}, the resulting $H$-band bar fraction ($58.0\\% \\pm 4.13\\%$ or 83 of 143)  is in excellent agreement with the $H$-band bar fraction of 60\\% reported by MJ07  based on ellipse fits of the OSUBSGS sample, with  a slightly more conservative inclination cut ($i\\leq 60^\\circ$).  \\subsection{Mass in Bulges, Disks, and Bars}\\label{smass} The  fractional $H$-band luminosities in the bulge, disk, and bar ($B/T$, $D/T$, Bar/$T$) of each galaxy can be considered as a fractional mass if we assume that the same mass-to-light ($M/L$) ratio can be used to convert the $H$-band luminosities  of both the numerator ($B$, $D$, or Bar) and the denominator ($T$) terms into a stellar mass. This is not an unreasonable assumption as the $H$-band $M/L$ ratio is not very sensitive to differences in dust or age that might exist between  the bulge, disk, and bar.  The uncertainties in  $M/L$ can  be estimated by looking at population synthesis models. Charlot, Worthey, \\& Bressan (1996) find that for idealized galaxies with a single generation of stars, the  uncertainties in $M/L$ ratio  due to different input stellar models and spectra are roughly $\\pm35\\%$ for a fixed metallicity and IMF.  Furthermore, as the age of a stellar population varies from $\\sim0.5$ Gyr to 10 Gyr, the $K$-band $M/L$ ratio rises by a factor of $\\sim$ 2 to 3  (Charlot 1996).  Asymptotic giant branch (AGB) stars dominate the  NIR light for ages between 0.1 and 1 Gyr, while  red giant branch (RGB) and super-giant branch (SGB) stars dominate between 1 Gyr and 10 Gyr.  In this paper, we convert the $B/T$ light ratio determined from $H$-band images to a $B/T$ mass ratio by assuming a constant mass-to-light ($M/L$) in the $H$-band for both the bulge and the rest of the galaxy. However, Moorthy \\& Holtzman (2006) present line strengths of bulges and inner disks for 38 spirals with Hubble types S0 to Sc.  They show 76\\% of spirals have negative metallicity gradients. $B-K$ color gradients are shown to largely match metallicity gradients and are likewise negative outward, indicating bulge $M/L$ is higher than in the disk and bar. If bulges are much older than the disks, then our prescription  would  underestimate  the true $B/T$ mass ratio. If we assume an extreme case where bulges are $\\sim12$ Gyr and the disk light is dominated  by a young 3 Gyr population, our assumption of a constant $H$-band $M/L$ ratio  would underestimate the true $B/T$ by a factor of $\\le2$ (Charlot 1996). In several sections of the paper (e.g., $\\S$~\\ref{sbulgenbt}, $\\S$~\\ref{smodel1}), we illustrate how our main results  would change if the true  $B/T$ was higher by up to a factor of two. On the other hand, central regions of galaxies may harbor intense episodes of star formation. If the bulge is younger than the disk and happens to have star formation and a significant young population of massive stars, then our prescription could overestimate the true $B/T$ mass ratio. This would make our current results on the high fraction of low $B/T$ bulges even stronger.  Using the total galaxy stellar mass from $\\S$~\\ref{stellmass}, the fractional masses can be converted into absolute masses. (We do not convert the $H$-band luminosity directly into a mass as the $H$-band images do not have photometric calibration). The results are shown in Table~\\ref{masstable}. For  our sample S1 of  143  bright ($M_B\\leq-19.3$) low-to-moderately inclined  ($i \\le 70^\\circ$) spirals with a mass-weighted mean Hubble type of Sab-Sb, we find that {\\it $\\sim70\\%$ of the stellar mass is in disks,  $\\sim$~10\\%  is in stellar bars and $\\sim20\\%$ is in bulges (with $\\sim11\\%$ in $n>2$ bulges and $\\sim9\\%$ in $n \\leq 2$ bulges).} Thus while bulges with $n\\le2$  are highly ubiquitous (see next section), they only account for a small fraction of the total stellar mass budget.  Figure~\\ref{3p-plot2} shows the stellar mass for bulges, disks, and bars  along the Hubble sequence.  It is useful to compare our results with those of Driver \\etal (2006), who performed bulge-disk decomposition of $B$-band images with GIM2D on 10,095 galaxies from the Millennium Galaxy Catalog  (Liske \\etal 2003; Driver \\etal 2005).  They found  68.6\\% of the stellar mass to be in disks, and 32.6\\% in bulges (with 30.8\\% in high $n$ bulges, and 1.8\\% in low $n\\le2$ bulges). Their study thus finds a  higher stellar mass fraction  in all bulges (32.6\\% vs  our 18.9\\%), and in high $n$ bulges (30.8\\% vs our 10.4\\%),  and a lower fraction in low $n\\le2$ bulges  (1.8\\% vs our  8.4\\%),  and disks+bars (68.6\\% vs our  71.6\\% + 9.6\\%). This difference can be attributed to the fact that  the  Driver \\etal (2006) study did not perform  bulge-disk-bar fits and thus, their $B/T$ ratios may be skewed to higher values.  \\subsection{Distribution of Bulge Index and $B/T$}\\label{sbulgenbt} Figure~\\ref{4p-bplot1}  shows the individual and mean  $B/T$ and bulge S\\'ersic index, plotted, as a function  of Hubble type and galaxy stellar mass. Barred and unbarred galaxies are shown separately. Figure~\\ref{2p-bplot2} shows the relationship between bulge index and $B/T$.  We first consider the $B/T$ values in  Figure~\\ref{4p-bplot1}. The mean $B/T$  in barred galaxies is lower than in unbarred galaxies, but there is a large overlap in the individual values. The offset in the mean $B/T$ of  barred and unbarred galaxies reported here,  agrees with the result of Laurikainen \\etal (2007; see $\\S$~\\ref{scomp1}) on the same sample. We also note that  $B/T$ does not correlate with Bar/$T$ (Fig.~\\ref{2p-brplot5}):  aside from the six galaxies with large Bar/$T$ ($>0.3$), most galaxies have low-to-moderate Bar/$T$ and a wide range of $B/T$ is seen at each Bar/$T$. This is reassuring and suggests that the bar fit is not arbitrarily biasing the $B/T$ values.  The distribution of Bar/$T$ is further discussed in $\\S$~\\ref{sbar2}.  How does  the $B/T$ vary as a function of Hubble type and galaxy stellar mass? Bulges with very high $B/T$ ($>0.4$)  exist primarily in galaxies with high mass ($M_\\star > 6 \\times 10^{10} M_\\odot$) and early types (S0/a to Sab). Bulges with very low  $B/T$ ($<0.1$) lie primarily in lower mass galaxies with later morphologies (Sb to Sc). {\\it It is striking that  $\\sim69\\%$   of  bright ($M_B\\leq-19.3$) low-to-moderately inclined  ($i \\le 70^\\circ$) spirals have $B/T \\leq 0.2$: these bulges are pervasive and exist across  the whole spectrum of S0/a to Sd. } The results are summarized in Table~\\ref{bulge1}. We shall return to this point in \\S~\\ref{smodel1}. We note again that these $B/T$ mass ratios were calculated assuming a constant $M/L$ ratio in the $H$-band for the bulge and disk components. As noted in $\\S$~\\ref{smass}, if the bulge in these high mass spirals is much younger (older) than the disk and bar, then the $B/T$ can be overestimated  (underestimated) by up to a factor of two, and the limiting value of  0.2 for the $B/T$ cited in the above fraction, would have to be modified in the extreme case to  0.1 and 0.4, respectively.  Some of the low $B/T \\le 0.2$  values for six barred S0/a and Sa galaxies in Figure~\\ref{4p-bplot1} may at first look suspicious. Balcells \\etal (2007) report the mean $B/T$ for S0 galaxies to be 0.25, so much smaller $B/T$ are potentially worrisome. OSUBSGS $H$-band images of these objects in Figure~\\ref{lowBT} show a smooth extended disk around the bulge. It should be noted that Hubble types were originally assigned on a combination of criteria including disk smoothness and spiral arm topology in addition to the prominence of the bulge. It is likely that these galaxies were assigned early Hubble types due to their smooth extended disks, in spite of their low bulge-to-disk ratio.  Similarly, some of the high  $B/T \\sim 0.4$ bulges in three of the Sc galaxies may at first seem odd. However, again, visual inspection  of their image (Figure~\\ref{lowBT}) reveals prominent spiral arms and clumpiness, which may explain why they were assigned late Hubble types.   How does  the bulge S\\'ersic index $n$ vary as a function of Hubble type, and galaxy stellar mass (Figure~\\ref{4p-bplot1}), as well as $B/T$ (Figure~\\ref{2p-bplot2})? The results are summarized in Table~\\ref{bulge1}. Only a small fraction ($\\sim1\\%$) of bright spirals have high $n\\ge4$ bulges; such bulges lie primarily in S0/a to Sab, and have a large $B/T > 0.2$. A  moderate fraction ($\\sim22\\%$)   have intermediate $2<n<4$ bulges; these exist in barred and unbarred S0/a to Sd,  and their $B/T$ spans a wide range  (0.03 to 0.5) with a mean of 0.29. {\\it A strikingly large fraction ($\\sim76\\%$) of bright  spirals have low $n \\leq 2$ bulges; such bulges exist  in barred and unbarred galaxies across a wide range of Hubble  types, and  their $B/T$  varies from 0.01 to 0.4, with most having $B/T \\leq0.2$. }  \\subsection{Comparison With Other Work}\\label{scomp1} As an independent check of our decomposition method, we compare our results with independently published decompositions.  We find our mean $H$-band $B/D$ (Figure~\\ref{log10BD}) ratios are comparable to the $K$-band $B/D$ derived with the $1D$ bulge-disk decompositions of Graham (2001) and Trujillo \\etal (2002). Like Graham (2001) and Trujillo \\etal (2002), we find $B/D$ is widely variable with Hubble type and that mean $B/D$ steadily declines from Sa through Scd galaxies.  Graham (2001) and Trujillo \\etal (2002) find bulge indices are widely variable within a Hubble type, but they are in general $>1$ for early types and $<1$ for late types. We likewise find wide scatter in bulge index with $n<1$ bulges existing in both early and late types. Figure~\\ref{log10BD} is also in good agreement with the more recent results of Graham \\& Worley (2008), who have calculated, as a function of morphology, inclination and dust-corrected $B/D$ and bulge S\\'ersic indices. They find $B/D$ values are typically $<1/3$.  Another meaningful comparison can be made with Laurikainen \\etal (2007)  who, using their own $2D$ decomposition code, fit a hybrid sample containing some OSUBSGS galaxies combined with additional S0 galaxies.  One difference between their work and ours is that they typically model bars with a Ferrers function, but may sometimes use a S\\'ersic profile, while we use only the latter. Also, they include additional components to model secondary bars or inner disks. They report a distinct offset in the mean $B/T$ between barred and unbarred galaxies,  which we confirm in Figure~\\ref{4p-bplot1}.  Their mean $B/T$ are similar to ours, and they conclude that pseudobulges exist throughout the Hubble sequence. The S\\'ersic indices derived by Laurikainen \\etal (2007) are likewise similar, on the mean, to ours for both barred and unbarred systems.  They likewise find $n\\le2$ bulges across early and late Hubble type galaxies.  Balcells \\etal (2003) emphasized that bulges typically have indices $\\sim3$ or lower. Our results in Figure~\\ref{4p-bplot1} are consistent.  We find a low frequency ($\\sim1\\%$) of high $n\\geq4$ bulges, with most bulges having $n\\leq3$.  \\subsection{Bar Strength}\\label{sbar2} Stellar bars exert gravitational torques on the gaseous component and are particularly efficient in driving  gas from the outer disk into the inner kiloparsec (see $\\S$~\\ref{sintro}).  Thus, it would be natural to have  a measure of bar strength that is  sensitive to the strength of the gravitational torque and  hence measures both the shape and mass of the bar.  Many measures of bar strength have been formulated. The $Q_b$ method of (Block \\etal 2002; Buta \\etal 2003; Buta \\etal 2005) measures directly the gravitational torque at a single point along the bar.  This method requires a value of scaleheight for the disk and a model of the potential to be made from the image. In the bar/inter-bar contrast method of Elmegreen \\& Elmegreen (1985) and Elmegreen \\etal (1996), bar strength is parameterized as the ratio between peak surface brightness in the bar region and the minimum surface brightness in the inter-bar region.  Elmegreen \\& Elmegreen (1985) and Elmegreen \\etal (1996) also characterize bar strength with the maximum amplitude of the $m=2$ mode from Fourier decomposition. When ellipse fitting is applied, the maximum ellipticity of the bar, $e_{bar}$, can be used to characterize bar strength (e.g. MJ07).  This constitutes a partial measure of bar strength only, however, as it offers no information about mass of the bar.  Bulge-disk-bar decomposition  in the $H$-band  provides another \\emph{partial} measure of bar strength through the $H$-band Bar/$T$ light ratio, which is a measure of the Bar/$T$ mass ratio under the assumption that the $H$-band $M/L$ ratio is the same for the bar and the rest of the galaxy, as discussed in $\\S$~\\ref{smass}. Figures~\\ref{fbar1} and \\ref{fbar2} explore the derived bar properties.  The upper left panel of Figure~\\ref{fbar1} plots the individual and mean Bar/$T$ against Hubble type. There is a wide range ($\\sim0.03$ to  $\\sim0.47$) in the  individual Bar/$T$ at a given Hubble type.  The mean Bar/$T$ remains fairly constant with Hubble type from Sa to Sb, but shows a possible weak decline by about 0.1 from Sb to Sc.    Their number statistics are too small to make any robust statement for later Hubble types. We also note that six systems have high  Bar/$T$ above $0.3$; these are displayed  in Figure~\\ref{highBrT}.  Bar S\\'ersic indices are mostly below unity. Neither the individual, nor the mean bar S\\'ersic index, show any trend with Hubble type or with stellar mass, for Sa to Sc galaxies (Fig.~\\ref{fbar1}). Thus, the steepness of the bar profile does not seem to depend on the Hubble type. Is the bar mass ratio and its mass profile related?  There is a wide range in the individual   Bar/$T$  at a given bar S\\'ersic index (Fig.~\\ref{fbar2}). The mean Bar/$T$  rises with bar index  out to a bar index of $\\sim0.6$, and then flattens out. This suggests that on the mean, bars of lower Bar/$T$ have flatter profiles.  Is there a relation between the bar strength and the bulge present in a galaxy? There is a wide range in the  individual   Bar/$T$  at a given $B/T$, and at a given bulge S\\'ersic index (Fig.~\\ref{fbar2}). The  mean  Bar/$T$  shows a weak decline for bulge S\\'ersic indices above two. Similarly  the mean  Bar/$T$  shows a weak rise from 0.1 to 0.25 as  $B/T$ rises out to 0.15, after which the trend flattens or reverses.  Both Bar/$T$ and maximum bar ellipticity $e_{bar}$ are partial measures of bar strength.  Figure~\\ref{fbar1} shows mean Bar/$T$ may scale weakly with Hubble type.  The bars with highest $e_{bar}$  (i.e, thin bars) are often termed strong bars, and  $e_{bar}$ has been shown to correlate with $Q_b$. Total bar strength should scale with both bar mass and bar ellipticity. Does bar strength have a dependence on Hubble type? The upper left panel of Figure~\\ref{fbar2} plots the product of Bar/$T$ and $e_{bar}$, as determined by MJ07 for galaxies mutually classified as barred, against Hubble type. There is a wide range in Bar/$T\\times e_{bar}$ in each bin, and mean bar strength shows no definite trend with Hubble type. We note that bars with high Bar/$T$ and high $e_{bar}$ should exert the largest  gravitational torque and be most effective at driving gas inflows. A nice example is the oval or lens  galaxy NGC 1317 (Figure~\\ref{highBrT}); the bar has a low ellipticity, but its $B/T$ is large as it is extended and massive. Such bars or lenses may exert significant gravitational torques although they are not very elongated.  \\subsection{Bar Fraction as a Function of $B/T$ and Bulge Index} \\label{sbar3} As outlined in $\\S$~\\ref{sbulgenbt}, we found that as many as  $\\sim76\\%$ of bright spirals have bulges with $n \\leq 2$; such bulges exist in barred and unbarred galaxies across a wide range of Hubble types, and  their spread in $B/T$ is from 0.01 to 0.4, with most having $B/T \\leq$ 0.2. The variation of the bar fraction as a function of $B/T$ and bulge $n$ can provide important constraints  on bulge formation scenarios ($\\S$~\\ref{smodel1}). Table~\\ref{barFraction} shows our results. The bar fraction declines with bulge index; $\\sim65\\%$ of the spirals with low $n\\le2$ bulges host bars while intermediate $2<n<4$ bulges have a lower bar fraction ($\\sim38\\%$). The high $n\\ge4$ bulges in the sample are unbarred, so the bar fraction is $0\\%$. Systems with low $B/T$ are more likely to be barred. For $B/T \\le 0.2$, the bar fraction is high ($\\sim68\\%$).  Systems with $0.2<B/T<0.4$ and $B/T \\ge 0.4$ have lower bar fraction ($\\sim42\\%$ and $\\sim17\\%$).  Overall, Table~\\ref{barFraction} shows bulges with low $n\\leq2$ and low $B/T \\le 0.2$ preferentially exist in barred galaxies. This is consistent with earlier work (Odewahn 1996; Barazza \\etal 2008; Marinova \\etal 2008; Aguerri \\etal 2008, in prep.) where an enhanced optical bar fraction is seen is galaxies with late Hubble types or low $B/D$. It may be tempting to infer this result to mean the formation pathway of two-thirds of low-$B/T$ bulges is related to bars in that spontaneous or/and tidally induced bars play a role in bulge formation (with the remaining one-third of such bulges may have been formed either by mechanisms like retrograde minor mergers or short-lived bars). We caution that this type of cause-effect relationship is not the only scenario consistent with this result.  It may also be possible that bar instabilities are favored in galaxies with low $B/T$ and no inner Lindblad resonances (ILR).  Under these conditions, the swing amplifier model with a feedback loop (Julian \\& Toomre 1966; Toomre 1981; Binney \\& Tremaine 1987) may be responsible for bar formation and partly account for the high bar fraction in  galaxies of low $B/T$.  \\subsection{Formation of Bulges}\\label{bulgeFormation} Our observational results provide some interesting challenges for models of galaxy evolution that try to  address the origin of present-day bulges. Any successful model must be able to account for the observed distribution of bulge $B/T$ and $n$ in  high mass ($M_\\star \\geq  1.0\\times 10^{10} M_\\odot$) spirals, as shown in  Table~\\ref{bulge1} and Table~\\ref{barFraction}. In particular, the following  results must be reproduced:  \\begin{enumerate} \\item In terms of the overall distribution of bulge $n$, as much as ($\\sim74\\%$) of high mass spirals have bulge $n\\leq2$: such bulges exist  in barred and unbarred galaxies and  their $B/T$ ranges from 0.01 to 0.4,  with most having $B/T \\leq0.2$ (Table~\\ref{bulge1}). A moderate  fraction   ($\\sim24\\%$) of high mass spirals have $2<n<4$, and just ($\\sim2\\%$) have $n\\geq4$.  \\item Theoretical models often make more robust predictions on the bulge-to-total mass ratio $B/T$  than on the bulge index $n$, so we consider the  empirical $B/T$ distribution in detail. We note  that as much as $\\sim66\\%$ of high mass spirals have bulges with $B/T \\leq 0.2$ (Table~\\ref{bulge1}). In terms of bar fraction, $\\sim68\\%$ are barred  (Table~\\ref{barFraction}).  \\item The fraction of bars rises among spirals with low bulge index $n$. About 63\\% of spirals with low $n\\leq2$ bulges host bars, while the bar fraction in spirals with $2<n<4$ bulges (44\\%) is two-thirds as large (Table~\\ref{barFraction}). \\end{enumerate}  In a hierarchical Universe, there are several physical processes that contribute to the assembly of bulges: major mergers, minor mergers, and secular evolution. We briefly describe these, expanding on our introduction in $\\S$~\\ref{sintro}.  Major mergers,  defined as those with mass ratio   $M_1/M_2 \\ge 1/4$, typically destroy the extended outer stellar disks during violent relaxation, leaving behind  a classical bulge.  Such bulges are associated with modest-to-high bulge  S\\'ersic indices, in the range $2<n<6$ (Hopkins \\etal 2008; Springel \\etal 2005; Robertson \\etal 2006; $\\S$~\\ref{smodel1}) in simulations. This trend is also consistent with the fact that among ellipticals, high luminosity ones tend to have  a S\\'ersic index $n>4$, while  low luminosity ones tend to have $2 \\le n \\le 3.5$  (Caon \\etal 1993; Kormendy et al. 2008). The final  S\\'ersic index depends on the amount of residual gas the settles into a somewhat disky component. Simulations by Hopkins \\etal (2008) find that the S\\'ersic indices  of remnants from 1:1 gas-rich major mergers lie in the range of $2<n<4$, with most above 2.5 (see Fig.~\\ref{fhopki1}). This body of evidence strongly suggests that many bulges with $n>2$ might have a major merger origin.  Minor mergers, typically defined as those with mass ratio $1/10 < M_1/M_2 < 1/4$, do not destroy the stellar disk of the primary system,  but can contribute to building bulges via three pathways. Firstly,  a fraction  $F_{sat}$  of the satellite's stellar mass can end up in the central region of the primary galaxy. The value of  $F_{sat}$ depends on how centrally concentrated the in-falling satellite is. Typically, the more diffuse outer stellar body is tidally stripped, while the  central core sinks by dynamical friction to the central region (e.g., Quinn et al. 1993; Walker et al. 1996). Secondly,  a non-axisymmetric feature (e.g., a stellar bar or bar-like feature) can be induced in the main disk, and gravitational torques exerted by the feature can drive gas into the inner kpc (e.g., Hernquist \\& Mihos 1995; Jogee 2006 and references therein), where subsequent SF forms a compact  high $v/\\sigma$  stellar component, or disky pseudobulge. Most of the gas inflow happens during the merger phase and large gas inflow rates (e.g., $\\gg1$  $M_\\odot$ per year) may  be generated. Thirdly,  gas inflow can also be caused by direct tidal torques from  the companion (e.g., Hernquist \\& Mihos 1995; Eliche-Moral \\etal 2006). It is to be noted that in the simulations by  Hernquist \\& Mihos (1995),  the gas inflow  generated by non-axisymmetric features (e.g., bar-like features) in the inner part of the disk is much larger than that caused by direct tidal torques from the satellite. In the recent work of Eliche-Moral \\etal (2006), N-body simulations of minor mergers followed by fits of 1D S\\'ersic+exponential models to the remnants, suggest that the bulge S\\'ersic index and $B/D$ ratio can grow as a result of the central re-concentration of stellar disk material in the primary system by tidal forces. Minor mergers are frequent under $\\Lambda$CDM, and the likelihood of multiple successive minor mergers occurring during the formation of a galaxy is high. Bournaud, Jog, \\& Combes (2007) study the effects of repeated minor mergers on galaxy structure.  They show that a disk galaxy undergoing successive minor mergers will eventually transform into an elliptical galaxy with an $r^{1/4}$-law profile and high $V/\\sigma$.  However, galaxy growth is not completely merger-driven and the efficiency of minor mergers at creating ellipticals must be regulated by other mechanisms (e.g., cold gas accretion).  In addition, the process of secular evolution can build  a disky bulge (pseudobulge) between merger events.  Here  a stellar bar or globally oval structure in a  {\\it non-interacting} galaxy drives  gas inflow into the inner kpc, where subsequent star formation forms a compact  high $v/\\sigma$  stellar component (e.g., Kormendy 1993; Jogee 1999; KK04;  Jogee, Scoville, \\& Kenney 2005;  Athanassoula 2005; Kormendy \\& Fisher 2005; Kormendy 2008). This process is different from that of minor mergers in the sense that it happens in the quiescent phase of the galaxy, between minor or major merger events. The prevalence of pseudobulges in galaxies of different Hubble types is discussed in KK04, and select examples of S0 galaxies with pseudobulges are also shown in Kormendy \\& Cornell (2004) and KK04.  The present-day bulge mass can be written as the sum of mass contributed from each process: \\begin{equation} \\label{bulgemass} M_{bulge}=M_{bulge}\\times(f_{\\rm maj} +  f_{\\rm min1} + f_{\\rm min2} + f_{\\rm min3} + f_{\\rm sec}), \\end{equation}  where  \\begin{itemize} \\item $f_{\\rm maj}$ is the percentage of the  bulge stellar mass, which is built by  major mergers, \\item $f_{\\rm min1}$ is the percentage of the bulge stellar mass, which is built during minor mergers from stars of the satellite. This depends on the  fraction  $F_{sat}$  of the satellite's stellar mass, which ends up in the central region of the primary galaxy during each minor merger. \\item $f_{\\rm min2}$ is the percentage of the bulge stellar mass, which is built during minor mergers from gas inflow  caused by a tidally induced bar. \\item $f_{\\rm min3}$ is the percentage of the stellar mass, which is built during minor mergers from gas inflow caused by tidal torques from the companion. \\item $f_{\\rm sec}$ is the percentage of the stellar mass, which is  built secularly between  merger events from gas inflow  caused by  bars or ovals \\end{itemize}  In \\S~\\ref{smodel1}, we compare our derived distribution of bulge $n$ and $B/T$ with hierarchical models that model major and minor mergers, but not secular evolution.  The main goal of the model is to see whether bulges built via major mergers can account for the large fraction of high mass spirals  with bulges of low $B/T$  or/and low $n$. A secondary goal is to see if a first order simplified prescription for minor mergers can broadly account for the observations. We stress here that bulge-building during minor mergers is  modeled in a very simple way: all the stars in the satellite are assumed to contribute to the bulge of the larger galaxy (i.e., $F_{sat}=100\\%$), and bulge-building via gas inflow driven through tidal torques and via gravitational torques from induced bars are ignored (i.e., $f_{\\rm min2}$~=~0, and $f_{\\rm min3}$~=~0).  Furthermore, the models entirely ignore secular evolution between mergers.  In a future paper, these extra terms will be addressed and a comprehensive picture built of the relative importance of  minor mergers and secular processes in making present-day bulges.   \\subsection{Comparison of $B/T$ With Hierarchical Models of Galaxy Evolution}\\label{smodel1} We compare our data with the  predictions from  cosmological semi-analytical models based on Khochfar \\& Burkert (2005) and Khochfar \\& Silk (2006).  We briefly describe the models first. The merger trees of dark matter (DM)  halos are derived by using the  extended Press-Schechter formalism (Press \\& Schechter 1974), as in Somerville \\& Kolatt (1999). When two DM halos merge, the merger time scale of the galaxies is calculated by considering the timescale it would take the satellite galaxies to reach the central galaxy at the center of the halo via dynamical friction (e.g., Kauffmann \\etal 1999;  Springel \\etal 2001). The baryonic physics, which includes radiative cooling, star formation, and feedback from supernovae, is treated via semi-analytic prescriptions (see Khochfar \\& Silk (2006) and references therein). Baryonic mass inside the dark matter halos is divided between hot gas, cold gas, and stars.  The hot gas is initially shock-heated to the halo virial temperature.  As the gas radiatively cools, it settles down into a rotationally supported disk at the halo center.  Cold disk gas is allowed to fragment and subsequently form stars according to the Schmidt-Kennicutt law (Kennicutt 1998). Star formation is regulated by feedback from supernovae using the prescription in Kauffmann \\etal (1999).  Major mergers are typically considered as those with stellar mass ratio $M_1/M_2 \\ge 1/4$. In the simulations, one assumes that during a major merger any  existing stellar disk is destroyed,  gas is converted to stars with some star formation efficiency (SFE), and all stars present undergo violent relaxation to form a bulge.  Therefore, the  bulge-to-total stellar mass ratio  ($B/T$)  of a bulge immediately after a major merger is always one. Note  that the SFE during a major merger is not assumed to be 100\\%  as there is mounting  evidence from  SPH simulations (Springel \\& Hernquist 2005; Cox \\etal 2008) that not all cold gas is converted to stars. Instead, the burst efficiency defined by Cox \\etal (2008) is applied to control the fraction of stars formed due to the interaction. This efficiency is dependent on the relative masses of merging galaxies and is expressed as \\begin{equation} \\label{cox} e=e_{1:1} {\\left(\\frac {M_{Satellite}} {M_{Primary}}\\right)}^\\gamma, \\end{equation} where $e_{1:1}$ is the burst efficiency for a 1:1 merger and $\\gamma$ fixes the dependence on mass ratio; Cox \\etal (2008) find $e_{1:1}=0.55$ and $\\gamma=0.69$. The remaining fraction (1-$e$) of gas is added to the gaseous disk and can start making stars.  As stated above, immediately after a major merger, the remnant is a bulge  with a $B/T$ equal to one. As time proceeds, $B/T$ falls because a stellar disk grows around the bulge as hot gas in the halo cools, settles into a disk, and forms stars. The formation of stars by any residual cold gas left at the end of the major merger also helps to grow the disk.  Thus $B/T$ falls until the next major merger happens, at which point $B/T$ is reset to one in the models.  The bulge may also grow in stellar mass due to minor mergers. Minor mergers are defined as mergers with mass ratio $ 1/10 < M_1/M_2 < 1/4$, and the stellar disk of the large companion is not destroyed during such mergers. The models assume that during minor mergers, all  the stars in the  satellite are added to the bulge of the host, while the gas settles in the disk. When DM halos grow by accretion or minor mergers, the  hot gas that comes in with a satellite is immediately stripped and added to the hot gas component of the host. The cold gas in the disk of the satellite is only added to the cold gas of the host if they merge.  Until they merge the satellite is using up its own cold gas to make stars.  Fig.~\\ref{mergehis} shows the relationship between the present-day $B/T$ of a a high mass ($M_\\star \\geq  1.0\\times 10^{10} M_\\odot$) spiral and the redshift $z_{\\rm last}$  of its last major merger. As expected,  systems where the last major merger occurred at earlier times, have had more time to grow a disk and have a lower $B/T$. The dispersion in the present-day $B/T$ at a given $z_{\\rm last}$ is  due to the different times spent by a galaxy in terms of being a satellite versus a central galaxy in a DM halo, since the cooling of gas and the growth of a  disk is stopped when a galaxy becomes a satellite. Thus, galaxies that became a satellite galaxy shortly after their last major merger  stayed at high $B/T$. Conversely,  galaxies that continued to be a central galaxy for a long time after their last major merger will have low $B/T$.   The present-day  $B/T$ of a high mass ($M_\\star \\geq  1.0\\times 10^{10} M_\\odot$) spiral depends on its major merger history. In particular, we note from Fig.~\\ref{mergehis} that a  high mass ($M_\\star \\geq  1.0\\times 10^{10} M_\\odot$) galaxy, which has undergone a past major merger since $z \\le 2$ will end up hosting a present-day  $B/T > 0.2$. In effect, {\\it a   high mass spiral can have  a  present-day $B/T \\le0.2$ only if its last major merger occurred at $z > 2$  (lookback times $>10$ Gyr).}   The predicted distribution of present-day $B/T$  depends on the galaxy merger history in the models and it is relevant to ask how well the latter is constrained  observationally. Over the redshift range $z\\sim0.24$ to 0.80 (lookback times of  3 to 7 Gyr), recent studies by Jogee \\etal (2008, 2009) find that   among  high mass ($M_\\star \\geq  2.5\\times 10^{10} M_\\odot$) galaxies, $\\sim$~10\\% of galaxies are undergoing mergers of  mass ratio $>1/10$, and $\\sim$~3\\% are undergoing major mergers of  mass ratio $>1/4$. These findings agree within a factor of  less than $\\sim2$  with the merger rates  from the models of Khochfar \\& Burkert (2001) over $z\\sim$~0.24 to 0.80. At higher redshifts, the empirical merger rate/fraction is uncertain due to relatively modest volumes and bandpass shifting effects, but there is  a general trend towards  higher merger fractions at higher redshifts (e.g., Conselice \\etal 2003). The models used here (Khochfar \\& Burkert 2001) agree with this trend and  predict  that $\\sim$~13.5\\% and $\\sim$~20\\%  of high mass ($M_\\star \\geq  1.0\\times 10^{10} M_\\odot$) spirals have undergone major mergers since $z\\le 2$ and  $z\\le 4$, respectively (see  Table~\\ref{tmodel1}).  The contribution of galaxies with different merger histories to the present-day $B/T$ distribution are  shown in  Table~\\ref{tmodel1}. The top and middle parts of the table describe systems with and without major mergers since $z\\le 2$ and  $z\\le 4$, respectively. In the model, $\\sim$~13.5\\% of present-day high mass ($M_\\star \\geq  1.0\\times 10^{10} M_\\odot$) spirals, experienced a major merger since $z\\le 2$, causing most of them  ($\\sim$~11.2\\%) to have a present-day  high $B/T > 0.4$ and a negligible fraction ($\\sim$~0.1\\%) to have a low present-day $B/T \\leq 0.2$. In contrast, the remaining $\\sim$~86.5\\% spirals experienced no major merger at $z\\le$~2, and most (67.2\\%) of them  have a present-day low $B/T \\leq 0.2$. If the comparisons are extended to systems without a major merger since $z\\le 4$, the numbers are very similar (see middle part of Table~\\ref{tmodel1}).  Table~\\ref{tmodel1}  shows that there is good agreement between the model and data for the fraction of high mass spirals with present-day low $B/T \\leq 0.2$ ($\\sim$~67\\% in the model versus $\\sim$~66\\% in the data). The model contribution to low $B/T \\leq 0.2$ comes almost entirely from galaxies, which have not had a major merger since  $z\\le 2$ (see column 4 in Table~\\ref{tmodel1}). In fact, most of these galaxies have  not even had a major merger since  $z\\le4$, as illustrated by the bottom part of Table~\\ref{tmodel1}. {\\it In the model, the fraction ($\\sim1.6\\%$; column 3 of Table~\\ref{tmodel1}) of high mass spirals, which have undergone a  major merger since $z\\le 4$ and host a bulge with a present-day $B/T \\le 0.2$, is {\\it a factor of over thirty smaller}  than the observed fraction  ($\\sim66\\%$) of high mass spirals  with  $B/T \\le 0.2$. } Thus, bulges built  via major mergers  since $z\\le 4$ seriously fail to account for the bulges present in $\\sim66\\%$ of  high  mass spirals. These results are also illustrated in Fig.~\\ref{pbulgeless1}, which shows the comparison between data and models for the cumulative fraction of high mass spirals as a function of  present-day $B/T$.   It is also interesting to note from Table~\\ref{tmodel1} that although the models reproduce well the frequency of bulges with present-day low $B/T \\leq 0.2$, they tend to over-produce the frequency of present-day  high $B/T > 0.4$ systems by nearly a factor of  two ($\\sim$~14\\% in the model versus $\\sim$~8\\% in the data; see columns 6 and 2 in middle part of Table 11). Most of this overproduction stems from major mergers at $z\\le4$ ($\\sim$~13\\%; column 4 in middle part  of  Table~\\ref{tmodel1}). This suggests that major mergers, as currently modeled here, are {\\it building  bulges too efficiently.}   One possible solution to this problem might relate to the suggestion  by Hopkins \\etal (2009) that the efficiency of bulge-building during the major merger of two spirals depends not only  on the mass ratio $M_{\\rm 1}/M_{\\rm 2}$ of the merger, but also depends on the cold gas mass fraction $f_{\\rm gas}$ in the disk.  In their semi-analytic  models,  the entire stellar mass of the satellite violently relaxes, but the fraction of stellar mass in the primary disk that violently relaxes and adds to the bulge is $M_1/M_2$. This differs with our models where the entire primary stellar disk is always destroyed in a major merger. Furthermore, the fraction $F$ of the total gas mass, which loses angular momentum, falls to the nucleus, and is transformed into stellar mass in a nuclear starburst, is  $\\sim(1-f_{\\rm gas})\\times(M_1/M_2)$. In particular,  $F$ is lower for more gas-rich disks, causing a suppression of the burst efficiency in gas-rich systems, and a reduction in the stellar mass that ends up in the bulge built during the  major merger. The predictions for the distribution of $B/T$ from the Hopkins \\etal (2009) models are shown in the bottom of Table~\\ref{tmodel1}, with major mergers considered as those with baryonic galaxy-galaxy mass ratio $>$~/1/3.  Due to the reduced stellar mass that ends up in the bulge after a major merger, the  Hopkins \\etal (2009) models tend to yield lower  $B/T$ after such a merger. Thus, the models produce overall fewer high $B/T > 0.4$ systems, and more intermediate  0.2~$< B/T <$~0.4 and low $B/T \\leq 0.2$ systems (see column 6 in Table~\\ref{tmodel1}). Nonetheless, the  predictions are not very different from the models by  Khochfar \\& Burkert (2005) and Khochfar \\& Silk (2006), which we use in this paper. In particular, the conclusion that the large frequency of high mass ($M_\\star \\geq  1.0\\times 10^{10} M_\\odot$) spirals  with low $B/T \\leq 0.2$ bulges can only be accounted for by spirals without a major merger since  $z\\le 2$ also holds with the Hopkins \\etal (2009) models.  Thus, we  conclude that {\\it  the observed large frequency ($\\sim$ 66\\%) of high mass ($M_\\star \\geq  1.0\\times 10^{10} M_\\odot$) spirals with low present-day $B/T \\leq 0.2$   can be accounted for in our and other hierarchical models by high mass spirals, which have not undergone a major merger since $z\\le 2$, and most of which have not even experienced a major merger over the last 12 Gyr since $z\\le 4$. } In the models, most of these present-day low $B/T \\leq 0.2$ bulges are built by  minor mergers  since  $z\\le4$. As noted earlier, our models explore bulge-building via minor and major mergers, but do not explicitly incorporate secular processes (see  $\\S$~\\ref{bulgeFormation}). In practice, secular processes may contribute to the building of  present-day low $B/T \\leq 0.2$ bulges, and are particularly relevant at $z\\le2$, where major mergers cannot build such bulges.  For completeness, we further explore how sensitive are our results to other  assumptions made in the data and models of  Khochfar \\& Burkert (2005) and Khochfar \\& Silk (2006):   \\begin{itemize} \\vspace{-3mm} \\item How sensitive are the results to the mass ratio used to separate major and minor mergers?  Fig.~\\ref{pbulgeless2} is similar to  Fig.~\\ref{pbulgeless1} except that the model now defines major mergers as those with  mass ratio $M_1/M_2 \\ge 1/6$.  In this case, about 30\\% of the model spirals undergo major mergers since $z\\le4$ rather than $\\sim20\\%$. The overall model $F$ (black dashed line) now under-predicts the data $F$ by  about 10\\%  for  $B/T >$~0.2. However, the main conclusion that bulges built by major mergers   since $z\\le4$ cannot account for most of the low $B/T\\le0.2$ bulges, present in a large percentage ($\\sim66\\%$) of spirals still holds.  \\vspace{-3mm} \\item How sensitive are the results to the $B/T$  cut used to define spirals? Fig.~\\ref{pbulgeless3} is similar to  Fig.~\\ref{pbulgeless1} except that here spirals are considered to be systems with a $B/T \\le0.55$ rather than 0.75 in the models, and a corresponding cut is applied to the data points. The results are similar to Fig.~\\ref{pbulgeless1}  \\vspace{-3mm} \\item How sensitive are the results to our assumed constant $H$-band mass-to-light ($M/L$)  for the bulge, disk, and bar? Fig.~\\ref{pbulgeless4} is similar to  Fig.~\\ref{pbulgeless1} except that the  $B/T$ of all the observed galaxies has been  multiplied by a factor of two, in order to test what would happen in the case where the $M/L$ ratio of the bulge in $H$-band is twice as high as that of the disk and bar.  This could happen in an extreme example where the dominant bulge stellar population was much older (e.g. 12 Gyr) than the age of the dominant disk stellar population (e.g., 3 Gyr). In such a case, the fraction of high mass spirals with $B/T\\le$~0.2 would change from $\\sim66\\%$ in Fig.~\\ref{pbulgeless1} to  $\\sim55\\%$, and deviate from the model overall model $F$ (black dashed line) by $\\sim20\\%$. However, the main conclusion that bulges built by major mergers since $z\\le4$ cannot account for most of the low $B/T\\le0.2$ bulges, present in a large percentage ($\\sim55\\%$) of spirals still holds. \\end{itemize}  Finally, it is important to note that so far we have compared the data and model only in terms of  bulge $B/T$, but not in terms of bulge index $n$ or in terms of bar fraction.  In effect, we have only shown that the models reproduce a subset of the results outlined in  points (1)-(3) of $\\S$~\\ref{bulgeFormation}.  Since the semi-analytic models do not produce a distribution of bulge index $n$, we resort to  presenting only an indirect comparison in Table~\\ref{tmodel2}. We assume that bulges, which form in major mergers have a bulge $n>2.5$. This assumption is based the evidence presented in $\\S$~\\ref{bulgeFormation}. Thus, in  Test 1 of Table~\\ref{tmodel2}, we compare the fraction ($\\sim$~66\\%) of galaxies in the semi-analytic models having $B/T \\le0.2$ and  no  major merger since $z\\le4$, to the observed fraction ($\\sim$~65\\%) of  galaxies  with $B/T \\le 0.2$ and bulge $n\\le 2.5$.  There is close agreement between the two values. In  Test 2 of Table~\\ref{tmodel2}, the fraction ($\\sim$~12.7\\%) of model galaxies with  $B/T>0.4$ and no  major merger since $z\\le4$ is a factor of $\\sim$~3 higher than the fraction ($\\sim$~3.5\\%) of high mass spirals with $B/T>0.4$  and  bulge $n>2.5$. Similarly, the fraction ($\\sim$~1.6\\%) of model galaxies with $B/T \\le 0.2$  and  a past major merger since $z\\le4$  is also a factor of  $\\sim$~2 higher than the observed fraction ($\\sim$~0.9\\%) of high  mass spirals with $B/T\\leq0.2$ and bulge $n>2.5$ (Test 3, Table~\\ref{tmodel2}). Thus, in terms of bulge $B/T$ and $n$, there is  good agreement between  data and model for Test 1 (involving model galaxies with no major mergers since $z\\le4$). This suggests that the vast majority of bulges with  $B/T \\le 0.2$ and $n \\le 2.5$ likely formed in galaxies having had no major merger since $z\\le4$. However, the agreement is less good for Tests 2 and 3 (involving model galaxies with major mergers  since  $z\\le4$) and this suggests that the models may be building bulges a little bit too efficiently during a major merger, in agreement with the conclusion reached earlier.      What about the role of  bars in the formation of these bulges of low $B/T$ and low $n$? A detailed direct comparison with the semi-analytic models is not possible as the role of bars is not yet modeled, but related comparisons are possible. First, it is important to note that bar-driven gas inflow into the inner kpc and the subsequent building of disky stellar components or `pseudobulges' (see $\\S$~\\ref{sintro}) can happen in both isolated galaxies and in minor mergers ($\\S$~\\ref{bulgeFormation}), since bars can be spontaneously induced in an isolated disk or tidally induced during an interaction or minor merger. The triggering of a bar is favored in a prograde interaction or minor merger. Thus, bulge-building via induced bars is more likely to happen in prograde rather than retrograde minor mergers. Statistically about half of minor mergers might be prograde or prograde-like, and half retrograde. Thus, one would expect bars to be induced in only half of the minor mergers. If this assumption is correct and if most  of the mass in bulges with  present-day $B/T \\le 0.2$ is formed in minor mergers, then one would expect only about half of these bulges to host bars. This  is close to what is observed, as  shown by  Table~\\ref{barFraction}. We see that   $\\sim63\\%$ of high mass spirals with low $n\\leq2$ bulges and $\\sim68\\%$ of spirals with low $B/T \\le 0.2$ bulges host bars. This suggests that  in high mass spirals, spontaneous and/or tidally induced bars \\emph{may} play a part in forming up to two-thirds of $B/T \\le 0.2$ or $n\\leq2$ bulges. The remaining one-third of such bulges may have been formed either by   mechanisms that  do  not involve bars (e.g., retrograde minor mergers) or by bars that are not long-lived."
    },
    "0807/0807.2860_arXiv.txt": {
        "abstract": "From Doppler velocity maps of active regions constructed from spectra obtained by the Extreme-ultraviolet Imaging Spectrometer (EIS) on the {\\it Hinode} spacecraft we observe large areas of outflow (20-50 km s$^{-1}$) that can persist for at least a day.  These outflows occur in areas of active regions that are faint in coronal spectral lines formed at typical quiet Sun and active region temperatures.  The outflows are positively correlated with non-thermal velocities in coronal plasmas.  The bulk mass motions and non-thermal velocities are derived from spectral line centroids and line widths, mostly from a strong line of \\ion{Fe}{12} at 195.12 \\AA.  The electron temperature of the outflow regions estimated from an \\ion{Fe}{13} to \\ion{Fe}{12} line intensity ratio is about $1.2-1.4\\times10^6$ K.  The electron density of the outflow regions derived from a density sensitive intensity ratio of \\ion{Fe}{12} lines is rather low for an active region.  Most regions average around $7\\times10^8$ cm$^{-3}$, but there are variations on pixel spatial scales of about a factor of 4. We discuss results in detail for two active regions observed by EIS. Images of active regions in line intensity, line width, and line centroid are obtained by rastering the regions.  We also discuss data from the active regions obtained from other orbiting spacecraft that support the conclusions obtained from analysis of the EIS spectra. The locations of the flows in the active regions with respect to the longitudinal photospheric magnetic fields suggest that these regions might be tracers of long loops and/or open magnetic fields that extend into the heliosphere, and thus the flows could possibly contribute significantly to the solar wind. ",
        "introduction": "The launch of the Extreme-ultraviolet Imaging Spectrometer (EIS) on the {\\it Hinode} spacecraft on September 23, 2006 has given us the opportunity to investigate the dynamics of coronal plasma, i.e, bulk plasma flows and non-thermal motions, in more detail and with higher spatial resolution than previously possible. In addition, the dynamics can be associated with plasma diagnostic measurements of electron temperature and density.  EIS can therefore help address fundamental unsolved problems in active region physics, such as heating in active region loops (e.g., Klimchuk 2006; Warren \\& Winebarger 2007), and the associations of active regions with the solar wind and heliospheric magnetic flux (e.g., Liewer, Neugebauer, \\& Zurbuchen 2004; Schrijver \\& DeRosa 2003).  Plasma motions might be particularly useful as tracers of either active region heating or sources of solar wind and heliospheric flux. For example, motions that are highly transient, spatially well-localized and comprised of both upflows and downflows, might signify active region loop heating and wave activity.  In contrast, motions which are uni-directional (either upflows or downflows), lasting over several days, and that cover an extensive spatial region might be expected to be associated with primarily long loops or open field regions, and are possible indicators of sufficiently open magnetic flux to extend into the heliosphere and thus the flows could contribute to the solar wind.  The association of flows with closed or open field can be facilitated by the creation of flow maps of active regions contructed by rastering the EIS spectrometer slit over an active region, and then correlating the maps with the underlying photospheric magnetic field.  Obtaining rasters of active regions and the creation of Doppler velocity flow maps was an early science goal of the EIS data analysis program.  Doschek et al. (2007) (and more recently Harra et al. 2008) showed that in active regions there are extensive areas of low intensity that are characterized by net bulk plasma outflows and enhanced non-thermal motions.  The appearance of these outflow regions in areas of very low intensity is striking when compared with maps of spectral line intensity, also constructed from the EIS rasters.  In this paper we expand the results of Doschek et al. (2007) and discuss in detail EIS rasters of Doppler shifts and non-thermal motions for two active regions.  We estimate electron temperatures and densities of the outflowing regions, discuss data from other {\\it Hinode} instruments and other satellites, and speculate on the types of magnetic structures that contain the flows and turbulent motions.  We also estimate the mass flux of the flows into the corona.  We will show that the locations of the flows in the active regions and the underlying morphology of the photospheric longitudinal magnetic flux suggest a link with heliospheric open fields that originate in active regions.  The flows seem more associated with the solar wind and heliospheric magnetic fields than with heating of active region loops.  Experimentally, flows are determined from measurements of a spectral line centroid, coupled with information on the rest wavelength of the line.  There are numerous observations of Doppler shifts in the literature, the most recent being obtained from the Solar Ultraviolet Measurements of Emitted Radiation (SUMER) spectrometer on the {\\it Solar and Heliospheric Observatory} (e.g., Teriaca et al. 1999a; Warren, Mariska, \\& Wilhelm 1997).  Mass motions can be related to heating in coronal flux tubes and the acceleration of the solar wind (e.g., Hassler et al. 1999; Tu et al. 2005; McIntosh et al. 2006). The solar wind work has centered mostly on Doppler shift measurements in coronal holes.  There are also numerous coronal flux tube dynamical numerical simulation models in the literature (e.g., Hansteen 1993; Warren \\& Winebarger 2007; and Patsourakos \\& Klimchuk 2006).  These models attempt to explain observed density and temperature measurements in coronal loops and relate the results to coronal heating mechanisms.  The origin of excess optically thin spectral line widths beyond their thermal Doppler widths is unclear.  Since a rocket flight by Boland et al. (1975) it has been known that spectral lines from ions present in the solar transition region and corona have wider profiles than expected from pure thermal Doppler broadening alone, based on the ionization equilibrium temperatures at which they are formed.  The profiles of the majority of the lines are Gaussian or close to Gaussian, and the excess broadening does not appear to be a strong function of position on the solar disk.  The bulk of our current knowledge of these motions in non-flaring plasma has been obtained from the S082-B slit spectrograph on {\\it Skylab} (e.g., Mariska 1992), the High Resolution Telescope Spectrograph (HRTS) rocket stigmatic spectrograph (e.g., Bartoe 1982), and SUMER (e.g., Chae, Sch\\\"uhle, \\& Lemaire 1998a).  The excess spectral line broadening is usually expressed as a non-thermal random mass motion component of unknown origin. It has been cited as an expected signature of waves in the corona (e.g., Mariska, Feldman, \\& Doschek 1978) and magnetic reconnection (e.g., Parker 1988).  It could also be due to different multiple plasma flow speeds in structures smaller than the spatial resolution of the spectrometer because the convolved Gaussian-like line profile would be wider than the widths of the individual Gaussians composing the overall profile.  Most recently, McIntosh et al. (2008) argue from {\\it Hinode} high spatial resolution Solar Optical Telescope (SOT) observations that Alfv\\,{e}nic waves observed in small-scale spicular-type structures are responsible for the excess broadening.  Although there are numerous published observations of Doppler shifts and non-thermal motions, there is little information, particularly for the corona, on the detailed relationship of the motions to actual structures, e.g., coronal loops.  There is far less information on the relationship of these dynamical parameters to other physical parameters such as electron temperature and density.  The apparent close relationship between large line widths and bulk plasma flows seen in the EIS spectra is one of the most interesting aspects of the new observations.  In Section 2 we present pertinent details of the EIS on {\\it Hinode}. In Section 3 we discuss the observations and data reduction.  Results are given in Section 4 and the implications of the results are discussed in section 5. ",
        "conclusions": "From the above analysis of active region dynamics with EIS we may conclude:  1. Extensive areas of active regions where the \\ion{Fe}{12} emission line intensity is weak exhibit higher Doppler shift outflows and higher non-thermal velocities than found in other areas of the active regions where the line intensity is much stronger.  This does not rule out large flows and line widths in high intensity areas of active regions; we are simply implying that there will be regions of low intensity in active regions that have significant Doppler shift and line width signals.  The outflows range from a few km s$^{-1}$ to as much as 50 km s$^{-1}$.  The non-thermal motions range from about 20 to 90 km s$^{-1}$.  2. The temperatures of the outflow regions obtained from the intensity ratio of an \\ion{Fe}{13} line to an \\ion{Fe}{12} line are about $1.4\\times10^6$ K, and electron densities vary from about $5\\times10^8$ to about $10^{10}$ cm$^{-3}$, depending on the particular region.  3. The outflow regions are concentrated primarily over or near magnetic regions of a single polarity, and can last for at least periods of a day and a half although there are variations within the flow region.  4. There is some evidence for variations of temperature within flow structures and multiplicity of flows, i.e., the excess widths of the lines need not be due to turbulence even in cases where the line profiles appear well-fit by single Gaussian profiles.  It is possible that a wide profile is composed of the line profiles of many outflowing regions, each slightly Doppler shifted with respect to each other due to slightly different Doppler shifts.  This scenario might produce a summed wide profile that could be fit with a single Gaussian.  A multi-component scenario definitely holds for the 11 December event.  Some of the outflow profiles show an obvious secondary outflowing component such that the overall profile departs significantly from a Gaussian.  5. The \\ion{Fe}{12} intensity in the outflow regions is much less than in the bright active region loops.  However, it is brighter than in a disk coronal hole that was observed by EIS on September 28, 2007 near 01:45 UT.  In this coronal hole the \\ion{Fe}{12} line is significantly fainter than in the active region outflow regions.  The faintness of the structures could be due to intermittency in the outflows.  That is, if the outflows are events with time scales significantly less that the exposure times, they will appear faint simply because they are bright for only part of the exposure time. For example, two hypothetical regions, equally bright per second, would appear different by a factor of, say, 5 in brightness if one region only lasted 5 seconds while the other lasted for a longer exposure time of 25 seconds.  Intermittent events of this nature would have to persist over time intervals of at least a day to explain the persistence of the outflows.  From inspection of context data from a number of spacecraft it appears that the outflow regions are associated with long magnetic flux tubes, and/or open field lines that extend into the heliosphere.  It is interesting to calculate the mass flux in the outflowing regions. For the December 11 event the outflow mass flux for all outflows greater than about 12 km s$^{-1}$ in the boxed region of Figure 4 is about $1\\times10^{-5}$ gm cm$^{-2}$ s$^{-1}$ and is about $1.6\\times10^{-6}$ g cm$^{-2}$ s$^{-1}$ for speeds exceeding 31 km s$^{-1}$.  For the August 23 event the mass outflow flux in the boxed region of Figure 1 is about $5.7\\times10^{-7}$ g cm$^{-2}$ s$^{-1}$ for speeds exceeding 12 km s$^{-1}$.  These quantities are obtained by multiplying the flow speed by the mass density for each pixel in the boxed regions with a outflow speed greater than the specified speed, and then summing the result over all pixels.  More investigation is needed to determine whether or not this mass flux, or a fraction of it, contributes to the solar wind. Alternatively, some of the mass flux might be confined within long closed flux tubes.  Liewer et al. (2004), Schrijver \\& DeRosa (2003), Luhmann et al. (2002), Neugebauer et al. (2002), and Wang \\& Sheeley (2002) have found from data and modeling that active regions can contribute substantially to the heliospheric magnetic field and solar wind.  Liewer et al. (2004) find that the open field lines are from the edges of active regions (see Figures 11 and 12 in their paper), in locations similar to where we find the flows discussed in this paper. Schrijver \\& DeRosa (2003) assert that connections with the heliospheric field to active regions are common, long-lived, and that during solar maximum the contributions of active regions to the interplanetary magnetic field can be as much as 50\\%.  As seen in Liewer et al. (2004), Figure 14 in the Schrijver \\& DeRosa paper shows open field lines connected to the edge of a plage region.  Also, a comparison of their simulation results with the {\\it TRACE} image in their figure shows that the open field regions emanate from a region that is largely dark in the {\\it TRACE} image, qualitatively similar to the locations of the EIS flows in the active regions we have discussed. It is therefore quite tempting to associate at least some of the EIS flows as sources of the solar wind confined to open field lines that extend into the heliosphere.  Some of the Doppler results of this paper are similar to those found by Harra et al. (2008), who also analyzed flows in an active region discussed previously by Sakao et al. (2007), and these authors suggest that the flows might contribute to the solar wind.  These authors found the same types of flows as reported earlier by Doschek et al (2007), but Doschek et al. (2007) did not make a connection with the solar wind.  In addition to the two active regions analyzed in detail in this paper, we have qualitatively examined five other active regions that exhibit similar behavior to the two regions we discussed in detail. However, more studies with EIS will be needed along with simulations of the solar magnetic field in order to test quantitatively the solar wind connection.  We are continuing to obtain high resolution rasters of active regions and are refining our knowledge of the zero velocity centroid wavelengths of lines in the EIS spectrum.  In particular, deeper exposures that reveal the flows in spectral lines formed over a broad temperature range should be useful in connecting the flows back to their loop footpoint origins.  More work also needs to be done on the line profiles in some of the flow regions, which exhibit obvious indications of multiple loop flows.  {\\it Hinode} is a Japanese mission developed and launched by ISAS/JAXA, collaborating with NAOJ as domestic partner, and NASA (USA) and STFC (UK) as international partners.  Scientific operation of the {\\it Hinode} mission is conducted by the {\\it Hinode} science team organized at ISAS/JAXA.  This team mainly consists of scientists from institutes in the partner countries.  Support for the post-launch operation is provided by JAXA and NAOJ, STFC, NASA, ESA (European Space Agency), and NSC (Norway).  We are grateful to the {\\it Hinode} team for all their efforts in the design, build, and operation of the mission.  The authors acknowledge support from the NASA {\\it Hinode} program and from ONR/NRL 6.1 basic research funds."
    },
    "0807/0807.3951.txt": {
        "abstract": "% % Text of abstract % %  We present the observational results of a survey designed to target and detect asteroids whose colors are similar to those of Vesta family members and thus may be considered as candidates for having a basaltic composition. Fifty basaltic candidates were selected with orbital elements that lie outside of the Vesta dynamical family. Optical and near-infrared spectra were used to assign a taxonomic type to 11 of the 50 candidates. Ten of these were spectroscopically confirmed as V-type asteroids, suggesting that most of the candidates are basaltic and can be used to constrain the distribution of basaltic material in the Main Belt.  Using our catalog of V-type candidates and the success rate of the survey, we calculate unbiased size-frequency and semi-major axis distributions of V-type asteroids. These distributions, in addition to an estimate for the total mass of basaltic material, suggest that Vesta was the predominant contributor to the basaltic asteroid inventory of the Main Belt, however scattered planetesimals from the inner Solar System ($a<2.0$ AU) and other partially/fully differentiated bodies likely contributed to this inventory. In particular, we infer the presence of basaltic fragments in the vicinity of asteroid 15 Eunomia, which may be derived from a differentiated parent body in the middle Main Belt ($2.5<a<2.8$). We find no asteroidal evidence for a large number of previously undiscovered basaltic asteroids, which agrees with previous theories suggesting that basaltic fragments from the $\\sim100$ differentiated parent bodies represented in meteorite collections have been \"battered to bits\" [Burbine, T.H., Meibom, A., Binzel, R.P., 1996. Mantle material in the Main Belt: Battered to bits? Met. \\& Planet. Sci. 31, 607]. ",
        "introduction": "}  The surface properties of asteroids as determined by spectroscopic observations are clues to the environmental conditions present in the early Solar System. Asteroids with basaltic material on their surface are interpreted as tracers of bodies whose interiors reached the melting temperature of silicate rock and subsequently differentiated \\citep{Gaffey93b,Gaffey02}. Until recently Vesta and the dynamically associated Vestoids were the only asteroids with observed surface basalt \\citep{McCord70,Binzel93}. A small number of basaltic asteroids have now been found that may derive from other differentiated parent bodies \\citep{Lazzaro00,Hammergren06,Binzel07,Roig08}. Other asteroids and asteroid families may also represent differentiated parent bodies [e.g. Flora \\citep{Gaffey84}, Eunomia \\citep{Reed97,Nathues05}, Eos \\citep{Mothe05}, Hebe \\citep{Gaffey98}, and Thetis, Agnia and Merxia \\citep{Sunshine04}], however it is not clear why so little basaltic material from these objects remains while the Vestoids are one of the most prominent families in the Main Belt. Two possibilities are that the Vestoids were produced relatively recently and that non-Vestoid basaltic fragments were heavily comminuted to sizes below our current observational limits \\citep{Davis89,Burbine96,Davis02,Bottke05b} or that some space weathering process has erased the diagnostic spectral signatures of surface basalts \\citep[e.g.][]{Wetherill88}.  In addition to asteroidal evidence for the widespread occurrence of differentiation, oxygen isotope analyses of the iron meteorites suggests that as many as $\\sim60$ differentiated parent bodies once existed \\citep{Chabot06,Burbine02}, some of which were more than 500 km in size \\citep{Yang07}. Furthermore, the HED basaltic achondrites are not isotopically homogenous and may represent more than one differentiated parent body \\citep{Yamaguchi01}.  These populations suggest that basaltic material was common throughout the inner Solar System, however it is unclear whether there currently exists any basaltic vestige of the numerous differentiated parent bodies that once existed. To provide insight on these issues we have determined an unbiased distribution of V-type asteroids in the Main Belt. Other studies have calculated unbiased distributions of asteroid types \\citep[e.g.][]{Gradie82,Gradie89,Bus02}, but they have been restricted to the largest and brightest asteroids and did not consider the presence of V-type asteroids beyond 2.5 AU. Our survey provides constraints on the number of basaltic asteroids up to an absolute magnitude of 17.2, which is an increase of roughly 3 orders of absolute magnitude compared to these previous studies and therefore probes for basaltic fragments that may represent the upper size limit to a heavily comminuted population.  Throughout this work we use specific terms related to the taxonomic classification, surface mineralogy and dynamical properties of asteroids. V-type is a taxonomic distinction based on spectral criteria defined by \\citet{Bus02}. Any object with photometric colors consistent with a V-type designation is a V-type candidate. Basaltic asteroids are differentiated bodies or fragments of differentiated bodies with basalt on at least part of their surface. Basaltic asteroids are characterized based on interpreted mineralogy from spectroscopic data. An object with photometric colors indicative of a basaltic composition is a basaltic candidate. For this work we consider a V-type taxonomy to be directly attributable to a basaltic composition, therefore the terms V-type and basaltic are interchangeable. Vestoids are a dynamically-defined sub-population of V-types thought to have originated from the Vesta parent body at least 1 billion years ago \\citep{Marzari96}. Both spectroscopic data and dynamical studies are required to classify an object as a Vestoid. Vestoids are typically interpreted as having a basaltic surface composition, however some asteroids with basaltic material on their surface are not Vestoids, e.g. Magnya \\citep{Lazzaro00}. An asteroid that is dynamically associated with Vesta but has not been spectroscopically observed is called a dynamical Vestoid.  In the early Solar System the likely heat source for metal-silicate differentiation and the production of basalt was the decay of short-lived radioactive isotopes such as $^{26}$Al and $^{60}$Fe \\citep{Grimm93,Goswami00,Tachibana03}. The half-lives of these isotopes are 0.73 and 1.5 Myr respectively, thus under this scenario a differentiated planetesimal or embryo must have accreted within a few half-lives ($< 3$ Myr after the onset of decay) to reach silicate melting temperatures ($>1000$ K). An alternative heat source for material in the early Solar System may have been electromagnetic inductive heating by an intense solar wind \\citep{Herbert91,Gaffey02}, however this mechanism is poorly constrained \\citep{McSween02} and may be incompatible with observations of T Tauri winds \\citep{Wood91}.  The increase of accretion time-scale with heliocentric distance implies that planetesimals from the terrestrial planet region of the Solar System incorporated a high abundance of live radionuclides and may have differentiated, whereas planetesimals that formed in the vicinity of the gas giants accreted too slowly to have melted. It is uncertain where the boundary between these two regions lay. \\citet{Bottke06} suggest that differentiation was primarily confined to the inner Solar System ($a<2.0$ AU) so that the iron meteorites that reach the Earth today trace differentiated planetesimals that were scattered into the Main Belt from smaller heliocentric distances. It is unclear whether the presence of non-Vestoid basaltic asteroids and partially melted asteroids in the outer belt [e.g.~Eos and Eunomia, \\citet{Mothe05,Reed97,Nathues05}] provide evidence for the in-situ differentiation of material in the Main Belt or can be explained by the \\citet{Bottke06} scenario.  Basaltic melts derived from a chondritic precursor will contain pyroxenes, olivine and plagioclase. Pyroxenes are spectroscopically characterized by deep and broad 1- and 2-$\\mu m$ absorption features. Figure \\ref{fig.basalt} shows optical and near-infrared spectra of basaltic asteroids 38070 (1999 GG2) and 33881 (2000 JK66), which are two of our basaltic candidates (\\S\\ref{sec.selection}). The visible spectrum was obtained with ESI \\citep{Sheinis02} on Keck II and the near-infrared with SpeX \\citep{Rayner03} on IRTF. Also plotted is a RELAB (Reflectance Experiment Laboratory) spectrum of a basaltic eucrite meteorite. These objects clearly show the absorption features at 1 and 2 $\\mu m$. The depth, center and shape of these features are dictated by temperature, grain size and calcium content \\citep{Hardersen04} and are attributable to Fe$^{2+}$ ionic transitions within the pyroxene crystal structure \\citep{Sunshine04}. Olivine and plagioclase also show Fe$^{2+}$ absorption features around 1 $\\mu m$.  Although the spectral features of basaltic material dominate in the near-infrared, optical data ($0.4-0.95~\\mu m$) are easier to obtain for faint asteroids and can be diagnostic of such surface material. Thus, we primarily relied on optical taxonomic classification as a way of selecting and characterizing our target asteroids.  Recent observational and theoretical work \\citep{Lazzaro00,Hammergren06,Bottke06,Binzel07, Nesvorny08,Roig08} along with the data presented here suggests that there exists non-Vestoid basaltic asteroids throughout the Main Belt that have yet to be discovered. We have placed new constraints on the overall population of these rare objects by investigating relatively faint asteroids that are not typically included in spectroscopic surveys. One of the goals of all research related to the extent of thermal alteration in the early Solar System is to understand how many differentiated parent bodies are represented by the Main Belt asteroids. This work is a small step towards that understanding.  The remaining sections of this paper are organized as follows: in \\S\\ref{sec.selection} we outline the procedure of selecting target asteroids for our observational campaign. In \\S\\ref{sec.obs} and \\S\\ref{sec.reduct} we discuss our observational procedure and data reduction techniques. In \\S\\ref{sec.dist} we use our observational results to derive a size-frequency distribution of basaltic asteroids in the Main Belt. We use this result to constrain the semi-major axis distribution and the total mass of basaltic asteroids. In \\S\\ref{sec.disc} we discuss and summarize these results. ",
        "conclusions": "}  We have identified 50 candidate basaltic asteroids from the SDSS MOC based on photometric and dynamical constraints. Based on observations of these basaltic candidates we have constrained the amount of basaltic material that is present outside of a region occupied by the dynamical Vestoids (defined by $a$-$e$-$i$ orbital elements, \\S\\ref{sec.selection}). We have determined the unbiased distribution of this basaltic material in absolute magnitude and semi-major axis.  \\subsection{Basaltic Asteroids with $a>2.5$ AU}  The prospect of multiple basaltic asteroids with semi-major axes greater than 2.5 AU is particularly compelling due to the unlikelihood of numerous such objects originating from the surface of Vesta and then crossing the 3:1 mean motion resonance \\citep{Nesvorny08,Roig08}.  Our selection process identified 12 basaltic candidates between the 3:1 and 5:2 mean motion resonances at 2.5 and 2.8 AU respectively. Six of these are clustered around the Eunomia family (Fig. \\ref{fig.orbits}) which is one of the largest families in the Main Belt \\citep{Zappala95} and may have derived from a partially or fully differentiated parent body. This claim is supported by a compositional gradient across the surface of Eunomia \\citep{Reed97,Nathues05} and a spread of surface compositions amongst members of the Eunomia family \\citep{Lazzaro99}, both of which are consistent with a differentiated mineralogy \\citep{Gaffey02}. Presumably the basaltic crust of the Eunomia parent body was removed in the large family forming collision and/or by subsequent collisions. Asteroid 21238 (1995 WV7) is a spectroscopically confirmed V-type that may represent some of this lost basaltic crust \\citep{Hammergren06,Binzel07,Carruba07}. One of the six candidates in the vicinity of Eunomia, asteroid 40521, is a spectroscopically confirmed V-type, however there is some uncertainty about whether this object originated from the Vesta or Eunomia parent bodies \\citep{Carruba07,Roig08}. Detailed numerical simulations will be necessary to investigate whether it is dynamically feasible for the remaining five candidates to have originated from Eunomia.  It is very unlikely that the selection of 6 candidates in the Eunomia family represents a statistical fluke related to the high density of objects in this region. Approximately 11\\% of MOC objects beyond the 3:1 resonance are part of the Eunomia family. Thus, assuming a binomial probability distribution for which there is an 11\\% chance of a single candidate being selected at random from the Eunomia family, there is only a 0.7\\% probability that 6 of the 16 candidates from beyond the 3:1 resonance would belong to the Eunomia family.  We have identified only 4 basaltic candidates with semi-major axes greater than 2.8 AU (Fig. \\ref{fig.orbits}, Table \\ref{tab.candidates}). One of these, asteroid 10537 (1991 RY16), may be classified as a V-type because of its deep 1 $\\mu m$ absorption feature which is typical of this class and indicative of a pyroxene-rich and/or olivine-rich mineralogy. However it also has an unusually deep absorption feature centered at $0.63~\\mu m$ that is uncommon to V-types (Fig. \\ref{fig.RY16}). These visible wavelength data have been confirmed by other observations [Duffard \\& Roig (2008), submitted to Planetary and Space Science]. Follow-up infrared observations provide additional insight on the surface composition of this unusual object \\citep{Moskovitz08}. Another outer belt candidate is unnumbered asteroid 2003 SG55. It is possible that this object originated from a differentiated Eos parent body \\citep{Mothe05}. The presence of multiple non-linear secular resonances that pass through the region \\citep{Milani90,Milani92} may have provided a pathway for the orbital migration of 2003 SG55 even though HCM analysis \\citep{Zappala90} shows that 2003 SG55 is outside of the Eos dynamical family. Future study of these outer belt basaltic candidates will have important implications for understanding the dynamical and thermal histories of these specific families.  \\subsection{Basaltic Asteroids with $a<2.5$ AU}  Defining the dynamical boundaries of the Vestoids is difficult due to effects such as Yarkovsky drift and resonance scattering \\citep{Carruba05,Carruba07}. Nevertheless we have attempted to identify non-Vestoid basaltic candidates inside of 2.5 AU because neglecting to do so could exclude basaltic asteroids of non-Vestian origin. The relatively uniform distribution of basaltic candidates in this region suggests two possibilities: (1) there exists non-Vestoid basaltic asteroids in the inner belt that represent a unique differentiated parent body, or (2) Vestian fragments actually populate the full dynamical extent of the inner Main Belt. The first possibility is testable by performing mineralogical analyses (e.g.~\\citet{Lawrence07,Gaffey93a}) on NIR spectra of basaltic candidates with large values of $\\Delta v$, for example 38070 (1999 GG2) with $\\Delta v=2.1$ km/s and 33881 (2000 JK66) with $\\Delta v=3.3$ km/s. If these inner belt basaltic candidates do represent a unique differentiated parent body then they may be compositionally distinguishable from Vesta and the Vestoids (e.g.~Magnya, \\citet{Hardersen04}). The second possibility can be investigated with numerical simulations. \\citet{Nesvorny08} show that the dynamical diffusion of Vestian material does in fact populate the semi-major axis extent of the inner Main Belt, but does not extend below orbital inclinations of 4$^\\circ$. In the future, the most effective test for inner belt, non-Vestoid basaltic asteroids will be a combination of dynamical studies and mineralogical analyses.   \\subsection{Total Basaltic Asteroid Inventory}  We find that the calculated size-frequency distribution of basaltic asteroids outside of the dynamical Vestoid population is heavily weighted towards small bodies. We have also shown that the mass of basaltic material outside of the dynamical Vestoid population is $M_{sur} < 5.9 \\times 10^{16}$ kg. This value is comparable to the mass of basaltic material inside of the dynamical Vestoid population. Due to uncertainties in the definition of dynamical families we do not attempt to make a statement about the fraction of $M_{sur}$ that is truly non-Vestoid in origin. We have attempted to place upper limits on $M_{sur}$ and $M_V$ and a conservative estimate on $M_{ex}$. Even so we note that $M_{sur} + M_V < M_{ex}$ by over an order of magnitude.  The size-frequency distribution ($N(H)$) and smaller total mass relative to $M_{ex}$, both inclusive and exclusive to the Vestoid dynamical family, are consistent with most (but not all!) of the basaltic asteroids deriving from the Vesta parent body. Vesta has several large craters on its surface \\citep{Thomas97} whose collisional formation would have produced fragments of varying size. Due to ejection velocities, the Yarkovsky effect, resonance interactions and close encounters, these fragments diffused away from the Vesta parent body \\citep{Bottke01,Carruba05,Nesvorny08}. The smallest fragments experienced the fastest dynamical evolution while larger objects would take longer to separate from Vesta due to the decreased significance of the Yarkovsky force for bodies larger than 10 km \\citep{Bottke00}. Over time, small Vestoids migrated to the edge of the family at the $\\nu_6$ and 3:1 resonances. Although many of the fragments that reached these resonances were scattered into planet crossing orbits or into the Sun, some crossed over into stable orbits \\citep{Gladman97,Roig08,Nesvorny08}. Thus, the size-frequency distribution of Vestoids shows a depletion of small bodies ($H>16$) relative to the SDSS MOC distribution. This is apparent in Fig. \\ref{fig.NvH} where  the slope of the Vestoid distribution (solid curve) for $H>16$ is slightly steeper than that of the MOC (dotted curve). This also holds true when compared to just the subset of MOC objects that are confined to the semi-major axis range of the Vestoids ($2.2 < a < 2.5$ AU).  The overabundance of small bodies (sizes $< 2$ km) amongst the basaltic asteroids outside of the dynamical Vestoid population (Fig. \\ref{fig.NH}) may represent the hiding place of these \"missing Vestoids\". If this is true then it suggests that the shallow slope of absolute magnitude distributions at large values of $H$ for HCM-defined collisional families \\citep{Morby03} may be attributed to the preferential loss of small bodies out of the families by dynamical processes such as the Yarkovsky effect and resonance scattering.  The semi-major axis distribution indicates that the full inventory of basaltic material in the Main Belt is not consistent with an origin exclusively coupled to the formation of the large Vestian crater. With an age of 1-3.5 Gyr \\citep{Bottke05a,Marzari96}, the primary Vesta crater was formed after the completion of planet formation. Thus, the asteroidal fragments generated in this impact evolved in a dynamical environment similar to the present day architecture of secular and mean motion resonances. In such an environment it is relatively easy for asteroids to migrate in semi-major axis because of the Yarkovsky effect, however it is more difficult to produce significant changes in inclination and eccentricity. This motivates the search for basaltic asteroids at very different $i$ and $e$ values than Vesta.  Our results suggest that basaltic asteroids with sizes in excess of 2 km do exist at low inclinations and beyond 2.5 AU (Fig. \\ref{fig.Na}). \\citet{Roig08} show that it is not possible to dynamically transport many hundreds of Vestoids with diameters greater than 2 km beyond the 3:1 resonance. This is further confirmed by \\citet{Nesvorny08} whose dynamical simulations do not produce any 3 km size bodies beyond this resonance or any Vestoids with low inclination orbits ($i<4^\\circ$).  We cite three explanations for the presence of basaltic material in regions of phase space that are dynamically inaccessible to the Vestoids. One possibility is that these basaltic asteroids originated from the surface of Vesta, not from its largest crater but sometime before the Late Heavy Bombardment (LHB) more that 3.8 Gyr ago. The higher density of objects in the pre-LHB asteroid belt \\citep{Petit02} would have made for a dynamical environment in which collisional and scattering events were greatly enhanced over present day rates. In addition, resonance sweeping during the LHB could have dispersed these asteroids' inclinations \\citep{Nesvorny08}.  A second possibility is that these objects represent scattered planetesimals from the terrestrial planet region of the Solar System \\citep{Bottke06}. In the inner Solar System solid-body accretion times would have been faster than in the Main Belt. These bodies would have accreted a higher quantity of live radioactive nuclides such as $^{26}$Al and been more likely to differentiate. The amount of basaltic material that was injected into the Main Belt from the inner Solar System is currently unconstrained.  A final explanation for the presence of basaltic asteroids at low inclinations and beyond 2.5 AU is that they derived from other differentiated parent bodies that formed in the Main Belt. Although we find no asteroidal evidence for a large number of undiscovered basaltic asteroids, we cannot rule out minor contributions to the basaltic inventory from non-Vestoid differentiated bodies whose basalt has been heavily depleted due to collisions and dynamical effects \\citep[i.e. a \"battered to bits\" scenario,][]{Burbine96}. It is possible that the observed distribution of basaltic asteroids in the Main Belt is derived from many different source populations including the Vestoids, inner Solar System planetesimals and a heavily depleted population of non-Vestoid differentiated bodies.  \\subsection{Comparison to Other Works}  We compare our calculated semi-major axis distribution to the simulation results of \\citet{Nesvorny08} and \\citet{Bottke06} (Fig. \\ref{fig.Na_comp}). This allows us to investigate the feasibility of two different scenarios for delivering basaltic asteroids to orbits outside of the dynamical Vestoid population. Figure \\ref{fig.Na_comp} shows the normalized distribution of non-Vestoid basaltic asteroids as a function of semi-major axis. These distributions have been normalized by their maximum values (all at 2.2 AU) because limitations inherent to each work precludes the comparison of an absolute number of basaltic asteroids. We confine these distributions to the inner Main Belt ($a<2.5$ AU) because of low number statistics for larger values of $a$ in our analysis and the \\citet{Bottke06} simulations and the lack of any Vestoids across the 3:1 resonance in the \\citet{Nesvorny08} simulations. The distributions from the numerical simulations exclude test particles that populate the dynamical Vestoid region of orbital element space as defined in \\S\\ref{sec.selection}. This produces distributions that are non-Vestoid based upon our specific dynamical criteria.  The \\citet{Nesvorny08} study was designed to understand the dynamical evolution of Vesta's ejected fragments. They modeled the diffusion of Vestoids over an estimated family lifetime of 2 Gyr. If our distribution of basaltic asteroids was exclusively derived from dynamically aberrant Vestoids then we would expect it to match the \\citet{Nesvorny08} result.  The \\citet{Bottke06} study models the injection of planetesimal-size bodies from the inner Solar System into the Main Belt during the epoch of planet formation. The specific source region of these planetesimals determines which are differentiated and contain basalt. Due to uncertainties in the time-scales of accretion and the initial abundance of radioactive isotopes, it is not clear where the boundary between differentiated and undifferentiated planetesimals existed. Thus we plot two distributions from the \\citet{Bottke06} results: one corresponding to all of the planetesimals that are scattered into the main belt from initial semi-major axes, $a_0$, less than 1.5 AU and those with $a_0<2.0$ AU. If our distribution was derived primarily from these scattered interlopers we would expect it to match the \\citet{Bottke06} result.  We calculate the chi-squared ($\\chi^2$) probability that these simulated distributions are derived from our observational result. The \\citet{Nesvorny08} distribution and the \\citet{Bottke06} distribution for a planetesimal source region of $a_0<1.5$ AU produce $\\chi^2$ probabilities equal to 99.6\\% and 98.5\\% respectively. Within the limits of our calculations these distributions fit the data equally well. The worst fit ($\\chi^2$ probability = 73\\%) is the \\citet{Bottke06} distribution for a planetesimal source region of $a_0<2.0$ AU. The difference in probabilities from the \\citet{Bottke06} simulations suggests that the contribution of basaltic asteroids from inner Solar System planetesimals is weighted towards objects originally separated in semi-major axis from the Main Belt rather than objects that formed at semi-major axes around 2.0 AU. This is not unexpected as faster accretion times at smaller heliocentric distances would have produced a greater number of differentiated planetesimals.  This work is the first survey specifically designed to constrain the global distribution of basaltic material in the Main Belt. Future all-sky surveys such as PanSTARRS \\citep{Kaiser02} and LSST \\citep{Claver04} will greatly increase the number of known asteroids and will allow a more complete description of the basaltic asteroids. The results presented here suggest that the total amount of basaltic material in the Main Belt is very much less than expected from the $\\sim100$ differentiated parent bodies traced by the iron meteorites and basaltic achondrites, a result which agrees with the suggestion that basaltic fragments have been comminuted to sizes below observational limits \\citep[e.g.][]{Burbine96} or that some unknown space weathering processes has obscured the spectral signature of basalt on asteroid surfaces \\citep{Wetherill88}. However, the inferred presence of a small amount of basaltic material throughout the Main Belt suggests that some vestige of crustal material from large differentiated bodies does exist and may represent the upper size limit of this population.  %% Using an acknowledgements command is not in the Elsevier template, %% but it can be used. \\ack Thanks to Greg Wirth for his patience and expert advice as we have strived for competency with ESI. We would like to acknowledge the tremendous help received from Greg Aldering, Rolin Thomas, Yannick Copin and the rest of the SNIFS team. Thanks to Bobby Bus for helpful suggestions regarding observing and data reduction. We are grateful to Richard Binzel and Fernando Roig for their insightful reviews. N.M. would like to acknowledge the support of NASA GSRP grant NNX06AI30H. R.J. would like to acknowledge the support of NSF planetary astronomy grant AST04-07134. E.G. acknowledges support from the NASA Astrobiology Institute. Some of the data presented herein were obtained at the W.M. Keck Observatory, which is operated as a scientific partnership among the California Institute of Technology, the University of California and the National Aeronautics and Space Administration. The Observatory was made possible by the generous financial support of the W.M. Keck Foundation. Part of the data utilized in this publication were obtained and made available by the The MIT-UH-IRTF Joint Campaign for NEO Reconnaissance. The IRTF is operated by the University of Hawaii under Cooperative Agreement no. NCC 5-538 with the National Aeronautics and Space Administration, Office of Space Science, Planetary Astronomy Program. The MIT component of this work is supported by the National Science Foundation under Grant No. 0506716. We wish to recognize and acknowledge the very significant cultural role and reverence that the summit of Mauna Kea has always had within the indigenous Hawaiian community.  We are most fortunate to have the opportunity to conduct observations from this mountain.  \\label{lastpage}  % Bibliographic references with the natbib package: % Parenthetical: \\citep{Bai92} produces (Bailyn 1992). % Textual: \\citet{Bai95} produces Bailyn et al. (1995). % An affix and part of a reference: %   \\citep[e.g.][Ch. 2]{Bar76} %   produces (e.g. Barnes et al. 1976, Ch. 2).-"
    },
    "0807/0807.2098_arXiv.txt": {
        "abstract": "% We carried out intensive spectroscopic observations of two WZ Sge-type dwarf novae, GW Lib, and V455 And during their superoutbursts in 2007, at 6 observatories.  The observations covered the whole of both superoutbursts from the very maximum to the fading tail.  We found evidence of the winds having a speed of $\\sim$1000 km s$^{-1}$ which blew in GW Lib during the rising phase.  The evolution of the hydrogen, helium, and carbon lines suggests flaring of the accretion disk and emergence of the temperature inversion layer on the disk. ",
        "introduction": "Dwarf novae are a class of cataclysmic variable stars, which show outbursts by the disk instabilities, and SU UMa-type dwarf novae are a subclass giving rise to two types of outburst: normal outbursts and superoutbursts \\citep[for a review][]{Osa96}.  There is a small group of SU UMa stars, being called WZ Sge-type dwarf novae.  Their peculiar feature is summarized as: the large outburst amplitude over 6 mag, the long recurrence cycle of the superoutburst (several years or more), and no (or few) normal outburst.  In addition, they show curious behavior in many points, such as ``early'' superhumps, a variety of rebrightening light curves after the main superoutburst, long fading tail, and so on \\citep[see][and references therein]{Uem08}.  The driving mechanisms working in WZ Sge stars are still in debate \\citep[see][]{nog07}, while the basic properties of normal SU UMa-type dwarf novae are generally explained by the thermal-tidal disk instability theory. For further investigation, we have prepared for coordinated spectroscopic observations during forthcoming superoutbursts, following a success of intensive spectroscopy during the 2001 superoutburst in WZ Sge itself \\citep{bab02,nog04}.  Under this circumstances, superoutbursts of two WZ Sge-type dwarf novae, GW Lib, and V455 And (=HS2331+3905) were discovered in the very early phase in 2007 by eager contributors to VSNET \\citep[about VSNET, see][]{kat04}. We succeeded in soon starting coordinated spectroscopic campaigns before or around the maximum in both stars, and here review the preliminary results and implications.  The details will be published in our forthcoming papers. ",
        "conclusions": ""
    },
    "0807/0807.0781_arXiv.txt": {
        "abstract": "\\ni The X-ray background (XRB) is produced by a large number of faint sources distributed over a wide range of redshifts. The XRB carries information on the spatial distribution and evolution of these sources. {The goals of the paper are: 1. to determine the redshift distribution of the soft X-ray background photons produced by all types of extragalactic sources, in order to relate fluctuations of the background to the large scale structures, 2. to determine the  redshift distribution of the soft XRB produced by AGN in order to calculate the evolution of the AGN X-ray luminosity density.} {A set of major X-ray surveys is used to determine the redshift distributions of the X-ray sources selected at various flux levels. Simple analytic fits to the data allow us to determine the smooth relationship between the redshift distribution and the source flux. The redshift distribution of the integral XRB flux is obtained  by averaging the fits over the source counts.} {It is shown that the distribution of extragalactic XRB photons in the $0.5-2$\\,keV band is adequately represented by the function: $d\\,n_{\\rm XRB} / d\\log z = 5.24\\:z^{1.52}\\,\\exp(-z/0.63)$. The huge voids postulated to explain the cold spots in the CMB maps create dips in the total XRB flux. However, the expected magnitude of the effect is comparable to the fluctuation amplitude of the XRB generated by the individual sources contributing to the background. The cosmic evolution of the AGN X-ray luminosity density up to redshift of $\\sim\\!5$ is calculated in an elegant and straightforward way. Systematic uncertainties of the present method are assessed and shown to be small. At redshift greater than one the present results could be compared directly with some recent estimates obtained in a standard way and the agreement between both methods is very good.}  \\vspace{2mm} \\ni   Keywords{X-rays: diffuse background  -- intergalactic medium   -- X-rays: galaxies } ",
        "introduction": "}  The X-ray background (XRB) is generated mostly by discrete extragalactic sources (e.g. \\citealt{lehmann01, kim07}, and references therein), predominantly by various types of Active Galactic Nuclei (AGN) and cluster of galaxies. A question of flux, luminosity and redshift distributions of these sources has been discussed in a great number of papers for the last $30$ years. One of the major outcome of these investigations is the conclusion that X-ray sources associated with the AGN are subject to strong cosmic evolution (e.g. \\citealt{miyaji00}, \\citealt{silverman07}, and references therein). As a result of the evolution, the redshift distribution of the XRB flux is wide. Thus, the integral XRB comprises the information on the large scale distribution of the X-ray sources over a wide redshift range.  In the present paper the redshift distribution of the XRB photons is investigated in detail. The analysis is based on an extensive observational data selected from several published X-ray sky surveys. A convenient analytic approximations are applied to model the observed redshift histograms of the extragalactic X-ray sources selected at several flux levels. These distributions are weighted by the source counts and summed up to obtain the redshift distributions of the integral XRB.  Next, the redshift distribution is used to define a relationship between the XRB signal and the large-scale fluctuations of the matter spatial distribution. The investigation has been raised by the recent report on the huge void generating a dip in the surface brightness of the radio background (\\citealt{rudnick07}). This void, responsible for the deficit of the radio surface brightness, allegedly generates also a cold spot in the CMB map via the late-time integrated Sachs-Wolfe effect. Although the careful statistical analysis by \\cite{smith08} has not confirmed the existence of this particular ``cold spot'' in the radio survey, a relationship between the large scale features of the matter distribution and the integrated sky brightness in various energy bands is a problem deserving some interest.  Apart of the question of the XRB fluctuations induced by voids, the XRB redshift distribution is interesting per se, as it allows to assess the evolution of AGN phenomenon. A standard way to estimate a rate and type of this evolution is based on the examination of the X-ray luminosity functions determined at the consecutive redshift bins.  Unfortunately, the X-ray surveys produce flux-limited rather than luminosity-limited samples of sources. In effect, luminosity functions at different redshifts cover different luminosity ranges. This in turn severely impedes estimates of the luminosity function over a wide range of luminosities and redshifts.  The total level of nuclear activity in galaxies within unit volume is given by the integral of the X-ray luminosity function. The question of the AGN cosmic evolution constitutes one of the central problems of observational cosmology, and has been investigated for the last forty years (this issue was for the first time recognized by \\citealt{schmidt68}).  Here a question of the AGN evolution is addressed without the  calculations of the X-ray luminosity function.  The available observational data on X-ray source counts and redshifts are used to evaluate the redshift distribution as a function of source flux. This relationship and the source counts allow to calculate the redshift distribution of the total XRB and the integral luminosity density generated by the AGN as a function of redshift.  The organization of the paper is following. First, I present the formulae used in calculations. Next, in Sec.\\ref{catalogs}, the basic information on the observational material extracted form the various archives is given. Since the comprehensive characteristics of the data and the source catalogs are described in the original papers, only the basic properties of the material are presented here. The numerical fits to the observed distributions are obtained in Sec.~\\ref{fits}. In that section the calculations of the redshift distribution of the XRB photons are described in details. These results are applied in the Sec.~\\ref{supervoids} to quantify the relationship between the voids and the XRB variations. In Sec.~\\ref{agn} the distribution of the XRB flux produced by AGN is used to calculate the evolution of the AGN activity. Finally, potential sources of errors inherent in the present method are discussed in the Sec.~\\ref{discussion}.  The `canonical' standard cosmology is assumed throughout, with $H_0 = 70$\\,km\\,s$^{-1}$Mpc$^{-1}$, $\\Omega_m = 0.3$, and $\\Omega_\\Lambda = 0.70$. ",
        "conclusions": "}  The main objectives of the present investigation, viz. estimates of the redshift distribution of the XRB photons, $b(z)$, and the evolution of the AGN luminosity density, $\\varepsilon(z)$, have been achieved using the smooth, analytic fits to the observed source redshift histograms. The present method is conceptionally simple and computationally straightforward. Unfortunately, it does not provide error estimates. The major sources of uncertainties have been indicated in the previous section. Here a quantitative estimate of the errors is discussed.  The errors of the present measurement of $b(z)$ are generated by the statistical nature of the investigated material and a chain of approximations applied to substitute the observed redshift distributions centered on a selected fluxes by an analytic function $f_S(z)$ continuous in both parameters, $z$ and $S$.  In fact, the visual inspection of the analytic fits displayed with the solid curves in Fig.~\\ref{z_distr} reveals some deviations from the redshift histograms.  To estimate the significance of these differences, the calculations have been performed using the actual histograms shown in Fig.~\\ref{z_distr} with broken solid lines instead of $f_S(z)$. For the each value of flux $S$ in the range  $10^{-17} - 10^{-11}$ \\cgs, the corresponding redshift distribution has been obtained by the linear interpolation between two histograms from the samples centered on the median fluxes nearest to $S$.  The results of this procedure are shown in the lower right panel of Fig.~\\ref{z_distr} with dots.  Generally good agreement between the distribution of points and the solid curve proves that the analytic approximations do not introduce perceptible systematic errors in the present investigation. It appears that the relatively large deviations for three data points (centered at redshifts: $0.087$, $0.65$, and $1.16$) result purely from the large scale structures. This is particularly likely for the first bin ($0.075 < z < 0.1$), where the discrepancy between the fits is produced entirely by the excess of sources in the localized NEP survey.  One should also notice, that the uncertainties of our main results are only weakly affected by the limited statistics of the individual samples and histograms.  This is because the final distributions are obtained by averaging the individual distributions and this procedure effectively reduces statistical fluctuations.  In the upper panel of Fig.~\\ref{evolution} the AGN emissivity calculated by \\cite{hasinger05} is shown.  The points with the error bars are redrawn here from their original paper. \\cite{hasinger05} apply more stringent criteria to select sources and use only well defined samples of type-1 AGN. Optically these objects are identified by the broad Balmer emission lines, while using the X-ray criteria, they have unabsorbed spectra indicating low intrinsic column densities. In the present analysis I have included all objects in which the X-ray emission originates in the active nuclei. Thus, our results cannot be compared directly to those by \\cite{hasinger05}. Nevertheless, despite entirely different method applied in the present paper, the distributions show good agreement at redshifts above $\\sim\\!1$.  Although most of the apparent discrepancies, which reach a factor of $3$ at $z\\approx 0.5$, are probably due to the distinct selection criteria of both investigations, one cannot exclude that some differences are caused by unrecognized systematic effects inherent in one or both methods.  The present method of the luminosity density calculations has also some disadvantages. In our approach the absolute luminosities of the individual objects are not determined. Consequently, only the integral luminosity density is obtained, and the cosmic evolution of any selected AGN luminosity class has to be studied by means of the standard methods.  \\vspace{2mm} \\ni ACKNOWLEDGEMENTS\\\\ This work has been partially supported by the Polish MNiSW grant N~N203~395934."
    },
    "0807/0807.0112_arXiv.txt": {
        "abstract": "We examine the claim of Babak and Grishchuk \\cite{BG} to have solved the problem of localising the energy and momentum of the gravitational field. After summarising Grishchuk's flat-space formulation of gravity, we demonstrate its equivalence to General Relativity at the level of the action. Two important transformations are described (diffeomorphisms applied to all fields, and diffeomorphisms applied to the flat-space metric alone) and we argue that both should be considered gauge transformations: they alter the mathematical representation of a physical system, but not the system itself. By examining the transformation properties of the Babak-Grishchuk gravitational energy-momentum tensor under these gauge transformations (infinitesimal and finite) we conclude that this object has no physical significance. ",
        "introduction": "Despite the central role played by the energy-momentum tensor of matter in General Relativity, there is no widely accepted way to localise the energy and momentum of the gravitational field itself. In the place of a genuine solution to this problem, we are forced to make do with an over-abundance of energy-momentum pseudotensors, objects designed to display some or other property befitting a measure of gravitational energy-momentum, but whose coordinate dependence renders them of little physical significance beyond giving the correct integrals at infinity in asymptotically flat spacetimes. Even for weak gravitational waves, the best measures at our disposal only become meaningful once we have averaged over many wavelengths.  The canonical response to the gravitational energy-momentum problem is to dismiss it as ``looking for the right answer to the wrong question''\\cite{MTW}; but while the well-known argument presented by Misner, Thorne and Wheeler is certainly compelling, it is far from watertight. They remind us that the equivalence principle ensures that all ``gravitational fields'' $\\Gamma^{\\alpha}_{\\phantom{\\alpha}\\beta \\gamma}$ can be made to vanish at a point by a suitable choice of coordinates, and conclude that because gravity is locally zero, there can be no energy density associated with it. However, this argument fails to consider tensors containing second derivatives of the metric, which unlike $\\Gamma^{\\alpha}_{\\phantom{\\alpha}\\beta \\gamma}$ cannot be made to vanish by choice of coordinates, and really do reflect the local curvature of spacetime: for example, the Riemann tensor can be used to construct objects such as the Bel-Robinson tensor \\cite{Bel}. Misner, Thorne and Wheeler also point out that, while the matter energy-momentum tensor derives its physical significance by curving space, a similar tensor for gravity would not be a source term for the field equations. However, this stance is based around a prejudice for writing the Einstein field equations as $G^{ab} = \\kappa T^{ab}$ with gravity on the left and matter on the right; there is nothing to stop us splitting up $G^{ab}$ in a covariant fashion, grouping one part with $T^{ab}$, and interpreting \\emph{this} as the total energy-momentum source, taking the remainder of $G^{ab}$ to be the gravitational `response'. Despite these reservations, the argument in \\cite{MTW} remains vindicated as yet by the failure of these escape-routes to yield anything which can be physically interpreted as an energy-momentum tensor.  It might appear that the only straightforward solution to the problem is to extend the definition of the matter energy-momentum tensor $T^{ab}$ (a functional derivative of the matter Lagrangian with respect to the metric) to the gravitational field, and conclude that the gravitational energy-momentum tensor is $-G^{ab}/\\kappa$, where $\\kappa= 8\\pi G/c^4 $. The Einstein field equations could then be interpreted as a constraint that everywhere sets to zero the sum of gravitational and matter energy-momentum. While one might claim this simple idea conveys some important physical insight, it suffers from numerous problems. Firstly, $-G^{ab}/\\kappa$ lacks the analytical power one expects from an energy-momentum tensor: the ability to split the set of all physical systems at a particular time into classes of \\emph{different} total energy and momenta, so that conservation laws alone can reveal that two particular systems could never be part of the same spacetime. Secondly, it leads us to conclude that the gravitational field only has energy where matter is also present, precluding the use of this prescription to describe the energetics of gravitational waves, or define a gravitational tension in the vacuum between massive bodies. Thirdly, the energy-momentum tensors for gravity and matter are conserved \\emph{separately} ($\\nabla_a G^{ab}=0$ and $\\nabla_a T^{ab}=0$) so that although there is a delicate balance that keeps their sum zero, it is not the case that energy or momentum simply `flows' between gravity and matter, as $\\nabla_a (T^{ab} - G^{ab}/\\kappa)=0$ \\emph{alone} would imply. Lastly, we note that the conservation law $\\nabla_aG^{ab}=0$ actually tells us nothing at all about the gravitational field; it is satisfied identically, without any need for the equations of motion to hold. Because of these drawbacks, if we are to regard $-G^{ab}/\\kappa$ as a solution to the gravitational energy-momentum problem, we consider it rather a trivial one. Clearly, the reason for this triviality is that we have over-worked the metric: we cannot use the functional derivative with respect to a dynamical field as a way of defining the energy-momentum tensor for that same field, as we will only end up writing down the equations of motion twice. This line of reasoning leads us to consider that one method of attack for this problem may be to separate the two roles played by $g^{ab}$ in General Relativity, that of dynamic field and spacetime metric.  In \\cite{G}, Grishchuk develops a ``field-theoretical'' approach to gravitation, which expresses the physical content of General Relativity (GR) in terms of a dynamical symmetric tensor field in flat Minkowski spacetime. Although this formulation has been carefully designed to agree with the empirical predictions of GR, in \\cite{BG} Babak and Grishchuk claim that the flat-space approach allows them to define a unique, symmetric, and non-trivial energy-momentum tensor for the gravitational field. The major purpose of this paper is to examine the extent to which this tensor is physically meaningful. ",
        "conclusions": "The formulation of gravity presented in \\cite{G} succeeds in recasting General Relativity as a flat-space theory of a symmetric tensor field. While we do not find fault with the formalism itself, we assert that care must by taken in its interpretation, as we believe we have demonstrated that only those quantities which can be defined solely in terms of GR tensors are of any physical importance. The physically insignificant content of the flat-space formalism is a consequence of an unmeasurable field $\\gamma^{ab}$ which is not uniquely determined by the requirement that it be a flat metric tensor.  There are in principle two ways to deal with the non-uniqueness of $\\gamma^{ab}$: 1.~Pick a particular flat metric and declare that this is the immutable `correct' choice, to be used in all situations; 2.~Allow $\\gamma^{ab}$ to depend somehow on the physical system we are describing, or how we have chosen to represent the system mathematically.  The problem with the first stance is that the theory still retains $\\gamma$-fixed gauge transformations. To see this, note that equation (53) of \\cite{BG} expresses the equivalence of Grishchuck's equations of motion ($r_{ab} \\equiv - P^c_{\\phantom{c}ab\\gamsem c} - P^c_{\\phantom{c}ad}P^d_{\\phantom{d}bc} + \\tfrac{1}{3}P_a P_b =0$) with the Einstein field equations: \\begin{eqnarray} R_{ab}[g]=\\breve{R}_{ab}[\\gamma] + r_{ab}[h,\\gamma]. \\end{eqnarray} Babak and Grishchuk interpret this relation as follows: given a flat-metric $\\gamma^{ab}$, an $h^{ab}$ that satisfies $r_{ab}=0$ will enforce $R_{ab}=0$, establishing the agreement with GR. However, one can always use this equation to make the converse argument: given a flat-metric $\\gamma^{ab}$, a $g^{ab}$ which solves $R_{ab}=0$ will enforce Grishchuk's equation $r_{ab}=0$. As $R_{ab}[\\phi^* g]=0$ if $R_{ab}[g]=0$, we can construct a whole range of solutions $\\{h^{\\prime ab}: r_{ab}[h^\\prime]=0\\}$ from $h^{ab}$ simply by applying diffeomorphisms to $g^{ab}$. Because we declared $\\gamma^{ab}$ to be immutable, these new solutions will correspond to $\\gamma$-fixed transformations of $h^{ab}$. Crucially, as $g^{\\prime ab}= \\phi^*g^{ab}$, no `geometric' experiment can tell any $h^{\\prime ab}$ apart from from $h^{ab}$. Thus, without a method to measure a non-GR quantity directly, we have to conclude that these new solutions represent physically equivalent systems, and that the $\\gamma$-fixed transformation is a gauge transformation of the theory.  The second stance appears to be able to dodge this argument, because one can claim that we should have applied the same diffeomorphism to $\\gamma^{ab}$ that we applied to $g^{ab}$, forcing us to perform a harmless DGT instead of a $\\gamma$-fixed transformation. However, if we take this view, we will need a heuristic for deriving $\\gamma^{ab}$ from measurable quantities, otherwise we will never know where to start with the `correct' pairing $(g^{ab}, \\gamma^{ab})$. In order that this heuristic be consistent with arbitrary DGTs (which are gauge transformations of any tensorial theory) any prescription for which  $T^{ab}=0 \\Rightarrow h^{ab}=0$ inevitably leads us to identify $\\gamma$-transformations as gauge transformations anyway, because we are free to represent the $T^{ab}=0$ limit in any diffeomorphism gauge we choose.  Accepting that $\\gamma$-transformations and $\\gamma$-fixed transformation are maps between different mathematical representations of the same physical system, we conclude that the exotic gauge transformation properties of $t^{ab}$ cannot allow us to interpret this tensor as a local measure of the energy and momentum content of the gravitational field. Although $t^{ab}$ is a perfectly legitimate mathematical construction, its dependence on the unmeasurable and non-unique tensor $\\gamma^{ab}$ renders it ill-defined, and devoid of physical meaning."
    },
    "0807/0807.0438_arXiv.txt": {
        "abstract": "{The search for B-mode polarization fluctuations in the Cosmic Microwave Background is one of the main challenges of modern cosmology. The expected level of the B-mode signal is very low and therefore requires the development of highly sensitive instruments with low systematic errors. An appealing possibility is bolometric interferometry.} {We compare in this article the sensitivity on the CMB angular power spectrum achieved with direct imaging, heterodyne and bolometric interferometry.} {Using a simple power spectrum estimator, we calculate its variance leading to the counterpart for bolometric interferometry of the well known Knox formula for direct imaging.} {We find that bolometric interferometry is less sensitive than direct imaging. However, as expected, it is finally more sensitive than heterodyne interferometry due to the low noise of the bolometers. It therefore appears as an alternative to direct imagers with different and possibly lower systematic errors, mainly due to the absence of an optical setup in front of the horns.} {} ",
        "introduction": " ",
        "conclusions": "} The sensitivities of the three different techniques only differ in the way the instrument filters the multipoles observed in the sky. An imager is affected by its resolution on the sky while an interferometer is affected by the ratio between the number of equivalent baselines and the number of horns as a function of multipoles. All of these filtering factors are less than one. However, imagers are usually operated in such a way that they are not limited by their angular resolution in the multipole region of interest, in that case $B_\\ell\\simeq 1$, and the imager is always more sensitive than an interferometer (bolometric or heterodyne). From the strict point of view of sensitivity, interferometers can therefore only compete with low angular resolution imagers.  There is a large difference in sensitivity between bolometric and heterodyne interferometers compared to an imager: the ratio $N_\\mathrm{eq}/N_h$ acts quadratically on the variance for a bolometric interferometer while it acts linearly for a heterodyne instrument. This is due to the fact that with a heterodyne interferometer, equivalent baselines are averaged after their measurement, resulting in a $1/N_\\mathrm{eq}$ factor on the variances. In a bolometric interferometer, the signals from all $N_h$ horns are added together multiplying the noise variance by $N_h$ while the coherent summation of equivalent baselines performs an efficient $1/N_\\mathrm{eq}^2$ reduction of the noise. This finally results in a factor $N_h/2N_\\mathrm{eq}$ for the variance of a bolometric interferometer relative to a heterodyne one. This is largely compensated by the difference in  NET between bolometric instruments and coherent ones. When comparing them, the ratio of their NET also appears quadratically and favours bolometric instruments that are dominated by the photon noise rather than by that of the amplifiers. This situation may change in the future with the improvements of the HEMT technologies but at frequencies around and above 100 GHz we are unlikely to face photon noise limited HEMTs in the near future. The difference between the NET would be even greater in space where the bolometers NET would drop as the background temperature while that of the coherent instruments would remain roughly constant.  With the present technologies of bolometers and coherent amplifiers, the hierarchy in terms of sensitivity between the three techniques (and layout) studied here is very clear for the multipole range $25<\\ell<200$ where the primordial B-mode signal is expected to be maximal. Imagers are the most sensitive, bolometric interferometers have a lower sensitivity, the ratio dropping quadratically with the multipole considered. Heterodyne interferometers have an even lower sensitivity but the ratio with an imager drops less rapidly. They remain however less sensitive than bolometric interferometers in the range of multipoles considered here, where the largest primordial B-mode signal is expected and where the lensing of the E-modes into B-modes is still low. At higher multipoles however, the bandwidth smearing effect and loss of coherence would be a real issue for bolometric interferometers while in a heterodyne interferometer, the separation into small bands would prevent the sensitivity from dropping.  The main remaining question is whether the gain in terms of systematic effects is worth the price of this sensitivity reduction if one builds an interferometer instead of an imager. In terms of optics for instance, an interferometer directly observes the sky. The primary beam is therefore only set by that of the horns, while in an imager, the telescope (mirror or lenses) produces sidelobes inducing poorly predictible  ground pickup that often prevent one from reaching the nominal sensitivity. An interferometer is also completely insensitive to spatially uniform polarized signals that vary with time such as polarized atmospheric contamination. These could also prevent an imager from reaching its nominal sensitivity by adding some spread in the noise. These examples mitigate the statistical sensitivity loss of an interferometer with respect to an imager. The differences in terms of systematic effects between imagers and bolometric and heterodyne interferometers are not obvious and deserve a detailed quantitative study in continuation of the work done by~(\\cite{bunn})."
    },
    "0807/0807.2926_arXiv.txt": {
        "abstract": "The all-flavor solar neutrino flux measured by the Sudbury Neutrino Observatory (SNO) constrains nonstandard energy losses to less than about 10\\% of the Sun's photon luminosity, superseding a helioseismological argument and providing new limits on the interaction strength of low-mass particles. For the axion-photon coupling strength we find $g_{a\\gamma}<7\\times10^{-10}~{\\rm GeV}^{-1}$. We also derive explicit limits on the Yukawa coupling to electrons of pseudoscalar, scalar and vector bosons with keV-scale masses. ",
        "introduction": "\\label{sec:introduction}  The interaction strength of new low-mass particles with photons, electrons or nucleons is severely constrained by the well-known requirement that stars not lose energy in excess of observational constraints~\\cite{Raffelt:1990yz, Raffelt:1996wa, Raffelt:2006cw}. Depending on the assumed particle mass and interaction structure, the most restrictive limits typically derive from population statistics in globular clusters, the white dwarf luminosity function, and the duration of the SN~1987A neutrino signal. Such constraints usually imply that the solar particle emission is small compared to its photon luminosity, but it can still serve as a powerful source, e.g.\\ for the ongoing solar axion searches~\\cite{Moriyama:1998kd, Inoue:2008zp, Andriamonje:2007ew, Arik:2008mq}.  Constraints based on the properties of the Sun remain of interest, even if they are less restrictive than other astrophysical arguments, because they are more direct and thus perhaps more comparable to laboratory experiments. Moreover, sometimes the Sun is enough to test a given hypothesis. In this case it is nice to have a simple argument at hand that does not require more complicated astrophysical reasoning.  Previously the most sensitive diagnostic for the solar interior was the helioseismological sound-speed profile, providing restrictive energy-loss limits~\\cite{Schlattl:1998fz}. Of course, the underlying chain of arguments is not simple and worse, a new determination of the solar element abundances~\\cite{Asplund:2004eu} has created an unresolved tension between solar modeling and helioseismology~\\cite{Bahcall:2004yr, PenaGaray:2008qe}. While the conservative limit of Ref.~\\cite{Schlattl:1998fz} likely remains unchanged, it is nice that the measured solar neutrino flux provides a somewhat more restrictive limit based on a simpler argument.  The recently completed SNO measurements of the all-flavor solar neutrino flux~\\cite{Ahmad:2002jz, Aharmim:2005gt, Aharmim:2008kc} probe directly the physical conditions of the particle-emitting solar core. The steep temperature dependence of the $^8$B neutrino production rate provides a sensitive test of the Sun's interior that would be hotter if a lot of ``invisible energy'' were produced. The main purpose of our note is to update solar energy-loss constraints using the SNO results.  The inner solar temperature is around 1~keV, so these limits always apply to sub-keV mass particles, notably axions. Recently the hypothesis of keV-scale bosons as possible dark matter candidates has received some attention~\\cite{Pospelov:2008jk}. Moreover, some time ago it was proposed~\\cite{Bernabei:2005ca} that keV-mass pseudoscalars explain the annual modulation observed in the DAMA/Libra experiment~\\cite{Bernabei:2008yi}. Unfortunately, this intriguing intepretation was based on the incorrect axio-electric absorption rate of Refs.~\\cite{Bernabei:2005ca, Dimopoulos:1986mi}, the correct rate being much larger but nearly independent of velocity~\\cite{Pospelov:2008jk}. While an earlier version of our manuscript was largely motivated by this now-dismissed interpretation, an evaluation of the solar limit for keV-mass bosons may still prove useful in future.  We present the new solar energy-loss constraint in the context of axions and apply it explicitly to the axion-photon interaction strength. This result is of interest for solar axion searches. We also treat explicitly keV-mass bosons $\\chi$ that couple to electrons. To this end we re-derive the $\\gamma+e^-\\to e^-+\\chi$ Compton cross section, correcting several errors in the literature. ",
        "conclusions": "\\label{sec:summary}  The SNO measurements of the all-flavor solar neutrino flux produced by the very temperature-dependent ${}^8$B reaction severely constrain anomalous solar energy losses. We have re-considered self-consistent solar models produced by Schlattl et~al.~\\cite{Schlattl:1998fz} who included axion losses by the Primakoff effect. We have observed that the predicted solar neutrino flux is nicely reproduced by a simple and intuitive power law as a function of the assumed anomalous solar energy loss. In this way the measured neutrino flux and the assumed energy loss are directly related in a simple form. The excellent agreement between the measured and predicted solar neutrino flux provides a restrictive limit on any new energy loss channel of the Sun. While constraints from other astrophysical arguments are usually more restrictive, the solar neutrino limit on new energy losses is complementary in that it is based on a direct diagnostic of the solar interior.  In particular, we have derived a new solar limit on the axion-photon interaction strength $g_{a\\gamma}$,  superseding an often-cited helioseismological result. Only the CAST experiment is sensitive enough to detect solar axions obeying our new constraint Eq.~(\\ref{eq:solarlimit}).  For bosons coupling to electrons, our limit extends to masses of almost 10~keV even though the solar inner temperature is around 1~keV. This would have been of interest to constrain the DAMA annual modulation in terms of keV-scale pseudoscalar dark matter particles. However, based on the corrected axio-electric absorption rate of Ref.~\\cite{Pospelov:2008jk} this interpretation is no longer viable. Instead, recent direct constraints on keV-scale pseudoscalar dark matter by CoGeNT~\\cite{Aalseth:2008rx} and CDMS~\\cite{Collaboration:2009ht} are more restrictive than the solar limit.  \\begin{figure}[t] \\includegraphics[width=0.9\\columnwidth]{fig1.eps} \\caption{Mass dependence of the solar $g_{ee\\chi}$ limit. The PS curve supersedes the one in our original preprint that was reproduced in Ref.~\\cite{Aalseth:2008rx}.\\label{fig:exclusion}} \\end{figure}"
    },
    "0807/0807.1925_arXiv.txt": {
        "abstract": "We present new metallicity measurements for 298 individual red giant branch stars in eight of the least luminous dwarf spheroidal galaxies (dSphs) in the Milky Way (MW) system.  Our technique is based on medium resolution Keck/DEIMOS spectroscopy coupled with spectral synthesis.  We present the first spectroscopic metallicities at $\\mathfeh < -3.0$ of stars in a dwarf galaxy, with individual stellar metallicities as low as $\\mathfeh = -3.3$.  Because our \\feh\\ measurements are not tied to empirical metallicity calibrators and are sensitive to arbitrarily low metallicities, we are able to probe this extremely metal-poor regime accurately.  The metallicity distribution of stars in these dSphs is similar to the MW halo at the metal-poor end.  We also demonstrate that the luminosity-metallicity relation previously seen in more luminous dSph galaxies ($M_V = -13.4$ to $-8.8$) extends smoothly down to an absolute magnitude of $M_V = -3.7$.  The discovery of extremely metal-poor stars in dSphs lends support to the $\\Lambda$CDM galaxy assembly paradigm wherein dwarf galaxies dissolve to form the stellar halo of the MW. ",
        "introduction": "\\label{sec:intro} In the hierarchical theory of the assembly of galaxy halos \\citep{whi78,sea78}, dwarf galaxies interact gravitationally with their hosts, shedding stars, losing gas, and eventually tidally dissolving into the diffuse halo.  Recent numerical simulations and semi-analytic models \\citep[e.g.,][]{bul05} suggest that a Milky Way-like halo can be explained entirely by disrupted satellites in the $\\Lambda$CDM paradigm.  Clearly, dwarf galaxies must play some role in building stellar halos of large galaxies because both the Milky Way (MW) and \\object{M31} exhibit tidal streams \\citep[e.g.,][]{iba01a,cho02,gil07} and dwarf galaxies in various stages of disruption \\citep[e.g.,][]{iba94, how08}.  However, the chemical abundances of individual stars in present-day MW dSphs do not match that of the MW halo.  \\citet{she01} found a lower \\afe\\ for dSph red giant branch (RGB) stars than for MW halo RGB stars.  The differences in chemical abundances led \\citet{tol03} to conclude that the bulk of the halo can not be composed of stars like those present in surviving dSphs.  In order to compare the bulk metallicities of stars in dSphs and the MW halo, \\citet[hereafter H06]{hel06} obtained medium-resolution spectra of the MW dSphs Sculptor, Fornax, Sextans, and Carina.  Using an empirical relation between the infrared \\ion{Ca}{2} triplet (CaT) equivalent width (EW) and \\feh\\ \\citep{tol01}, H06 find a lack of $\\rm{[Fe/H]} < -3.0$ stars in these four MW dSphs. Given this absence, they concluded that the MW halo, which contains numerous such stars, could not have formed from present-day dSphs.  Nevertheless, dSphs could still be the building blocks if they once contained a population of extremely metal-poor stars. However, it remains difficult to understand how they could have lost \\emph{all} of those stars by now.  In this Letter, we revisit the absence of metal-poor stars in MW dSphs by targeting lower luminosity galaxies and by using a more direct technique to measure \\feh\\ (based on Fe lines) than has been used before on low or moderate resolution spectra.  In \\S\\,\\ref{sec:measure}, we describe new metallicity measurements from the \\citet[hereafter SG07]{sim07} data set of eight of the least luminous MW dSphs.  As a result, we report for the first time the discovery of extremely metal-poor stars ($\\mathfeh < -3.0$) in MW dSphs.  In \\S\\,\\ref{sec:mdf}, we compare the ultra-faint dSph MDF to the halo MDF.  We also present the luminosity-metallicity relation for the full range of MW dSph luminosities.  In \\S\\,\\ref{sec:concl}, we briefly summarize our findings and discuss work on dSph chemical abundances beyond [Fe/H]. ",
        "conclusions": "\\label{sec:concl}  We have presented metallicity measurements for eight of the least luminous known galaxies in the Universe.  We also discover, for the first time, stars outside the MW field halo population with $\\mathfeh < -3.0$.  Furthermore, we have shown that the distribution of [Fe/H] in present-day dSphs reaches nearly as low as that of the MW stellar halo, and that the dSph luminosity-metallicity relation is well-defined for nearly 4~dex in luminosity.  There are two main differences between our study and previous works that might have contributed to our discovery of extremely metal-poor stars in MW dSphs.  First, our spectral synthesis approach is valid for any metallicity and is not restricted to calibrated ranges, like the CaT technique.  Second, we explored very faint dSphs whereas H06 examined more luminous dSphs. We have found extremely metal-poor stars only in the faintest dSphs. It remains to be seen whether any of the brighter dSphs also harbor extremely metal-poor stars.  [Fe/H] is just one abundance puzzle in the role of dSphs in building the stellar halo.  Additional elements will need to be examined to obtain further clues. Most notably, [$\\alpha$/Fe] ratios in the more luminous dSphs are on average lower than in the halo \\citep[e.g.,][]{she03,gei07}.  However, cosmologically motivated models including star formation and chemical enrichment \\citep{rob05,fon06} may explain the difference.  These models along with the discovery of extremely metal-poor stars in long-lived dSphs support the original hierarchical paradigm of galaxy formation \\citep[e.g.,][]{sea78,whi78}."
    },
    "0807/0807.1114_arXiv.txt": {
        "abstract": "This pedagogical review covers an unsolved problem in the theory of protoplanetary disks: the growth of dust grains into planetesimals, solids at least a kilometer in size.  I  summarize timescale constraints imposed on  planetesimal formation by circumstellar disk observations, analysis of meteorites, and aerodynamic radial migration.  The  infall of $\\lesssim$ meter-sized solids in a hundred years is the most stringent constraint.  I review proposed mechanisms for planetesimal formation.  Collisional coagulation models are informed by laboratory studies of microgravity collisions. The gravitational collapse (or Safronov-Goldreich-Ward) hypothesis involves detailed study of the interaction between solid particles and turbulent gas.  I cover the basics of aerodynamic drag in protoplanetary disks, including radial drift and vertical sedimentation.  I describe various mechanisms for particle concentration in gas disks -- including turbulent pressure maxima, drag instabilities and long-lived anticylonic vortices.  I derive a general result for the minimum size for a vortex to trap particles in a sub-Keplerian disk.  Recent numerical simulations  demonstrate that particle clumping in turbulent protoplanetary disks can trigger gravitational collapse.  I discuss several outstanding issues in the field. ",
        "introduction": "This chapter serves as a pedagogical introduction to planetesimal formation and related processes in protoplanetary disks.  I defer to \\citet{dbcw07} for a comprehensive review of research in this rapidly progressing field.  % Section \\ref{s:times} serves as an extended introduction by discussing timescale constraints from disk observations, meteoritics, and dynamics.  The purposeful overlap with other chapters in this volume  aims to reinforce the overlap between sub-fields.  Section \\ref{s:mech} describes formation mechanisms: collisional coagulation (\\S\\ref{s:coag}) and gravitational collapse (\\S\\ref{s:grav}).  Section \\ref{s:drag} covers the dynamics of gas drag in protoplanetary disks, including the basic form of drag forces (\\S\\ref{s:tstop}), vertical settling and radial migration (\\S\\ref{s:drsett}), and particle trapping mechanisms (\\S\\ref{s:trap}). ",
        "conclusions": "\\label{s:concl} Planetesimal formation remains unsolved, despite considerable progress on several fronts.  Laboratory studies of collisions place constraints on sticking assumptions \\citep{bw08}.  Numerical simulations coupled with analytic theory provide insights to physical processes in disks.  The response of particles to -- and the effect of particles on -- turbulent gas is an especially promising area.  Contrary to previous expectations, simulations show that gravitational collapse of $\\gtrsim 10$ cm solids into planetesimals occurs in  turbulent protoplanetary disks \\citep{nature07}.   Forming planetesimals directly by gravitational collapse from chondrule-sized, $\\lesssim 1$ mm, solids would be an elegant solution.  But such small particles are more tightly coupled to the gas, making it difficult to understand (analytically or numerically) how they might collapse in standard disk models.  Sadly numerical (and even laboratory) experiments lack direct observational confirmation.  But the initial conditions and basic ingredients of theoretical models can be compared to the meteoritic record and observations of protoplanetary disks.  Spitzer is delivering exciting, detailed results.  Soon Herschel and ALMA will further expand our knowledge of the conditions in -- and structure of -- disks."
    },
    "0807/0807.3920_arXiv.txt": {
        "abstract": "We present a system model for optical and far UV spectra of the nova-like variable UX~UMa\tinvolving a white dwarf, secondary star, gas stream, hot spot and accretion disk using our code BINSYN and based on an initially adopted system distance. Calculated SED intensity data successfully fit successive tomographically-extracted annuli longward of the Balmer limit but require a postulated `iron curtain' shortward of the Balmer limit that is applied to the annulus section closest to the secondary star, while postulated recombination emission fills in the model SED shortward of the Balmer limit and is applied to the annulus section more remote from the secondary star. The same model fits $UBV$ 1954 light curves by Walker and Herbig. Fits to $HST$ $FOS$ spectra are approximate but require assumed time-variable changes in the SED. Comparable effects, possibly involving variable absorption, afflict $FUSE$ spectra. Fits to $IUE$ spectra by the model show time-dependent residuals that indicate changes in the accretion disk temperature profile, possibly indicative of a slightly variable $\\dot{M}$ from the secondary star. Using model-based component light contributions and the improvement on the Bailey relation by Knigge we determine the system distance and mass transfer rate. ",
        "introduction": "\\label{s-intro}  Cataclysmic variables (CVs) are semi-detached binary stars in which a late-type typically main sequence star loses mass onto a white dwarf (WD) via Roche lobe overflow and accretion proceeds through a viscous disk. In the nova-like (NL) subclass, of which UX UMa is the prototype, the mass transfer rate from the secondary star is large enough that dwarf nova outbursts do not occur. (See \\citet{warner95} for a history of UX UMa studies and a thorough review of CV types and properties.) Important early studies of UX UMa by \\citet{wh1954,jph1954} and \\citet{kw1963} established the system photometric properties, while \\citet{nr1974} discovered coherent oscillations. $IUE$ spectra identified UX UMa as a UV emission line variable \\citep{holm1982,king1983}. X-ray observations \\citep{becker1981,wood1995} detected a soft source, further refined \\citep{pratt2004} to a noneclipsed soft component and a deeply eclipsed hard component. An accretion disk wind was discovered by \\citet{mason1995} and modeled by \\citet{knigge1997}.  \\citet{knigge1998a} had mixed success in modeling the spectral energy distribution (SED) of UX UMa with a combination of stellar atmosphere synthetic spectra. A similar approach by \\citet{froning2003} (hereafter FR2003) in a fit to the $FUSE$ spectrum produced an estimated $\\dot{M}=1.6{\\times}10^{-9}~M_{\\odot}~{\\rm yr}^{-1}$. However, a mass transfer rate of $\\dot{M}=1.6{\\times}10^{-9}~M_{\\odot}~{\\rm yr}^{-1}$ is small enough that the system would be unstable against outburst, contrary to observation. Neither of these two studies included a SED contribution from the WD; on the other hand, tomographic reconstruction of the UX UMa disk \\citep{ru1992} is consistent with a standard model accretion disk with $\\dot{M}=5{\\times}10^{-9}~M_{\\odot}~{\\rm yr}^{-1}$. A comparable study by \\citet{bap1995} hereafter BA1995, with a different estimated distance to the system, finds $\\dot{M}=1{\\times}10^{-8}~M_{\\odot}~{\\rm yr}^{-1}$. \\citet{frank1981} modeled IR and optical light curves and determined system parameters. \\citet{smak1994a} modeled optical region light curves and derived bright spot parameters. \\citet{mason1997} argue that a more complicated structure than a bright spot on an otherwise unmodified accretion disk rim is necessary. BA1995 derived a WD $T_{\\rm eff}$ of 52,000K to 70,000K, but FR2003 points out that a $T_{\\rm eff}$ that large is inconsistent with FUV spectra. Neither of the tomographic reconstruction studies has been the basis of an effort to model the spectra. These continuing inconsistencies are the motivation to attempt to derive a single model that represents a substantial part of the disparite observational data. ",
        "conclusions": "An important part of the system analysis for UX UMa depends on knowledge of its distance. BA1995 calculate a value of 345pc with an estimated error of 34pc. The 345pc determination depends on fits in a color-magnitude diagram, with one wavelength at 1523\\AA. The fits use a theoretical color-magnitude relation derived on assumed radiation characteristics of the source. If the source can be represented by a black body, the derived distance is 401pc, and if standard Kurucz model atmospheres, 312pc. We feel the residuals in the SED fits found in the present study are large enough that an assumed radiation characteristic for the source should be treated with caution, and that the estimated error in the 345pc determination could be larger. Other distance determinations are in \\S3. The \\citet{knigge2006} method is an important improvement on the \\citet{bailey1981} relation; as cited in \\S3., this method leads to a distance of 215pc if the secondary provides all of the system K-flux and 376pc if the secondary provides 1/3 of the system flux. Based on the parameters of our \\S4. model, and adopting a secondary $T_{\\rm eff}=3575$K \\citep[Table~3]{knigge2006} the secondary provides 0.48 of the system flux at 2.2$\\mu$, leading to a calculated distance of 312pc. Consider the sensitivity of the calculated distance to variation of system parameters. The flux ratio secondary/(secondary+disk) is sensitive to the accretion disk flux which in turn depends on $\\dot{M}$. (The contribution of the adopted 20,000K WD is only a few tenths of a percent at 2.2$\\mu$.) From Figure~1 we estimate that the scatter of the observational data constrains $\\dot{M}$ within a factor of about 2 for the adopted system distance of 250pc. Increasing $\\dot{M}$ to $1.0{\\times}10^{-8}M_{\\odot}~{\\rm yr}^{-1}$ produced a new calculated distance of 332pc. We increased $M_{\\rm wd}$ by $10\\%$, while maintaining $M_{\\rm sec}$ fixed, to test the calculated distance sensitivity to $M_{\\rm wd}$ variation. The new distance was 323pc. Variation in the observed K magnitude of a few times 0.01 produces variation of only a few parsecs in the calculated distance. We propose a distance of $312\\pm30$pc as the best currently available distance determination to UX UMa and note the accordance with the value found by BA1995 using a fit to Kurucz model atmospheres.  This study finds that a standard model ${\\dot{M}}=5.0{\\times}10^{-9}M_{\\odot}~{\\rm yr}^{-1}$ accretion disk surrounding a 20,000K (but see below), $0.47M_{\\odot}$ WD at a distance of 250pc provides a model that reasonably fits spectral intensity data (Figure~4) and spectral flux data (Figure~8 through Figure~11). All of the DS1 and DS2 data can be equally well represented if the distance is 345pc and ${\\dot{M}}=1.0{\\times}10^{-8}M_{\\odot}~{\\rm yr}^{-1}$. In particular, the observational data of Figure~1 depend only on the adopted distance. Interpolating between the 250pc and 345pc calibrations, for our preferred distance of 312pc, the corresponding $\\dot{M}=8{\\times}10^{-9}M_{\\odot}~{\\rm yr}^{-1}$. We take this value to be our final result for the mass transfer rate. In principle it would be possible to iterate our solution, starting over with the new distance and $\\dot{M}$ determination but, based on our study of the ${\\dot{M}}=1.0{\\times}10^{-8}M_{\\odot}~{\\rm yr}^{-1}$ case at 345pc, we believe there would be no improvement in any of the plots presented in this paper.  \\citet{puebla2007} (hereafter P2007) use a separate method to study accretion rates and, for UX~UMa (their Table~2) find ${\\dot{M}}=1.7{\\times}10^{-8}M_{\\odot}~{\\rm yr}^{-1}$ from black body fits and ${\\dot{M}}=1.4{\\times}10^{-8}M_{\\odot}~{\\rm yr}^{-1}$ for their model accretion disk fits, adopting a distance to UX UMa of 340pc and $M_{\\rm wd}/M_{\\odot}$ in the range 0.4-0.8. This result is in approximate agreement with our results. In their Figure~6 P2007 fit $IUE$ spectra of UX~UMa with two models. Their second model, for $M_{\\rm wd}/M_{\\odot}=0.4$, finds ${\\dot{M}}=5.5{\\times}10^{-9}M_{\\odot}~{\\rm yr}^{-1}$, in close agreement with our model. P2007 parameterize the WD contribution with ${\\zeta}=f_{\\rm wd}/f_{\\rm disk}$ and define a ``disk-dominated'' system as one in which ${\\zeta}<0.1$, where the flux values are integrated contributions from 1500\\AA~to 3250\\AA. We would prefer to define ``disk-dominated'' in terms of an integration extending to 950\\AA~since a large part of the flux from a hot WD can occur shortward of 1500\\AA. Because of the ambiguity of the WD $T_{\\rm eff}$ in UX~UMa (see below), we do not estimate a value for $\\zeta$ except to note that, from Figure~4a, the accretion disk supplies well over $90\\%$ of the system flux if the WD $T_{\\rm eff}=20,000$K (compare with the zero temperature WD where the accretion disk supplies $100\\%$ of the system flux, excluding the secondary).  Several studies (\\S3) support the adopted $M_{\\rm sec}=0.47M_{\\odot}$. The adoption of $q=1.0$ places the system at the boundary of instability against dynamical scale mass transfer and a $M_{\\rm wd}$ larger than $0.47M_{\\odot}$ seems likely. \\citet{bap1995} state that adoption of the \\citet{smak1994a} $M_{\\rm wd}=0.70M_{\\odot}$ increases the system distance from 345pc to 386pc; the effect of a smaller $q$ is to increase the size of the accretion disk and so make it more luminous. A WD mass greater than $0.47M_{\\odot}$ would be smaller and would produce a deeper potential well so, for the same $\\dot{M}$, the accretion disk would be hotter and the spectral gradient would be steeper. Preservation of a fit to the calibrated temperatures of Figure~1 would require either a reduction in $\\dot{M}$ or a nonstandard model temperature profile.  The evidence concerning the WD $T_{\\rm eff}$ is ambiguous. BINSYN requires specification of a WD $T_{\\rm eff}$ and our adopted \\S4. model includes a 20,000K WD but, as Figure~4a shows, this model differs almost negligibly from a zero temperature WD; the hotter inner annuli, Table~4, contributions dominate the system synthetic spectrum. A 40,000K or hotter WD, directly visible to the observer, would be strongly inconsistent with the Figure~3 intensity data while the SED data, Figure~8 through Figure~11, all are inconsistent with a hot WD. The only data supporting a hot WD are Figure~4a and Figure~4b. Although the DS1 data plotted in Figure~4a and Figure~4b superficially support a 40,000K WD, we find it more attractive to attribute the excess flux source to something like a boundary layer in common with the other unmodeled hot sources described in connection with Figure~4a. We argue that, in Figure~4a, the central region is seen through an absorbing layer, likely the ``transition region'' between the accretion disk and the fast wind \\citep{knigge1997}, which produces the deep absorption features.  Studies of other cataclysmic variable systems generally support a WD $T_{\\rm eff}$ hotter than 34,000K \\citep{sion1999,knigge2000,sion2008} at the orbital period of UX UMa. It is informative that the WD temperature determination for DW UMa obtained during a low state of accretion when the WD was clearly visible \\citep{knigge2000} were much higher than during the normal high accretion state, indicating that the accretion disk can occult the WD. Since the orbital period of UX UMa is outside the range for systems that undergo low states, there is no opportunity to determine how much disk occultation occurs. \\citet[Table~3]{sion1999} lists CV systems for which absorbing curtains obscuring the WDs have been calculated. Note that our application would be more complex since, in addition to obscuring the central region, we require obscuration of the `front' of particular annular regions, separate from the WD.  As we have demonstrated, an accretion disk with constant $\\dot{M}$ fits both the DS2 and DS1 observations, and this datum (a fixed, unchanging accretion disk) might imply that a changed WD explains the difference between Figure~3 and Figure~4b. But if $\\dot{M}$ has remained nearly constant we see no mechanism to heat the WD (and we do not believe a changed $\\dot{M}$ could heat the WD from 20,000K to 40,000K), while there are documented changes in the accretion disk\t(Figure~8 to Figure~13) and one such change might be formation of a hot emitting region between the times of the two data sets (June 4, 6, and 7, 1992 for DS2 and Nov. 11, 1994 for DS1).  Our relatively simple bright spot model, \\S4.1, represents the $UBV$ light curves of Figure~5 through Figure~7, including the luminosity maximum just before eclipse. These simulations are the basis for our determination of the system orbital inclination. UX UMa occasionally shows anomalous light curves that can not be represented by even a sophisticated model like that of HO1994. An example is in Figure~15. Less extreme cases show a depressed light level from phase 0.2 to 0.8, with a downward tilt to phase 0.8, followed by a rapid rise to light maximum at phase 0.9. These features become more pronounced from $V$ to $B$ to $U$. An example is (our) Figure~6 or Figure~3 of \\citet{jph1954}. \\citet{mason1997} attribute this effect to a disk `bulge' upstream from the bright spot and tie it to a similar effect in UV spectra. In Figure~15 we suggest that, except for a temporally enlarged `bulge' or associated disturbance, the system would show a brightness peak at orbital phase 0.9 with a brightness reduction to the light level seen at phase 0.15. The same `dips' are seen in LMXRBs, discussed by \\citet{livio1993} and with the same proposed explanation. This subject is discussed further by \\citet{knigge1998a} and FR2003. This explanation for anomalous light curves differs from the \\citet{smak1994b} proposal for circumdisk absorbing material.  An important result of this study is that, starting with an adopted distance of 250pc to UX UMa, intensity-based observed spectra, flux-based observed spectra, and photometric data all can be approximately represented by a standard model with ${\\dot{M}}=5.0{\\times}10^{-9}M_{\\odot}~{\\rm yr}^{-1}$ and with an elevated temperature in the accretion disk quadrant that includes the stream impact region. Application of the \\citet{knigge2006} method then leads to an improved distance estimate of 312pc followed by a revised estimate of the (average) mass transfer rate (at the times of the DS1 and DS2 data sets). In spite of this success, SED fits show departures from the standard model in agreement with BINSYN studies of IX Vel \\citep{linn2007} and QU Car \\citep{linnell2008} which also exhibit departures from a standard model. We postpone attempting to model the IC and the recombination spectrum that fills in the Balmer continuum for the `back' spectra, Figure~4c to Figure~4j, to a subsequent publication.  We thank the referee for a prompt report; responding to it substantially improved this paper. We are grateful to Dr. Baptista for supplying data from BA1998 for this investigation. PG wishes to thank Mario Livio for his kind hospitality at the Space Telescope Science Institute where part of thi work was carried out. Support for this work was provided by NASA through grant number HST-AR-10657.01-A  to Villanova University (P. Godon) from the Space Telescope Science Institute, which is operated by the Association of Universities for Research in Astronomy, Incorporated, under NASA contact NAS5-26555. PS is supported by HST grant GO-09724.06A.  This research was partly based on observations made with the NASA/ESA Hubble Space Telescope, obtained at the Space Telescope Science Institute, which is operated by the Association of Universities for Research in Astronomy, Inc. under NASA contract NAS5-26555, and the NASA-CNES-CSA {\\it Far Ultraviolet Explorer}, which is operated for NASA by the Johns Hopkins University under NASA contract NAS5-32985."
    },
    "0807/0807.1264_arXiv.txt": {
        "abstract": " ",
        "introduction": "The collisional process plays a key role in all stages of a planetary system history, from the phase of planetary formation through collisional accretion to the late phases where the populations of small bodies evolve collisionally in a disruptive way. So far, these problems have been addressed using models appropriate to solid bodies represented as brittle rocky materials in which porosity was neglected or modeled at a macroscopic (i.e. resolved) level. Numerical codes, called {\\it hydrocodes} (see, e.g., Benz and Asphaug, 1994) have been developed to compute the impact induced fragmentation of such solid bodies by solving the elastic-plastic conservation equations associated with a model of brittle failure to account for the fracture of the solid material. This already allowed to improve our understanding of the impact response of small bodies such as basalt-like material and asteroids belonging to the taxonomic class S, supposed to be composed of material with negligible porosity. However, several evidence point to the presence of a high degree of porosity in some small body classes. Asteroids belonging to the C taxonomic class are now believed to be highly porous, as indicated by the low bulk density ($\\approx 1.3$ g$/$cm$^3$) estimated for some of them, such as the asteroid 253 Mathilde encountered by the NEAR Shoemaker spacecraft (Yeomans et al., 1997), and as inferred from meteorite analysis (Britt et al., 2006). Small body populations evolving at larger distances from the Sun, i.e. the Jupiter-Family Comets, the Kuiper Belt Objects, and the other classes of comets contain also a high level of porosity, as indicated by the estimated low bulk densities (below $1$ g$/$cm$^3$, e.g., Rickman, 1998) and the analysis of interplanetary dust particles collected on Earth.  In parallel, the dissipative properties of porous media are invoked more and more as playing an important role in the formation of early planetesimals (e.g., Wurm et al., 2005). Hence, porosity emerges slowly as playing a major role from the time of the formation of the planets to the collisional evolution of the present day Solar System.  Despite the growing focus on porosity, our ability to model its effect on the outcome of impacts and collisions remains limited. To address the question of modeling porosity in this context it is necessary to first define its scale in comparison with the other relevant dimensions involved in the problem such as the size of the projectile and/or the crater, etc.  In particular we define microscopic porosity as a type of porosity characterized by pores sufficiently small that one can reasonably assume that they are distributed uniformly and isotropically over these relevant scales. Macroscopic porosity on the other hand would be characterized by pores with sizes such that the medium can no longer be assumed to have homogeneous and isotropic characteristics over the scales of interest. In this case, pores have to be modeled explicitly and the current hydrocodes developed for the modeling of non-porous brittle solids  can still be used. The presence of these large macroscopic voids will only affect the transfer efficiency and the geometry of the shock wave resulting from the impact which can be computed using this existing software. On the other hand, a body containing microporosity (i.e. porosity on a scale much smaller than can be numerically resolved) can be crushable: cratering on a microporous asteroid might be an event involving compaction rather than ejection (Housen et al., 1999). In an impact in microporous material, a part of kinetic energy is dissipated by compaction which leads to less ejection and lower velocities of the ejected material. These effects cannot be reproduced using hydrocodes developed for the modeling of non-porous solids. Therefore, a model is needed which takes pore compaction into account.  Sirono (2004) proposed a model which can be used to study low velocity collisions of porous aggregates. Using this model, the author showed that the energy dissipation by compaction can lead to the sticking of dust aggregates. However, this model is not appropriate for high velocity impacts because of the simplicity of the equation of state (by construction, the pressure only depends on density).  Recently, Wuennemann et al. (2006) proposed a so-called $\\epsilon-alpha$ model suitable to model porous material and for the use in hydrocodes. In this paper, we propose an alternative approach based on the so-called $P-alpha$ model (Herrmann, 1969) which we found to be more appropriate to use with our numerical method (see Sec.~\\ref{sec:altmod}). Our 3D Smooth Particle Hydrodynamics (SPH) code, originally developed by Benz and Asphaug (1994), includes a model of brittle failure of non-porous material and successfully reproduced impact experiments on non-porous basalt targets. Moreover, associated to the N-body code {\\it pkdgrav} to account for the effect of gravity (see e.g., Richardson et al., 2000), it reproduced successfully for the first time the formation of S-type asteroid families resulting from the catastrophic disruption of non-porous parent bodies (see, e.g., Michel et al., 2001; Michel et al. 2003). A first implementation of a porosity model has been made in this code, by Benz and Jutzi (2006) using a simplified version of the $P-alpha$ model (the density $\\rho$ was used instead of the pressure $P$) but was found to be inappropriate for materials with a high degree of porosity (see Sec.~\\ref{sec:highpor}). Here, we improve on this first work by implementing the full $P-alpha$ model which offers several advantages and has a more general application.  In the following, we begin by describing our porosity model and comparing it to alternative models. We discuss some (counter intuitive) effects which occur when dealing with highly porous material. In the case of a high porosity, the energy dissipation during compaction leads to a large thermal pressure which can even become large enough to cause a decrease of the matrix density (Sec.~\\ref{sec:highpor}). We show that our model is suitable to model these effects. In Sec.~\\ref{sec:nummod}, we recall the equations used in our modeling of brittle solids and then we show how porosity is related to these equations (Sec.~\\ref{sec:distsolid}). In Sec.~\\ref{sec:tests}, simple test cases are presented: first, a simulation of a 1D shock wave in porous media is compared to the analytical solution; in the second test, we simulate the compression of porous pumice and measure the so-called crush-curve. As a first application, we perform a simulation of a Deep Impact-like impact and compare simulations with and without porosity model. We also compare different compaction models using a porous basalt target. Conclusions are then exposed in Sec.~\\ref{sec:concl}. A detailed comparison of numerical simulations with high-velocity impact experiments on porous material is presented in a forthcoming paper (Jutzi et al., in preparation). ",
        "conclusions": "\\label{sec:concl}  In this paper, we have presented a new approach to model sub-resolution porosity in brittle solids that can be coupled to a 3D SPH hydrocode in order to simulate impacts and collisions involving porous bodies. Such bodies are believed to be present in all the populations of small bodies in our Solar System. Therefore, understanding their impact response is crucial to determine the collisional evolution of those populations during all stages, and to assess the accretion efficiency of small bodies during planetary formation. Moreover, this can help defining efficient mitigation strategies against the impact of a porous body on Earth.  In practice, the implementation of our model does not consume excessive CPU time and is easily implemented in a parallel code (porosity is a local property) so that simulations involving multi-million particles can readily be performed. In fact, extensive testing has shown that high resolution is really needed to obtain converged solutions in the case of simulations involving fracturing and/or porosity.  We presented two test cases. In the first test, we compared the simulation of a 1D compaction wave in porous material with the theoretical solution given by the Hugoniot equations. The amplitude and velocity of the resulting waves (elastic precursor and shock) in the simulation agree with the theoretical solution. The second test is a simulation of an experiment consisting in compressing a porous sample. We show that an advantage of the $P-alpha$ model is that an experimentally measured crush-curve can be directly used to define the relation $\\alpha(P)$. Therefore, if a crush-curve has been measured for a given material, we can then simulate an impact on this material using its very crush-curve and not an arbitrary one which may lead to different outcomes.  As an application of our model, we presented the simulation of a Deep Impact-like impact. The main conclusion of this first application is that to model the behavior of a porous material during an impact, it is not enough to use the small bulk density expected for a porous material and compute its response from an impact by using a classical model of brittle failure of non-porous rock material, as one might be tempted to do. Actually, the cratering or disruption of porous material involves different processes than the ones involved in a non-porous brittle material (e.g. the crushing of pores), and this makes a huge difference in the outcome.  We also investigated the compaction behavior in an impact in basalt using different compaction models. We showed that using the $\\rho-alpha$ or $\\epsilon-alpha$ models (without modifications) can lead to difficulties to reach full compaction ($\\alpha=1$) in the case of highly porous material. Using these models, the compaction behavior depends much on the EOS parameters which relate energy and pressure. On the other hand, the $P-alpha$ model can be used without modifications to simulate the compaction of highly porous material and the compaction behavior is not sensitive to these EOS parameters.  Our next step will be to validate our code in more details and in a context adapted to its future applications by confronting it with real high-velocity impact experiments on porous targets. Such a validation will mainly consist of reproducing the size distribution of the fragments (and their velocities when they are measured), using as inputs all measured material parameters. Indeed, it is important that matching the results does not rely on a fine tuning of all parameters but rather on the validity of the physical model. This will be the subject of a next paper. Comparisons with impact experiments in laboratory will have to be performed in the long run for a wide variety of materials (porous and non porous) to improve our understanding of the impact process as a function of material properties and impact conditions. The formation and evolution of planetary systems involves bodies with a wide range of properties and colliding with each other in different regimes of impact energies, leading either to their accretion and the formation of planets, or to their disruption as in the current stage of our Solar System. It is then important to have the possibility to simulate the collisional process in these different regimes and between bodies with different degrees of porosity. Collisional evolution models, which need constraints to characterize the outcomes of collisions, will certainly benefit from these investigations.  \\newpage"
    },
    "0807/0807.1787_arXiv.txt": {
        "abstract": "Double peaked broad emission lines in active galactic nuclei are generally considered to be formed in an accretion disc. In this paper, we compute the profiles of reprocessing emission lines from a relativistic, warped accretion disc around a black hole in order to explore the possibility that certain asymmetries in the double-peaked emission line profile which can not be explained by a circular Keplerian disc may be induced by disc warping. The disc warping also provides a solution for the energy budget in the emission line region because it increases the solid angle of the outer disc portion subtended to the inner portion of the disc. We adopted a parametrized disc geometry and a central point-like source of ionizing radiation to capture the main characteristics of the emission line profile from such discs. We find that the ratio between the blue and red peaks of the line profiles becoming less than unity can be naturally predicted by a twisted warped disc, and a third peak can be produced in some cases.  We show that disc warping can reproduce the main features of multi-peaked line profiles of four active galactic nuclei from the Sloan Digital Sky Survey. ",
        "introduction": "\\label{intro}  A small fraction of active galactic nuclei~(AGN) show double-peaked broad emission line profiles \\citep*{era94,era03,str03,gez07}. The possibility has been considered for a long time that at least some of these lines arise directly from the accretion discs assumed to feed the central supermassive black holes. The H$\\alpha$ profile observed in the spectrum of Arp 102B was modelled as arising in a relativistic disc by \\citet{che89} and \\citet*{chen89} in an attempt to explain the asymmetric double-peaked profiles. In these stationary circular relativistic disc models, Doppler boosting will make the blue peak of the profile higher than the red one. However, \\citet{mil90} questioned the relativistic disc explanation for Arp 102B, by showing that at least at some epochs the profile asymmetry is reversed, with the red peak higher than the blue one in contrary to the prediction of homogeneous, circular relativistic disc models. This type of profile asymmetry has been observed in some other double-peaked sources \\citep{str03,gez07}. Thus, other scenarios for the profile asymmetry were considered later on, including various physically plausible processes in the accretion flow, such as spiral shocks, eccentric discs, bipolar outflows \\citep{cha93,cha94,era95,sto97,zhe90,vei91}, as well as  a binary black hole system \\citep{beg80,gas96,zha07b}.  In addition to the line profile asymmetry, the energetic budget for the line emission region has not been solved so far for the disc model.  It is demonstrated that the local viscous heating in the accretion disc is not sufficient to explain the observed H$\\alpha$ intensity \\citep{chen89,era97}, thus reprocessing of ultraviolet(UV)/X-ray continuum from the inner accretion disc is required. However, only a very small fraction of radiation from inner disc is expected to intercept the outer part of the disc in the case of a flat geometrically thin disc. Two type of processes have been proposed to increase the fraction of light incident on to the disc emission-line region.  \\citet{chen89} proposed that the inner part  of the accretion disc is hot and becomes geometrically thick due to insufficient radiative cooling, and the X-ray emission from such an ion-supported torus is responsible for such energy input.  However, double peak emitters do not always have a low accretion rate \\citep{zha07A,bian07}.  Furthermore, a detailed calculation of advection dominated accretion flow (ADAF) models suggests that only a small portion of X-rays actually hit the emission-line portion of the disc \\citep{cao06}. On the other hand, \\citet{cao06} proposed that a slow moving jet can scatter a substantial fraction of UV/X-ray light from the inner accretion disc back to the outer part of the disc.  While this may be the likely case for radio-loud double-peaked emitters, the majority of double-peaked emitters are radio quiet.  On the other hand, warped accretion discs are believed to exist in a variety of astrophysical systems, including X-ray binaries, T Tauri stars, and AGN. The dynamics of warping accretion discs has been discussed by a number of authors \\citep{pap83,pal95,ogi99,nel99,nel00,nay05}. Four main mechanisms for exciting/maintaining warping in accretion discs have been proposed, including: tidally induced warping by a companion in a binary system \\citep{teb93,teb96,lar96}, radiation driven or self-inducing warping \\citep{pri96,pri97,mal96,mal97,mal98}, magnetically driven disc warping \\citep{lai99,lai03,pfl04}, and frame dragging driven warping \\citep{bap75,kum85,arm99}. As a result, the accretion disc in some AGN should be nonplanar. Evidence for the existence of warped discs in AGN has been found by observations \\citep{her05,cap07}. A warped disc makes it possible for the radiation from the inner disc to reach the outer parts, enhancing the reprocessing emission lines. Thus, this warping can solve the long-standing energy-budget problem because the subtending angle of the outer disc portion to the inner part has been increased by the warp.  It can also provide the asymmetry required for the variations of the emission lines. The effect of disc warping on iron line profiles treated relativistically have been made by \\citet{har00} and \\citet{cad03}. \\citet{bac99} studied the broad-line H$\\beta$ profiles from a warped disc. Nevertheless, his treatment is non-relativistic. A relativistic treatment of warped disc on the Balmer lines which can be compared with observations is still not available in the literature.  In this paper, the broad Balmer emission due to reprocessing of the central high-energy radiation by a warped accretion disc is investigated. Our main purpose here is to describe the effect of disc warping on the line profiles. A parametrized disc model is adopted. In \\S\\ref{method} we summarize the assumptions behind our model, and present the basic equations relevant to our problem. We present our results and the comparison with observations in \\S\\ref{result}, and conclusion and discussion in \\S\\ref{sum}. ",
        "conclusions": "\\label{sum} We compute the reprocessing emission-line profiles from a warped relativistic accretion disc around a central black hole by including all relativistic effects.  A parametrized disc model is used to obtain an insight into the impact of disc warping on the double-peaked Balmer emission-line profiles. For simplicity, we assumed that the disc is illuminated by a point-like central source, which is a good approximation, the line emissivity is proportional to the continuum light intercepted by the accretion disc, and line emission is isotropic.   We find the following. \\begin{enumerate} \\item For twist-free or low phase amplitude~(described by $n_1$) warped disc, the asymmetrical and frequency-shifted single-peaked line profiles as expected are produced in most cases. The rarity of such sources suggests that a warping disc is usually twisted. \\item The flux ratio of the blue peak to red one becoming less than unity can be predicted by a twisted warp disc with high phase amplitude. The phase amplitude has a significant influence on the line profiles. \\item A third peak and the variation in blue/red peak position which have been found by observations \\citep{vei91,str03} are produced in our model. \\item The influence from retrograde disc on line profiles has a inverse effect compared with a prograde one. \\item The fraction of the radiation incident to the outer disc from the inner part can be enhanced by disc warping. The results are shown in Fig.~\\ref{sang}. \\end{enumerate}  We showed that warped disc model is flexible enough to reproduce a variety of line profiles including triplet-peaked line profiles, or double peaked profiles with additional plateaus in the line wing as observed in the SDSS spectra of some AGN, while the models have the same number of free parameters as eccentric disc models. While we leave the detailed fit to a future paper, the essential characteristics in the line profiles is in good agreement with observed line profiles. Future monitoring of line profile variability in these triplet sources can provide a critical test of the warped disc model."
    },
    "0807/0807.3762.txt": {
        "abstract": "Asteroid families, traditionally defined as clusters of objects in orbital parameter space, often have distinctive optical colors. We show that the separation of family members from background interlopers can be improved with the aid of SDSS colors as a qualifier for family membership.  Based on an $\\sim$88,000 object subset of the Sloan Digital Sky Survey Moving Object Catalog 4 with available proper orbital elements, we define 37 statistically robust asteroid families with at least 100 members (12 families have over 1000 members) using a simple Gaussian distribution model in both orbital and color space. The interloper rejection rate based on colors is typically $\\sim$10\\% for a given orbital family definition, with four families that can be reliably isolated only with the aid of colors. About 50\\% of all objects in this data set belong to families, and this fraction varies from about 35\\% for objects brighter than an $H$ magnitude of 13 and rises to 60\\% for objects fainter than this. The fraction of C-type objects in families decreases with increasing $H$ magnitude for $H > 13$, while the fraction of S-type objects above this limit remains effectively constant. This suggests that S-type objects require a shorter timescale for  equilibrating the background and family size distributions via collisional processing. The size distribution varies significantly among  families, and is typically different from size distributions for background populations. The size distributions for 15 families display a well-defined change of slope and can be modeled as a ``broken'' double power-law. Such ``broken'' size distributions are twice as likely for S-type familes than for C-type families (73\\% vs. 36\\%), and are dominated by dynamically old families. The remaining families with size distributions that can be modeled as a single power law are dominated by young families ($<$1 Gyr). When size distribution requires a double power-law model, the two slopes are correlated and are steeper for S-type families. No such slope--color correlation is discernible for families whose size distribution follows a single power law. For several very populous families, we find that the size distribution varies with the distance from the core in orbital-color space, such that small objects are more prevalent in the family outskirts. This ``size sorting'' is consistent with predictions based on the Yarkovsky effect. ",
        "introduction": "The size distribution of asteroids is one of most significant observational constraints on their history and is considered to be the ``planetary holy grail'' (Jedicke \\& Metcalfe 1998, and references therein). It is also one of the hardest quantities to determine observationally because of strong selection effects. Recently, Ivezi\\'{c} et al. (2001, hereafter I01) determined the asteroid size distribution to a sub-km limit using multi-color photometry obtained by the Sloan Digital Sky Survey (York et al. 2000; hereafter SDSS). Here we extend their work by using an updated ($4^{th}$) version of the SDSS Moving Object Catalog (Ivezi\\'{c} et al. 2002a, hereafter I02a).  The main goal of this paper is to study size distributions of asteroid families. Asteroid dynamical families are groups of asteroids in orbital element space (Gradie, Chapman \\& Williams 1979, Gradie, Chapman \\& Tedesco 1989, Valsecchi {\\em et al.} 1989). This clustering was first discovered by Hirayama (1918, for a review see Binzel 1994), who also proposed that families may be the remnants of parent bodies that broke into fragments. About half of all known asteroids are believed to belong to families; Zappal\\'{a} {\\em et al.} 1995 (hereafter Z95), applying a hierarchical clustering method to a sample of 12,487 asteroids, finds over 30 families. Using the same method and a larger sample of $\\sim$106,000 objects, Nesvorn\\'{y} et al. (2005, hereafter N05) identify $\\sim$50 statistically robust asteroid families.  The size distributions of asteroid families encode information about their formation and evolution, and constrain the properties of the families' parent bodies (e.g., Marzari, Farinella \\& Davis 1999; Tanga et al. 1999; Campo Bagatin \\& Petit 2001; Michel et al. 2002; de Elia \\& Brunini 2007; Durda et al. 2007; and references therein). Motivated by this rich information content, as well as the availability of new massive datasets, here we address the following questions \\begin{enumerate} \\item What is the fraction of objects associated with families? \\item Do objects that are not associated with families show any heliocentric color gradient? \\item Do objects that are not associated with families have uniform size distribution independent of heliocentric distance? \\item Do objects associated with families have a different size distribution than those that are not in families? \\item Do different families have similar size distributions? \\item Is the size distribution related to family color and age? \\end{enumerate}  These questions have already been addressed numerous times (e.g. Mikami \\& Ishida 1990; Cellino, Zappal\\'{a} \\& Farinella 1991; Marzari, Davis \\& Vanzani 1995; Z95; Morbidelli et al. 2003; N05). The main advantages of the size distribution analysis presented here, when compared to previous work, are \\begin{itemize} \\item The large sample size: we use a set of $\\sim$88,000 objects for which both SDSS colors and proper orbital elements computed by Milani \\& Kne\\v{z}evi\\'{c} (1994) are available \\item Simple and well-understood selection effects: the SDSS sample is $>$90\\% complete without a strong dependence on magnitude (Juri\\'{c} et al. 2002, hereafter J02) \\item Improved faint limit: the sample of known objects listed in the latest ASTORB \u00dele from January 2008 to which SDSS observations are matched is now essentially complete to $r\\sim19.5$ (corresponding to $H\\sim17$ in the inner belt, and to $H\\sim15$ in the outer belt) \\item Improved family definitions due to color constraints (rejection of interlopers and separation of families overlapping in orbital space) \\item Improved accuracy of absolute magnitudes derived using SDSS photometry, as described below. \\end{itemize}  The SDSS asteroid data are described in Section 2, and in Section 3 we describe a novel method for defining asteroid families using both orbital parameters and colors. Analysis of the size distribution for families and background objects is presented in Section 4, and we summarize our results in Section 5. ",
        "conclusions": "We have used a large sample of asteroids ($\\sim$88,000) for which both orbital elements and SDSS colors are available to derive improved membership for 37 asteroid families. The addition of colors typically rejects about 10\\% of all dynamically identified candidate members due to mismatched colors. Four families can be reliably isolated only with the aid of colors. About 50\\% of objects in this data set belong to families, with this fraction representing a lower limit due to a conservative requirement that a candidate family must include at least 100 members. The resulting family definitions are in good agreement with previous work (e.g. Z95, N05) and all the discrepancies are well understood. Although SDSS has observed only about 1/3 of all known asteroids, it is remarkable that the sample discussed here provides color information for more than an order of magnitude more objects associated with families than analyzed in the published literature.  This data set enables the determination of absolute magnitude (size) distributions for individual families to a very faint limit without a need to account for complex selection effects. We verify that size distribution varies significantly among families, and is typically different from size distributions for background populations. Consequently, the asteroid size distribution cannot be described by a universal function that is valid throughout the main belt (e.g. Jedicke \\& Metcalfe 1998, Ivezi\\'{c} et al. 2001, and reference therein). This finding will have an influence on conclusions derived from modeling the size distribution under this assumption (e.g. Bottke et al. 2005, and references therein). In particular, it is not clear how to interpret a detailed dependence of the critical specific energy (energy per unit mass required to fragment an asteroid and disperse the fragments to infinity) on asteroid size derived from such models, when the starting observational constraint on size distribution is an average over multiple families with significantly varying size distributions.  We show that for objects with $H_{corr} < 13$, the background population dominates (family fraction decreases toward lower $H_{corr}$, indicating a shallower size distribution for large objects), while for objects with $H_{corr} > 13$, the red family fraction remains effectively constant to our completeness limit while the blue family fraction falls off. This indicates that the time to collisionally equilibrate the family and background populations (see e.g., Morbidelli et al. 2003) is shorter for red objects than blue.  The size distributions for 15 families display a well-defined change of slope and can be modeled as a ``broken'' double power-law. The first evidence for this effect and a discussion of its significance are presented by Morbidelli et al. (2003). Using a data set with much simpler correction for the observational selection effects, we confirm their result in a statistically more robust way. We also find such ``broken'' size distributions are twice as likely for S-type familes than for C-type families (73\\% vs. 36\\%), and are dominated by dynamically old families. The remaining families with size distributions that can be modeled as a single power law are dominated by young families ($<$1 Gyr).  The eight largest families all show a change of size distribution slope to much smaller values at the faint end (see Table 5). This result has a direct consequence when prediciting the number of very small objects ($D\\sim1$ km). In particular, it could explain why the Statistical Asteroid Model developed by Tedesco, Cellino \\& Zappal\\'{a} (2005) predicts too many objects: the data presented here are inconsistent with the SAM assumptions for the number of objects in its most populous families such as Eunomia and Themis.  We find that when size distribution requires a double power-law model, the two slopes are correlated and are steeper for S-type families. No such slope--color correlation is discernible for families whose size distribution follows a single power law. While beyond the scope of this work, the modeling of such correlations may shed light on the internal structure and material properties of asteroids.  For several very populous families, we find that the size distribution varies with the distance from the core in orbital-color space, such that small objects are more prevalent in the family outskirts. As discussed by N05 (and references therein), this ``size sorting'' is consistent with predictions based on the Yarkovsky/YORP effect.  While these results provide significant new observational constraints for the properties of main-belt asteroids, very soon the observations will further improve. The upcoming large-scale sky surveys, such as Pan-STARRS (Kaiser et al. 2002) and LSST (Tyson 2002), will obtain even more impressive samples, both in size, diversity of measurements and their accuracy. For example, LSST will scan the whole observable sky every three nights in two bands to a $5\\sigma$ depth equivalent to $V=24.7$. These data will enable much improved analysis due to several factors \\begin{itemize} \\item Due to hundreds of observations, the orbits will be determined directly, instead of relying on external data, resulting in a sample about 30-40 times larger than discussed here \\item The effective faint limit will be extended by about 5 magnitudes, correspoding to ten times smaller size limit (diameters of several hundred meters) \\item Due to many photometric observations obtained with the same well-calibrated system, the uncertainties in absolute magnitudes will be an order of magnitude smaller \\item The addition of the $y$ band (at $\\sim$1 $\\mu$m) will improve the color classification due to better sensitivity to the $\\sim$1 $\\mu$m absorption feature present in spectra of many asteroids. \\end{itemize} These new data will undoubtely reinvigorate both observational and theoretical studies of main-belt asteroids."
    },
    "0807/0807.2447_arXiv.txt": {
        "abstract": "We use a combination of analytic and numerical arguments to consider the impact of quasar photo-heating during \\HeII reionisation on the thermal evolution of the intergalactic medium (IGM).  We demonstrate that rapid ($\\Delta z< 0.1$--$0.2$), strong ($\\Delta T > 10^4\\rm~K$) photo-heating is difficult to achieve across the entire IGM unless quasar spectra are significantly harder than implied by current observational constraints.  Although filtering of intrinsic quasar radiation through dense regions in the IGM does increase the mean excess energy per \\HeII photo-ionisation, it also weakens the radiation intensity and lowers the photo-ionisation rate, preventing rapid heating over time intervals shorter than the local photo-ionisation timescale.  Moreover, the hard photons responsible for the strongest heating are more likely to deposit their energy inside dense clumps, which cool rapidly and are furthermore invisible to most observational probes of the IGM temperature.  The abundance of such clumps is, however, uncertain and model-dependent, leading to a fairly large uncertainty in the photo-heating rates.  Nevertheless, although some of the IGM may be exposed to a hardened and weakened ionising background for long periods, most of the IGM must instead be reionised by the more abundant, softer photons and  with accordingly modest heating rates ($\\Delta T \\la 10^{4}\\rm~K$), although localised patches of much higher temperature are still possible.  The repeated ionisation of fossil quasar \\HeIII regions does not increase the net heating because the recombination times in these regions typically exceed the IGM cooling times and the average time lag between successive rounds of quasar activity.  Detailed line-of-sight radiative transfer simulations confirm these expectations and predict a rich thermal structure in the IGM during \\HeII reionisation.  The resulting complex relationship between temperature and density may help resolve discrepancies between optically thin simulations of the \\Lya forest and recent observations. ",
        "introduction": "Helium, the second most abundant chemical element in the Universe, is expected to be singly ionised at the same time as hydrogen due to its similar ionisation potential ($24.6\\rm~eV$), larger photo-ionisation cross-section and lower abundance relative to hydrogen.  Even for ionising sources with relatively soft, stellar like spectra, the reionisation of \\HI and \\HeI in the IGM is expected to be concurrent ({\\it e.g.}  \\citealt{GirouxShapiro96}).  In contrast, the much larger ionisation threshold ($54.4 \\rm~eV$) and smaller photo-ionisation cross-section of singly ionised helium requires that \\HeII reionisation proceeds only when sources with sufficiently hard spectra become numerous.    A putative population of massive, metal free stars (\\citealt{Oh01b,Venkatesan03}), high redshift X-rays (\\citealt{Oh01,Venkatesan01,RicottiOstriker04}) or thermal emission from shock heated gas (\\citealt{Miniati04}) may provide the requisite ionising photons.  The stacked spectra of Lyman break galaxies at $z\\sim 3$ (\\citealt{Shapley03}) also exhibit an \\HeII recombination line strong enough to suggest that galaxies may contribute significantly to \\HeII reionisation (\\citealt{FurlanettoOh07b}). However, quasars are considered to be the most likely candidates for completing \\HeII reionisation (\\citealt{Madau99,MiraldaEscude00,Sokasian02,Bolton06,FurlanettoOh07b,Paschos07}), although bright quasars are rare at high redshift and unlikely to contribute substantially to either the \\HI or \\HeII ionising photon budget at $z>5$ (\\citealt{Dijkstra04,Meiksin05,Srbinovsky07,BoltonHaehnelt07b,Jiang08}). Consequently, unless \\HeII reionisation was completed by an as yet unidentified population of sources with hard spectra at $z>5$, its latter stages are expected to coincide with the observed peak in quasar activity around $z\\simeq 2-3$.  There are several pieces of observational evidence which may support this picture.  Observations of the \\HeII \\Lya Gunn-Peterson trough and patchy absorption in the \\HeII \\Lya forest at $z \\simeq 3$ indicate a situation analogous to that implied by the \\HI opacity data at $z\\simeq 6$ (\\citealt{Jakobsen94,Davidsen96,Heap00,Smette02,Zheng04,Shull04,Reimers05b,Fechner06}). However, like the \\HI measurements at higher redshift ({\\it e.g.} \\citealt{Fan06,Becker07}), the interpretation of the \\HeII data is hampered by its insensitivity to the \\HeII fraction once the \\Lya absorption becomes saturated.  Other indirect evidence includes an apparent boost to the IGM temperature at $z\\simeq 3.3$ (\\citealt{Schaye00,Ricotti00}), measurements of metal line ratios which indicate the ultraviolet (UV) background is hardening around the same redshifts (\\citealt{Songaila98,Vladilo03}), and the detection of a sharp dip in the \\HI \\Lya forest effective optical depth at $z\\simeq 3.2$ (\\citealt{Bernardi03,Faucher08}).  The latter result may be associated with the observed boost in the IGM temperature at the same redshifts (\\citealt{Theuns02d}). On the other hand, conflicting evidence exists for all of these cases ({\\it e.g.} \\citealt{Zaldarriaga01,Kim02,McDonald01,McDonald05}).  Thus, the observational evidence for \\HeII reionisation remains controversial.   Nevertheless, from a theoretical standpoint it is still expected \\HeII reionisation will have a significant impact on the thermal state of the IGM (\\citealt{MiraldaRees94,AbelHaehnelt99,Theuns02d,Paschos07,FurlanettoOh07}). Firstly, energy injected by the hard sources which reionise \\HeII will raise the IGM temperature. This is consistent with some of the observational constraints derived from \\Lya forest line widths which imply a boost of $\\Delta T \\simeq 1$--$2\\times 10^{4}\\rm~K$ around $z\\sim 3$ (\\citealt{Schaye00,Ricotti00}). The magnitude of the increase will depend on the spectrum of the intrinsic ionising radiation, and thus also on the amount it is subsequently filtered by the intervening IGM (\\citealt{AbelHaehnelt99}).  Secondly, the reionisation of \\HeII will alter the temperature-density relation of the IGM.  In the optically thin limit, the IGM will follow a tight polytropic temperature-density relation, $T=T_{0}\\Delta^{\\gamma-1}$, expected to arise for normalised densities $\\Delta = \\rho/\\langle \\rho \\rangle \\leq 10$ when photo-heating by a spatially uniform UV background is balanced by cooling due to adiabatic expansion (\\citealt{HuiGnedin97,Valageas02}). However, the assumption of a spatially uniform UV background ionising an optically thin IGM will break down during \\HeII reionisation.  Instead, radiative transfer effects and/or the inhomogeneous distribution of the ionising sources will produce a complex, multi-valued and perhaps even inverted ($\\gamma<1$) temperature-density relation (\\citealt{Gleser05,Tittley07,FurlanettoOh07,Bolton08}).   Despite all of these studies, very little work has been done to establish a physical argument for exactly how large the temperature boost in the IGM during \\HeII reionisation will be.  Furthermore, it is unclear over what timescale any global temperature boost should occur.  This depends on understanding how photo-heating in an optically thick IGM proceeds, and thus on the details of radiative transfer through the inhomogeneous IGM. The observational evidence provides valuable hints, but the data are still too poorly constrained to distinguish between different models in detail.  A wide variety of values for the temperature boost during \\HeII reionisation, $5~000 \\rm~K < \\Delta T < 30~000\\rm~K$, have therefore been adopted in the literature, with the timescales for the temperature boost varying from instantaneous to periods spanning several redshift units ({\\it e.g.} \\citealt{AbelHaehnelt99,Theuns02d,Inoue03,HuiHaiman03,FurlanettoOh07}). The main reason for this is that photo-heating in an optically thick IGM during \\HeII reionisation is almost always modelled in an approximate way.   Only a handful of studies have explicitly quantified the expected IGM temperature boost during \\HeII reionisation using radiative transfer calculations. \\cite{AbelHaehnelt99} were the first to point out the importance of radiative transfer through an optically thick IGM.  Using one dimensional radiative transfer simulations, they found the filtering of ionising radiation through the IGM could increase \\HeII photo-heating rates by a factor of $1.5$--$2.5$ above the optically thin values, leading to $\\Delta T \\sim 10^{4}\\rm~K$.  More recently, \\cite{Tittley07} demonstrated that the temperature boost depends sensitively on the typical spectra of the sources which drive \\HeII reionisation.  Finally, \\cite{Paschos07} have presented a detailed, three dimensional radiative transfer simulation of \\HeII reionisation. They found that radiative transfer effects could boost the IGM temperature by around a factor of two.  However, none of these studies explicitly discuss over what timescale this heat input may reasonably occur, or how the temperature boost depends in detail on the density distribution of the intervening IGM.  A clear understanding of photo-heating during \\HeII reionisation is thus vital for interpreting the thermal history of the IGM.  In this paper we critically re-examine the photo-heating of the IGM by quasars during \\HeII reionisation.  In particular, using a combination of analytical arguments and line-of-sight radiative transfer calculations, we shall address what the likely temperature boost in the IGM during \\HeII reionisation is, and how rapidly this temperature boost should occur.   Furthermore, we consider what the fate of the hard photons responsible for heating the IGM is likely to be: where do they end up in the IGM, and how important is the filtering of ionising radiation through the IGM in setting the temperature during \\HeII reionisation?  We begin in \\S2, where we  briefly outline the relationship between \\HeII photo-heating and the photo-ionisation timescale.  This argument is key to establishing the amount of heat which may be injected into the IGM within a given time.  In \\S3 we present a model for estimating the maximum temperature achievable during \\HeII reionisation in an optically thick IGM.  A more detailed treatment of the filtering of \\HeII ionising photons through a clumpy medium is presented in \\S4, and in \\S5 we consider the impact of \\HeII photo-heating in recombining fossil \\HeIII regions.  We then compare our analytical results with detailed line-of-sight radiative transfer simulations in \\S6, and use these results to discuss the impact of \\HeII reionisation on the IGM temperature density relation in \\S7.  We summarise our results, compare them to current observational constraints and conclude in \\S8.  Throughout this paper we shall assume $\\Omega_{\\rm m}=0.26$, $\\Omega_{\\Lambda}=0.74$, $\\Omega_{\\rm b}h^{2}=0.024$, $h=0.72$, $\\sigma_{8}=0.85$, $n_{\\rm s}=0.95$ and an IGM of primordial composition with a helium mass fraction of $Y=0.24$ ({\\it e.g.} \\citealt{OliveSkillman04}). The cosmological parameters are consistent with the most recent WMAP results (\\citealt{Dunkley08}) aside from a slightly larger normalisation for the power spectrum.  Unless otherwise stated, all distances are expressed as comoving quantities. ",
        "conclusions": "In this work we have performed a detailed analysis of the photo-heating of the IGM during \\HeII reionisation using a combination of analytical arguments and numerical radiative transfer models.  Our aim was to critically examine the physical processes involved in modelling the fate of hard, \\HeII ionising photons in the IGM and their impact on the thermal evolution of the IGM.  The main results of this work are as follows.  \\begin{enumerate}  \\item{Filtering through the IGM weakens as well as hardens ionising radiation (\\citealt{AbelHaehnelt99,Bolton04}).  However, significant photo-heating of the IGM during \\HeII reionisation can only occur over timescales comparable to or shorter than the local \\HeII photo-ionisation timescale,  $t_{\\rm ion} = \\Gamma_{\\rm HeII}^{-1}$, irrespective of the mean excess energy per \\HeII photo-ionisation, $\\langle E \\rangle_{\\rm HeII}$.  Consequently, this limits the temperature boost, $\\Delta T$, attainable over a given timescale within a fixed volume of the IGM.   For an $L_{*}$ quasar at $z=3$ with a typical EUV spectra index of $\\alpha_{\\rm s}=1.5$ and a lifetime $t_{\\rm s} = 10^{8} \\rm~yrs$, $\\Delta T$ will not exceed $10^{4}\\rm~K$ at distances $R>40\\rm~Mpc$ from the quasar, and will furthermore only do so at $R<40\\rm~Mpc$ if the softer photons near the \\HeII ionisation edge do not contribute towards photo-ionisation.  For an average separation between $L_{*}$ quasars of $\\sim 100\\rm~Mpc$, this implies a global temperature boost of $\\Delta T \\ga 10^{4}\\rm~K$ over $\\Delta z=0.1-0.2$ is difficult to achieve unless quasars with much harder EUV spectra are common.}   \\item{Hard photons are preferentially absorbed in the densest regions of the IGM where the largest temperature boosts will thus occur. However, current constraints on the IGM temperature based on observations of the low column density \\Lya forest (\\citealt{Schaye00,Ricotti00,McDonald01}) are mainly sensitive to densities close to the cosmic mean or below (\\citealt{Schaye01}). The fate of hard photons and expected heating rates thus depends on the relative abundance of high column density $N_{\\rm HeII}$ absorbers: if most of the heat is deposited in dense systems, we have no way of inferring this from the \\Lya forest. Unfortunately, the abundance of such systems is rather model-dependent, and varies with the unknown radiation field $\\Gamma_{\\rm HeII}$ during reionisation and the correspondence between physical density $n_{\\rm H}$ and observed column densities $N_{\\rm HI}$. This translates into considerable uncertainties in the photo-heating rates and the resulting IGM temperature boost.  A plausible range of values, given model uncertainties, is $\\Delta T \\approx 7,000-15,000$K, somewhat intermediate between the optical thin and thick values of $\\Delta T \\approx 4,000$K and $\\Delta T \\approx 30,000$K respectively, for a $J_{\\nu} \\propto \\nu^{-1.5}$ unprocessed source spectrum.}  \\item{Even though the mean separation between \\HeII Lyman limit systems is comparable to the expected size of \\HeIII regions around $L_{*}$ quasars, most of the IGM could potentially still be ionised primarily by soft photons, at least if quasars are randomly distributed. If this were the case, the average amount of heating over the entire IGM would be relatively modest. In general, only unusual regions that are far from the ionising sources will be exposed to a highly-filtered spectrum for a long enough period to heat substantially.  This can only happen toward the end of  \\HeII reionisation, because the photo-ionisation timescale is so long for such heavily filtered spectra.}  \\item{The reheating of partially recombined fossil \\HeIII regions is unlikely to substantially increase the IGM temperature beyond that achieved during the first round of \\HeII reionisation.  In general, recombination timescales in these fossil regions typically exceed cooling timescales and the average time lag between ionisations by successive generations of quasars.  The net heating is therefore not increased.}  \\item{We run a detailed set of $270$ line-of-sight radiative transfer simulations using a variety of initial conditions to confirm these analytic predictions.  We further demonstrate that the filtering of ionising radiation through the IGM will produce a rich thermal structure.   This is likely to be further complicated by the inhomogeneous distribution of the ionising sources and the typically short lifetimes of quasars (\\citealt{FurlanettoOh07}).  This complex, multi-valued temperature-density relation may help to resolve discrepancies between optically thin hydrodynamical simulations of the \\Lya forest and recent detailed observations (\\citealt{Bolton08})}  \\end{enumerate}   How do these conclusions compare to the current observational constraints on the IGM thermal state at $z=3$?  In the case of the magnitude of the IGM temperature boost, measurements of the Doppler widths of absorption lines in the \\Lya forest currently offer the best constraints.  The data of \\cite{Schaye00} are consistent with a relatively sharp ($\\Delta z \\simeq 0.2$), large ($\\Delta T>10^{4}\\rm~K$) temperature jump occurring around $z=3.3$ at densities around the cosmic mean. However, the error bars on these measurements are still large, and a more gradual, moderate temperature boost, as suggested by this work, is also consistent with the data.  Similarly, the results of \\cite{Ricotti00} are also consistent with a significant and sudden temperature boost at $z\\sim 3$, although large statistical errors again leave a more gentle thermal evolution possible.  On the other hand, \\cite{McDonald01} find no evidence for a rapid boost in the IGM temperature at $z\\sim 3$.  Their constraints, corrected to correspond to measurements at mean density, at most allow only for a modest evolution in the IGM temperature ($\\Delta T \\sim 4~000\\rm~K$ between $z=2-4$), and are also consistent with no temperature evolution at all.   A detailed analysis of a large set of high resolution \\Lya forest spectra, perhaps using alternative temperature sensitive statistics, is necessary to obtain improved constraints on the thermal state of the IGM at these redshifts.   Another intriguing result which may indirectly favour a rapid heat input scenario, somewhat at odds with the results presented here, is the evidence for a sharp dip ($\\Delta z=0.2$) of around $10$ per cent in the effective optical depth of the \\Lya forest at $z=3.2$ (\\citealt{Bernardi03,Faucher08}).   \\cite{Theuns02d} used hydrodynamical simulations to interpret this feature as being due to a sudden increase in the IGM temperature at $z\\simeq 3.2$, such that $\\Delta T\\sim 10^{4}\\rm~K$.  However, our results suggest that it is unlikely this transition could occur globally over such a short period.  Alternatively, this feature may instead be due to a rapid change in the hydrogen photo-ionisation rate at $z=3.2$, although again there is no current evidence to suggest this may be the case. Further investigation into this feature using detailed hydrodynamical simulations of the IGM is thus desirable.  Our results therefore suggest a general picture for the thermal evolution of the IGM during \\HeII reionisation by quasars where: (i) the temperature boost during \\HeII reionisation  is significantly uncertain, and could be considerably less than has been commonly assumed in the optically thick heating limit, (ii) any such temperature boost must be achieved over a timescale longer than $\\Delta z >0.1-0.2$ and (iii) the resulting temperature-density relation in the IGM will be much more complex than the tight, power-law temperature density relation expected in an optically thin IGM.   Finally, we note that these results will not apply if the \\HeII in the IGM is predominantly reionised by sources other than quasars. If the \\HeII in the IGM is reionised early on by stellar sources, high redshift X-rays, thermal emission from shocked gas or some other, as yet unidentified source, then the thermal evolution of the IGM during \\HeII reionisation may be rather different."
    },
    "0807/0807.0442_arXiv.txt": {
        "abstract": "The frequencies of solar oscillations are known to change with solar activity.  We use Principal Component Analysis to examine these changes with high precision.  In addition to the well-documented changes in solar normal mode oscillations with activity as a function of frequency, which originate in the surface layers of the Sun, we find a small but statistically significant change in frequencies with an origin at and below the base of the convection zone.  We find that at $r=(0.712^{+0.0097}_{-0.0029})R_\\odot$, the change in sound speed is $\\delta c^2 / c^2 = (7.23 \\pm 2.08) \\times 10^{-5}$ between high and low activity. This change is very tightly correlated with solar activity.  In addition, we use the splitting coefficients to examine the latitudinal structure of these changes.  We find changes in sound speed correlated with surface activity for $r \\gtrsim 0.9R_\\odot$. ",
        "introduction": "\\label{sec:intro} Normal modes of oscillation of the Sun have provided a powerful tool to peer into the solar interior.  In particular, modern experiments, both ground- and space-based, have measured the intermediate degree global oscillation spectrum with high precision since the beginning of solar cycle 23.  Accurate determinations of interior structure and dynamics are now possible (see, e.g., review by \\citealt{C-D02}).  These measurements contain a wealth of information about the fundamental causes of solar variability.  It is generally believed that the seat of the solar dynamo is located at the base of the convection zone \\citep[e.g., review by][]{Charbonneau2005}. Because helioseismology provides the only direct measurements of this region of the solar interior, these results can play an important role in constraining dynamo theories.  In particular, a number of authors have attempted to use global and local helioseismic techniques to determine limits on the strength of the magnetic field at the base of the convection zone \\citep[e.g.,][and references therein]{Chouetal03}.  In this paper, we attempt to improve helioseismic measurements of changes in this region.  Global modes of solar oscillation are described by three numbers that characterize the spherical harmonics that are used to define the horizontal structure of the mode. These are (1) radial order $n$ that related to the number of nodes in the radial direction, (2) the degree $\\ell$ that is related to the horizontal wavelength of the mode, and (3) the azimuthal order $m$ that defines the number of nodes along the equator.  In a spherically symmetric star, the $2\\ell + 1$ modes of an $(n,\\ell)$ multiplet are degenerate, but effects that break spherical symmetry such as magnetic fields or rotation lift the degeneracy and results in frequency splittings. The frequencies $\\nu_{n\\ell m}$ of the modes within a multiplet can be expressed as an expansion in orthogonal polynomials: \\begin{equation} \\nu_{n\\ell m} = \\nu_{n\\ell} + \\sum_{j=1}^{j_{\\rm max}} a_j (n,\\ell) \\, {\\cal P}_j^{(\\ell)}(m). \\label{eq:acoefs} \\end{equation} Early investigators \\citep[e.g.,][]{Duvalletal86} commonly used Legendre polynomials, whereas now one often uses the Ritzwoller-Lavely formulation of the Clebsch-Gordan expansion \\citep{RL91} where the basis functions are polynomials related to the Clebsch-Gordan coefficients. In either case, the coefficients $a_j$ are referred to as $a$-coefficients or splitting coefficients.  Solar structure is determined by inverting the mean frequency $\\nu_{n\\ell}$, while the odd-order coefficients $a_1, a_3, \\ldots$ depend principally on the rotation rate \\citep{Durneyetal88} and reflect the advective, latitudinally symmetric part of the perturbations caused by rotation.  Hence, these are used to determine the rate of rotation inside the Sun. The even order $a$ coefficients on the other hand result from magnetic fields and asphericities in solar structure, and the second order effects of rotation \\citep[e.g.,][]{GT90,DG91}.  Solar oscillation frequencies are known to vary on timescales related to the solar activity cycle.  This was first suggested by \\citet{WN85} and confirmed soon after by \\citet{Elsworthetal90} and \\citet{LW90}.  It was quickly established that the frequency shifts were strongly correlated with surface activity \\citep[etc.]{Woodardetal91,BB93,Elsworthetal94,Reguloetal94}.  \\citet{LW90} observed that the frequency shifts depended very strongly on mode frequency $\\nu$, and very weakly on degree $\\ell$ of the mode, and \\citet{A-Getal92} and \\citet{Elsworthetal94} confirmed these results. These authors concluded that all or most of the physical changes responsible for the changes in frequency were confined to the shallow layers of the Sun.  In general, this picture has been confirmed in more recent studies (e.g., observational results:  \\citealt{HKH99,HKH02,BA00,Verneretal04,DG05}, etc., and theoretical results:  \\citealt{Goldetal91,Balmforthetal96,Lietal03}, etc.).  A change in the second helium ionization zone at $r=0.98R_\\odot$, first suggested by \\citet{Goldetal91} and \\citet{Gough02}, has been confirmed by \\citet{BM04} and \\citet{Verneretal06}.  The even-order mode splitting parameters sample effects of structural asphericities on the mode frequency.  \\citet{K88} suggested that they were correlated with observed changes in surface temperature.  Subsequent work has shown that the aspherical components of the mode frequencies are tightly correlated with surface magnetic activity \\citep{HKH99,Antiaetal01}.  This high correlation lends further credence to the idea that frequency shits are caused by surface and/or near-surface effects.  This can be tested directly with high degree modes that sample the near-surface layers of the Sun.  However, as the degree $l$ increases, global modes become increasingly hard to measure precisely due to the decrease in mode lifetimes \\citep{Rhodesetal98,R-Setal01,Korzenniketal04}. The lack of reliable measurements of these modes has led some authors to use ring diagrams to measure high degree modes and measure changes in the shallow layers of the solar convection zone.  These studies have confirmed that structural changes do occur in the near-surface layers of the Sun \\citep{Betal07}.  Direct inversions of changes in the structure of the solar interior probed by the spherically symmetric global modes have not yielded any measurable differences in the deep interior \\citep{Basu02,E-Detal02} and there have been upper limits placed on the changes at the base of the convection zone \\citep{E-Detal02}.  \\citet{CS05} and \\citet{SC05} have presented evidence of a possible change in mode frequencies with lower turning points near the base of the convection zone.  Internal dynamics, on the other hand, show clear and unequivocal evidence of change over the course of the solar cycle. In the convection zone, bands of different rotational velocities (called zonal flows) have been shown to migrate poleward at high latitudes and equator-ward at low latitudes \\citep{Schou99,Howeetal00,AB00,AB01}.  Temporal variations in dynamics have been shown to penetrate throughout the entire convection zone, and even below \\citep{Vorontsovetal02,BA06,Howeetal05,Howeetal06}.  One conclusion that can be drawn unequivocally from previous studies of the changes in solar structure is that any changes deeper than those in the outermost layers of the Sun are very small, and hence very difficult to detect through their signatures in oscillation frequencies.  \\citet{CS05} used a smoothing technique to attempt to remove the effect of surface variations, and found that the scaled frequency differences showed evidence of change near the base of the convection zone, but could not say more about the physical nature of the changes.  Attempts to invert the frequencies directly have never shown any changes larger than the inversion errors \\citep[e.g.,][]{Basu02,E-Detal02}. Therefore, although current helioseismic instruments have determined the solar oscillation frequencies with tremendous precision, statistical errors in those frequencies are still too large to make any direct detections of structure changes in the deep interior.  In this work, we take a somewhat different approach to attempting to detect changes at the base of the convection. We use Principal Component Analysis (PCA) to separate the frequency differences over the last solar cycle into a linear combination of different time-dependent components.  This has the effect of decreasing the effects of measurement errors in the measured helioseismic frequencies, which allows us to isolate as precisely as possible the changes in frequency over time.  In section \\ref{sec:data}, we describe the data used, and the methods employed to analyze them.  In section \\ref{sec:results}, we present the results in detail. Finally, in section \\ref{sec:conc} we discuss the significance of the results and presents our conclusions. ",
        "conclusions": "\\label{sec:conc} We have analyzed the changes in solar oscillation frequencies over the course of solar cycle 23.  In order to reduce the effects of measurement errors and detect the faintest signatures of solar variability possible, we have employed a Principal Component Analysis of the frequencies.  The mean frequencies are known to vary over the solar cycle, and this variation is known to be tightly correlated with surface activity.  We have found this correlation as well.  In addition to the frequency dependent change which these earlier authors have detected, we have found a small but statistically significant change in modes with turning points at or below the convection zone.  This confirms the result of \\citet{CS05}.  This signature is present in both the MDI and the GONG data sets.  We have inverted these results to obtain the difference in sound speed, and we have confirmed that there is a change in solar structure at the base of the convection zone over the course of the solar cycle.  The measured change at a radius of $r=(0.712^{+0.0097}_{-0.0029})R_\\odot$ is $\\delta c^2 / c^2 = (7.23 \\pm 2.08) \\times 10^{-5}$, where the errors in radius are a measure of the resolution of the inversion taken from the first and third quartile points of the inversion kernel.  We have also used the splitting coefficients to investigate how the changes in structure vary over latitude.  We have found that the changes in the solar interior are tightly correlated with the latitudinal distribution of surface activity.  The most statistically significant changes detected in the analysis are changes in sound speed in the outer ten percent (by radius) of the solar interior."
    },
    "0807/0807.0397_arXiv.txt": {
        "abstract": "We present new fully self-consistent models of the formation and evolution of isolated dwarf galaxies. We have used the publicly available N-body/SPH code \\mbox{HYDRA}, to which we have added a set of star formation criteria, and prescriptions for chemical enrichment (taking into account contributions from both SNIa and SNII), supernova feedback, and gas cooling. We extensively tested the soundness of these prescriptions and the numerical convergence of the models. The models follow the evolution of an initially homogeneous gas cloud collapsing in a pre-existing dark-matter halo. These simplified initial conditions are supported by the merger trees of isolated dwarf galaxies extracted from the milli-Millennium Simulation. The star-formation histories of the model galaxies exhibit burst-like behaviour. These bursts are a consequence of the blow-out and subsequent in-fall of gas. The amount of gas that leaves the galaxy for good is found to be small, in absolute numbers, ranging between $3\\times10^7\\ M_{\\sun}$ and $6 \\times 10^7\\ M_{\\sun}$. For the least massive models, however, this is over 80 per cent of their initial gas mass. The local fluctuations in gas density are strong enough to trigger star-bursts in the massive models, or to inhibit anything more than small residual star formation for the less massive models. Between these star-bursts there can be time intervals of several Gyrs.  The models' surface brightness profiles are well fitted by S\\'ersic profiles and the correlations between the models' S\\'ersic parameters and luminosity agree with the observations. We have also compared model predictions for the half-light radius \\re, central velocity dispersion $\\sigma\\rs{c}$, broad band colour $B-V$, metallicity $[Z/Z_{\\sun}]$ versus luminosity relations and for the location relative to the fundamental plane with the available data. The properties of the model dwarf galaxies agree quite well with those of observed dwarf galaxies. However, the properties of the most massive models deviate from those of observed galaxies. This most likely signals that galaxy mergers are starting to affect the galaxies' star-formation histories in this mass regime ($M \\gtrsim 10^9\\ M_{\\sun}$).  We found that a good way to assess the soundness of models is provided by the combination of \\re and $\\sigma\\rs{c}$. The demand that these are reproduced simultaneously places a stringent constraint on the spatial distribution of star formation and on the shape and extent of the dark matter halo relative to that of the stars. ",
        "introduction": "Dwarf galaxies (DGs), galaxies with blue absolute magnitude $-14>M_\\mathrm{B}>-18$ mag, are the most numerous type of galaxies in the nearby universe \\citep{ferguson1994}. Besides their number there are various reasons why DGs are very important in modern astronomy, most notably their cosmological relevance.  In the Lambda cold dark matter ($\\Lambda$CDM) cosmological model, present day galaxies are the product of a series of mergers. The early universe consisted of a uniform density, superimposed with small Gaussian fluctuations. As their amplitude increases monotonically toward smaller scales, they give, after inflation, rise to a myriad of small virialized dark matter (DM) haloes. These small haloes collapse first under influence of their own gravity, followed by the collapse of haloes with larger masses, swallowing the smaller haloes. We thus have hierarchical clustering, where massive galaxies, such as our own Milky Way, have formed through the merger of DGs. Dwarf galaxies are therefore an important link in the chain of hierarchical clustering. Moreover due to their relatively low total mass ($\\mathcal{O}(10^7-10^9\\ M_{\\sun}$)) they are quite sensitive to the effects of gravitational interactions, star formation (SF) and supernova (SN) explosions. On the one hand DGs are therefore the ideal objects to study when researching SF and SNe in a galaxy environment, and dynamical evolution of galaxies in general. Understanding the birth and life of DGs on the other hand could prove crucial when trying to understand the formation and evolution of more massive stellar systems. And even more: cosmological simulations of galaxy formation exhibit very strong dependence on the way star formation, feedback (FB) through supernovae and stellar winds (SW) are modelled within stellar systems \\citep[see e.g.][]{white1991,kauffmann1999}. Therefore realistic simulations of single dwarf galaxies can supply information on how to implement star formation and feedback processes in large cosmological simulations \\citep[e.g.][]{springel2005}, where a fully consistent implementation of SF and FB should be possible for moderate simulation volumes.  From the theoretical point of view various approaches for the study of DGs are possible, most notably there are semi-analytical models (SAMs), consisting of a limited set of analytical equations solved numerically, and fully numerical simulations, Lagrangian (e.g. N--body/SPH) or Eulerian (grid codes). SAMs have the advantage of computational efficiency, due to their speed it is possible to scan a large part of the parameter space in a reasonable period of time. In numerical simulations however a fully self consistent implementation of SF, FB, chemical enrichment, gas cooling, ... is possible, at the price of high computational effort.  \\citet{dekel1986} were the first to model Dwarf Galaxies (DGs), emphasising the impact of supernova-driven winds on the formation of dwarfs. \\citet{yoshii1987} found good agreement between observational properties of DGs and those predicted by their SAM. These authors calculated the properties of DGs (and globular clusters), using a model that takes into account the dynamical response to a supernova-driven wind. An evolutionary population synthesis method is employed to bring the models into the observational domain.  According to these models, spheroidal galaxies are a one-parameter family of birth mass (initial total mass $M_\\mathrm{t,i}$). These models were later extended in \\citet{naga2004}, who used a Monte Carlo technique to construct the merger tree leading up to the formation of a DG. \\citet{naga2005} incorporated N--body simulations instead of a Monte Carlo method to construct the merger tree. They found good agreement between galaxies produced by their SAM, ranging from dwarf elliptical to elliptical galaxies, and observations. SAMs are also used to investigate certain processes involved in the birth and life of dwarf galaxies rather than trying to reproduce dwarf galaxies, e.g. the model of \\citet{burkert1997}, investigating the role of SNIa in delaying star formation and of \\citet{ferrara2000} who investigated the role of stellar feedback and dark matter.  Although plenty of numerical research has already been carried out concerning the formation and evolution of massive elliptical galaxies (EG) \\citep[to name a few:][]{katz1991, kawata2001, chiosi2002, kawata2003, kobayashi2005, merlin2006}, starting from various cosmological initial conditions, only few people have tried directly modelling DGs with fully numerical simulations. One problem is that DGs are too small to be adequately resolved in large cosmological simulations. A first fully numerical study of dwarfs was performed by \\citet{mori1997,mori1999}, who explored the role of gas cooling in dwarf galaxy formation. \\citet{chiosi2002} used an N--body/SPH code to simulate the formation of galaxies, including dwarf spheroidal, dwarf elliptical and normal elliptical galaxies. They started from cosmologically inspired initial conditions. Assuming that star formation starts after an initial merging period, the simulations start with the collapse of an initially homogeneous gas cloud in a dark matter halo \\citep{white1978}. Observable properties of DGs were broadly reproduced by their models. More recently, \\citet{stinson2007} conducted simulations of dwarf spheroidal galaxies (dSph). They used initial conditions similar to those of \\citet{chiosi2002}, except that a degree of rotation was added to the gas. Like \\citet{chiosi2002}, these authors find burst-like star formation. The work of \\citet{stinson2007} is based on the work by \\citet{stinson2006}, who investigated the influence of star formation criteria on the formation of a galaxy, as well as the dependence of the processes involved on the number of particles. \\citet{mashchenko2007} performed fully cosmological simulations of the formation of a few dwarf galaxies, focusing on the gravitational heating of the dark matter, induced by gas flows. They argue that this heating can effectively convert a cusped dark matter halo into a DM halo with a central core. Other studies of dSphs include those of \\citet{read2005}, who looked at the effect of baryonic mass loss from a two component galaxy (dark and baryonic matter) and \\citet{read2006}, who used cosmological simulations to investigate the smallest baryonic building blocks for galaxy formation. \\citet{mashchenko2005} employed a simple model for the formation of dwarf spheroidal galaxies. A single star-burst was imposed on a gas distribution, employing a density cutoff, after which the remaining gas was removed. They find reasonable agreement with data from local group dwarf spheroidal galaxies.  A study of dwarf galaxies in a cosmological context was performed by \\citet{scannapieco2001}, investigating the effect of gas outflows on the formation of dwarf galaxies. Other studies using simulations, not directly aimed at reproducing dwarf galaxies, include those of \\citet{maclow1999}, who used a two-dimensional grid code to model blowout/blow-away of gas in dwarf galaxies, and more recently \\citet{marcolini2006}, who used a three-dimensional grid code to study the ISM in dwarf spheroidal galaxies by imposing a burst-like Star formation history (SFH) and comparing the results to the Draco dSph.  In this paper we present the results of our models of dwarf galaxies, with initial conditions taken from the $\\Lambda$CDM formalism. As we construct spheroidal, non-rotating galaxies, we expect them to resemble dwarf spheroidal and dwarf elliptical galaxies. We took great care to avoid the introduction of a direct dependence on the number of particles or on the resolution into the formalism. We also investigate the role of the different star formation criteria in detail. We then compare our model data to observations using the central velocity dispersion, the half-light radius, $B-V$ broad band colours, total metallicity and the fundamental plane.  The outline of this paper is as follows. In \\S~\\ref{sec:code}, we summarise the most important features of the modelling code. The initial conditions of our models are discussed in \\S~\\ref{sec:incond}. In \\S~\\ref{sec:results}, we compare our models with observational data. We end with a summary and our conclusions in \\S~\\ref{sec:sum}. ",
        "conclusions": "\\label{sec:sum} We have presented N--body/SPH models of dwarf galaxies. The basic SPH scheme was modified to include star formation, supernova feedback, chemical enrichment and gas cooling. The simulations start from a homogeneous gas cloud collapsing onto a DM halo. A first series of models (A models) was used to investigate the influence of the gas particle number on the overall evolution of the system. We found that a minimum of 25$\\,$000 particles is required to resolve the evolution accurately. Using less particles leads to a reduced first star formation peak. Because of this, the ISM is injected with less energy (less SNe), resulting in higher subsequent SF peaks and hence a larger final stellar mass (Table \\ref{tabnpart}). Based on these experiments, we decided to use 30$\\,$000 gas particles and 15$\\,$000 DM particles in the simulations.  All our model galaxies exhibit burst-like star formation. This is a direct consequence of star formation in a shallow potential well. Supernovae heat the gas, causing it to expand. The gas density consequently drops below the threshold $\\rho\\rs{c}$ for star formation. After the lapse of an adiabatic cooling period, the gas is again allowed to cool and contract. Star formation thus appears to be a self-regulating mechanism. % Since substantial amounts of gas remain in and around the model galaxies throughout the duration of the simulations (see fig. \\ref{cummassfig}), one expects that tidal interactions and ram-pressure stripping in dense environments will have a significant effect on their SFHs.  We compared the model galaxies with observational data in terms of morphology (S\\'ersic $n$, $R\\rs{e}$, $\\mu_0$), kinematics ($\\sigma\\rs{1D,0}$), colour, metallicity and location with respect to the FP ($\\kappa_1,\\kappa_2,\\kappa_3$). Despite their simplicity, the models do very well in the range $-8 > M\\rs{B} > -15$ mag. The models reproduce the slopes of the size, colour, metallicity, and velocity dispersion vs. luminosity relations very well. The zero-points of these relations are reproduced to within a factor of 2. The position of the model dwarf galaxies in the three-dimensional space spanned by $(R\\rs{e},\\sigma\\rs{1D,0},I\\rs{e})$ agrees excellently with the observations. Dwarf galaxies trace a sequence within the FP (i.e. in the $(\\kappa_1,\\kappa_2)$ projection) that is almost orthogonal to that of massive galaxies. In the dwarf regime, the height above the FP (in the $(\\kappa_1,\\kappa_3)$ projection) is a decreasing function of galaxy mass. This is a consequence of less massive galaxies having shallower potentials than more massive ones, decreasing the tendency of the gas to fall into the potential well and increasing the capability of supernovae to support the gas by heating it.  Furthermore we found that a good check for the validity of models is the combination of \\re and $\\sigma\\rs{c}$. \\re is very sensitive to the location of star formation and hence to the time evolution of the density of the ISM.  $\\sigma\\rs{c}$ reflects, through the virial theorem, the total amount and distribution of (dark) matter. We demonstrate this using our B and C models, where it can be clearly seen (figs. \\ref{MB_re} and \\ref{MB_disp}) that using a more diffuse DM halo indeed raises $R\\rs{e}$, and at the same time reduces $\\sigma\\rs{c}$."
    },
    "0807/0807.0504_arXiv.txt": {
        "abstract": "AstraLux is a Lucky Imaging camera for the Calar Alto 2.2-m telescope, based on an electron-multiplying high speed CCD. By selecting only the best 1-10\\% of several thousand short exposure frames, AstraLux provides nearly diffraction limited imaging capabilities in the SDSS {\\em i'} and {\\em z'}\\ filters over a field of view of 24$\\times$24 arcseconds. \\\\ By choosing commercially available components wherever possible, the instrument could be built in short time and at comparably low cost. We briefly present the instrument design, the data reduction pipeline, and summarise the performance and characteristics\\footnote{Based on observations collected at the Centro Astron\\'omico Hispano Alem\\'an (CAHA) at Calar Alto, operated jointly by the Max-Planck Institut f\\\"ur Astronomie and the Instituto de Astrof\\'isica de Andaluc\\'ia (CSIC).}. ",
        "introduction": "The recovery of the full theoretical resolution of large ground-based optical telescopes in the presence of atmospheric turbulence has been a major goal of technological developments in the field of astronomical instrumentation in the past decades. The de-facto standard nowadays is adaptive optics, the active real-time correction of the incoming distorted wavefronts by means of deformable mirrors. While diffraction limited image quality can be routinely achieved in the near infrared at large telescopes, this is only possible with considerable technical effort.  Besides speckle techniques in the Fourier domain, a relatively simple approach is the selection of images based on their quality: instead of using all images of a large set of short-time exposures, only those showing little image distortion due to the variable strength of atmospheric turbulence are combined to a high-resolution result. Like in the speckle interferometric approach, this requires only a detector with fast readout capability and moderate computational effort. The application of this technique to fainter astronomical targets, e.g. low-mass double stars, was limited by the readout noise of the available detectors at visible wavelengths until a few years ago. At typical exposure times of few ten milliseconds -- necessary to ``freeze'' the effects of atmospheric turbulence -- photon noise limited detectors are obligatory to use this so-called ``Lucky Imaging'' technique on targets fainter than $\\approx$10\\,mag. Since the image quality of each single frame has to be determined, e.g. by measuring the Strehl ratio of a reference star, the readout noise of the detector sets stringent limits on the minimum brightness of this reference.  The advent of electron multiplying CCDs (EMCCD) with single-photon detection capabilities led to considerable interest in the Lucky Imaging technique on the part of professional astronomers \\cite{Jerram:2001,Hynecek:2002,Mackay:2001}. First experiments at the Nordic Optical Telescope (NOT) with LuckyCam proved that Lucky Imaging is a very promising alternative to adaptive optics, allowing diffraction limited imaging at telescopes in the 2--3\\,m class at wavelengths below 1$\\mu$m \\cite{Baldwin:2001,Tubbs:2002,Tubbs:2003,Law:2006}.  These results triggered the development of a similar instrument for the 2.2-m telescope at Calar Alto by our group. The Lucky Imaging camera AstraLux was built in less than 5 months thanks to the availability of most parts as off-the-shelf equipment.  In the following we describe the instrument design and data processing pipeline and summarise the instrument's performance and key characteristics. Interested readers looking for more comprehensive information are referred to the diploma thesis, which covers AstraLux and its performance in full detail \\cite{Hormuth:2007a}\\footnote{Available online at \\texttt{http://www.mpia.de/ASTRALUX/}}. ",
        "conclusions": "Within less than one year it was possible to design, build, and characterise a Lucky Imaging instrument for the Calar Alto 2.2-m telescope. As a common user instrument from 2007 on, AstraLux has become the standard tool for diffraction limited imaging at the Calar Alto observatory. It is currently mostly used for binarity surveys among stars and minor planets.  AstraLux is able to reach Strehl ratios as high as 25\\% in the {\\em z'} band. In general, Lucky Imaging provides an improvement of the Strehl ratio by a factor of 10, corresponding to an increase of the signal-to-noise ratio for point sources by a factor of 10$-$20, depending on atmospheric conditions. Thus a selection of only the best 5$-$10\\% of all images does definitely not have a negative effect on the detection limit for point sources.  The requirements for the reference star magnitude are similar as for observations with adaptive optics. The performance starts to significantly decrease at $I$=14\\,mag, but image quality improvements are still possible with stars as faint as 15$-$16\\,mag. The measured isoplanatic angle in the $z'$-band is with $\\approx$40\\arcsec\\ as large as typical values in $K$-band for adaptive optics observations.  The measured close companion detection limit at an angular separation of 1\\arcsec\\ is on average 6\\,mag, worse than what adaptive optics can provide. But: adaptive optics at 8-m class telescopes currently has this capability only in the $H$ and $K$-band at wavelengths $>$1.5\\,$\\mu$m. The achievable contrast ratio in speckle imaging observations is typically two magnitudes less than for Lucky Imaging.  The encouraging results regarding both development speed as well as scientific output triggered the start of the project ``AstraLux Sur'', an almost identical copy of the Calar Alto version to be used as visitor instrument at ESO's New Technology Telescope (NTT) at La Silla, Chile. AstraLux Sur is currently on its way to Chile after a development time of less than 2 months and will have first light in mid-July 2008."
    },
    "0807/0807.2737_arXiv.txt": {
        "abstract": "In this paper we report on the spectral evolution of 4U 1608--522 performed as part of the long Galactic Bulge monitoring with {\\it INTEGRAL\\/}. The data set include the April 2005 outburst. {\\it BeppoSAX} archival data (two observations, in 1998 and 2000) have been also analysed and compared  with the  {\\it INTEGRAL\\/} ones. Three different spectral states have been identified from the hard Color--Intensity diagram derived from {\\it INTEGRAL\\/}: the canonical Hard and Soft as well as an Intermediate state. The hard state spectrum is well described by a weak black body component plus a Comptonised plasma component with high electron temperature ($kT_{\\rm e}\\simeq 60$ keV) extending up to 200 keV without any additional cut--off.  % The soft spectra are characterized by a cold Comptonised plasma ($kT_{\\rm e}$= 2--3 keV, and 7 keV for the intermediate state) and a strong disk black body component. A reflection component, indicating reflection  of the X-ray radiation from the accretion disc, is also present in the soft state revealed by  {\\it BeppoSAX} in 1998. The 2000 {\\it BeppoSAX} observation revealed the source in quiescent state modelled by a neutron star atmosphere (assuming a neutron star with radius 10 km and mass 1.4 M$_\\odot$) with an effective temperature, $kT_{\\rm eff}$ of 0.1 keV plus a power law component with $\\Gamma \\sim$ 3 detected for the first time for this source. This spectrum can also be modelled with a simple black body compatible with emission originating from a small fraction of the NS surface of radius of 0.4 km. ",
        "introduction": "A Low Mass X-ray Binary (LMXB) is a stellar system composed by a late type star and a compact object which accretes material from the companion star via Roche lobe overflow. The Neutron star (NS) LMXBs  are classified as Z or Atoll source based on the track of the colour-colour diagram \\citep{hasinger} corresponding to different X-ray spectral and variability properties. The Atoll sources move from the Island branch that corresponds to the Hard spectral state to the Banana branch corresponding to the Soft state.  Many LMXBs are Soft X-ray Transient sources (SXRTs) showing spectral state changes driven possibly by accretion instability phenomena. In fact, they are often in a quiescent state with a low luminosity of 10$^{33}$ erg s$^{-1}$ and abruptly show X-ray outbursts that consist of fast flux increase reaching a $L_X \\sim 10^{37-38}$ erg s$^{-1}$ followed by a slower nearly exponential decay (lasting weeks or months) \\citep{campanareview}. During the outburst the sources reveal spectral states transitions from the Hard (onset the outburst) to the Soft (at the peak of the outburst) passing through Intermediate spectral states. The mass accretion rate is the main parameter responsible for the spectral changes (hardening and softening).  NS in quiescent state on average are more luminous than black holes (BH) in quiescence. This observational property may be used to distinguish the NS or BH nature of the compact object (Rutledge et al. 2000, Campana et al. 2001). Unfortunately so far, few sources have been detected during the quiescent state namely Aql~X-1, Cen~X-4, EXO~0748--676, 4U~2129+47 \\citep{campanareview}, XTE~J2123-058 \\citep{tomsick04} and 4U~1608-522 itself \\citep{asaiquie}.   The Atoll 4U~1608--522 is a LMXB with a NS as compact object as derived by the presence of Type I X-ray bursts and superbursts (Nakamura et al. 1989, Remillard et al. 2006). It is one of the SXRTs showing long periods of quiescence (such us EXO~0748--676 and 4U~2129+47) spaced out by periodic outbursts (Lochner et al. 1994, Simon et al. 2004) like Aql~X-1 \\citep{tanaka}.  In the past decade, the {\\it BeppoSAX\\/} and {\\it RXTE\\/} broad band capability provided a step forward in the study of the NS LMXBs establishing their high energy behaviour well above the standard X-ray band ($\\le$ 20 keV). More recently, reflection components and high energy tails have been reported, e.g. 4U~1820--30, 4U~1705--44  (Tarana et. al. 2007, Fiocchi et al. 2006, Piraino et al. 2007) and 4U~1608--522 itself \\cite{gie}. For \\source\\/ a radio flux density upper limit of $<$0.19 mJy at 8.5 GHz is available \\cite{migliari06}.  We report here on recent data from {\\it INTEGRAL\\/} to characterise the high energy emission and spectral variation. For this purpose we performed the scientific analysis of the data in the 4-200 keV energy range with JEM-X \\citep{lund} and IBIS \\citep{uber}. Two non contemporaneous observations with {\\it BeppoSAX\\/} \\citep{boella} have been also analysed.  We combined three methods of analysis: {\\it Temporal}, using the light curves in different energy bands; {\\it Photometric}, with the hardness--instensity Diagram; and {\\it Spectral}, by spectral modelling. ",
        "conclusions": "We reported on the study of the spectral behaviour, as derived with {\\it INTEGRAL} and {\\it BeppoSAX\\/} satellites, of the SXRT source 4U 1608--522 in different spectral states characterised by luminosity ranging from 4 $\\times 10^{33}$ erg s$^{-1}$ to 6 $\\times 10^{37}$ erg s$^{-1}$ from the quiescent up to the soft state. The soft and hard spectra have been modelled with a soft component, described with a disk black body emission or black body emission, plus a Comptonization component which is described by a Comptonised corona. The {\\it BeppoSAX\\/} soft state spectra show also a reflection emission from the disk.  During the hardening (from the banana to the island branch of the color-intensity diagram) the inner disk black body temperature decreases from 0.7 to 0.4 keV % while the Comptonised corona temperature increases and the optical depth decreases, becoming an optically thin and very hot corona in the hard state ($kT_{\\rm e}$ changes from 2 to 60 keV).  In spite of the difficulty to precisely determine the soft parameters  (because of the limited instrumentation band-width often sensitive at E$>$ 4 keV) and the need to always freeze the seed photons temperature in our fits, we can still draw some remarks and explanations on the source.  From the diskbb normalization model we have estimated the inner disk radius\\footnote{estimated with formula given in Gienrli{\\'n}ski \\& Done 2002.}. For the soft spectra a radius ranging from 20 to 60 km has been obtained, indicating that the disk extend near the NS. On the contrary, for the hard spectrum the inner radius extends up to 120 km, as expected to allow the formation of hot plasma corona. This imply that change in the spectral state corresponds to a different disk-corona geometry  of the system, as already observed in NS and also BH transient systems \\citep{zdz2004}.  Moreover, we observed a seed input photons temperature higher ($T_{\\rm 0}$ ranging from 0.6 to 1.2 keV) than the inner disk temperature. In view of the fact that the $kT_{\\rm in}$ is representative of the latest/inner part of the accretion disk,  the seed photons temperature  must be provided from a hotter inner region that could be either close the neutron star i.e. an optically thick  boundary layer heated by friction effect (as suggests possibly for the source GRS 1724-30, Barret et al. 2000) or from the neutron star surface itself.  The boundary layer model predicts a temperature ranging from 0.5 to 2.0 keV (Lin et al. 2007 and references therein) whereas the predicted temperature for the \"spreading boundary layer model\" of a NSLMXB (with NS radius of 15 km and mass of 1.4 M$_{\\odot}$) is of 2.4 keV (Lin et al. 2007, Revnivtsev et al. 2006, Suleimanov \\& Poutanen 2006) well above the temperature of 0.6-1.2 keV observed for 4U~1608--522.  Another possibility is that the seed photons could be provided by the stellar surface of the NS. In this case the temperature is higher than the NS temperature estimated during quiescence due to the NS cooling during the quiescent state at low accretion rate \\cite{Done_review}.  In the hard state the energy spectrum doesn't show any high energy cut--off. This behaviour is usually attributed either to non-thermal emission processes probably related to jet formation, as for 4U~0614+091 \\cite{migliari06b}, or to hybrid thermal-non thermal emission, as for 4U~1820--30 (Tarana et al. 2007). In any case, the spectra characteristics and models of \\source\\/ are different from the 4U~1820--30 ones. Also the large uncertainties on the value for the corona electron temperature do not allows to derive the presence of a cut-off. The lack of any firm detection in the radio band support this scenario though non-thermal process could not be ruled out. Indeed in any case, 4U~0614+091 also only an upper limit were derived in the radio band but later presence of jet was detected (Migliari et al. 2006, 2006b).   We have also investigated the spectrum of 4U~1608--522 during a quiescent state with a luminosity higher than previous reported ones. The quiescent state has been modelled by a thermal NSA emission coming from the cooling NS and, for the first time, a power law with $\\Gamma \\simeq$ 3 was required. The thermal emission is due to the cooling NS (with R=10 km and M=1.4M$_\\odot$) with an effective temperature of 0.12 keV, comparable to previous observed temperature of 0.17 keV \\citep{rut99}. The nature of the power law tail is still unclear. This component was also observed as the dominant one above 2-3 keV in X-ray quiescent spectrum of transient NSs such as Aql~X-1, Cen~X-4 (Asai et al. 2006, Campana et al. 2000, Rutledge et al. 2000), KS~1731--260 \\citep{rut2002}. It could be originated by a residual low level accretion onto the NS magnetosphere also in the quiescence state \\citep{campanareview} or to an interaction of pulsar wind/surrounding gas such as for PSR~1259-63 \\citep{tavani97}. This last case could be not applied to \\source\\/ that is not a pulsar like system. This power law component was never detected during previous quiescent spectrum of 4U~1608-522 suggesting that it appears only occasionally and suggesting a variability in the quiescent emission."
    },
    "0807/0807.1502_arXiv.txt": {
        "abstract": "{}{To account for the observed differential metallicity distribution (DMD) of the Milky Way halo, a semi-analytical model is presented in the framework of the hierarchical merging paradigm for structure formation.} {It is assumed that the Milky Way halo is composed of a number of sub-haloes with properties either as observed in the dwarf satellite galaxies of the Local group (shape of metallicity distribution, effective yield) or derived from calculations of structure formation (sub-halo distribution function).} {With reasonable assumptions for the parameters involved, we find that the overall shape and effective yield of the Galactic halo DMD can be reproduced in the framework of such a simple model. The low metallicity tail of  the DMD presents a defficiency of stars with respect to the simple model predictions (akin to the G-dwarf problem in the solar neighborhood); it is suggested that an early infall phase can account for that problem, as well as for the observed DMDs of dwarf satellite galaxies.} {Accretion of galaxies similar (but not identical) to the progenitors of present day dwarf satellites of the Milky Way may well have formed the Galactic halo. } ",
        "introduction": "The metallicity distribution of long-lived stars is one of the most powerful probes of galactic chemical evolution. Indeed, the shape of the metallicity distribution is independent of the rate at which stars are formed. It depends essentially on the stellar initial mass function (IMF), which fixes the ``intrinsic'' (true) yield of the stellar population and the loss or gain of (gaseous) mass from the galaxian system (outflow or inflow, see next section and Prantzos 2007a). In the case of the solar neighborhood, for instance, the shape of the metallicity distribution suggests that the local disk was constructed by infall on long timescales (e.g. Boissier and Prantzos 1999).  The regular shape of the metallicity distribution of the Milky Way (MW) halo can readily be explained by the simple model of galactic chemical evolution with outflow, as suggested by Hartwick (1976, see also Sec. 2.1). However, that explanation lies within the framework of the monolithic collapse scenario for the formation of the MW (Eggen, Lynden-Bell and Sandage, 1962). It seems dificult to keep that interpretation in the currently dominant paradigm of hierarchical galaxy formation. A coherent interpretation should be sought in the framework of the new paradigm, according to which the MW halo was formed from the hierarchical merging of smaller components (sub-haloes).  Several attempts to account for the metallicity distribution of the MW halo in the modern framework were undertaken in recent years, starting with the  work of Bekki and Chiba (2001). They model a galaxy closely ressembling the MW and  despite the low spatial resolution of their simulations (2 kpc), they are able to provide a halo metallicity distribution with a peak at the observed metallicity ([Fe/H]=--1.6), but with a shape which is much narrower than the observed one ; theyobtain a double peaked metallicity distribution (which is not observed), probably reflecting late major mergers due to the small mass resolution. In their five MW-type simulations (in a cosmological framework and using various prescriptions for ouflow/feedback) Scanapieco and Broadhurst (2001) find no ressemblance of the resulting halo metallicity distributions with the observed one. This is also the case in Font et al. (2006), who find systematically too metal rich haloes with respect to the MW halo (see their Fig. 6). On the contrary, using semi-analytical models (starting with the extended Press-Schecter formalism and using various prescriptions for star formation, feedback and mixing of the stellar ejecta) Tumlinson (2006) and Salvadori, Schneider and Ferrara (2007) find rather good agreement between the model and observed halo metallicity distribution.   Independently of their success or failure in reproducing the observations, those recent models provide little or no physical insight into the physical processes that shaped the metallicity distribution of the MW halo. Indeed, if the halo was built from a  large number of sucessive mergers of sub-haloes, why is its metallicity distribution so well described by the simple model with outflow (which refers to a single system)? And what determines the peak of the metallicity distribution at [Fe/H]=$\\sim$--1.6, which is (successfully) interpreted in the simple model by a single parameter (the outflow rate) ? Even in the succesfull models cited in the previous paragraph, no clear answer is provided to these questions.    In this work we present an attempt to built the halo metallicity distribution analytically, in the framework of the hierarchical merging paradigm. Preliminary results have already been presented in Prantzos (2007b). It is assumed that the building blocks were galaxies with properties similar to (but not identical with) those of the Local group dwarf galaxies that we observe today. In that way, the physics of the whole process become (hopefully) more clear. The plan of the paper is as follows: In Sec. 2 we present the simple model of the outflow and  its overall successfull description of observations; we discuss, in particular, some issues  related to its low- and high- metallicity tails. In Sec. 3 we present the properties of the sub-haloes that formed the MW halo. These are derived either from observations of the Local froup dwarf galaxies (the shape of the metallicity distribution, Sec. 3.1 and the effective yield, Sec. 3.2) or from the theory of large scale structure formation (the distribution function of baryonic sub-haloes, Sec. 3.3). The results of the model are presented in Sec. 3.4 and summarized in Sec. 4. ",
        "conclusions": "In this work we present an analytical model for the metallicity distribution of the MW halo, in the framework of the hierarchical merging scenario for galaxy formation. It is assumed that the MW halo was composed of a number of sub-haloes with stellar masses ranging from $M_1$=2 10$^6$ \\ms \\ to $M_2$=2 10$^8$ \\ms \\ and distributed as in Eq. (9). Furthermore, each of  the stellar sub-haloes has a DMD given by the simple model (Eq. 2) with a mass dependent effective yield given by Eq. (4). Those assumptions are based on observations of  dSph satellites of the MW galaxy, while the distribution function is obtained from numerical simulations of dark matter halo formation (after explicitly accounting for  modification due to outflows).  It is found that, under those assumptions, the bulk of the MW halo DMD is reproduced in a fairly satisfactory way. The advantage of that model  over the simple outflow model (proposed by Hartwick in 1976) is that it treats the MW halo not as a single entity but as a composite one, in line with the hierarchical merging scenario for galaxy formation. Furthermore, the proposed composite model  avoids the \"paradox\" noticed in Sec. 3.2, namely that the effective yield of the MW halo appears to be substantially lower than the  effective yields of dwarf galaxies of comparable mass.   Observations suggest that the low metallicity tail of both the MW halo and local dSphs is defficient in stars w.r.t. the predictions of the simple outflow model. Prantzos (2003) suggested that early infall may account for that defficiency in the case of the MW halo, while Helmi et al. (2006) prefer pre-enrichment (at a level of [Fe/H]=$\\sim$--2.8) for the dSphs. It is shown here, with a simple toy model based on observations of dSphs, that an early infall phase can account for the low metallicity tails of all small galaxian systems and, consequently, of the  MW halo conceived as a composite system.  The results obtained here with semi-analytical models should be substantiated with  detailed numerical simulations of the formation of a stellar halo in a MW-type galaxy. Such simulations suffer at present from several drawbacks, which make it difficult for them to provide robust predictions. For instance, the issue of star formation vs feedback is far from being settled. Furthermore, as shown (and stressed) in Prantzos (2003), accounting for the finite lifetimes of stars produces substantially different results than the commonly adopted Instantaneous recycling approximation (IRA) at the earliest epochs (and lowest metallicities); since almost all numerical models  adopt IRA (exceptions being Samland et al. 1997 or Lia et al. 2002), their predictions cannot be considered as accurate, even if the issue of feedback is satisfactorily resolved.  Finally, as emphasized in Sec. 2, an in-depth discussion of the important implications of the low metallicity tail of the DMD should rely on the final release of the results of the Hamburg/ESO survey (Sch\\\" orck et al. in preparation); the toy model presented here relies on preliminary results and should only illustrate the potential of the proposed analytical model to probe the physics of the early galaxy."
    },
    "0807/0807.4782_arXiv.txt": {
        "abstract": "The fraction of ionizing photons which escape their host galaxy and so are able to ionize hydrogen in the inter-galactic medium (IGM) is a critical parameter in studies of the reionization era and early galaxy formation. In this paper we combine observations of Ly$\\alpha$ absorption towards high redshift quasars with the measured UV luminosity function of high redshift galaxies to constrain the escape fraction ($\\fesc$) of ionizing photons from galaxies at $z\\sim5.5-6$. We employ an $N$-body simulation to describe the density and peculiar velocity fields and to identify virialized halos. To model the ionizing background (IBG) we associate these halos with star-bursting galaxies using a semi-analytic prescription. We extract ensembles of mock {\\Lya} absorption spectra for a range of values of $\\fesc$ assuming all resolved ($M\\gtrsim10^9$~\\msun) simulation halos are able to contribute to the IBG. The observed Ly$\\alpha$ transmission constrains the escape fraction to lie in the range $\\fesc\\sim10-25\\%$ (at $z\\sim5.5-6$). Excluding halos with $M< 10^{10} M_{\\odot}$ (as might be expected if galaxy formation is suppressed due to the reionization of the IGM) implies a larger escape fraction of $\\fesc\\sim20-45\\%$. Using the numerical results to calibrate an analytic relation between the escape fraction and minimum galaxy halo mass we also extrapolate our results to a mass ($M\\sim10^8$~\\msun) corresponding to the hydrogen cooling threshold. In this case we find $\\fesc\\sim5-10\\%$, consistent with observed estimates at lower redshift. We find that the escape fraction of high redshift galaxies must be greater than $5\\%$ irrespepctive of galaxy mass. Based on these results we use a semi-analytic description to model the reionization history of the IGM, assuming ionizing sources with escape fractions suggested by our numerical simulations. We find that the IBG observed at $z\\sim5.5-6$ implies a sufficient number of ionizing photons to have reionized the Universe by \\zsix. However, if the minimum mass for star-formation were $\\gtrsim10^9$~\\msun, the IBG would be over-produced at $z\\lesssim5$. In summary, our results support a scenario in which the IGM was reionized by low mass galaxies. ",
        "introduction": "\\label{introduction_section}  The sources thought to have produced the UV radiation which reionized the hydrogen gas in the inter-galactic medium (IGM) are star-bursting galaxies and quasars [e.g. \\citet{Barkana2001}]. However, the number  density of the quasar population is observed to decline exponentially with redshift at $z\\gtrsim 2.5$, implying that galaxies probably contribute the bulk of UV photons which drive reionization at $z\\gtrsim 6$ and also the evolution of the ionizing background post-overlap between $ z \\sim 3-6$ \\citep{Madau1999,Fan2002,Srbinovsky2007,Bolton2007a}. The potential contribution of galaxies to the UV radiation field is a function of their number density, the star formation rate (SFR) within the galaxies and the spectral energy distribution (SED) of the stellar population. However, the fraction of ionizing photons which escape their host galaxy into the IGM, is limited by intervening absorption in the inter-stellar medium (ISM). The value of this escape fraction ($\\fesc$) is therefore a critical parameter in studies of reionization, and remains relatively uncertain.  Owing to its importance, constraints on the value of the escape fraction feature extensively in the literature. However, the definition of $\\fesc$ varies depending on the study. A tendency exists in observational studies to quote values of $\\fesc$ as the escape fraction of Lyman-limit photons, relative to the escape fraction at $1500\\AA$ ($\\flyescplus$). In this paper, we define $\\fesc$ as the absolute escape fraction of photons at the Lyman-limit, since this is the parameter of physical importance for the reionization history of hydrogen. Where necessary, we convert values of the escape fraction from published studies to be consistent with this definition, assuming $\\flyescplus=0.24$ and an extinction with a median value of E(B-V) $=0.15$ at $z \\lesssim 3$ \\citep{Siana2007}.  Early theoretical studies predicted values for the escape fraction of between $\\fesc \\sim 3-15\\%$ \\citep{Dove1994, Dove2000,Ciardi2002, Fujita2003}, however \\citet{Clarke2002} suggested that $\\fesc$ could be much higher than this if the feedback dependent porosity of the ISM is considered. At high redshift, \\citet{Fujita2003} argue that increased supernovae activity create escape tunnels through the ISM for ionizing radiation, and estimate $\\fesc>20\\%$ by $z>5$. Similarly, \\citet{Razoumov2006} find that $\\fesc$ evolves from $\\sim1-2\\%$ at $z=2.39$ to $\\sim6-10\\%$ at $z=3.6$. In contrast, \\citet{Wood2000} suggest that $\\fesc$ decreases with redshift (becoming $\\sim 1\\%$ at $z\\sim 10$) due to the increasing density of the galactic disk. At $z\\sim20$, \\citet{Whalen2004} and \\citet{Alvarez2006}, have derived escape fractions as high as $\\fesc\\sim 95\\%$, however in these cases the large value can be attributed to the assumed population-III stellar population. Finally, \\citet{Kitayama2004} studied the evolution of \\HII~regions surrounding a central population-III star embedded in a galaxy and find that $\\fesc$ increases with the mass of the star and inversely with the mass of the halo.   In the local Universe, observational estimates of $\\fesc \\sim 2-11\\%$, have been derived\\footnote{In this range for $\\fesc$ we have excluded the finding of \\citet{Hurwitz1997} for one of the four galaxies in their studied sample in which $\\fesc \\lesssim 57\\%$.} from measurements of emission line strengths relative to the Lyman-limit flux [\\citet{Leitherer1995,Hurwitz1997,Heckman2001,Bergvall2006,Grimes2007}]. An alternative approach measures the flux in the Lyman-continuum relative to the flux at a more readily detectable frequency above the Lyman-limit (usually $1500\\AA$). Comparison is then made to the intrinic flux ratio predicted by an assumed  SED of galaxies, which can be characterized by the flux decrement across the Lyman break ($LB$). Assuming $LB=3$, \\citet{Deharveng2001} estimates values of $\\fesc\\lesssim 10\\%$ in nearby galaxies, and \\citet{Malkan2003} estimates values of $\\fesc\\lesssim 1-6\\%$ from a sample of $11$ blue galaxies at $z\\sim 1$. At $z\\sim3$, \\citet{Steidel2001} compile a composite spectrum from $29$ Lyman-break galaxies and find an average escape fraction of $\\fesc>10\\%$, and from observations of 2 galaxies  \\citet{Giallongo2002} determine values of $\\fesc \\lesssim 4\\%$. From a sample of 14 star-forming galaxies, \\citet{Shapley2006} determines the escape fraction to be  $\\fesc\\sim 3\\%$, however, in two of these galaxies detection of Lyman continuum flux suggests an escape fraction of the order of $10\\%$. A less stringent constraint is found by \\citet{Inoue2005}, who find the escape fraction to be $\\fesc \\lesssim 20-40 \\%$ at $z\\sim 3$ using narrow band images of two galaxies. Finally, the escape fraction has also been constrained by \\citet{Schaye2006} who found $\\fesc\\sim10\\%$ at $z\\sim3$, by comparing the absorption systems near an ionizing source with those in the average IGM.     Alternatively, if an SED is available for the galaxy, then the the effects of intrinsic and intervening absorption can be modelled. Using this method \\citet{Fernandez2003} find $\\fesc \\leq 4\\%$, by fitting to the observed SEDs of 27 galaxies at $1.9\\leq z\\leq 3.5$. Similarly, at $z\\sim1$, \\citet{Siana2007} fits optical/near-infrared model SEDs to the observed spectra of 21 galaxies and finds the average, intrinsic Lyman-break factor to be $\\sim 6$, from which they conclude the global escape fraction to be, $\\fesc \\lesssim 4\\%$. Recently, \\citet{Chen2007} estimate $\\fesc\\sim2\\%$ at $z\\sim2$ from the observed distribution of neutral hydrogen column densities in the after-glow spectra of long duration GRBs. This approach alleviates the dependence on largely unknown intrinsic spectral properties of the galaxy, however presents another issue of a possible GRB environment bias. Finally, combining estimates of $\\fesc$ from published studies with values derived from measurements of the ionizing background, \\citet{Inoue2006} find that the escape fraction increases from $\\fesc \\sim 1-10\\%$ in the redshift interval $1 \\lesssim z \\lesssim 4$.  While there is as yet no consensus on the value of $\\fesc$, the implications for the reionization history have been addressed by several authors. Recently, detailed numerical simulations \\citep{Gnedin2007a} have predicted a value for $\\fesc$ of between $1$ and $3\\%$, from halos of mass $M\\gtrsim5\\times10^{\\scs 10}M_{\\odot}$ within the redshift range $3<z<9$. In addition to a small escape fraction in massive galaxies \\citet{Gnedin2007a} further predict that halos with $M\\lesssim 5\\times10^{\\scs 10}M_{\\odot}$ have an escape fraction which is negligibly small. This very low efficiency of reionization would have profound effects on the reionization history. Indeed, \\citet{Gnedin2007b} argues that the small escape fraction for lower mass galaxies implies that the observed galaxy population at $z\\sim6$ could not have produced enough UV photons to reionize the IGM. This suggests that the process of reionization takes place in a ``photon starved\" environment. This conclusion has been previously reached by \\citet{Bolton2007a} who studied the ionizing background (IBG) based on the {\\Lya} absorption properties towards high redshift quasars. However, \\citet{Bolton2007a} inferred a higher a value for the escape fraction of $\\fesc \\sim 20-30\\%$ (at \\zsix).     In this paper we employ an approach that is similar to \\citet{Bolton2007a}, and constrain the escape fraction of ionizing photons (at $z\\sim5.5-6$) by combining the observed luminosity function (LF) of galaxies with the observed opacity of the high redshift IGM. However, with respect to the study of the escape fraction our modelling includes several improvements. We employ an $N$-body simulation to describe the density and peculiar velocity fields and to identify virialized halos. We determine the UV luminosity of each halo using a semi-analytic description and constrain model parameters by fitting to the observed galaxy LF. We then generate an inhomogeneous ionizing background using these discrete galaxies, which depends on the escape fraction of ionizing radiation. For a range of escape fractions we extract an ensemble of mock {\\Lya} absorption spectra along 3072 sight-lines through our simulation box and determine the value of $\\fesc$ which best models the observed {\\Lya} transmission. We note that in our model the escape fraction is not explicitly dependent on galaxy mass or angular direction.  The paper is organized as follows. In \\S~\\ref{model_section} we outline our models for the IBG due to discrete galaxies. We also describe our computation of the {\\Lya} photon transmission. In \\S~\\ref{Gamma_and_escape_fraction_section} we show our results for $\\fesc$ at several redshifts assuming MFPs which are consistent with the observationally derived values from \\citet{Fan2006a}. In \\S~\\ref{galaxy_suppression_section} we repeat our analysis to constrain $\\fesc$ assuming the suppression of galaxy formation in low mass halos ($M<10^{10}M_{\\odot}$), while in  \\S~\\ref{reionization_massive_galaxies_section} we exclude halos with $M<5\\times10^{10}M_{\\odot}$. In \\S~\\ref{extrapolation_section} we extrapolate our numerical results to other minmum galaxy masses. In \\S~\\ref{small_escape_fraction_models_section} we discuss the implications of our results using a semi-analytic model of reionization history. Finally we make some concluding remarks in \\S~\\ref{conclusions_section}. Throughout the paper we assume a $\\Lambda$CDM cosmology with parameters ($\\sigma_{\\scs 8}, h, \\omegab, \\omegam, \\omegal$)=$(0.8,0.7,0.04,0.3,0.7)$. All distances are in co-moving units unless otherwise stated. ",
        "conclusions": "\\label{conclusions_section}  In this paper we have employed an $N$-body simulation to describe the density and peculiar velocity fields in a $(65.6$~Mpc/$h)^{\\sss 3}$ volume of the IGM at $z=5.5,5.7,6.0$. The simulation also identified the locations and masses of virialized halos, with the minimum halo mass accurately resolved by the simulation being $2\\times10^9M_{\\odot}$. We associate the identified halos with galaxies and assign to them a UV luminosity using models for the SFR and the SED. By fitting to the observed UV galaxy LF \\citep{Bouwens2007} we constrained the best fit free parameters to be ${\\fstar}~=0.11$ and ${\\tdc}~=0.16$, where $\\fstar$ is the fraction of baryons which participate in star formation and $\\tdc$ is the duty-cycle of star-bursting galaxies. We then constructed an ionizing background which reflects the spatial distribution of UV luminous galaxies. Our model made the assumption of a universal MFP at each redshift.  Given the density, velocity and ionizing radiation fields we extracted an ensemble of {\\Lya} absorption spectra along 3072 sight-lines. From this sample we computed the mean effective transmission as a function of the fraction of ionizing photons produced that escape their host galaxy ({$\\fesc$}). In our simulations $\\fesc$ effectively adjusts the intensity of the ionizing background. The value of $\\fesc$ was determined so that the mean effective transmission predicted by the model matched the observed value. Assuming halos above the mass resolution of our simulation to host galaxies, we find values for the escape fraction at $z\\sim5.5-6$ of $\\fesc \\sim 10-25\\%$ assuming values for the MFP which are suggested by \\citep{Fan2006a}. In models where we assume the formation of galaxies to be suppressed in halos with masses below $10^{10}M_{\\odot}$ (possibly due to radiation feedback following reionization), we find the escape fraction to be in the range $\\fesc \\sim 20-40\\%$. Our results for the escape fraction are consistent with those found in \\citet{Bolton2007a} who constrain the escape fraction to be $\\fesc \\sim 20-30\\%$, using a similar approach to the one adopted in this paper. This is true both for scenarios in which the minimum halo mass is set by the simulation resolution ($2\\times 10^{\\scs 9}M_{\\odot}$) and for the larger value of $10^{\\scs 10} M_{\\odot}$.   We also considered the scenario suggested by \\citet{Gnedin2007a} in which the escape fraction of ionizing photons is negligible in halos with masses below $\\sim5\\times10^{10}M_\\odot$. In this case we find the escape fraction is required to be $\\fesc \\sim 46-86\\%$ in order to reproduce the observed optical depth of the IGM. This is inconsistent with theoretical modelling of the escape fraction by \\citet{Gnedin2007a} who suggest $\\fesc\\sim1-3\\%$. Using the numerical results to calibrate an analytic relation between the escape fraction and minimum galaxy halo mass we extrapolate our results to a mass ($M\\sim10^8$~\\msun) corresponding to the hydrogen cooling threshold. In this case we find $\\fesc\\sim5-10\\%$, consistent with observed estimates at lower redshift.     Since the SFR is estimated at $1350\\AA$, while $\\fesc$ is estimated at $912\\AA$, the escape fraction is degenerate with the adopted galaxy SED. The value of the Lyman-break factor ($LB$), which is the intrinsic flux decrement measured across the Lyman break, is particularly sensitive to to the age of the stellar population. Assuming that the normalization of the adopted SED at wavelengths above the Lyman-limit, and the slope of the SED at wavelengths below the Lyman-limit are both correct, we can express this degeneracy with respect to LB as $\\fesc=\\fesc^{\\sssrm{LB}=3}\\left({{LB}}/3 \\right)$. In this expression $\\fesc^{\\sssrm{LB}=3}$ is the escape fraction which we have constrained in this paper, and $\\fesc$ the value obtained for an alternate $LB$.   We use a semi-analytic model that describes the evolution of the ionized state of the post-reionization IGM to study whether the escape fraction implied by our numerical modelling is consistent with reionization at $z\\gtrsim6$. We find that the measured ionization rates imply a history that is consistent with reionization at \\zsix~for a range of these models, including those in which the minimum halo mass of galaxies contributing to the ionizing emissivity is $M\\sim10^8,10^{9}$ and $10^{10}$ solar masses respectively. However, we find that models with minimum masses above $10^9$~\\msun~over-estimate the ionization rate in the post-overlap IGM. These models imply that the Universe was reionized by low mass galaxies with an escape fraction of $\\sim 5\\%$.  In summary, the transmission of \\Lya~photons through the IGM at $z\\sim5.5-6$ implies an ionization rate that requires the escape fraction of ionizing photons to be in excess of $5\\%$. Our results for the escape fraction and star-formation efficiency based on our numerical simulations combined with the model reionization histories suggest that the observed transmission and luminosity function are compatible with reionization at $z\\sim6$."
    },
    "0807/0807.4261_arXiv.txt": {
        "abstract": "We present a semi-analytical kinetic calculation of the process of non-linear diffusive shock acceleration (NLDSA) which includes the magnetic field amplification due to cosmic ray induced streaming instability, the dynamical reaction of the amplified magnetic field and the possible effects of turbulent heating. The approach is specialized to parallel shock waves and the parameters we chose are the ones appropriate to forward shocks in Supernova Remnants. Our calculation allows us to show that the net effect of the amplified magnetic field is to enhance the maximum momentum of accelerated particles while reducing the concavity of the spectra, with respect to the standard predictions of NLDSA. This is mainly due to the dynamical reaction of the amplified field on the shock, which noticeably reduces the modification of the shock precursor. The total compression factors which are obtained for parameters typical of supernova remnants are $R_{tot}\\sim 7-10$, in good agreement with the values inferred from observations. The strength of the magnetic field produced through excitation of streaming instability is found in good agreement with the values inferred for several remnants if the thickness of the X-ray rims are interpreted as due to severe synchrotron losses of high energy electrons. We also discuss the relative role of turbulent heating and magnetic dynamical reaction in driving the reduction of the precursor modification. ",
        "introduction": "The supernova remnant (SNR) paradigm for the origin of galactic cosmic rays heavily relies on the mechanism for particle acceleration being the \\emph{Diffusive Shock Acceleration} (DSA) at the shock front generated in the supernova blast. This mechanism, also known as \\emph{first order Fermi acceleration}, has been studied in great detail and a two decades long work has led to the development of a kinetic non-linear theory that allows us to assess the importance of the dynamical reaction of the accelerated particles on the shock itself. The nonlinear effects described by the theory turn out to be not just corrections. They rather reveal the profound reasons why the mechanism works in the first place as a cosmic ray accelerator: the large efficiency required to explain the energetics ($10-20\\%$ of the kinetic energy of the blast wave going into cosmic rays) and the large magnetic fields required to explain the maximum energies observed in cosmic rays are the two main motivations for developing such a nonlinear theory, and, not surprisingly, are among the most successful predictions of the theory.  The initial attempt at building a nonlinear theory led to two-fluid models \\citep{dr_v80,dr_v81} which provided information on the dynamics and thermodynamics of the shocked gas and cosmic ray gas but not on the spectrum of the accelerated particles. A satisfactory kinetic approach, able to predict the spectrum of accelerated particles was later proposed by \\cite{malkov1,malkov2} and \\cite{blasi1,blasi2}. More recently a kinetic model which takes into account the possibility of arbitrary diffusion coefficients was put forward by \\cite{amato1}. The self-excitation of unstable modes leading to magnetic field amplification was then also introduced by \\cite{amato2}.  Parallel to these analytical approaches, developed primarily in the assumption of quasi-stationarity of the acceleration process, numerical approaches following the temporal evolution were also developed \\citep{bell87,je91,elli90,ebj95,ebj96,kj97,kj05,jones02}. These have been of crucial importance for the description of some aspects of the phenomenology connected with the acceleration process in SNRs, especially during the Sedov-Taylor phase. Numerical methods for the solution of the transport equation for cosmic rays and of the conservation equations for the plasma in the shock region have been extensively used to reproduce the observed multifrequency observations from single SNRs \\citep[e.g.][]{V+05,cc+05}.  A typical result of all approaches to nonlinear shock acceleration is that the spectra of accelerated particles are far from being power laws. The concave shape leads to spectra as flat as $p^{-3.2}$ close to the maximum momentum, corresponding to total compression factors that may exceed $\\sim 50-100$ \\citep{amato1}. In general the total compression factor $R_{tot}$ is found to scale with the Mach number as $R_{tot}\\propto M_0^{3/4}$ \\citep{be99}. Such levels of shock modification do not compare well with some observations, which suggest $R_{tot}\\sim 7-10$ \\citep[see e.g.][]{V+05}.  From the phenomenological point of view this problem has been faced by invoking some sort of turbulent heating: part (or all) of the energy in the form of Alfv\\'en waves which are responsible for the scattering of charged particles is assumed to be damped on the thermal gas, thereby causing its heating in the upstream precursor region \\citep{v-mck81,mck-v82}. This process, originally investigated as a possible mechanism to limit the magnetic field amplification, keeping the turbulent field in the linear regime (e.g. $\\delta B/B\\ll 1$), is currently called upon also in situations in which $\\delta B/B\\gg 1$.  \\begin{table*} \\begin{minipage}{150mm}\\label{tab:SNR} \\caption{Parameters for 5 well studied SNRs.} \\begin{tabular}{ccccc} \\hline SNR & $u_0 (\\rmn{km/s})$ & $B_2(\\umu\\rmn{G})$ & $\\alpha_2\\times 10^3$\\\\ \\hline Cas A& 5200 (2500)& 250--390 & 32 (36) \\\\ Kepler & 5400 (4500)& 210--340 &  23 (25)\\\\ Tycho & 4600 (3100)& 300--530 &  27 (31)\\\\ SN 1006 & 2900 (3200)& 91--110 & 40 (42) \\\\ RCW 86 & (800) & 75--145 & 14-35 (16-42)\\\\ \\hline \\end{tabular}  \\medskip The value are from \\cite{V+05} (in parentheses) and from \\cite{P+06}. In order to estimate the normalized downstream magnetic pressure $\\alpha_{2}=B_2^2/(8\\pi\\rho_0u_0^2)$, for the SNRs discussed by \\cite{P+06} we used $\\rho_0=0.1\\,m_p/\\rmn{cm}^3$ (SN 1006) and $\\rho_0=0.5\\,m_p/\\rmn{cm}^3$ (in the other cases). \\cite{V+05}, instead, provide directly $\\alpha_2$. \\end{minipage} \\end{table*}  An important piece of information has been recently added to the debate on whether SNRs can be the sources of galactic cosmic rays: Chandra X-ray observations of some remnants showed narrow filaments of non-thermal origin \\cite[see][and references therein for a review]{vink}. If the thickness of the rims is assumed to be due to severe synchrotron losses limiting the lifetime of high energy electrons, then one can infer the strength of the magnetic field in the downstream region, which turns out to be $\\sim 100$ times stronger than magnetic fields in the ISM.  On the other hand, it has been proposed that the narrow rims may reflect the damping scale of the magnetic turbulence rather than the loss length of electrons \\citep{pohl05}. This interpretation is at odds with the morphology of the radio emission, as discussed by \\cite{radio}  \\citep[see also the discussion in][]{morlino}, but at present it is not possible to rule it out. If it turns out to be correct, then there would be no observational constraint on the magnetic field in the shock region.  If on the other hand the interpretation based on strong synchrotron losses is confirmed, the inferred levels of magnetization can be interpreted as a result of streaming instability induced by cosmic rays efficiently accelerated at a parallel shock front \\citep{bell78,bl2001,bell2004}, although alternative mechanisms have also been put forward \\citep[e.g.][]{gj07}. In Table~\\ref{tab:SNR} we list the SNRs where evidence has been collected, from X-ray observations, for a strong magnetic field: $u_0$ is the shock velocity, $B_2$ is the value of the magnetic field downstream of the shock as inferred from the X-ray brightness profile and $\\alpha_2$ is the magnetic energy density immediately downstream of the forward shock, in units of the total kinetic pressure at upstream infinity ($\\alpha_{2}=B_2^2/(8\\pi\\rho_0u_0^2$)). The data are from \\cite{V+05} (numbers in parentheses) and from \\cite{P+06}.  The efficient acceleration and the magnetic field amplification are the two most impressive manifestations of nonlinear diffusive acceleration at SNR shocks. Both these aspects have been included in a Monte Carlo scheme by \\cite{elli06}, with the magnetic field amplification described accordingly to the phenomenological scenario proposed by \\cite{bl2001}. \\cite{elli06} find that, when amplification is efficient, the wave pressure makes the plasma upstream of the shock less compressible and the change in the energy density of the magnetic turbulence across the subshock strongly affects $\\Rs$, which in turn affects the injection efficiency and the entire process of cosmic ray acceleration.  In a previous paper \\citep{jumpl} we used a three-fluid approach (gas, cosmic rays and Alfv\\'en waves), to show  the very general nature of the magnetic dynamical feedback: the shock dynamics is significantly affected by the turbulence backreaction as soon as the magnetic pressure becomes comparable to that of the gas upstream of the subshock. In particular, for the magnetization levels inferred from available observations (see Tab.~\\ref{tab:SNR}) and the saturation values that are expected from different amplification mechanisms proposed in the literature, we found that $\\Rt$ would naturally lead to values in the range 6-10, in good agreement with the values currently inferred from observations.  In the present paper we use a kinetic model for particle acceleration at a parallel shock in the non-linear regime, together with the modified conservation equations in the precursor and at the subshock in order to describe particle acceleration, the dynamical reaction of accelerated particles, the generation of magnetic field through streaming instability and the dynamical reaction of the magnetic field on the shock, which also results in modified cosmic ray spectra.  The paper is structured as follows: in Sec.~\\ref{sec:jumpcond}  we write down the correct jump conditions for a parallel, modified shock when magnetic turbulence is excited by the cosmic ray streaming, as in \\cite{jumpl}. In Sec.~\\ref{sec:model} we summarize our kinetic model based on that by \\cite{amato2}, including the self-consistent treatment of the magnetic field amplification via resonant streaming instability. The latter is described in detail in Sec.~\\ref{sec:resSI}. In Sec.~\\ref{sec:results} we present our main results, namely the solutions for DSA with resonant streaming instability: in particular we discuss the reduced modification of the precursor and the consequent steepening of the spectrum near the (increased) maximum momentum. In Sec.~\\ref{sec:TH} we investigate the combined effect of magnetic reaction and turbulent heating and we show that the dominant effect on the precursor modification is likely to be that of the magnetic reaction. Our conclusions are in Sec.~\\ref{sec:conclusions}. ",
        "conclusions": "\\label{sec:conclusions}  The possibility that the narrow filaments detected in X-rays may result from severe synchrotron losses of high energy electrons has sparked much attention on the issue of magnetic field amplification at shock fronts, since the inferred fields are $\\sim 100$ times larger than the typical interstellar ones. The importance of such magnetic fields for the origin of cosmic rays is immediate: if the turbulent field leads to Bohm diffusion, it is easy to show that particle acceleration at shock fronts may lead to the generation of protons with energy close to the knee.  In a previous paper \\citep{jumpl} we investigated the hydrodynamical consequences of having a large magnetic field amplified by cosmic ray streaming upstream of the shock and we demonstrated that when the magnetic pressure equals or overcomes the pressure of the background plasma, a condition easy to realize, the compressibility of the plasma is reduced, so to decrease the compression factor of the cosmic ray modified shock and smoothen its precursor.  In the present paper we went beyond the hydrodynamical picture and we solved the combined system of conservation equations, cosmic ray diffusion-convection equation, and equation for the magnetic field amplification in order to investigate the effects of the precursor smoothening on the spectrum of accelerated particles. We find that resonant streaming instability is sufficient to amplify a parallel pre-existing magnetic field to levels which compare well with X-ray observations, if the latter are interpreted as a result of strong synchrotron losses.  In these circumstances, we also confirm the crucial role of the dynamical feedback of the magnetic field, which leads to total compression factors around $\\sim 7-10$ (to be compared with the typical predictions of standard NLDSA, which gives $R_{tot}\\sim 20-100$), in good agreement with the values suggested by observations and by fits to the multifrequency spectrum of several SNRs \\citep{V+05,warren}. The reduced compression in the precursor does not inhibit particle acceleration, in fact a fraction 50-60\\% of the ram pressure at the shock is converted into accelerated particles, for parameters typical of a SNR in the Sedov phase.  The kinetic calculation of particle acceleration in the nonlinear regime was carried out using the approach of \\cite{amato1,amato2}. The effects of resonant streaming instability were treated as in \\cite{skillingc,bell78}, but including the presence of the precursor which implies magnetic field compression in addition to amplification. We neglected here the possibility of having non-resonant amplification \\citep{bell2004}, for several reasons: 1) the relation between pressure and energy flux of these waves is not yet well defined; 2) the diffusion properties of charged particles in turbulence at wavelengths much shorter than the Larmor radius of particles are not known as yet, and need a dedicated effort of investigation; 3) as showed by \\cite{amato3}, the non resonant modes are likely to be especially important in the free expansion phase of a SNR and in the early stages of the Sedov phase.  Perhaps not surprisingly, we find that the effect of magnetic feedback on the spectrum of accelerated particles is that of reducing the concavity. This result can be easily explained based on the reduced compressibility in the shock precursor, caused by the magnetic pressure. The slope of the spectrum typically remains close to $4$ at energies smaller than $\\sim 1$ TeV, while being flatter at higher energies (Fig. \\ref{fig:Ucsispettro}). We also discuss the effect of the amplified field on the velocity of the scattering centers, which causes a further steepening of the spectrum, leading it to get closer to a power law. We find that the spectrum of accelerated particles becomes even steeper than $p^{-4}$ if the velocity of the scattering centers is assumed to be the Alfv\\'en speed in the amplified field. Although this maybe an unphysical assumption, it serves the scope of showing that the results of calculations carried out in this strongly nonlinear regime may well be affected by details which at the present time we are unable to control.  The smoother precursor induced by the magnetic feedback also implies a larger value of the maximum achievable momentum for the accelerated particles. We recall that, as showed by \\cite{bac07}, a strong precursor leads to lower the maximum momentum since particles feel a smaller mean fluid speed and therefore a longer acceleration time. Using the acceleration time for modified shocks as calculated by \\cite{bac07}, we find that for a $1000\\rmn{yr}$ old remnant the maximum energy can be as high as $10^6\\rmn{GeV}$, a factor 5-10 larger than one could find without including the magnetic backreaction. In principle even larger maximum energies can be achieved by taking into account the non-resonant streaming instability in the early phases of the SNR evolution.  Our entire analysis was carried out for the case of a shock propagating in the direction parallel to the ambient magnetic field. The amplification of the background field due to cosmic ray streaming soon changes this situation, in that the magnetic field upstream becomes highly turbulent. Changing the field obliquity should not change much the picture discussed in the present paper as long as the projection of the shock velocity along the field lines remains supersonic. If the supernova shock propagates in a medium in which the background field is mainly oriented in a given direction, there will be locations at which the shock is quasi-parallel and others at which it is quasi-perpendicular. The implications for the morphology of the non-thermal emission require a dedicated investigation, in that the quasi-perpendicular regions could be expected to be bright because of the more efficient perpendicular configuration, but quasi-parallel regions could be expected to be bright because of more efficient magnetic field amplification. Clearly the degree to which these expectations may be true depends on many details which at present are not under control, although observations of specific astrophysical objects (e.g. SN 1006) will provide us with important information on this issue.   Finally, we investigated the effect of turbulent heating on the shock modification. We generalized the formalism introduced by \\cite{v-mck81} and \\cite{mck-v82} in order to account for the possibility that a finite fraction of the wave energy is damped on the background plasma. We find that, although in general the heating of the upstream plasma results in a decrease of the compressibility, the effect is at most competitive, if not subdominant, compared with the dynamical feedback of the amplified magnetic field on the plasma. When turbulent heating is indeed dominant, a fraction very close to unity of the wave energy is transformed in thermal energy, and the main effect is that of suppressing the growth of the magnetic field, so that it becomes challenging to explain X-ray observations in terms of severe synchrotron losses. Moreover in this case the maximum energy of accelerated particles is drastically reduced and it becomes much smaller than the knee.  It is worth stressing that while the magnetic feedback on the background plasma is based on well known physics, the turbulent heating is treated at best in a very phenomenological way, with little attention to the specific physical processes that may be responsible of the non-adiabatic heating and on how modes with different wavelengths are damped. In these circumstances we think that the effect of magnetic feedback is much more solidly assessed and should be considered as the chief mechanism for reducing the compression in cosmic ray modified shock waves with evidences for amplified magnetic fields.  On the other hand, the role of turbulent heating may become much more important for older remnants, when the shock velocity drops below $\\sim 2000 \\rmn{km/s}$. At these times, the amplification of the magnetic field may have become less important from the dynamical point of view, so that the magnetic feedback is also less important. These stages, as discussed by \\cite{pz05} play an important role in determining the spectrum of cosmic rays observed at the Earth, at least at energies much lower than the knee energy."
    },
    "0807/0807.1996_arXiv.txt": {
        "abstract": "We present MUSE, a software framework for combining existing computational tools for different astrophysical domains into a single multiphysics, multiscale application. MUSE facilitates the coupling of existing codes written in different languages by providing inter-language tools and by specifying an interface between each module and the framework that represents a balance between generality and computational efficiency.  This approach allows scientists to use combinations of codes to solve highly-coupled problems without the need to write new codes for other domains or significantly alter their existing codes.  MUSE currently incorporates the domains of stellar dynamics, stellar evolution and stellar hydrodynamics for studying generalized stellar systems.  We have now reached a ``Noah's Ark'' milestone, with (at least) two available numerical solvers for each domain.  MUSE can treat multi-scale and multi-physics systems in which the time- and size-scales are well separated, like simulating the evolution of planetary systems, small stellar associations, dense stellar clusters, galaxies and galactic nuclei. In this paper we describe three examples calculated using MUSE: the merger of two galaxies, the merger of two evolving stars, and a hybrid $N$-body simulation.  In addition, we demonstrate an implementation of MUSE on a distributed computer which may also include special-purpose hardware, such as GRAPEs or GPUs, to accelerate computations.  The current MUSE code base is publicly available as open source at {\\tt http://muse.li}. ",
        "introduction": "The Universe is a multi-physics environment in which, from an astrophysical point of view, Newton's gravitational force law, radiative processes, nuclear reactions and hydrodynamics mutually interact.  The astrophysical problems which are relevant to this study generally are multi-scale, with spatial and temporal scales ranging from $10^4$ m and $10^{-3}$ seconds on the small end to $10^{20}$m and $10^{17}$s on the large end. The combined multi-physics, multi-scale environment presents a tremendous theoretical challenge for modern science.  While observational astronomy fills important gaps in our knowledge by harvesting ever-wider spectral coverage with continuously increasing resolution and sensitivity, our theoretical understanding lags behind these exciting developments in instrumentation.  In many ways, computational astrophysics lies intermediate between observations and theory.  Simulations generally cover a wider range of physical phenomena than observations with individual telescopes, whereas purely theoretical studies are often tailored to solving sets of linearized differential equations. As soon as these equations show emergent behavior in which the mutual coupling of non-linear processes result in complex behavior, the computer provides an enormous resource for addressing these problems.  Furthermore simulations can support observational astronomy by mimicking observations and aiding their interpretation by enabling parameter-space studies.  They can elucidate the often complex consequences of even simple physical theories, like the non-linear behavior of many-body gravitational systems.  But in order to deepen our knowledge of the physics, extensive computer simulations require large programming efforts and a thorough fundamental understanding of many aspects of the underlying theory.  From a management perspective, the design of a typical simulation package differs from construction of a telescope in one very important respect.  Whereas modern astronomical instrumentation is generally built by teams of tens or hundreds of people, theoretical models are usually one-person endeavors.  Pure theory lends itself well to this relatively individualistic approach, but scientific computing is in a less favorable position.  So long as the physical scope of a problem remains relatively limited, the necessary software can be built and maintained by a single numerically educated astronomer or scientific programmer.  However, these programs are often ``single-author, single-use'', and thus single-purpose: recycling of scientific software within astronomy is still rare.  More complex computer models often entail non-linear couplings between many distinct, and traditionally separate, physical domains. Developing a simulation environment suitable for multi-physics scientific research is not a simple task.  Problems which encompass multiple temporal or spatial scales are often coded by small teams of astronomers.  Many recent successful projects have been carried out in this way; examples are GADGET \\citep{2001NewA....6...79S}, and starlab \\citep{2001MNRAS.321..199P}.  In all these cases, a team of scientists collaborated in writing a large-scale simulation environment.  The resulting software has a broad user base, and has been applied to a wide variety of problems.  These packages, however, each address one quite specific task, and their use is limited to a rather narrow physical domain. In addition, considerable expertise is needed to use them and expanding these packages to allow the addition of a new physical domain is hampered by early design choices.  In this paper we describe a software framework that targets multi-scale, multi-physics problems in a hierarchical and internally consistent implementation.  Its development is based on the philosophy of ``open knowledge'' \\footnote{See for example {\\tt http://www.artcompsci.org/ok/}.}.  We call this environment MUSE, for MUltiphysics Software Environment. ",
        "conclusions": "The Multiscale Software Environment presented in this paper provides a diverse and flexible framework for numerical studies of stellar systems.  Now that the Noah's Ark milestone has been reached, one can ask what new challenges MUSE has to offer.  Many of the existing modules have been adapted for grid use and, as demonstrated in \\S\\,\\ref{Sect:GRID}, MUSE can be used effectively to connect various computers around the world.  However, there are currently a number of limitations in its use, and in its range of applications, which will be addressed in the future.  Most of the current application modules remain unsuitable for large-scale scientific production simulations. The stellar dynamics codes do not yet efficiently deal with close binaries and multiples, although modules are under development, and external potentials, though relatively easy to implement, have not yet been incorporated.  Binary evolution is not implemented, and the diagnostics available to study the output of the various modules remain quite limited.  Many improvements can be made to the environment, and we expect to include many new modules, some similar to existing ones, others completely different in nature.  The current framework has no method for simulating interstellar gas, although this would be an extremely valuable addition to the framework, enabling study of gas-rich star clusters, galaxy collisions, colliding-wind binary systems, etc.  In addition, radiation transfer is currently not implemented, nor are radiative feedback mechanisms between stars and gas.  Both would greatly increase the applicability base of the framework.  However, both are likely to challenge the interface paradigm on which MUSE is based.  The current MUSE setup, in which the individual modules are largely decoupled, has a number of attractive advantages over a model in which we allow direct memory access. The downside is that MUSE in its present form works efficiently only for systems in which the various scales are well separated.  Communication between the various modules, even of the same type, is currently all done via the top interface layer. For small studies, this poses relatively little overhead, but for more extensive calculations, or those in which more detailed data must be shared, it is desirable to minimize this overhead. One way to achieve this would be by allowing direct data access between modules. However, for such cases, the unit conversion modules could not be used, and consistency in the units between the modules cannot be guaranteed. As a result, each module would be required to maintain consistent units throughout, which may be hard to maintain and prone to bugs. In addition, the general problem of sharing data structures between modules written in different languages, currently resolved by the use of the glue language, resurfaces."
    },
    "0807/0807.3993_arXiv.txt": {
        "abstract": "{ The infrared is a key wavelength regime for probing the dusty, obscured nuclear regions of active galaxies. We present results from an infrared study of 87 nearby Seyfert galaxies using the Spitzer Space Telescope and ground-based telescopes.  Combining detailed modelling of the 3 - 100 micron spectral energy distributions with mid-IR spectral diagnostics and near-infrared observations, we find broad support for the unified model of AGNs.  The IR emission of Seyfert 1s and 2s is consistent with their having the same type of central engine viewed at a different orientation.  The nature of the putative torus is becoming clearer; in particular we present evidence that it is likely a clumpy medium.  Mid-infrared correlations between tracers of star formation and AGN ionizing luminosity reveal the starburst-AGN connection implied by the black hole/bulge mass relation, however it is not yet clear if this is due to feedback. ",
        "introduction": "The infrared wavelength regime provides a unique probe into the nuclear regions of active galaxies (AGNs), as the circumnuclear dust absorbs and reradiates higher energy photons from the nucleus. We are conducting a study of the infrared emission from nearby Seyfert galaxies, using the Spitzer Space Telescope and ground-based telescopes \\citep{buc06}. The goal is to construct matched-aperture infrared (IR) spectral energy distributions (SEDs) of the central few kiloparsecs of a large sample of nearby Seyfert galaxies, and to decompose the SEDs into the contributions to the dust heating from the active nucleus and any circumnuclear starburst. ",
        "conclusions": "We have constructed 20\\arcsec\\ IR SEDs of a sample of 85 nearby Seyfert galaxies.  Decomposition of the SEDs into AGN and starburst contributions provides support for clumpy torus models.  The SEDs are consistent with the unified scheme for AGN.  We find that the Seyfert 1s are AGN-dominated at 12~\\micron\\ while the Seyfert 2s are starburst-dominated, which is likely to be a selection effect. We find evidence that the cool dust is primarily heated by star formation in these Seyferts."
    },
    "0807/0807.1281_arXiv.txt": {
        "abstract": "{} {We study the generation of transversal oscillations in coronal loops represented as a straight thin flux tube under the effect of an external driver modelling the global coronal EIT wave. We investigate how the generated oscillations depend on the nature of the driver, and the type of interaction between the two systems. } {We consider the oscillations of a magnetic straight cylinder with fixed-ends under the influence of an external driver modelling the force due to the global EIT wave. Given the uncertainties related to the nature of EIT waves, we first approximate the driver by an oscillatory force in time and later by a shock with a finite width. }{Results show that for a harmonic driver the dominant period in the generated oscillation belongs to the driver. Depending on the period of driver, compared to the natural periods of the loop, a mixture of standing modes harmonics can be initiated. In the case of a non-harmonic driver (modelling a shock wave), the generated oscillations in the loop are the natural periods only. The amplitude of oscillations is determined by the position of the driver along the tube. The full diagnosis of generated oscillations is achieved using simple numerical methods.}{} ",
        "introduction": "Latest high-resolution coronal observations have shown that coronal structures are very dynamic entities with flows and waves propagating along them. Waves and oscillations in coronal loops have received increased attention in the last few years due to the possibility to use the observed properties to diagnose not only the magnetic field in these structures, but also the sub-resolution space distribution of loops, plasma properties, etc. (e.g. Roberts, Edwin and Benz 1984; Aschwanden et al. 1999; and Nakariakov et al. 1999; Banerjee et al. 2007, Verth et al. 2007, Arregui et al. 2008, Verth and Er\\'elyi 2008). One specific type of oscillations observed in coronal loops is the fast kink mode which disturbs the symmetry axis of the loop (for a full description of these waves see, e.g. Edwin and Roberts 1983) and it is nearly incompressible.  Many of kink waves and oscillations have their origin in the interaction of coronal loops with various external sources and drivers (see, e.g. Hindman and Jain 2008, Erd\\'elyi and Hargreaves 2008). One of the possible explanations of oscillations in coronal loops and/or prominence fibrils is that they have their origin in the interaction of these loops with global coronal waves, e.g. EIT waves. EIT waves (Thompson et al. 1999) are waves generated by sudden energy releases (flares, coronal mass ejections, etc) and they are able to propagate over very large distances in the solar low corona.  Observational evidence for large-scale coronal impulses initiated during the early stage of a flare and/or CME has been provided by the EIT instrument onboard SOHO, TRACE/EUV, STEREO/EUVI. EIT waves propagate in the quiet Sun with speeds of 250\\,--\\,400 km s$^{-1}$ at an almost constant altitude. At a later stage in their propagation EIT waves can be considered as a freely propagating wavefront which is observed to interact with coronal loops (see, e.g. Wills-Davey and Thompson, 1999). Using TRACE/EUV 195 \\AA\\ observations, Ballai, Erd\\'elyi and Pint\\'er (2005) have shown that EIT waves - seen in this wavelength - are waves with average periods of the order of 400 seconds. Since at the height where the EUV lines are formed, the magnetic field can be considered vertical, EIT waves were interpreted as fast magnetohydrodynamic (MHD) waves. This interpretation was confirmed using multi-wavelength STEREO/EUVI observations by Long et al. (2008). Recently, Attrill et al. (2007a,b) proposed that the diffuse EIT coronal bright fronts are due to driven magnetic reconnections between the skirt of the expanding CME magnetic field and the favorably orientated quiet Sun magnetic field. According to this latter model, the propagation process of the front consists of a sequence of successive reconnection events.  Although a large consensus was reached on the trigger mechanism of these global coronal waves and effects EIT waves can generate, the nature of these large scale disturbances is still unknown, despite the multitudes of models. The main reason of this uncertainty is the lack of high temporal and spatial resolution as well as the limited field of view (in the case of TRACE/EUV).  The present paper investigates the temporal and spatial variations of transversal oscillations in a coronal loop under the influence of an external driver representing the coronal global EIT waves. {\\bf Strictly speaking a coronal loop application would require the consideration of an external magnetic field. However, we consider this study as a starting point in a much more complex analysis.} Due to the uncertainties in resolving the nature of EIT waves, we will discuss separately the cases of a harmonic driver and a driver of a finite width with a pulse-like temporal distribution. It should be noted here that the study of the present papers apply not only to EIT waves as external driver, but it could be applied to any external source.  In the next section we introduce the working model and derive the governing equation of transversal oscillations based on the principle of force equilibrium in conjunction with the continuity of mass and magnetic flux. The governing equation is solved for the a {\\bf driving force when its particular form is not specified}. Section 3 is devoted to the study of the periodical motion of the coronal loop under the influence of a few particular drivers. Finally, in Section 4 we summarize our results and discuss some possible key extensions which were neglected here but they could be added to this model in subsequent studies. ",
        "conclusions": "The generation and propagation of oscillations in coronal loops modelled by a straight magnetic cylinder with fixed ends is studied when the coronal loop is under the effect of an EIT wave, as a driver. We found that if all forces acting on the flux tube are taken into account, the governing equation describing the propagation of standing transversal (kink) waves is described by an inhomogeneous Klein-Gordon-type equation and the inhomogeneous part of the equation is represented by the force (over a unit volume) by EIT waves. The evolutionary equation contains information about the propagation speed of waves (here the kink speed) and the cut-off frequency of kink modes. The cut-off value is determined by the densities inside/outside the loop and the Alfv\\'en speed (i.e. magnetic field).  Using the combined Laplace and Fourier sine transform techniques, the governing equation is solved, such that the solution takes into account general initial and boundary conditions.  Particular solutions have been found in the case of an EIT wave considered first as a wave with a frequency $\\omega_{EIT}$, and later as a shock wave with finite front thickness (a non-harmonic driver). The results show that in case of a non-harmonic driver the periods of generated modes always belong to the natural periods of the loop. On the other hand, in the case of a periodic driver - for an arbitrary period of the driver - there will be a mixture of standing modes which could explain on the observed period ratio. The analysis carried out here for different type of drivers show that the generated oscillations will carry predominantly information about the driver rather than the loop itself.  The oscillations described in this paper were all modelled in the framework of ideal MHD. In reality coronal loop oscillations are observed to decay relatively rapidly and several mechanisms have been proposed to explain this damping (Ruderman and Roberts 2005, Terradas et al. 2005, 2007, Selwa et al. 2007). The inclusion of a dissipative (or energy lost) mechanism in the present model will be addressed in a future study. It could be possible that in the case of a loop oscillating as a whole in the kink mode, the friction with the environment could be also an important factor whose inclusion in the model could result in a possible explanation of the damping of these oscillations.  The present study can be further extended to the case when the external EIT wave acts not only on a single magnetic loop but on a system of adjacent loops. In this case the primary displacement of the first coronal loop (generated by the incoming EIT wave) will be the driver for the oscillations in the second loop (and so on), leading to coupled loop oscillations. {\\bf In order to describe a realistic loop, the present model is going to be expanded to consider the effect of an external magnetic field. It is expected that the presence of this field will generate an additional force which will tend to suppress the oscillation of the tube.}"
    },
    "0807/0807.3767_arXiv.txt": {
        "abstract": "{CoRoT, the first space-based transit search, provides ultra-high precision light curves with continuous time-sampling over periods of up to 5 months. This allows the detection of transiting planets with relatively long periods, and the simultaneous study of the host star's photometric variability. % In this letter, we report on the discovery of the transiting giant planet CoRoT-Exo-4b and use the CoRoT light curve to perform a detailed analysis of the transit and to determine the stellar rotation period. % The CoRoT light curve was pre-processed to remove outliers and correct for orbital residuals and artefacts due to hot pixels on the detector. After removing stellar variability around each transit, the transit light curve was analysed to determine the transit parameters. A discrete auto-correlation function method was used to derive the rotation period of the star from the out-of-transit light curve. % We derive periods for the planet's orbit and star's rotation of $9.20205 \\pm 0.00037$ and $8.87\\pm 1.12$ days respectively, consistent with a synchronised system. We also derive the inclination, $i=90.00_{-0.085}^{+0.000}$ in degrees, the ratio of the orbital distance to the stellar radius, $a/R_{\\rm s}=17.36_{-0.25}^{+0.05}$, and the planet to star radius ratio $R_{\\rm p}/R_{\\rm s}=0.1047_{-0.0022}^{+0.0041}$. % We discuss briefly the coincidence between the orbital period of the planet and the stellar rotation period and its possible implications for the system's migration and star-planet interaction history.} ",
        "introduction": "\\label{intro}  Transits provide unique insights into fundamental aspects of extra-solar planets that are currently beyond the reach of other techniques: mean density (through the determination of true masses and radii), atmospheres (through transmission spectroscopy and secondary transit observations), and formation and evolution mechanisms (through the statistics of the orbital parameters, including the true mass). All but two of the transiting planets published so far have been very short-period planets ($<5$\\,d), whose properties are bound to be affected by extreme proximity to their host star. The exceptions are HD\\,147506b \\citep{bkt+07} and HD\\,17156b \\citep{bal+07}, with periods of 5.6 and 21.2\\,d, both in very eccentric orbits ($e \\ge 0.5$).  CoRoT is the first space-based transit survey, and over its lifetime will survey 120\\,000 stars for up to 5\\,months with precisions down to 0.1\\,mmag per hour. This letter is the $4^{\\rm th}$ in a series presenting new extra-solar planets discovered thanks to the CoRoT observations and associated ground-based follow-up program. Papers I \\citep{bba+08} and II \\citep{aab+08} present the discoveries of CoRoT-Exo-1b and CoRoT-Exo-2b, both very close-in giant planets. Paper III \\citep{bqd+08} reports on the observation of the spectroscopic transit of CoRoT-Exo-2b. Here we report on the detection of CoRoT-Exo-4b, a relatively long period (9.2\\,d) gas-giant transiting planet, and on the analysis of the CoRoT light curve (the production of which is described in Section~\\ref{obs}) to determine the transit parameters (see Section~\\ref{transit}) and the stellar rotation period (see Section~\\ref{star}). This analysis makes no assumptions aside from the planetary nature of the companion, which was confirmed through ground-based follow-up observations which are reported in a companion letter (\\citealt{mbg+08} hereafter Paper V). In Paper V, the full system parameters are determined using the transit parameters reported in the present work. These parameters are used in Section~\\ref{concl}, where we discuss the implications of the relatively long period and of its proximity to the stellar rotation period. A further detection, CoRoT-Exo-3b, will be presented in \\citet[][Paper VI]{dha+08}. ",
        "conclusions": "\\label{concl}  CoRoT-Exo-4b's orbital period is the second longest of any transiting planet to date, following HD\\,17156b \\citep{bal+07}. To our knowledge, it is only the second transiting system where the star's photometric rotation period has been measured photometrically (the other being HD\\,189733, \\citealt{hw08}). There are other systems where the rotation period and planet orbital period are very similar, particularly $\\tau$~Boo \\citep{cds+07,dmf+08} and XO-3 \\citep{jmb+08,wht+08} (the rotation periods are derived from Doppler imaging and $v \\sin i$ respectively, rather than measured photometrically). In both cases, the period of the orbit and the stellar rotation period are not exactly the same, but is compatible with a synchronised outer envelope when allowing for differential surface rotation. Both planets, which orbit F-type stars like CoRoT-Exo-4, are massive and have very short periods ($\\sim 3$\\,d), and thus may induce strong enough tides in the star to have synchronised its convective envelope.  A simple calculation, based on the work of \\citet{dlm04} and \\citet{jgb08}, shows that CoRoT-Exo-4b, which has a smaller mass and longer period, would not exert a significant tidal torque on its host star unless one invokes an unphysically high rate of tidal energy dissipation in the star. However, \\citet{mbv+08} discuss the case of two other transiting systems containing F-type stars, XO-4 and HAT-P-6, both with orbital periods longer than 7 days which are roughly half the (spectroscopically determined) stellar rotation period. They argue that this may indicate the existence of resonant interactions between the planet's orbit and its rotating host star, the absence of such resonances in systems containing cooler stars suggesting that Jupiter-mass planets can only interact effectively with stars with very shallow outer convective zones.  Similarly, if some other factor brought the CoRoT-Exo-4 system to the 1:1 ratio between stellar rotation and orbital period that is observed today, resonant interaction between star and planet may have maintained this ratio thereafter. Any subsequent evolution of the star's rotation rate is likely to have been modest: 9\\,d is close to the peak of the rotation period distribution for F-type stars both during the TTauri phase and in the field \\citep[see e.g. Fig.~1 of][]{bar03}. The initial resonance may have occured naturally if the proto-planetary disk was truncated near the co-rotation radius, and the planets' migration halted close to the inner edge of the disk. This hypothesis is supported by the absence of detectable eccentricity (see Paper V).  This system clearly warrants further observational and theoretical investigation to pin down its tidal and rotational evolution status. For example, more detailed analysis of the out-of-transit light curve should enable the active regions on the stellar surface to be mapped in a time-resolved fashion \\citep{lbr07} to search for signs of star-planet magnetic interaction."
    },
    "0807/0807.3284_arXiv.txt": {
        "abstract": "Despite the known general properties of the solar cycles, a reliable forecast of the 11-year sunspot number variations is still a problem. The difficulties are caused by the apparent chaotic behavior of the sunspot numbers from cycle to cycle and by the influence of various turbulent dynamo processes, which are far from understanding. For predicting the solar cycle properties we make an initial attempt to use the Ensemble Kalman Filter (EnKF), a data assimilation method, which takes into account uncertainties of a dynamo model and measurements, and allows to estimate future observational data. We present the results of forecasting of the solar cycles obtained by the EnKF method in application to a low-mode nonlinear dynamical system modeling the solar $\\alpha\\Omega$-dynamo process with variable magnetic helicity. Calculations of the predictions for the previous sunspot cycles show a reasonable agreement with the actual data. This forecast model predicts that the next sunspot cycle will be significantly weaker (by $\\sim 30\\%$) than the previous cycle, continuing the trend of low solar activity. ",
        "introduction": "Investigation of solar activity is one of the oldest solar physics problems. Here we consider this phenomenon in the context of the sunspot number variations, which have detailed observational data during the past 23 solar cycles. There is no doubt that the 11-year cyclic variations of the sunspot number are connected to the dynamo process inside the Sun. Therefore it is natural to search for a solution by investigating nonlinear dynamo models, and use them for predicting the solar cycles. However, our current understanding of the solar dynamo is quite poor. The current predictions of the next solar cycle, number 24, show a wide range of the expected sunspot number. For example, using an axisymmetric, mean-field model of a flux transport dynamo Dikpati and Gilman (2006) predicted that the cycle 24 will be about 30\\% - 50\\% stronger than the previous cycle 23 with the sunspot number reaching 155 - 180 at the maximum. Using polar magnetic field measurements and assuming a direct relationship between the polar field strength and the future toroidal magnetic field Svalgaard et al. (2005) predicted that the cycle 24 will have a peak sunspot number of $75 \\pm 8$. A weak cycle 24 with the sunspot number of $74 \\pm 10$ was also proposed by Javaraiah (2007) from analysis of sums of the areas of sunspot groups in the  latitudinal interval of $0^\\circ - 10^\\circ$ in both hemispheres. There are many other predictions for the upcoming cycle(s), reviewed by Kane (2007) and Obridko and Shelting (2008).  The great variety of the predictions is caused by uncertainties in the dynamo models and model parameters, and errors in both models and observations. In the paper, we present initial results of application of a data assimilation method to a simple nonlinear dynamo model \\citep{kitikos08}. This model reproduces the basic properties of a solar cycle, such as the shape of the sunspot number profile. The advantage of the data assimilation methods is in their ability to combine the observational data and the models for possible efficient and accurate estimations of the physical properties, which cannot be observed directly. Here we consider an implementation of the Ensemble Kalman Filter method, EnKF, \\citep{evensen07}, which is effective for investigation of nonlinear dynamical models. The method is useful in several aspects: it supports estimations of past, present, and even future states of a system, and it can do so even when the precise nature of a modeled system is unknown. ",
        "conclusions": "The results of assimilation of the annual sunspot number data into the solar dynamo model and the prediction of the previous solar cycles (Fig.~\\ref{fig3}) demonstrate a new method of forecasting the solar activity cycles. Using the EnKF method and a simple dynamo model we obtained reasonable predictions usually for the first halfs of the sunspot cycles with the error $\\sim 8 - 12$\\%, and in some cases also for the declining phase of the cycles. This method predicts a weak solar cycle 24 (Fig.~\\ref{fig3}i) with the smoothed annual sunspot number at the maximum of approximately 80. It is interesting that the simulations show that the previous cycle does not finish in 2007 as was expected, but it still continues. According to the estimates the maximum of the next cycle will be approximately in 2013.  The application of the data assimilation method, EnKF, for modeling and predicting the solar cycles shows the power of this approach and encourages further development. It also reveals significant uncertainties in the model and the data. Among these are the uncertainties in the determination of the start of a solar cycle from the sunspot number series (in particular, when the cycles overlap), leading to the uncertainty in the phase relation between the model solution and the data. Also, there are significant uncertainties in the relationship between the sunspot number data and the physical properties of the solar magnetic field, in the absence of magnetic field and helicity data, and, of course, in the dynamo model. Our conclusion is that for robust and accurate predictions of the solar cycles the information contained in the sunspot number data is insufficient. For further development we plan to apply the data assimilation method to more complete 2D dynamo models, which describe the latitudinal distribution of the solar magnetic field, and use the magnetograph data available for the past 3 cycles.  This work was supported by the Center for Turbulence Research (Stanford) and the International Space Science Institute (Bern)."
    },
    "0807/0807.3551_arXiv.txt": {
        "abstract": "We study the clustering of LRG galaxies in the latest spectroscopic SDSS data releases, DR6 and DR7, which sample over 1 ${\\,\\rm Gpc^3/h^3}$ to z=0.47. The 2-point correlation function $\\xisp$ is estimated as a function of perpendicular $\\sigma$ and line-of-sight $\\pi$ (radial) directions. We find a significant detection of a peak at $r\\simeq 110$Mpc/h, which shows as a circular ring in the $\\sigma-\\pi$ plane. There is also significant evidence for a peak along the radial direction whose shape is consistent with its originating from the recombination-epoch baryon acoustic oscillations (BAO). A   $\\xisp$ model with no radial BAO peak is disfavored at $3.2\\sigma$, whereas a model with no magnification bias is disfavored at $2\\sigma$. The radial data enable, for the first time, a direct measurement of the Hubble parameter $H(z)$ as a function of redshift. This is independent from earlier BAO measurements which used the spherically averaged (monopole) correlation to constrain an integral of $H(z)$. Using the BAO peak position as a standard ruler in the radial direction, we find: $H(z=0.24)= 79.69 \\pm 2.32 (\\pm 1.29)$ km/s/Mpc for z=0.15-0.30 and $H(z=0.43)= 86.45 \\pm 3.27 (\\pm 1.69)$ km/s/Mpc for $z=0.40-0.47$. The first error is a model independent statistical estimation and the second accounts for systematics both in the measurements and in the model. For the full sample, $z=0.15-0.47$, we find $H(z=0.34)= 83.80 \\pm 2.96 (\\pm 1.59)$ km/s/Mpc. ",
        "introduction": "\\label{intro}  Luminous red galaxies (LRG's) are selected by color and magnitude to obtain intrinsically red galaxies in SDSS (Eisenstein etal 2001). % These galaxies trace a big volume, around $1 {\\rm \\, Gpc^3 h^{-3}}$, which makes them ideal for studying clustering on large scales \\citep[see][]{Hogg2005}. Attention has been paid especially to the baryon acoustic peak around a scale of $100$ Mpc/h, because of its value as a standard ruler. In Eisenstein etal (2005) the baryon acoustic peak was detected in the spherically averaged two-point correlation function (i.e. the monopole) using LRG's from an earlier SDSS release (about half the size of the final data). Both 2dFGRS and SDSS spectroscopic redshift surveys have been used to constrain cosmological parameters via the galaxy power spectrum \\citep{clusteringlrg,sanchez}, including information from the baryon acoustic feature \\citep{hutsia, hutsib, percival, Sanchez09}. Photometric LRG catalogs cover a larger volume and larger densities and have also been used to obtain cosmological constraints \\citep{padmanabhan2007,blake}.  In this paper, as in Papers I \\citep{paper1} and II \\citep{paper2} of this series, we focus on the LRG's anisotropic redshift space correlation function $\\xi(\\sigma, \\pi)$, where $\\pi$ is the line-of-sight (LOS) or radial separation and $\\sigma$ is the transverse separation. There are three sources of anisotropy at scales larger than 40Mpc/h. It is well known that peculiar motion distorts the correlation function anisotropically \\citep{Kaiser}. It is also well known that assuming the wrong background cosmology will lead to an anisotropic $\\xi$ \\citep{AP}. What is under-appreciated is that gravitational lensing also introduces an anisotropy of its own to the galaxy correlation function \\citep{matsulen,hui1,hui2}. In this paper, we will present evidence of this lensing distortion in the LOS direction.  A distinct focus of this paper in the series is on the baryon acoustic oscillation (BAO) feature in $\\sigma-\\pi$ plane.  In principle, it can be used as a standard ruler to measure both the Hubble expansion rate $H(z)$, via its observed redshift span in the radial direction, and the angular diameter distance $D_A (z)$, via its observed angular size in the transverse direction \\citep{glazebrook,seoeisenstein}. We will implement this idea to measure $H(z)$. We are aided in this endeavor by two effects. Redshift distortions in the LOS direction make the $\\xi$ negative on intermediate scales ($\\sim 50 - 90$ Mpc/h), while magnification bias gives $\\xi$ a positive boost especially on large scales ($\\gsim 100$ Mpc/h). This combination  enhanced the contrast of the radial BAO peak with respect to the BAO in the perpendicular direction. Because the noise is shot-noise dominated any increase in the signal, such as bias or magnification bias, increase the signal-to-noise and helps the BAO detection. We will show that the shape of the correlation function in the radial direction is in good agreement with the predictions. We will also validate the interpretation of the feature as baryon acoustic in nature by studying its appearance in directions away from the radial direction, and performing a parametric fit to the monopole.  It is useful to point out some related earlier work. Redshift distortions in the LRG's and quasars at $z\\sim 0.55$ have been studied using the 2dF-SDSS LRG and QSO Survey \\citep[2SLAQ][]{cannon,ross,daangela,wake}. \\citet{okumura} measured $\\xisp$ away from the LOS direction in the SDSS using a sample similar to ours: they used about 47000 LRG's over a redshift range of $0.16 - 0.47$ while we use 75000 LRG's over $z$ of $0.15 - 0.47$. None of the earlier papers attempted a direct measurement of $H(z)$ from the baryon feature in the LOS direction.  This paper is organized as follows. Section \\ref{sec:theory} gives a summary of the theory, including a brief introduction to magnification bias. In section \\ref{sec:analysis} we perform an analysis of the clustering signal around the position of the BAO peak, both in the monopole and in the $\\sigma-\\pi$ plane, especially in the radial direction, from which we obtain the Hubble parameter H(z) in section \\ref{sec:interpretation}. We end with a conclusion in section \\ref{sec:conclusions}, where we deduce implications for the dark energy equation of state, and  where we emphasize the need for further theoretical work on the nonlinear coupling between peculiar motion and magnification bias. ",
        "conclusions": "\\label{sec:conclusions}   In this paper, we have studied in detail a peak in the LRG correlation function, in the monopole, in the circular ring on the $\\sigma-\\pi$ plane, and especially in the LOS direction. Its location around $110$ Mpc/h is consistent with the interpretation that it originates from baryon acoustic oscillations. Its significance can be assessed in several ways. We have performed a parametric fit to the monopole: the data requires a non-zero $\\Omega_b$ with high significance. In fact, we find that data appears to favor a value that is slightly higher than the standard WMAP5 value. However, the difference is less than $2-\\sigma$ at both the full sample and the low redshift slice $z = 0.15 - 0.30$, but is larger at the high redshift slice $z = 0.40 - 0.47$ (see Fig. \\ref{fig:fit.slice}). We have also found this tendency to high $\\Omega_b$ in paper III \\citep{paper3} of this series, where we analyze the BAO scale in the  3-point correlation function, which tests non-linear growth of perturbations, in contrast to the 2-point correlation function, that tests linear perturbations. Plots of the signal and the signal-to-noise of $\\xi$ in the $\\sigma-\\pi$ plane (Fig. \\ref{fig:pisigma} and \\ref{fig:s2nps.11}) provide a reassuring view of the reality of the signal: a negative valley (in blue) of negative correlations at $\\pi \\sim 50 - 90$ Mpc/h together with a ring (in red) of positive correlations at $\\sim 110$ Mpc/h. The BAO peak in the LOS direction is pronounced, and detected with significance (Fig. \\ref{fig:losdirectionmodel2} and \\ref{fig:baos2nall}).  It would be hard to explain such a peak in the 2-point monopole, in the plane $\\sigma-\\pi$, in the LOS and in the 3-point function with systematic effects. We have checked for possible systematics by varying the angular mask and the radial selection function, and find the detection robust using either DR6 or DR7 SDSS data (see \\S \\ref{sec:SystematicsTests}). The data on scales $\\pi = 40 - 140$ Mpc/h is well fit by a model that includes linear redshift space distortions, magnification bias and BAO. In fact, a model with no BAO is disfavored by the data at the $3 \\sigma$ level, while a model with no magnification bias is disfavored at the $2 \\sigma$ level. Our detection of the BAO modulation in the LOS direction is helped by the combination of two effects: redshift space distortions make $\\xi$ negative on scales of $50 - 90$ Mpc/h while magnification bias gives a significant positive boost on larger scales.  We thus have significant evidence that the LOS correlation on scales $40 <\\pi <140$ Mpc/h reproduces well the expected BAO signal. Consequently we use the LOS data to infer $H(z)$ in two ways which we have labeled the shape and the peak method. In the shape method we have used the shape of the LOS correlation to fit $H(z)/H_0$ marginalized over other cosmological parameters. The result is shown as the first entry to Table \\ref{tab:rbao2}. In the peak method we find the location of the peak position and use it as a standard ruler using Eq.\\ref{eq:obtentionh}. This is a more direct, geometric test in the sense that the BAO is used strictly as a standard ruler (i.e. we are not using model-dependent shape of the correlation function), and in the sense that we constrain $H(z)$ rather than integrals thereof. The results are shown as 3 bottom entries of Table \\ref{tab:rbao2}.   There are two parts to our error analysis. One is the errors on $\\xi(\\sigma, \\pi)$ itself, and the other is the errors on our measurement of the peak location. Both have been extensively tested with simulations. The error model for $\\xi$ is obtained from what is to our knowledge the largest ever cosmological simulation run to date, MICEL7680 with 453 ${\\,\\rm Gpc^3/h^3}$, $2048^3$ dark matter particles and 107 million halos in a single box \\citep{mice,mice2}. We are able to create out of this simulation 216 independent mock LRG catalogs, each of which has a similar size to our SDSS DR6 sample. Our mock catalogs have the same number density and very similar 2- and 3-point functions compared to data. Fig.\\ref{fig:xigmu3} shows, for the first time on BAO scales, a comparison of the redshift space LOS linear model with non-linear measurements from simulations. Both model and simulation show very similar shapes and a prominent BAO peak which validates both our modeling and our simulated mocks. The error model constructed from these simulations is further validated by a comparison with jack-knife errors obtained from the data itself. Further details can be found in \\S \\ref{sec:nonstandard} and Fig.\\ref{fig:errorvalid2}.  The other part of our error analysis consists of using Monte Carlo simulations to simulate our peak measurement process. This is described in \\S \\ref{sec:staterror}. This method involves a further systematic uncertainty: because the data is quite noisy we can not be certain of what is the true model for the mean of the Monte Carlo simulations. We have used the data itself (or a smooth version of it) to measure both the statistical and systematic errors on these measurements. Thus, these errors are more model independent than the ones in the shape method, which depend strongly on the range of scales used for the fit.  \\begin{figure} \\centering{ \\epsfysize=6.5cm\\epsfbox{fitwOlHz.eps}} \\caption{ The dotted line and continuous (black) open contours across the figure show the best fit and 1-sigma contour from $q[z=0.34]$ in Eq.\\ref{eq:qz}, ie using only our two independent $H(z)$ measurements. Long-dashed line contour show $\\Delta\\chi^2=6.18$ (2-sigma for 2 dof) constraints from combining our $H(z)$ measurements with $H_0$ from HST and our $\\Omega_m$ estimate from redshifts distortions (Paper I). Short-dashed contours (blue) show the  corresponding $\\Delta\\chi^2=6.18$  constraints from the external data-set WMAP5, SNIa and BAO monopole. The inner red continuous contour shows  $\\Delta\\chi^2=1$ when combining both measurements. \\label{fig:fitwOl}} \\end{figure}  Given our measurements of $H(z)$, it is natural to ask: what constraint do they put on the dark energy abundance and equation of state? Our measurements of $H(z)$ at $z = 0.24$ and $z = 0.43$ by themselves only weakly constrain $\\Omega_{\\rm DE}$ and $w$. The constraint comes essentially from the acceleration parameter $q(z)$ in Eq.\\ref{eq:qz}. The mean and 1-sigma constrains on $\\Omega_{\\rm DE} - w$ from $q(z)$ are shown as  dotted and continuous black open contours across Fig.\\ref{fig:fitwOl}. These constraints are very weak as compared to the constraints from  WMAP5 + SN + monopole BAO \\cite{wmap5}, which are shown for $\\Delta\\chi^2=6.18$ as a long-dashed (blue) line contour in Fig.\\ref{fig:fitwOl}. The large short-dashed (black) line contour corresponds to the 2-sigma constraints when we combine our $H(z)$ measurements with $H_0 = 70 \\pm 5$ km/s/Mpc, based on the HST analysis \\cite{HST2007}, and our best fit $\\Omega_m=0.245 \\pm 0.020$ using the same LRG data in Paper I. The 2-sigma contours from the two independent data sets agree well, which is not trivial specially given the different assumptions and measurements involved in each data set. Combining both data sets, we find \\begin{eqnarray} w &=& -0.957 \\pm 0.053 \\\\ \\Omega_{DE} &=& 0.734 \\pm 0.023 \\end{eqnarray} which is shown as red continuous contour in Fig.\\ref{fig:fitwOl}. This is in excellent agreement with the cosmological constant model. The improvement in the error with respect to the 1-sigma WMAP5+SN+BAO value of $w = -0.97 \\pm 0.06$ and is quite modest. A separate paper  presents the implication on cosmological parameters when using our radial BAO scale measurements in a way that is independent of $H_0$ \\citep{paper5}.   \\begin{figure} \\centering{\\epsfysize=6.5cm\\epsfbox{sigmapiPAU.eps}} \\caption{ $\\xips$ in data for the z=0.15-0.30 slice recalculated from real LRG positions with an added photo-z error of $0.003(1+z)$ as predicted for the PAU survey. It should be compared with middle panel of Fig.\\ref{fig:pisigma} which uses spectroscopic redshifts. \\label{fig:pisigmaPAUb} } \\end{figure}   To illustrate what might be achievable in the near future, we show in Fig. \\ref{fig:pisigmaPAUb} the expected $\\xi(\\sigma,\\pi)$ for the PAU survey \\citep{PAU} which has a photo-z accuracy of $0.003(1+z)$. We redo the analysis using current SDSS data (slice z=0.15-0.30) but we disperse the LRG positions by this photo-z error (we assume a Gaussian distribution with $0.003(1+z)$ dispersion). The BAO signal does not appear to be washed away, as shown in Fig. \\ref{fig:pisigmaPAUb}. This is expected because the BAO peak is broader than the scale corresponding to the photo-z dilution. The PAU Survey proposes to map over 10 times the SDSS DR6 volume, ie to z=0.9 (this is a factor of 3 in sampling variance error) with 20 times the LRG number density (ie for $L>L_*$) so shot noise will be negligible. All this should increase the signal-to-noise over SDSS DR6 by better than a factor of four.  Thus, there is potential for a substantial improvement in the near future for the $H(z)$ measurements using the techniques presented here."
    },
    "0807/0807.4237_arXiv.txt": {
        "abstract": "We present our recent results on the cosmic evolution of the outskirst of disk galaxies. In particular we focus on disk--like galaxies with stellar disk truncations. Using UDF, GOODS and SDSS data we show how the position of the break (i.e. a direct estimator of the size of the stellar disk) evolves with time since z$\\sim$1. Our findings agree with an evolution on the radial position of the break by a factor of 1.3$\\pm$0.1 in the last 8 Gyr for galaxies with similar stellar masses. We also present radial color gradients and how they evolve with time. At all redshift we find  a radial inside-out bluing reaching a minimum  at the position of the break radius, this minimum is followed by a reddening outwards. Our results constraint several galaxy disk formation models and favour a scenario where stars are formed inside the break radius and are relocated in the outskirts of galaxies through secular processes. ",
        "introduction": "Early studies of the disks of spiral galaxies (\\cite[Patterson 1940]{Patterson40}, \\cite[de Vaucouleurs 1959]{deVaucouleurs59}, \\cite[Freeman 1970]{Freeman70}) showed that this component generally follows an exponential radial surface-brightness profile, with a certain scale length, usually taken as the characteristic size of the disk. \\cite[Freeman (1970)]{Freeman70} pointed out, though, that not all disks follow this simple exponential law. In fact, a repeatedly reported feature of disks for a representative fraction of the spiral galaxies is that of a truncation of the stellar population at large radii, typically 2-4 exponential scale lengths (see e.g. the review by \\cite[Pohlen et al. 2004]{Pohlen04}).  Several possible break-forming mechanisms have been investigated to explain the truncations. There have been ideas based on maximum angular momentum distribution: \\cite[van der Kruit (1987)]{Kruit87} proposed that angular momentum conservation in a collapsing, uniformly rotating cloud naturally gives rise to disk breaks at roughly 4.5 scale radii. \\cite[van den Bosch (2001)]{Bosch01} suggested that the breaks are due to angular momentum cut-offs of the cooled gas. On the other hand, breaks have also been attributed to a threshold for star formation (SF), due to changes in the gas density \\cite[Kennicutt (1989)]{Kennicutt89}, or to an absence of equilibrium in the cool Interstellar Medium phase (\\cite[Elmegreen \\& Parravano 1994]{ElmegreenParravano94}, \\cite[Schaye 2004]{Schaye04}). Magnetic fields have been also considered (Battaner et al. 2002) as responsible of the truncations. More recent models using collisionless N-body simulations, such as that by \\cite[Debattista et al. (2006)]{Debattista06}, demonstrated that the redistribution of angular momentum by spirals during bar formation also produces realistic breaks. In a further elaboration of this idea, \\cite[Ro\\v{s}kar et al. (2008)]{Roskar08} have performed high resolution simulations of the formation of a galaxy embedded in a dark matter halo. In these models, breaks are the result of the interplay between a radial star formation cut-off and redistribution of stellar mass by secular processes. A natural prediction of these models is that the stellar populations present an age minimum in the break position. This prediction could be probed by exploring the color profiles of the galaxies.   Furthermore, addressing the question of how the radial truncation evolves with z is strongly linked to our understanding of how the galactic disks grow and where star formation takes place. \\cite[P\\'erez (2004)]{Perez04} showed that it is possible to detect stellar truncations even out to z$\\sim$1. Using the radial position of the truncation  as a direct estimator of the size of the stellar disk, \\cite[Trujillo \\& Pohlen (2005)]{TP05} inferred a moderate ($\\sim $25\\%) inside-out growth of disk galaxies since z$\\sim$1. An important point, however was missing in the previous analyses: the evolution with redshift of the radial position of the break at a given stellar mass. The stellar mass is a much better parameter to explore the growth of galaxies, since the luminosity evolution of the stellar populations can mimic a size evolution (\\cite[Trujillo et al. 2004]{Trujillo04}, \\cite[Trujillo et al. 2006]{Trujillo06}). We present in this contribution a quick summary of our recent findings on the stellar disk truncation origin and its evolution with redshift. The results presented here are based on the following publications: \\cite[Azzollini et al. (2008a)]{Azzollini08a}, \\cite[Azzollini et al. (2008b)]{Azzollini08b} and \\cite[Bakos et al. (2008)]{Bakos08}. Throughout, we assume a flat $\\Lambda$-dominated cosmology ($\\Omega_{M}$ = 0.30, $\\Omega_{\\Lambda}$=0.70, and $H_{0}$=70 km $s^{-1}$ $Mpc^{-1}$). ",
        "conclusions": ""
    },
    "0807/0807.3417_arXiv.txt": {
        "abstract": "{Thermal regime of the baryons behind shock waves arising in the process of virialization of dark matter halos is governed at cetrain conditions by radiation of HD lines. A small fraction of the shocked gas can cool down to the temperature of the cosmic microwave background (CMB). We estimate an upper limit for this fraction: at $z=10$ it increases sharply from about $q_{_T}\\sim 10^{-3}$ for dark halos of $M=5\\times 10^7\\msun$ to $\\sim 0.1$ for halos with $M=10^8\\msun$. Further increase of the halo mass does not lead however to a significant growth of $q_T$ -- the asymptotic value for $M\\gg 10^8\\msun$ is of 0.3: We estimate star formation rate associated with such shock waves, and show that they can provide a small but not negligible fraction of the star formation. We argue that extremely metal-poor low-mass stars in the Milky Way may have been formed from primordial gas behind such shocks. } ",
        "introduction": "Formation of the first stars is determined by energy losses in ro-vibrational lines of molecular hydrogen H$_2$ and its deuterated analogue HD (Lepp \\& Shull 1983; Shchekinov 1986; Puy et al. 1993; Palla, Galli \\& Silk 1995; Galli \\& Palla 1998, 2002; Stancil, Lepp \\& Dalgarno 1998; Tegmark et al. 1997; Puy \\& Signore 1997, 1998; Uehara \\& Inutsuka 2000; Flower 2002; Nakamura \\& Umemura 2002; Machida et al. 2005, Nagakura \\& Omukai 2005). In turn, the amount of H$_2$ and HD and their ability to cool gas greatly depend on dynamical and thermal regime of the gas. In particular, shock waves strongly enhance the rate of conversion of atomic hydrogen to its molecular form (Shchekinov \\& Entel 1983; Suchkov, Shchekinov \\& Edelman 1983; Shapiro \\& Kang 1987; Kang \\& Shapiro 1992; Shchekinov 1991; Ferrara 1998; Yamada \\& Nishi 1998; Uehara \\& Inutsuka 2000; Cen 2005; Machida et al. 2005; Shchekinov \\& Va\\-si\\-liev 2005; Johnson \\& Bromm 2005). On the other hand, the first stars have formed from the gas processed by the shock waves unavoidable in the process of virialization of dark matter halos (Shapiro 1993; Haiman, Thoul \\& Loeb 1996; Tegmark et al. 1997; Abel, Bryan \\& Norman 2000, 2002). Therefore possible enhancement of H$_2$ and HD in these conditions can have important consequences for characteristics of the first stars (Oh \\& Haiman 2002; Johnson \\& Bromm 2005; Shchekinov \\& Vasiliev 2005).  When dark matter halos merge, shock waves form and compress the baryons (see discussion in (Barkana \\& Loeb 2001; Ciardi \\& Ferrara 2004)). At sufficiently large velocities of colliding flows ($v>8$ km s$^{-1}$) fractional ionization in shocked gas increases above the frozen cosmological value, and speeds up chemical kinetics of H$_2$ molecules. Collisions with velocities above $v>8.6[(1+z)/20]^{-1/6}$ km s$^{-1}$ may lead to a rapid formation of HD and an efficient cooling down to the minimum temperature $T=T_{\\rm CMB}=2.7(1+z)$ (Vasiliev \\& Shchekinov 2005; Johnson \\& Bromm 2005; Shchekinov \\& Vasiliev 2005). It is worth noting that HD can also play major role in central regions of collapsing dark halos, where a high-density environment favours transition of D into molecular form, so that even in low-mass halos HD cooling can dominate (Ripamonti 2007).  Explicit answer of how large is the fraction of cold gas in merging halos can be found only in hydrodynamic simulations. Very recently (Maio et al. 2007) performed 3D SPH simulation of structure formation with accounting of HD chemistry. They found that incorporation of effects from HD cooling increases clumpiness of gas, i.e. clouds are den\\-ser and more compact with respect to the case when only H$_2$ cooling is considered. However, full 3D simulations are time consuming, and always are made within a particular realization of a random hydrodynamic field. They therefore represent only very limited regions in the space of possible random hydrodynamic fields of a given spectrum. This circumstance restricts the final thermodynamic state of bary\\-ons within the domain corresponding to the chosen interelations between the amplitudes of different wave modes. As a result, estimates of the fraction of the cold baryons able to form stars are biased by such limitations. At such circumstances 1D hydrodynamical simulations of chemistry and thermodynamics of shocked primordial gas can play an auxiliary role for understanding of what in principle can be expected in the conditions preceeding formation of the first stars. In this paper we use 1D planar computations to model chemical and thermal regime of gas behind the shock waves after a head-on collision of two clouds of equal sizes. This approach is justified by the fact that in supersonic cloud collisions the rarefaction motion transverse to the symmetry axis takes longer  than the collision time $t_c=3R/2v_c$ (Gilden 1984). On the other hand, the estimates of the fraction of baryons cooled down by HD molecules we obtaine within this approximation are obviously upper limits.  At the stages when the halos are close to merge their velocities are of the order of the virial value corresponding to the total mass of the halos. We therefore connect the collisional velocities of the clouds to the mass of the halo $M$ formed in this collision \\begin{equation} \\label{vel} v_c=(3^4\\pi^3\\Omega_m\\rho_c)^{1/6}G^{1/2}M^{1/3}(1+z)^{1/2}, \\end{equation} $\\Omega_m$ is the matter closure parameter, $\\rho_c=3H_0^2/8\\pi G$, the critical density.  In Section 2 we describe dominant thermal and chemical processes in shocked gas; in Section 3 we discuss consequences of HD cooling for formation of the first stars; summary is given in Section 4.  Throughout the paper we assume a $\\Lambda$CDM cosmology with the parameters $(\\Omega_0, \\Omega_{\\Lambda}, \\Omega_m, \\Omega_b, h ) = (1.0,\\ 0.7,\\ 0.3,\\\\ 0.045,\\ 0.7 )$ and deuterium abundance $2.6\\times 10^{-5}$, consistent with the most recent measurements (Spergel et al. 2007). ",
        "conclusions": "In this paper we have shown that  \\begin{itemize}  \\item  A small fraction of baryons in merging halos can cool down to very low temperatures close to the temperature of the cosmic microwave background. This fraction increases with the halo mass, and can reach $\\simeq 0.2$ for masses $M>2\\times 10^7~M_\\odot$ at $z=20$ and for $M>10^8~M_\\odot$ at $z=10$; averaged over the halo mass spectrum this fraction varies from 10$^{-3}$ at $z=20$ to 0.1 at $z=10$.  \\item  Such cold gas is unstable against gravitational fragmentation, with the mass of the primary fragments decreasing for mergers of higher masses $M$: $M_J\\leq 1.3\\times 10^4 \\\\ M^{-1/3}M_\\odot$. Masses of protostars formed in gas cool\\-ed by HD molecules are in general lower than those formed in conditions when H$_2$ cooling dominates.  \\item  The contribution to the cosmic star formation rate of the mergers with a predominance of HD cooling, and therefore with presumably low mass first stellar objects, increases from less than 0.5\\% at redshift $z=18-20$ to 10-30 \\% at $z=10$. Extremally metal-poor low mass stars in the Milky Way may have been formed in mergers dominated by HD cooling.  \\end{itemize}"
    },
    "0807/0807.0414_arXiv.txt": {
        "abstract": "{We suggest a scenario where the three light quark flavors are sequentially deconfined under increasing pressure in cold asymmetric nuclear matter as found, e.g., in neutron stars.} The basis for our analysis is a chiral quark matter model of Nambu--Jona-Lasinio (NJL) type with diquark pairing in the spin-1 single flavor (CSL), spin-0 two flavor (2SC) and three flavor (CFL) channels. We find that nucleon dissociation sets in at about the saturation density, $n_0$, when the down-quark Fermi sea is populated (d-quark dripline) due to the flavor asymmetry induced by $\\beta$-equilibrium and charge neutrality. At about $3n_0$ u-quarks appear and a two-flavor color superconducting (2SC) phase is formed. The s-quark Fermi sea is populated only at still higher baryon density, when the quark chemical potential is of the order of the dynamically generated strange quark mass. We construct two different hybrid equations of state (EoS) using the Dirac-Brueckner Hartree-Fock (DBHF) approach and the EoS by Shen \\etal in the nuclear matter sector. The corresponding hybrid star sequences have maximum masses of, respectively, 2.1 and 2.0 M$_\\odot$. Two- and three-flavor quark-matter phases exist only in gravitationally unstable hybrid star solutions in the DBHF case, while the Shen-based EoS produce stable configurations with a 2SC phase-component in the core of massive stars. {Nucleon dissociation via d-quark drip could act as a deep crustal heating process, which apparently is required to explain superbusts and cooling of X-ray transients.} ",
        "introduction": "The phenomenology of compact stars is intimately connected to the EoS of matter at densities well beyond the nuclear saturation density, $n_0=0.16$ fm$^{-3}$. Compact stars are therefore natural laboratories for the exploration of baryonic matter under extreme conditions, complementary to those created in terrestrial experiments with atomic nuclei and heavy-ion collisions. Recent results derived from observations of compact stars provide serious constraints on the nuclear EoS, see \\cite{Klahn:2006ir} and references therein. A stiff EoS at high density is needed to explain the high compact-star masses {% and radii, for which there is growing evidence from recent observations. A mass of $M\\sim 2.0$~M$_\\odot$ has been reported for some low-mass X-ray binaries (LMXBs)\\eg 4U 1636-536 \\cite{Barret:2005wd}, based on the assumption that the abrupt drop in the coherence of the lower kilohertz quasiperiodic oscillation (QPO) may be related to the innermost stable circular orbit, see also  \\cite{Barret:2007df}. From observations of the bright isolated neutron star RX J1856.5-3754 (shorthand: RX J1856) in the optical and X-ray frequency range, a conservative lower limit of the apparent neutron star radius of $R_\\infty=16.5$ km is derived \\cite{Trumper:2003we}. This corresponds to a true (de-redshifted) radius of $R=14$ km for a 1.4 M$_\\odot$ neutron star, or equivalently, to a star mass of at least 2.1 M$_\\odot$ when the radius does not exceed 12 km \\cite{Klahn:2006ir}. } Another example is EXO 0748-676, an LMXB for which the compact-star mass {\\it and} radius have been constrained to $M\\ge 2.10\\pm 0.28$~M$_\\odot$ and $R \\ge 13.8 \\pm 0.18$~km \\cite{Ozel:2006km} {% by a simultaneous measurement of the Eddington limit, the gravitational redshift, and the flux of thermal radiation.} However, the status of the results for the latter object is unclear, because the gravitational redshift $z=0.35$ observed in the X-ray burst spectra \\cite{Cottam:2002} has not been confirmed, despite numerous attempts. {% Further constraints on the masses and radii of compact stars have been reported (see \\eg \\cite{Leahy:2007fb,Leahy:2008cq}), but they deserve a careful discussion which is beyond the scope of the present paper.} While compact-star phenomenology apparently points towards a stiff EoS at high density, heavy-ion collision data for kaon production \\cite{Fuchs:2005zg} and elliptic flow \\cite{Danielewicz:2002pu} set an upper limit on the stiffness of the EoS \\cite{Klahn:2006ir}.  A key question regarding the structure of matter at high density is whether a phase transition to quark matter occurs inside compact stars, and whether it is accompanied by unambiguous observable signatures. It has been argued  that the observation of a compact star with high mass and large radius, as reported for EXO 0748-676, would be incompatible with a quark core \\cite{Ozel:2006km}, because quark deconfinement softens the EoS and lowers the maximum mass and corresponding radius. % However, Alford \\etal have demonstrated \\cite{Alford:2006vz} with a few counter examples that quark matter cannot be excluded by this argument. In particular, for a recently developed hybrid star EoS \\cite{Klahn:2006iw}, based on the DBHF approach in the nuclear sector and a three-flavor chiral quark model \\cite{Blaschke:2005uj}, stable hybrid stars with masses ranging from 1.2 M$_\\odot$ to 2.1 M$_\\odot$ is obtained, in accordance with modern mass-radius constraints, see also \\cite{Blaschke:2007ri}. In this model, a sufficiently low critical density for quark deconfinement has been achieved by a strong diquark coupling, while a repulsive vector meanfield in the quark matter sector resulted in sufficient stiffness to achieve a high maximum mass of the compact star sequence. The corresponding hybrid EoS for symmetric matter was shown to fulfill the constraints derived from elliptic flow in heavy-ion collisions. In the present work we discuss a new scenario that comprises a sequential transition from nuclear matter to deconfined quark matter, which could play an important role in asymmetric matter, in particular for the phenomenology of compact stars.  Chiral quark models of NJL type with dynamical chiral-symmetry breaking have the property that the symmetry is restored (and quarks are deconfined) separately for each flavor. When solving the gap and charge-neutrality equations self-consistently, the chiral-symmetry restoration for a given flavor occurs when the chemical potential of that flavor reaches a critical value that is approximately equal to the dynamically generated quark mass, $\\mu_{f}=\\mu_{c,f} \\approx m_{f}$,  where $f=u, d, s$. In asymmetric matter the quark chemical potentials are different. Consequently, the NJL model behavior suggests that the critical density of deconfinement is flavor dependent, see Fig.~\\ref{f:mucrit}. \\begin{figure} % \\includegraphics[angle=0,width=0.9\\linewidth,clip=]{d-drip3.eps} \\caption{(Color online) Schematic picture of chemical potentials (columns) and sequential deconfinement of quarks with increasing baryon density (from left to right). The flavor dependent thresholds for chiral symmetry restoration (deconfinement) are approximately given by the dynamically generated quark masses $m_f$, $f=u, d, s$ (dashed lines). With increasing quark chemical potential, $\\mu=(\\mu_u+\\mu_d)/2$, the d-quark chemical potential is the first to reach the threshold in isospin asymmetric matter. Nucleon dissociation therefore sets in as d-quarks are deconfined. Still higher $\\mu$ is needed to form 2-flavor and 3-flavor quark matter phases. \\label{f:mucrit}} \\end{figure} In this scenario the down quark flavor is the first to drip out of nucleons when the density increases, followed by the up quark flavor and eventually also by strange quarks. This behavior is absent in simple and commonly applied % bag-model equations of state, because they are essentially flavor blind.  Under the $\\beta$-equilibrium condition in compact stars the chemical potentials of quarks and electrons are related by $\\mu_d=\\mu_s$ and $\\mu_d=\\mu_u+\\mu_e$. The mass difference between the strange and the light quark flavors $m_s \\gg m_u, m_d$ has two consequences: 1) the down and strange quark densities are different, so charge neutrality requires a finite electron density and, consequently, 2) $\\mu_d>\\mu_u$. When increasing the baryochemical potential the d-quark chemical potential is therefore the first to reach the critical value, $\\mu_{c,d}\\approx m_d$, where the chiral symmetry gets (approximately) restored in a first-order transition and deconfined d-quarks appear. Due to the finite value of $\\mu_e$ the u-quark chemical potential is still below $\\mu_{c,u}\\approx m_u$, while the s-quark density is zero due to  the high s-quark mass. A {\\it single-flavor} d-quark phase therefore forms in co-existence with the positively charged nuclear-matter medium.  Why has this interesting scenario been left unnoticed? One reason is that bag models, which are commonly used to describe quark matter in compact star interiors cannot address sequential deconfinement. Another reason is that the single-flavor d-quark phase is negatively charged and cannot be neutralized with leptons. % It was therefore disregarded in dynamical approaches like NJL models which in practice are used mainly to describe the deconfined and ``pure'' quark matter phase only. In the following we discuss the single-flavor phase for the first time, under the natural assumption that the neutralizing background is nuclear matter. Since nucleons are bound states of quarks, a mixed phase of nucleons and free d-quarks could naturally arise when nucleonic bound states dissociate (Mott effect). ",
        "conclusions": "In this paper we suggest a new quark-nuclear hybrid model for compact star applications that fulfills modern constraints from observations of compact stars. Due to isospin asymmetry, down-quarks may ``drip out'' from nucleons and form a single-flavor color superconducting (CSL) phase that is mixed with nuclear matter already at the crust-core boundary in compact stars. The CSL phase has interesting cooling and transport properties that are in accordance with constraints from the thermal and rotational evolution of compact stars \\cite{Blaschke:2007bv}. It remains to be investigated whether this new compact star composition could lead to unambiguous observational consequences, % and whether it is thermodynamically favored also when effects like Coulomb screening and surface tension are accounted for. We conjecture that the d-quark drip may serve as an effective deep crustal heating mechanism for the explanation of the puzzling superburst phenomenon and the cooling of X-ray transients."
    },
    "0807/0807.0622_arXiv.txt": {
        "abstract": "The transfer function $T(k)$ of dark matter (DM) perturbations during matter domination is obtained by solving  the linearized collisionless Boltzmann-Vlasov equation.  We provide an \\emph{exact} expression for $T(k)$ for \\emph{arbitrary distribution functions of decoupled particles and initial conditions}, which can be systematically expanded in a Fredholm series. An exhaustive numerical study of thermal relics for different initial conditions  reveals that the first \\emph{two} terms in the expansion of $T(k)$ provide a remarkably accurate and simple approximation valid on all scales of cosmological relevance for structure formation in the linear regime. The natural scale of suppression is the free streaming wavevector at matter-radiation equality, $   k_{fs}(t_{eq}) = \\left[{4\\pi\\rho_{0M}}/{[\\langle \\vec{V}^2\\rangle\\,(1+z_{eq})]} \\right]^\\frac12 $. An important ingredient is a non-local kernel determined by the distribution functions of the decoupled particles which describes the \\emph{memory of the initial conditions and gravitational clustering} and yields a correction to the fluid description. This correction is negligible at large scales $k \\ll k_{fs}(t_{eq})$ but it becomes important at small scales $k \\geq k_{fs}(t_{eq})$. Distribution functions that favor the small momentum region yield longer-range memory kernels and lead to an \\emph{enhancement of power at small scales} $ k > k_{fs}(t_{eq}) $. Fermi-Dirac and Bose-Einstein statistics lead to long-range memory kernels, with longer range for bosons, both resulting in enhancement of $T(k)$ at small scales. For DM thermal relics that decoupled while ultrarelativistic we find $ k_{fs}(t_{eq}) \\simeq  0.003 \\; (g_d/2)^\\frac13 \\; (m/\\mathrm{keV}) \\; [\\mathrm{kpc}]^{-1} $, where $ g_d $ is the number of degrees of freedom at decoupling.   For WIMPS we obtain $ k_{fs}(t_{eq}) = 5.88 \\; (g_d/2 \\; )^\\frac13 \\; (m/100\\,\\mathrm{GeV} )^\\frac12 \\; (T_d/10 \\,\\mathrm{MeV})^\\frac12 \\; [\\mathrm{pc}]^{-1} $. For $k\\ll k_{fs}(t_{eq})$, $T(k) \\sim 1- \\mathrm{C}[k/k_{fs}(t_{eq})]^2 $ where $C =\\mathrm{O}(1)$ and independent of statistics for thermal relics. We provide simple and accurate fits for $T(k)$ in a wide range of \\emph{small} scales $ k > k_{fs}(t_{eq})$ for thermal relics and different initial conditions.   The   numerical and analytic results    for {\\it arbitrary} distribution functions and  initial conditions   allow an   assessment of  DM candidates through their impact on structure formation. ",
        "introduction": "The \\emph{concordance} $\\Lambda\\mathrm{CDM}$ standard cosmological model successfully explains   a wide range of highly precise astrophysical and cosmological observations. Main ingredients are an early stage of accelerated expansion   (inflation) a more recent stage of accelerated expansion driven by dark energy,  and the presence of dark matter (DM) composed of primordial particles which are cold and collisionless at the time when the first structures formed\\cite{dodelson,kt}.  \\medskip  In the cold dark matter (CDM) model, structure formation proceeds in a hierarchical bottom up approach: small scales become non-linear and collapse first and their merger and accretion leads to structure on larger scales. This is a consequence of the fact that CDM features negligible small velocity dispersions leading to a power spectrum that favors small scales. In this hierarchical scenario, dense clumps that survive the merger process form satellite galaxies.  Large scale $\\Lambda\\mathrm{CDM}$ simulations seemingly lead to an overprediction of satellite galaxies \\cite{moore2} by almost an order of magnitude over the number of satellites that have been observed in Milky-Way sized galaxies \\cite{kauff,moore,moore2,klyp,will}. These simulations   also yield a distinct prediction: virialized DM halos should feature a  density profile that increases monotonically towards the center \\cite{frenk,dubi,moore2,bullock,cusps} such as the Navarro-Frenk-White (NFW) profile \\cite{frenk} or more general central density profiles $ \\rho(r) \\sim r^{-\\beta} $ with $ 1\\leq \\beta \\lesssim 1.5 $ \\cite{moore,frenk,cusps}. These profiles accurately describe clusters of galaxies but indicate a divergent cusp at the center of the halo.  There is, however, an accumulating body of observational evidence \\cite{dalcanton1,van,swat,gilmore,salucci,battaglia,cen,wojtak} that seem to indicate that the density profile in the central regions of dwarf galaxies is smooth leading to the suggestion that the central regions feature smooth cores instead of cusps as predicted by CDM.  \\medskip  Warm dark matter (WDM) particles were invoked \\cite{mooreWDM,turokWDM,avila} as possible solutions to the above mentioned discrepancies both in the overabundance of satellite galaxies and as a mechanism to smooth out the cusped density profiles predicted by CDM simulations into the  cored profiles seemingly observed in dwarf-spheroidal satellite galaxies (dShps). WDM particles feature a range of velocity dispersion in between the CDM and hot dark matter (HDM) leading to free streaming scales that smoothes out small scale features and could be consistent with core radii of the dSphs.   Even if these problems of the CDM model find other astrophysical solutions, it remains an important and intrinsically interesting question to assess the clustering properties of WDM  candidates, since these may emerge from extensions beyond the standard model of particle physics. Furthermore, it remains an important aspect to understand possible departures of the CDM paradigm at small scales provided by alternative particle physics candidates.    \\medskip  The gravitational clustering properties of collisionless DM in the linear regime are described by the   power spectrum of gravitational perturbations and in particular the transfer function that converts the primordial power spectrum into the power spectrum later in the matter dominated era, which in turn,  is the input in the large scale numerical simulations of galaxy formation. Free streaming \\cite{kt} of collisionless DM leads to a suppression of the transfer function on length scales smaller than the free streaming scale via Landau damping \\cite{bondszalay}.  \\medskip  Perturbations in a collisionless system of particles with gravitational interactions is fundamentally different from fluid perturbations in the presence of gravity. The (perfect) fluid equations correspond to the limit of vanishing mean free path. In a gravitating fluid, pressure gradients tend to restore hydrostatic equilibrium with the speed of sound in the medium, and short wavelength fluctuations are simple acoustic waves. For large wavelengths, the propagation of pressure waves cannot halt gravitational collapse on a dynamical time scale. The dividing line is the Jeans length: perturbations with  shorter wavelengths oscillate as sound waves, while perturbations with longer wavelengths undergo gravitational collapse.  In a gas of collisionless particles with gravitational interaction the situation is different since the mean free path is much larger than the size of the system (Hubble radius) and the fluid description is not valid. Instead, the collisionless Boltzmann-Vlasov  equation for the distribution function must be solved to obtain the dynamics of perturbations \\cite{peebles,bert,dodelson}.  Just as in the case of plasma physics, the linearized Boltzmann-Vlasov equation describes \\emph{collective} excitations \\cite{plasmas}. In the case of a collisionless gas with gravitational interactions these collective excitations describe particles free-streaming in and out of the gravitational potential wells of which they are the source. The damping of short wavelength \\emph{collective} excitations is akin to \\emph{Landau damping in plasmas} \\cite{plasmas}: it is a result of \\emph{dephasing} via phase mixing when the particles are out of phase with the potential wells that they produce \\cite{bt} and leads to the collisionless damping of the collective modes \\cite{trema1}.  \\medskip  Gilbert \\cite{gilbert} studied the linearized Boltzmann-Vlasov equation in a matter dominated cosmology for non-relativistic particles described by an (unperturbed) Maxwell-Boltzmann distribution function. This equation was cast as a Volterra integral equation whose numerical solution revealed a limiting value of the wavevector above which density perturbations are Landau  damped and below which they grow via gravitational (Jeans) instability \\cite{gilbert}.   The results were consistent with replacing the speed of sound by the Maxwellian velocity dispersion in the Jeans length [up to a normalization factor of $\\mathcal{O}(1)$]. Gilbert's equations were solved numerically to study: (i) the collisionless damping of density fluctuations in an expanding cosmology with massive neutrinos with an approximated Fermi-Dirac distribution function \\cite{bondszalay}; (ii) the dissipationless clustering of neutrinos with a truncation of the Fermi-Dirac distribution function and an  analytic fit to the numerical solution of the integral equation \\cite{bran}. These results were also used to analyze the linear regime \\cite{turokWDM}; (iii) cosmological perturbations and massive light neutrinos \\cite{bertwatts,bert}, and more recently  similar integral  equations were solved  approximately for thermal neutrino relics \\cite{ringwald}.  \\medskip  DM may be composed of several species \\cite{palazzo}, a possibility that can be   accommodated in most extensions of the Standard Model, with candidates ranging from weakly interacting massive particles (WIMPs) \\cite{kt} to axions \\cite{kt} and sterile neutrinos \\cite{dw,este,kusenko}. DM candidates may be produced in the early Universe via different mechanisms and some of them are conjectured to decouple and freeze-out with non-thermal distribution functions. Sterile neutrinos are an example of this DM candidate \\cite{dw,este,kusenko}. It is important to understand the clustering properties of possible DM particles that decoupled \\emph{in or out} of local thermal equilibrium. Predictions for the effective free-streaming length and transfer function are a necessary ingredient in the study of structure formation. The free-streaming length in such \\emph{mixture} was recently investigated in ref. \\cite{freestream} for \\emph{arbitrary} distribution functions of relic particles that decoupled in or out local thermodynamic equilibrium (LTE). This study analyzed the Boltzmann-Vlasov equation in Minkowski space-time focusing on the marginally unstable collective modes. It revealed important (and intriguing) aspects of the free streaming lengths associated with the distribution function of the decoupled  particles.  \\medskip  A program that bridges the microphysics of the production, evolution and decoupling of DM particles with large scale structure formation begins by obtaining the distribution function of the DM particles from the dynamics of production and decoupling. The DM distribution function determines the abundance \\cite{kt}, primordial phase space properties and generalized constraints on the DM candidate \\cite{hogan,coldmatter} and their free streaming lengths \\cite{freestream}. Results in this direction have been reported \\cite{coldmatter,kupesha}. The primordial phase space density of particles  $ \\mathcal D $  permits to obtain deep insights on DM \\cite{coldmatter}.  The phase space density can only decrease during gravitational dynamics via mergers and ``violent relaxation''\\cite{tg}. Recent photometric and kinematic data on dwarf spheroidal satellite galaxies in the Milky Way (dShps) and the observed DM density today yield upper and lower bounds on the mass, primordial phase space densities and velocity dispersion of the DM candidates\\cite{coldmatter}. Combining these constraints with recent results from $N$-body simulations yield estimates on the mass of the DM particles   in the {\\bf range of a few keV} \\cite{coldmatter}.  \\medskip  The DM distribution function after freeze-out is the main ingredient to obtain the transfer function and the power spectrum in the linear theory. In principle the DM transfer function may be obtained by modifying the publicly available CMB anisotropy codes\\cite{cmbfast,lewis} that treat baryons, photons and dark matter \\cite{ma},  to account for the different distribution functions of several species and range of parameters, masses and couplings.  In practice, this is a computationally daunting problem: fairly complicated non-equilibrium distribution functions for a variety of possible DM species (axions, sterile neutrinos, etc) must be input in the codes (with substantial modifications of the standard codes). The exploration of the parameter space (masses, couplings and mixing angles) and their impact on the transfer function would require an enormous computational effort.  After matter-radiation equality, the contribution from the baryon-photon fluid to the DM transfer function can only be a few percent\\cite{dodelson}. DM only couples to the baryon-photon fluid via gravitational interaction and the most prominent feature of this coupling on the DM transfer function are baryon acoustic oscillations (BAO) on a scale corresponding to the sound horizon at recombination, corresponding today to about $150\\,\\mathrm{Mpc}$\\cite{EH,HS,eisenstein}. Our goal is to understand the (DM) transfer function at much \\emph{smaller} scales $\\lambda \\lesssim \\,\\mathrm{Mpc}$, at which the effect of the acoustic oscillations of the baryon-photon fluid (BAO) can be safely neglected.    \\medskip  {\\bf Objectives:} The main goals of this article are the following:  {\\bf (a):}    to provide an \\emph{analytic, accurate and simple} framework  to obtain the DM transfer function during matter domination for general initial conditions and \\emph{ arbitrary distribution functions} of relic DM particles that decoupled in or out of LTE and that are \\emph{non-relativistic} after matter-radiation equality. Neglecting the gravitational coupling to the baryon-photon fluid entails that the DM transfer function will eventually require corrections  at the  few percent level. However,  we seek to gain understanding of robust features of $T(k)$ at small scales.     {\\bf(b):} to understand the   impact of the statistics of the relic particle on the small scale properties of the transfer function. More precisely, we study  which features of the distribution function affect more prominently the power spectrum at small scales.     \\medskip      \\medskip  {\\bf Summary of Results:} We study the linearized collisionless Boltzmann-Vlasov equation in the non-relativistic limit for particles that decoupled in or out of LTE with arbitrary (but isotropic) distribution functions  and general  initial conditions.  It is transformed into an integro-differential equation for density perturbations which features non-local kernels that include \\emph{memory of gravitational clustering} and represent corrections to the fluid description.  The influence of this memory becomes more important at \\emph{small scales}. Fermi-Dirac (FD) and Bose-Einstein (BE) statistics result in \\emph{long range} memory, with the longer range for Bose-Einstein statistics. Distribution functions that favor small momenta lead to longer range memory and result in an \\emph{enhancement of the transfer function at small scales}.    \\medskip  We obtain an \\emph{exact} expression   for $T(k)$ for \\emph{arbitrary initial conditions and distribution functions} and provide a   Fredholm expansion   that systematically includes the effects of memory of gravitational clustering and higher moments of the distribution function, which are corrections to the fluid approximation.  An exhaustive numerical study of thermal relics that decoupled relativistically, (Fermions (FD) and Bosons, (BE)) or non-relativistically [WIMPs, (MB)] for different initial conditions reveal that the  first two terms of the Fredholm series, given by eqn. (\\ref{Tint2ndord}) provide an accurate and simple approximation to the exact transfer function. These include explicitly the corrections to the fluid description and the influence of statistics on small scales. This approximation furnishes a useful tool to extract the broad properties of the transfer functions for arbitrary initial conditions and distribution functions.    \\medskip  The approximate expression for the transfer function  given by eq.(\\ref{Tint2ndord}) features  three  main ingredients: {\\bf (i):} the solutions of Jeans' fluid equations, {\\bf(ii:)} the free streaming solution of the Boltzmann equation \\emph{in absence of gravitational perturbations}, {\\bf (iii:)} a non-local kernel that depends on the distribution function of the decoupled particles, it defines the second contribution in eq.(\\ref{Tint2ndord}), includes \\emph{memory} of gravitational clustering and represents a correction beyond the fluid approximation.  The scale of suppression of  $T(k)$ is the comoving free streaming wave-vector at matter-radiation equality, $k_{fs}(t_{eq})$, where \\be \\displaystyle k_{fs}(t) = \\sqrt{\\frac{4 \\, \\pi \\; \\rho_{0M} \\; a(t)}{\\left\\langle \\displaystyle \\frac{{\\vec{p}}^{\\, 2}}{m^2} \\right\\rangle}} \\;  \\label{kfs1} , \\ee and the angular brackets refer to the average with the distribution function of the decoupled particle, $ \\rho_{0M} $ is the DM matter density today and $ a(t) $ the scale factor.  We find, \\bea k_{fs}(t_{eq}) = \\left\\{ \\begin{array}{l} \\displaystyle{ \\frac{5.88}{ \\mathrm{pc}} \\; \\Big( \\frac{g_d}2 \\; \\Big)^\\frac13 \\; \\Big(\\frac{m}{100\\,\\mathrm{GeV}} \\Big)^\\frac12 \\; \\Big(\\frac{T_d}{10 \\,\\mathrm{MeV}}  \\; \\Big)^\\frac12} ~~\\mathrm{WIMPs} \\\\\\displaystyle{ 0.00284\\, \\Big( \\frac{g_d}2\\Big)^\\frac13 \\; \\frac{m}{\\mathrm{keV}} \\; [\\mathrm{kpc}]^{-1} ~~\\mathrm{FD~thermal~relics}} \\\\ \\displaystyle{   0.00317\\, \\Big( \\frac{g_d}2\\Big)^\\frac13 \\; \\frac{m}{\\mathrm{keV}} \\; [\\mathrm{kpc}]^{-1} ~~\\mathrm{BE~thermal~relics}} \\end{array} \\right. \\eea where $ g_d $ is the number of relativistic species at decoupling.  \\medskip  For large scales  $ k \\lesssim k_{fs}(t_{eq}) $, the contribution of the non-local memory kernel is subleading, whereas for scales $ k > k_{fs}(t_{eq}) $  the memory of gravitational clustering  and corrections beyond the fluid approximation    become important.  FD and BE distribution functions lead to long range memory kernels with  Bose-Einstein statistics leading to the longest range. Longer range memory leads to an enhancement of the transfer function (and power spectrum) at \\emph{small scales} as depicted in fig. \\ref{fig:gamagrande}. For $ k > k_{fs}(t_{eq}) , \\; (\\gamma > \\sqrt2) $ the difference in statistics of the decoupled particles becomes very important  \\medskip  At large scales $k \\ll k_{fs}(t_{eq})$  the transfer function   $ T(k) $ in terms of the variable $ \\gamma = \\sqrt2 \\; k/k_{fs}(t_{eq}) $ has the  following behaviour for all cases considered \\be \\label{taluniI} T(k) = 1 - \\left( \\frac{\\gamma}{\\gamma_0}\\right)^2 + {\\cal O}(\\gamma^4) \\ee where $ \\gamma_0 $ is    the same for Fermi-Dirac, Bose-Einstein and Maxwell-Boltzmann distribution functions,  and  is given by eq.(\\ref{gamacero}).  The transfer function for the three different cases (MB, FD and BE) can be compared   by parametrizing them in terms of $ \\gamma $. This comparison for different initial conditions is displayed in    fig. \\ref{fig:gamagrande}. Simple functional fits to $T(k)$ are provided on different regions of $k$ for various initial conditions for thermal relics. ",
        "conclusions": "In this article we studied the evolution of gravitational and DM density perturbations from the collisionless Boltzmann-Vlasov during matter domination, and obtained an \\emph{exact} expression for the transfer function $T(k)$ for \\emph{arbitrary distribution function of the decoupled DM particle and initial conditions}.  We have transformed the non-relativistic Boltzmann-Vlasov equation into an integro-differential equation which features a non-local kernel that describes the memory of gravitational clustering and yields corrections to the fluid description. This formulation lends itself to a systematic Fredholm expansion for the evolution of DM density and gravitational perturbations and  $T(k)$, and makes explicit the influence of the distribution function on $T(k)$.  Distribution functions that favor the small momentum region lead to \\emph{longer range} kernels, a persistence of the memory of the initial conditions and gravitational clustering resulting in an \\emph{enhancement} of $T(k)$ at \\emph{small scales}.   The natural scale of suppression of $T(k)$ is determined by  the free-streaming wave vector at matter-radiation equality,   $k_{fs}(t_{eq}) = \\sqrt{6}/l_{fs}(0)$ where the comoving free streaming wavevector is \\be k_{fs}(t) = \\left[\\frac{4\\pi \\rho_{0M} a(t)}{\\langle \\left(\\frac{\\vec{p}}{m} \\right)^2 \\rangle} \\right]^\\frac{1}{2}\\, ,\\label{kfs1} \\ee   the angular brackets refer to the average with the distribution function of the decoupled particle, and $l_{fs}(0)$ is the comoving free streaming distance traveled by the decoupled particle from the time of matter-radiation equality $t_{eq}$ until today.      We find \\be k_{fs}(t_{eq}) = \\frac{k_{fs}(0)}{\\sqrt{1+z_{eq}}} \\ee with \\bea k_{fs}(0) = \\left\\{ \\begin{array}{l} \\frac{0.325}{\\big[10^{-3}\\mathrm{pc}\\big]}\\, \\Big( \\frac{g_d}{2} \\Big)^\\frac{1}{3} \\Big(\\frac{m}{100\\,\\mathrm{GeV}} \\Big)^\\frac{1}{2}\\Big(\\frac{T_d}{10 \\,\\mathrm{MeV}} \\Big)^\\frac{1}{2} ~~\\mathrm{WIMPs} \\\\ 0.157\\, \\Big( \\frac{g_d}{2}\\Big)^\\frac{1}{3}\\Big(\\frac{m}{\\mathrm{keV}} \\Big)\\, [\\mathrm{kpc}]^{-1} ~~\\mathrm{FD~thermal~relic} \\\\ 0.175\\, \\Big( \\frac{g_d}{2}\\Big)^\\frac{1}{3}\\Big(\\frac{m}{\\mathrm{keV}} \\Big)\\, [\\mathrm{kpc}]^{-1} ~~\\mathrm{BE~thermal~relic} \\end{array} \\right.\\eea where $g_d$ is the number of relativistic species at decoupling.   We provided a detailed numerical study for thermal relics, WIMPs and fermionic and bosonic particles that decoupled while relativistically. The result of the numerical study of $T(k)$  as a function of $ \\gamma = \\sqrt{2}k/k_{fs}(t_{eq})$ for the three cases and different initial conditions is displayed in fig.  \\ref{fig:gamagrande}.  This study reveals that the first \\emph{two} terms in the Fredholm expansion yield a remarkable accurate approximation to $T(k)$ in the range of scales of cosmological relevance for structure formation. In all cases considered we find that the exact solution differs from the second order approximation in less than $5\\%$ in the range $0<\\gamma\\lesssim 6$ corresponding to scales down to  a fraction of the free-streaming length and much less than this on large scales for $0 < \\gamma \\lesssim 1$.  \\medskip  A simple and accurate approximation to $T(k)$ is given by eqn. (\\ref{Tint2ndord}). The second term in this expression includes corrections beyond the fluid approximation and higher order moments in the Boltzmann hierarchy and describes the memory of initial conditions and gravitational clustering. It explicitly depends on the distribution function of the decoupled particle. Non-local kernels with longer range lead to an \\emph{enhancement} of $T(k)$ at small scales. FD and BE thermal relics feature kernels with longer range than WIMPs, with BE statistics (for relics that decoupled relativistically) leading to longer range and more power at small scales.  This behavior is   clearly exhibited   in fig. \\ref{fig:gamagrande}.  For long scales $k \\ll k_{fs}(t_{eq})$ we find the behavior \\be T(k) = 1- \\Bigg( \\frac{\\gamma}{\\gamma_0}\\Bigg)^2+\\cdots \\ee where  $\\gamma_0 $ is given by (\\ref{gamacero}) for the initial conditions considered here independent of the statistics for thermal relics.  We provide   fits of $T(k)$, within a wide range of (small) scales for $k>k_{fs}(t_{eq})$ in the intervals where it exhibits a simple power-like or exponential behavior for MB, FD and BE statistics: see eqs.(\\ref{fitMB}), (\\ref{fitfg}), (\\ref{feradia}), (\\ref{fa2}) and (\\ref{fitbo}).   Although we do not attempt to provide a global fit with \\emph{one} function, it is clear that in the small scale region, the functional forms of $T(k)$ found in this article by exact numerical solution of the Boltzmann-Vlasov equation (in Gilbert's form) are very different from the  fits by Bardeen \\emph{et.al.}\\cite{BBKS} often quoted in the literature.   \\medskip  An important consequence of this study is a distinct imprint of the particle statistics and its distribution function on the transfer function at \\emph{small scales}, $ \\lambda \\ll l_{fs}(0) $: a distribution function that enhances the \\emph{small} momentum region yields a \\emph{longer range memory} of the initial conditions and gravitational clustering and an \\emph{enhancement} of the transfer function at small scales as depicted in fig. \\ref{fig:gamagrande}. This result may prove important in the elucidation of the small scale structure of DM halos and perhaps lead to an explanation of the filamentary structures found in numerical simulations in ref.\\cite{gao}.  The tools provided in this article to study the transfer function for arbitrary distribution functions and general initial conditions, in particular the simple and remarkably accurate approximation to $T(k)$ given by eqn. (\\ref{Tint2ndord}) allows a systematic and robust assessment of different DM candidates."
    },
    "0807/0807.3371_arXiv.txt": {
        "abstract": "{ To study the relevance of mergers for the fueling of QSOs, we are currently conducting an HST imaging campaign of a sample of QSO host galaxies classified as ellipticals in the literature. Here, we present results from a study of the first five QSO host galaxies imaged with HST/ACS. For the majority of objects, strong signs of interactions such as tidal tails, shells, and other fine structure are revealed. We estimate the nature and age of the merger by comparing the images with numerical simulations. The merger ages range between a few hundred Myr up to a Gyr. These timescales are comparable to starburst ages in the QSO hosts previously inferred from Keck spectroscopy, but longer than theoretical estimates of AGN duty cycles. A possible scenario emerging from our results is that most QSO host galaxies experienced mergers with accompanying starbursts but that the activity is triggered with a delay of several hundreds Myr after the merger. To probe whether there is indeed a causal connection between the merger and the QSO activity, we study a control sample of inactive ellipticals. Our preliminary results do not reveal comparable fine structure. ",
        "introduction": "The ubiquity of supermassive black holes (BHs) in the center of galaxies shows that more than the mere presence of a massive BH is needed to trigger the activity observed in Active Galactic Nuclei (AGNs). Moreover, the steep evolution of activity with redshift indicates that the accretion onto the BH must have been more common in the earlier universe and thus also the triggering mechanism. Mergers have been suggested to induce the sudden inflow of gas to the center needed to feed the AGN \\citep[e.g.][]{too72,sto82,san88}.  A close connection between mergers and AGN activity has been found for ultra-luminous infra-red galaxies (ULIRGs). Observations of these galaxies are consistent with an evolutionary scheme in which (at least some) ULIRGs form through mergers of gas-rich galaxies and represent the initial dust-enshrouded stage in the evolution of optically selected QSOs \\citep[e.g.][]{san88,can01,vei06}. Finally, the QSO hosts may end up as inactive elliptical galaxies.  However, a general connection between mergers and AGN activity is still being debated. While there is little doubt that mergers are helpful to provide the gas and remove angular momentum, they are certainly not sufficient, considering the numerous examples of inactive interacting galaxies. Also, mergers may be necessary for QSOs only. For their low-luminosity cousins, Seyfert galaxies, often residing in spiral host galaxies, there is little direct evidence for unusually high rates of interaction \\citep[e.g.][]{mal98}. Secular evolution through processes such as bar instabilities may be the dominant effect in the evolution of these galaxies \\citep[e.g.][]{com06}. But even for QSOs, the role of mergers for the activity remains unclear. Recent high-resolution imaging studies showed that many QSOs reside in elliptical hosts \\citep[e.g.][]{dis95,bah97,flo04}. From HST/WFPC2 images of 33 AGNs (radio-loud QSOs, radio-quiet QSOs, and radio galaxies) at a redshift of $z \\simeq 0.2$, \\citet{dun03} concluded that ``for nuclear luminosities $M_V$ $<$ -23.5, the hosts of both radio-loud {\\em and} radio-quiet AGN are virtually all massive elliptical galaxies with basic properties that are indistinguishable from those of quiescent, evolved, low-redshift ellipticals of comparable mass''.  To study the relevance of mergers for the fueling of classical QSOs, we are currently conducting an HST imaging campaign of a sample of QSO host galaxies classified as ellipticals by \\citet{dun03}. Here we present results from a pilot study of five QSO host galaxies ($z \\simeq 0.2$) using very deep  (5 orbits) HST/ACS images (F606W). ",
        "conclusions": "For four of the five QSOs, our images reveal dramatic signs of interactions such as tidal tails, shells, and other fine structure in the hosts (Figs.~\\ref{fourqsos},\\ref{mc}), suggesting that a large fraction of QSO host galaxies may have experienced a relatively recent merger event \\citep{ben08}.  One spectacular example of regular inner shell structure is MC2 1635+119 \\citep{can07}. In $N$-body simulations, the observed shells can be produced in a minor merger event. Assuming that the outermost shell is at 12.5 kpc (Fig.~\\ref{mc}), we estimate a merger age of $\\sim$30-400 Myr, depending on the type of profile of the giant elliptical, its effective radius, and the amount of dark matter. Taking into account a larger tidal feature at 65 kpc (Fig.~\\ref{mc}) that may be an older shell formed during the same encounter, the merger age can be up to 1.7 Gyr \\citep{can07}. However, while the inner shell structure can be produced by a radial minor merger, we cannot exclude other scenarios such as a major merger. Indeed, the total light contribution from the shells ($\\sim$6\\%) and extended structures are more indicative of a major merger. In this case, the inner shell structure might have been formed by material ``raining'' back into the central regions of the merger remnant \\citep{can07}.  For all objects, deep Keck spectroscopy revealed major starburst episodes ($\\sim$1-2 Gyr; \\citealt{can06,can07}). These timescales are comparable to the merger ages, but significantly longer than theoretical estimates for QSO duty cycles \\citep[e.g.][]{yu02}. Our results indicate that while most QSO host galaxies experienced mergers with accompanying starburst, there is a time delay of several hundred Myr between the tidal interaction and the actual fueling of the central BH. This is in agreement with recent hydrodynamic simulations \\citep[e.g.][]{spr05,hop07}.  \\begin{figure*}[t!] \\begin{center} {\\includegraphics[clip=true,scale=0.5]{bennert_1.eps}} \\caption{\\footnotesize Deep HST/ACS images of four of the five early-type QSO host galaxies \\citep{ben08}. North is up, east is to the left. {\\bf PHL\\,909} ({\\it top left}): A ring-like structure is seen around the QSO nucleus. Diffuse outer material form another ring and tidal tails to both sides of the host. {\\bf PKS\\,0736+01} ({\\it top right}): A large ($r$ $\\sim$ 50\\,kpc) but faint spiral-like structure surrounds the QSO. The irregular structure and changes in pitch angle may indicate spatial wrapping of material from a (minor) merger event rather than a spiral disk seen face on. {\\bf PG\\,0923+201} ({\\it bottom left}): No fine structure can be seen in the host galaxy, but it lies in an environment with several interacting companions. {\\bf OX\\,169} ({\\it bottom right}): The extended linear structure is likely a tidal tail seen nearly edge on. Note the extended shell-like features east of the nucleus. }\\label{fourqsos} \\end{center} \\end{figure*}  \\begin{figure*}[t!] \\begin{center} {\\includegraphics[clip=true,scale=0.5]{bennert_2.eps}} \\caption{\\footnotesize The same as in Fig.~\\ref{fourqsos} for MC2\\,1635+119 \\citep{can07}. Spectacular interleaved shells occur at $r$ $\\sim$ 5-12\\,kpc ({\\it top left}). The shell structure is seen more prominently when subtracting a PSF+host galaxy model as fitted by GALFIT (\\citealt{pen02}; {\\it top right}). An arc-like feature extends out to $\\sim$ 32\\,kpc ({\\it bottom left}). An even larger, faint and diffuse structure can be seen $\\sim$65 kpc west of the center ({\\it bottom right}). }\\label{mc} \\end{center} \\end{figure*}"
    },
    "0807/0807.3692_arXiv.txt": {
        "abstract": "{The cosmological constant problem is studied in a two component cosmological model. The universe contains a cosmological constant of an arbitrary size and sign and an additional component with an inhomogeneous equation of state. It is shown that, in a proper parameter regime, the expansion of the universe with a large absolute value of the cosmological constant may asymptotically tend to de Sitter space corresponding to a small effective positive cosmological constant. It is argued that such a behavior can be regarded as a solution of the cosmological constant problem in this model. The mechanism behind the relaxation of the cosmological constant is discussed. A connection with modified gravity theories is discussed and an example of a possible realization of the cosmological constant relaxation in $f(R)$ modified gravity is described.}  \\vspace{2cm} ",
        "introduction": "\\label{int}  The state of accelerated expansion of the present Universe seems to be better and better confirmed by the cosmological observations \\cite{SN,WMAP,LSS}. The question of the dynamical mechanism responsible for the accelerated expansion, however, still lacks its definite answer. The last decade has witnessed the arrival of numerous models of the accelerated expansion including dark energy, braneworld models, modified gravity and many others \\cite{Rev}. An interesting fact is that when confronted against the observational data, a simple $\\Lambda$CDM model, in which the cosmological constant (CC) is the cause of acceleration, fits the data very well. When the conceptual simplicity of the $\\Lambda$CDM model is taken into account, it is easy to understand why it is a benchmark model for the analysis of cosmological observations. While it is quite clear why the $\\Lambda$CDM model is so appealing from the observational side, its status from the fundamental theoretical perspective is much more problematic. It is by now a notorious fact that the CC value predicted in Quantum Field Theory (QFT) differs from the observed value by an embarrassingly large number of orders of magnitude \\cite{Wein}. The problem of explaining the observed value of the cosmological constant is therefore one of the largest challenges in theoretical physics \\cite{Wein,Stra,Nob}. This problem is sometimes refered to as {\\em the old cosmological constant problem}.  It is further exacerbated by the fact that in all other approaches to the problem of the accelerated expansion of the universe it is assumed that the CC problem is somehow solved.  Many attempts to solve the CC problem have been made during several last decades \\cite{Wein,Stra,Nob}. However, so far none of them has provided a fully satisfactory solution of the CC enigma. The most frequent problem that models of various sorts encounter is the necessity of {\\em fine-tuning}. The parameters of the model have to be chosen with extraordinary precision in order for the model to lead to the resolution of the CC problem. Even a very small deviation from these fine-tuned values disrupts the efficiency of the proposed mechanisms.  In this paper we propose a dynamical cosmological model with a specific regime in which it is possible to contemplate the resolution of the CC problem. The proposed model is simplified insofar that it does not contain all the (matter or radiation) components that naturally participate in the evolution of the universe. However, given the difficulty of the CC problem and its resilience to different attempts of solution, it seems preferable to first concentrate on the very mechanism which could produce the observed value of the effective cosmological constant for a universe with values of $\\Lambda$ comparable to those predicted in QFT.  There exists additional problem related to the size of the cosmological constant energy density (or more generally the present dark energy density). Namely, the observational data reveal that the energy density of matter (which at present epoch comes predominantly from nonrelativistic matter) is of the same order of magnitude as the CC (present DE) energy density. Dark energy and nonrelativistic matter scale differently with the expansion of the universe and it is quite remarkable that at present epoch these two energy densities are comparable. This problem is also called {\\em the cosmic coincidence problem}. In this paper we are primarily concerned with the problem of the size of the cosmological constant, whereas the cosmic coincidence problem is not addressed.  The principal aim of this paper is to study a two component cosmological model which possesses a cosmological constant of a large absolute value. We investigate conditions under which the asymptotic expansion in this model is of de Sitter type where the asymptotic value of the Hubble function corresponds to a small value of the effective cosmological constant ($H_{asym}^2=\\Lambda_{eff}/3$). In the studied models the solution of the cosmological constant is understood as a situation in which the universe ends up in an asymptotic de Sitter regime at large scale factor values characterized by a small cosmological constant (small in a sense that $\\Lambda_{eff} \\ll |\\Lambda|$). The absence of fine-tuning in the model studied in this paper is achieved if the parameters of the model do not have to cancel to many decimal places to lead to the solution of the CC problem. The dynamical process in which the universe in the studied model tends to de Sitter space with a small $\\Lambda_{eff}$ is also referred to as the relaxation of the cosmological constant. In the following sections we present a two component model in which the relaxation of the cosmological constant is realized. ",
        "conclusions": "\\label{conc}  The cosmological constant problem is a spot in theoretical physics landscape where the inadequacy of standard theoretical approaches is evident. An unconventional new ingredient is clearly called for. It is unclear, however, how big a deviation from the standard principles this new ingredient should represent. In this paper we have presented a simple approach based on a cosmological component with an inhomogeneous equation of state. The new ingredient is an inhomogeneous term in (\\ref{eq:p}) which can be interpreted as nonlinear viscosity or the effect of modified gravity. In a particular parameter regime, which by itself requires no specific fine-tuning, the universe with a cosmological constant of any sign and an arbitrarily large absolute value asymptotically ends up in de Sitter regime with a small value of the effective cosmological constant. This result provides a solution of the CC problem without the need of fine-tuning. A preliminary analysis of a $f(R)$ modified gravity theory lends support to the claim that the mechanism of the CC relaxation studied for a model with EOS (\\ref{eq:p}) also functions for modified gravity theories. The main results of this paper exemplify a scenario in which a large $\\Lambda$ coming naturally from QFT calculations coexists with a small asymptotic value of the effective cosmological constant. The said results further open a possibility that the measured value of the cosmological constant is not the value coming from QFT, but it is determined by the QFT value. Apart from the intrinsic value of these results, they also allow us a bit different perspective on cosmological parameter puzzles: maybe instead of devising complex nonstandard ways of understanding the cosmological parameter values we should try to understand how the values of these cosmological parameters influence simple nonstandard dynamics. The relaxation mechanism for the cosmological constant presented in this paper hopefully follows the latter route.   {\\bf Acknowledgements.} The author would like to thank N. Bili\\'{c}, B. Guberina, R. Horvat and H. Nikoli\\'c for useful comments on the manuscript. This work was supported by the Ministry of Education, Science and Sports of the Republic of Croatia under the contract No. 098-0982930-2864."
    },
    "0807/0807.1374_arXiv.txt": {
        "abstract": "The heating of the ion-neutral (or ambipolar) diffusion may affect the thermal phases of the molecular clouds. We present an investigation on the effect of this heating mechanism in the thermal instability of the molecular clouds. A weakly ionized one dimensional slab geometry, which is allowed for self-gravity and ambipolar diffusion, is chosen to study its thermal phases. We use the thermodynamic evolution of the slab to obtain the regions where slab cloud becomes thermally unstable. We investigate this evolution using the model of ambipolar diffusion with two-fluid smoothed particle hydrodynamics, as outlined by Hosking \\& Whitworth. Firstly, some parts of the technique are improved to test the pioneer works on behavior of the ambipolar diffusion in an isothermal self-gravitating slab. Afterwards, the improved two-fluid technique is used for thermal evolution of the slab. The results show that the thermal instability may persist inhomogeneities with a large density contrast at the intermediate parts of the cloud. We suggest that this feature may be responsible for the planet formation in the intermediate regions of a collapsing molecular cloud and/or may also be relevant to the formation of star forming dense cores in the clumps. ",
        "introduction": "Observations of molecular clouds show the existence of small density fluctuations threaded by magnetic fields. The recent progress of the molecular cloud observations has established that the small and tiny scale structures are very ubiquitous. For example, Langer et al. (1995) observed small condensations with size from $0.007$ to $0.021\\mathrm{pc}$ in Taurus Molecular Cloud~1. The mass of these fragments is estimated to be $< 0.01-0.15 M_\\odot$. Studies of the time variability of absorption lines indicates the presence of fluctuations on scales of $10^{-4} \\mathrm{pc}$ (5-50 AU) and masses of $10^{-9} \\mathrm{M_\\odot}$ (Boiss\\'{e} et al. 2005). At larger scales (about $10,000$~AU), Pan et al. (2001) find significant differences in $\\mathrm{CN}$, $\\mathrm{CH}$, and $\\mathrm{CH^+}$ absorption lines. The \\textsc{spitzer} experiment has begun producing higher spatial resolution mid-infrared maps (Churchwell et al. 2004) and revealed the fine structure of the star forming regions. The advent of major new facilities in the coming years should yield several breakthroughs in this field (Andr\\'{e} et al.~2008). A higher resolution observation in future is assuredly expected to reveal hidden, small structures, as well as sub-stellar objects with very small mass (Atacama Large Millimeter Array, \\textsc{alma}, becoming partly available in 2011, fully operational at 2013).  The classical problems of cloud fragmentation are relevant to star formation because they show that there is a preferred scale (Jeans length) of gravitational instability as soon as one considers an initial structured region with some length-scale (see, e.g., Larson~1985). By the way, it is not easy for the small and tiny scale structure to form as a result of gravitational instability, since their size are much smaller than the Jeans length. According to Langer et al. (1995) indeed, the small scale structures appear to be gravitationally unbound. This suggests that some fragmentation mechanisms other than pure Jeans gravitational instability may be important in the clouds. Molecular clouds are thought to be supersonically turbulent, since they exhibit supersonic linewidths (Zuckerman \\& Palmer~1974). Base on the idea of turbulence, some mechanisms to generate substructures inside molecular clouds have been proposed (Elmegreen 2007). As well, a further process based on magnetohydrodynamic waves have been suggested to produce inhomogeneities in the clouds (Folini et al.~2004). Another basic physical process that could trigger formation of density fluctuations is thermal instability, because it acts on time-scales that can be much shorter than the duration of the turbulent motions (Gilden 1984). The effect of thermal instability in fragmentation of the clouds is clearly a competitive process relative to other mechanisms, although all could operate.  In molecular clouds, a neutral molecular gas is intermixed with an ionized component that is tied directly to the magnetic field. Birk~(2000) has used the two-fluid technique to find the thermal condensation modes in weakly ionized hydrogen plasma. Nejad-Asghar \\& Ghanbari~(2003) studied the effect of linear thermal instability in a weakly ionized magnetic molecular cloud within the one-fluid description. Nejad-Asghar (2007) has recently made the assumption that the molecular cloud is initially an uniform ensemble which then fragments due to thermal instability. He find that ambipolar drift heating is inversely proportional to density and its value, in outer parts of the cloud, can be significantly larger than the average heating rates of cosmic rays and turbulent motions. His results show that the isobaric thermal instability can occur in intermediate regions of the cloud; therefore it may produce the cloud fragmentation and formation of the condensations.  The study under the one-fluid approach did not precisely consider the effect of the ion-neutral friction, since they did not treat the flow as two fluids which are composed of ions and neutrals. Then, Fukue \\& Kamaya~(2007) revisited the effect of the ion-neutral friction of the two-fluid on the growth of the linear thermal instability. Their results indicate that the friction with the magnetic field affects the morphology and evolution of the interstellar matter. Falle et al.~(2006) described numerical calculations using an \\textsc{amr} magnetohydrodynamic code that show that thermal instability may have an important role to play in the formation of the hierarchical structure of molecular clouds. Here, we use the two-fluid numerical simulation to study the evolution of weakly ionized molecular clouds in the nonlinear thermal instability regime.  Many authors have developed computer codes that attempt to model ion-neutral diffusion. Black \\& Scott~(1982) used a two-dimensional, deformable-grid algorithm to follow the collapse of isothermal, non-rotating magnetized cloud. The three-dimensional work of MacLow et al. (1995) treats the two-fluid model in a version of the \\textsc{zeus} magnetohydrodynamic code. An algorithm capable of using the smoothed particle hydrodynamics (SPH) to implement the ambipolar diffusion in a fully three-dimensional, self-gravitating system was developed by Hosking \\& Whitworth (2004, hereafter HW). They described the SPH implementation of two-fluid technique that was tested by modeling the evolution of a dense core, which is initially thermally supercritical but magnetically subcritical.  In this paper, we firstly re-formulate and improve the two-fluid SPH implementation of the ambipolar diffusion in an isothermal self-gravitating one dimensional slab. Afterwards, we use this two-fluid SPH technique to investigate the nonlinear thermal evolution of the slab. In this way, the continuum equations of the self-gravitating slab and its isothermal case are given is section~2. The two-fluid SPH technique is presented in section~3. Section~4 devoted to the chosen physical scales, initial conditions, and the computer experiments to explore the temperature profile and density fluctuations in the contracting one dimensional molecular cloud. Finally, section~5 is allocated to the summary and conclusions with some prospects. ",
        "conclusions": "Molecular clouds have a hierarchical structure that extends from the scale of the cloud down to much smaller masses for unbound structures. A weakly ionized self-gravitating one dimensional slab geometry is assumed in order to gain some insight into the fragmentation process. The two-fluid SPH technique was used to investigate the nonlinear thermal evolution of the slab. The initial conditions for this simulation are an inverse square cosine-hyperbolic profile for initial density and a parallel magnetic field directed perpendicular to the slab.  The initial values of the cooling and heating functions are shown in Figure~\\ref{coolheat0}. According to this figure, the isobaric thermal instability criterion is satisfied in the outer parts of the cloud. In course of time, the thermal instability is shifted to the intermediate regions of the slab so that the condensation mechanism is occurred there, thus, these regions are thermally unstable while the outer part is stable. The evolution were carried out to time $17.2[t]$ so that the fragmentation of the cloud via thermal instability is revealed.  The cooling and heating rates cause to increase the temperature of the clouds at outer regions while decrease the temperature at the inner part of it. It is obvious that the instability of the cloud at the intermediate parts, causes the formation of two relative cool regions in those areas, which is shown in Figure~\\ref{tempfig}. The rapid growth of thermal instability results a density imbalance between the cloud and the surroundings. The isothermal density profile and density contrast, relative to the isothermal case, are shown in Figure~\\ref{denfig}. In this run, the thermal instability can fully develop and makes density fluctuations in the intermediate parts of the cloud.  This feature may be responsible for the planet formation in the intermediate parts of a collapsing molecular cloud and/or may also be liable for the formation of star forming dense cores in the clumps. The macro-velocity fields in the molecular clouds are highly turbulent and supersonic. It is then of uppermost importance to consider both collision and merger of the formed condensations. Merging is possibly the main onset mechanism to form dense cores, which likely evolve to star formation."
    },
    "0807/0807.3147_arXiv.txt": {
        "abstract": "In the T2K near neutrino detectors, about 60 000 Hamamatsu Multi-Pixel Photon Counters (MPPCs) will be used. The mass production of MPPC has started in February 2008. In order to perform quality assurance and to characterize each device, we have developed an MPPC test system. For each MPPC, gain, breakdown voltage, noise rate, photo detection efficiency, and cross-talk and after-pulse rate are measured as functions of the bias voltage and temperature. The design of the test system and the measurement procedure are described. ",
        "introduction": "The T2K (Tokai-to-Kamioka)~\\cite{T2K} is a long baseline neutrino oscillation experiment in Japan. In the near detector system~\\cite{T2K-ND280}, we plan to use about 60~000 Hamamatsu Multi-Pixel Photon Counters (MPPCs)~\\cite{MPPC, MPPC2}.  \\begin{table}[htbp] \\begin{center} \\caption{Specifications of T2K-MPPC (S10362-13-050C).} \\begin{tabular}{ccc} \\hline \\hline \\multicolumn{2}{c}{Item} & Spec. \\\\ \\hline \\multicolumn{2}{c}{Active area} & 1.3$\\times$1.3~mm$^2$ \\\\ \\multicolumn{2}{c}{Pixel size} & 50$\\times$50~$\\mu$m$^2$ \\\\ \\multicolumn{2}{c}{Number of pixels} & 667 \\\\ \\multicolumn{2}{c}{Operation voltage} & 70~V (typ.) \\\\ \\multicolumn{2}{c}{PDE @ 550~nm} & $>$15\\% \\\\ Dark count &($>$0.5~pe) & $<$1.35~Mcps\\\\ $[$ @ 25$^\\circ$C $]$  & ($>$1.2~pe) & $<$0.135~Mcps \\\\ \\hline \\hline \\end{tabular} \\label{tab:spec} \\end{center} \\end{table}  \\begin{figure}[tbp] \\begin{center} \\includegraphics[width=0.4\\textwidth]{T2K-MPPC.eps}% \\caption{Picture of T2K-MPPC.} \\label{fig:MPPC} \\end{center} \\end{figure}%  We have worked with Hamamatsu to develop the MPPC for T2K in the past a few years. The major specifications of T2K-MPPC are summarized in Tab.~\\ref{tab:spec}. A picture of T2K-MPPC is shown in Fig.~\\ref{fig:MPPC}. The mass production started in February 2008. In order to perform quality assurance and to characterize each device, we have developed an MPPC test system. By August 2008, 15~000 MPPCs will be tested with our system at Kyoto University, while the remainder will be tested at other testing facilities abroad.  In this paper, the design of the test system, testing procedure and performance of T2K-MPPC are presented. ",
        "conclusions": "We have developed an MPPC test system that can simultaneously characterize 64 MPPCs. The method to evaluate the basic performance of MPPC such as gain, breakdown voltage, dark noise rate, photon detection efficiency, cross-talk and afterpulse rate, has been established. With our test system, these parameters as functions of temperature and bias voltage are automatically measured. We have measured the performance of T2K-MPPC and found it to satisfy our requirements."
    },
    "0807/0807.1973_arXiv.txt": {
        "abstract": "I show that the predicted densities of the inner dark matter halos in \\LCDM\\ models of structure formation appear to be higher than estimates from real galaxies and constraints from dynamical friction on bars.  This inconsistency would not be a problem for the \\LCDM\\ model if physical processes that are omitted in the collisionless collapse simulations were able to reduce the dark matter density in the inner halos.  I review the mechanisms proposed to achieve the needed density reduction. ",
        "introduction": "I was invited to review secular evolution in disk galaxies.  Rather than attempt a very superficial review of this vast topic, I here focus on dynamical friction.  Several other possible topics could be included in a review of secular evolution, such as: scattering of disk stars, which I reviewed only recently (Sellwood 2008a); mixing and spreading of disks (\\eg~Sellwood \\& Binney 2002; Ro\\u skar \\etal\\ 2008; Freeman, these proceedings); and the formation of pseudo-bulges (\\eg\\ Kormendy \\& Kennicutt 2004; Binney, these proceedings).  The current \\LCDM\\ paradigm for galaxy formation (\\eg~White, these proceedings) makes specific predictions for the dark matter (DM) densities in halos of galaxies.  I first argue that halos of some barred galaxies are inconsistent with this prediction, and then consider whether DM halo densities could be lowered by internal galaxy evolution. ",
        "conclusions": "The inner densities of DM halos of galaxies today continue to challenge the \\LCDM\\ model for galaxy formation.  Both direct estimates in a few galaxies, and dynamical friction constraints from bar pattern speeds require inner densities lower than predicted in DM-only simulations by at least a factor of a few.  Several processes that are omitted from DM only simulations can reduce the inner halo density.  Feedback has a very slight effect unless a large mass of gas is blasted out from the deepest point in the peotential well.  Bar friction requires an extreme bar and removes a large fraction of the angular momentum in the baryons.  Dynamical friction on massive gas clumps is too slow, unless moving gas clumps exceed $\\sim 1\\%$ of total baryonic mass.  Scattering by merging or recoiling BHs affects only very center.  Thus any of the four suggested mechanism needs to be stretched if it is to cause a significant reduction."
    },
    "0807/0807.3237_arXiv.txt": {
        "abstract": "It has been suggested recently that the change in cosmological redshift (the Sandage test of expansion) could be observed in the next generation of large telescopes and ultra-stable spectrographs. In a recent paper we estimated the change of peculiar velocity, i.e.\\ the peculiar acceleration, in nearby galaxies and clusters and shown it to be of the same order of magnitude as the typical cosmological signal. Mapping the acceleration field allows for a reconstruction of the galactic gravitational potential without assuming virialization. In this paper we focus on the peculiar acceleration in our own Galaxy, modeled as a Kuzmin disc and a dark matter spherical halo. We estimate the peculiar acceleration for all known Galactic globular clusters and find some cases with an expected velocity shift  in excess of 20 cm/sec for  observations fifteen years apart, well above the typical cosmological acceleration. We then compare the predicted signal for a MOND (modified Newtonian dynamics)  model in which the spherical dark matter halo is absent. We find that the signal pattern is qualitatively different, showing that the peculiar acceleration field could be employed to test competing theories of gravity. However the difference seems too small to be detectable in the near future. ",
        "introduction": "\\label{intro} Present cosmology has plenty of accurate data on the cosmic expansion, from standard candles to big bang thermal relics, but still lacks a direct observation of expansion which does not rely on interpreting the redshift as due the expansion itself. Such a direct test of expansion has been suggested in the 60's by Sandage \\citep{sandage62} and recently reproposed by various authors \\citep*{loeb98,pasquini06,2007PhRvD..76f3508L,2007MNRAS.382.1623B,2008PhRvL.100s1303U,2008PhRvD..77b1301U,2007astro.ph..1433C} in view of tremendous technological advances in spectroscopy and telescope collecting power. The expansion can be directly observed by measuring the redshift of sources at cosmological distances (e.g.\\ quasars) at two different epochs in time, obtaining for a homogeneous and isotropic Universe the relation \\begin{equation} \\dot{z}=(1+z)H_{0}-H(z)\\end{equation} The expected effect at $z\\approx1$ amounts to less than a centimeter per second per year. Today's spectroscopy has reached already a sensitivity of few meters per second: by using spectra populated by many sharp lines, some authors have shown that a statistical sensitivity of the required level can be reached with the next generation of optical telescopes \\citep{2008MNRAS.386.1192L}.  The total signal $\\dot{z}$ results clearly from the combination of the cosmological acceleration and of the line-of-sight component of the local (peculiar) acceleration field, just as the redshift itself combines global and local velocities. At large scales ($z\\ge0.5$) the peculiar acceleration signal acts as a noise over the cosmological signal and it might be hard to discriminate between the two contributions. At smaller depths however the peculiar field dominates and can be observed directly. In a recent paper \\citep*{2008PhLB..660...81A} we found that the peculiar acceleration in nearby clusters and galaxies is in fact of the same order of magnitude of the cosmological signal at larger distances and could be measured with the same instrumentation.  In this paper we focus instead on the peculiar acceleration of our own Galaxy and see how this can help reconstructing the Milky Way gravitational potential. We model our Galaxy as a Kuzmin disk plus a spherical dark halo and calculate the expected peculiar acceleration as a function of source distance and Galactic coordinates. Then we apply the general expression to all known Galactic globular clusters, finding signals up to 20 cm/sec in 15 years, quite larger than the typical expected cosmological signal.  The peculiar acceleration field maps directly the gravitational potential without the need to rely on virialization assumptions. One is therefore led to ask whether the peculiar acceleration field could help in discriminating among competing gravitational theories. A well-known model of non-Newtonian gravity is the modified Newtonian dynamics (MOND) theory \\citep{Brada:1995}: in this theory, the galaxy can be modeled as a baryonic disk, without any spherical halo, with parameters chosen to obtain the same asymptotic disk rotational velocities. We estimated the Milky Way peculiar acceleration field in MOND and compared it to the Newtonian field for sources outside the disk, where the geometry of the two models should maximize the difference. We find qualitatively distinct patterns in the two cases but we find rather small quantitative differences, not larger than 1 cm/sec, probably too low to be within the reach of foreseeable observations. It is of course possible that peculiar acceleration can be used to discriminate Newtonian dynamics from other models of gravity in our own Galaxy, thereby extending the test of gravity far beyond the present solar system limits. Furthermore, measuring peculiar accelerations could be a way to determine the gravitational potential parameters.   \\section[]{Peculiar acceleration} \\label{pec}  We start by investigating the peculiar acceleration signal expected in the case of a generic spiral galaxy. Later, in Section \\ref{milkyway}, we will address the specific case of our own Galaxy.  The signal we are after is the velocity shift $\\Delta v$ given by \\begin{equation} \\Delta v=a_{s} \\Delta T, \\end{equation} where $a_s$ is the line of sight component of the acceleration and $\\Delta T$ is the time interval between two observations of the velocity of a test particle.  Let us first consider the disc component of the galaxy. If a test particle (e.g.\\ stars, gas) orbits on the disc, a mathematical theorem states that, due to the symmetry of the distribution, its peculiar acceleration is affected only by the mass embedded within its distance from the centre, and it is the same as if this mass was totally concentrated in the centre.  In order to model the disc component, we will consider the so-called Kuzmin disc, namely a disc with superficial density \\begin{equation} \\label{densi} \\Sigma(r)=\\frac{hM}{2\\pi(R^{2}+h^{2})^{3/2}}, \\end{equation} where M is the total disc mass, $R$ and $z$ are cylindrical coordinates and $h$ is the scale length of the disc. The two-parameter Newtonian gravitational potential outside the disc is \\begin{equation} \\label{kuzpot} \\phi_{K}=-\\frac{MG}{[R^{2}+(|z|+h)^{2}]^{1/2}}, \\end{equation} and the equipotential surfaces are concentric spheres centered at $\\pm h$ (see Fig.~\\ref{fig:ill1}). The test particle acceleration then reads \\begin{equation} \\label{acc2} a_{K}=-\\frac{MG}{[R^{2}+(|z|+h)^{2}]}, \\end{equation} with $|z|$ accounting for the acceleration field above and below the disc.  The force field of a Kuzmin disc no longer converges towards the origin of axis (see Fig.~\\ref{fig:ill1}). Consequently, the projection angle $\\theta'$ must be corrected by $\\gamma$, namely the angle between the acceleration and the radial direction $r$.  On the other hand the peculiar acceleration of a test particle outside the disc, e.g.\\ a globular cluster, is not only affected by the  gravitational potential generated by the disc, but also by the possible presence of a dark matter halo needed to explain the observed rotation curves. An alternative description of the gravitational potential can be given within the MOND framework: in this case, only baryonic matter in the disc is present, and the gravitational law is modified. The peculiar acceleration of a test particle outside the disc will then be affected by the axisymmetrical gravitational potential of the disc.  We will now calculate the expected signal in the two scenarios: first, the standard configuration of a disc plus a cold dark matter halo; second, the alternative description in terms of a disc plus MOND. The bulge component will be neglected in both scenarios, since its spherical symmetry allows one to treat it simply as an additional contribution to the total mass for scales outside the bulge.  \\subsection{Disc + CDM halo} \\label{CDM} In standard Newtonian mechanics we model the galaxy as a Kuzmin disc embedded in a CDM halo, described by a spherical logarithmic potential, following  \\cite{Read:2005}. Thus we adopt the Kuzmin potential for the disc (as in Eq.~\\ref{kuzpot}) and we model the halo potential as logarithmic, \\begin{equation} \\label{halo} \\phi_{L}=\\frac{1}{2}v_{0}^{2}\\log{\\Big(R_{c}^{2}+R^{2}+\\frac{z^{2}}{q}\\Big)}, \\end{equation} where $R_{c}$ is the scalelength, $q$ is the halo  flattening ($q=1$ recovers  spherical symmetry) and $v_{0}$ is the asymptotic value of the velocity at large radii.  The Kuzmin and the halo accelerations must be projected along the line of sight and then added together. While for spherical symmetric logarithmic potential the acceleration is radial and the angle between the line of sight and $r$ is  simply $\\theta'$ (see Fig.~\\ref{fig:ill1}), as mentioned before, the projection angle for the  Kuzmin acceleration does not point towards the origin. One clearly has $\\cos\\theta'\\approx \\sin\\beta$. Using $r\\cos\\beta\\simeq R_g \\theta$, $z\\simeq\\pm\\sqrt{r^2-R^2}$ and $R\\simeq R_g\\theta$, the two line of sight accelerations read \\begin{eqnarray} a_{s,K}&=&\\frac{MG}{R_{g}^{2}\\theta^{2}\\Big[1+\\Big(|\\tan{\\beta}|+\\frac{h}{R_{g}\\theta}\\Big)^{2}\\Big]}\\sin{(\\beta\\mp\\gamma)}\\\\ a_{s,L}&=&\\frac{v_{0}^{2}R_{g}\\theta\\sqrt{1+\\frac{\\tan^{2}{\\beta}}{q^{2}}}}{R_{c}^{2}+R_{g}^{2}\\theta^{2}(1+\\frac{\\tan^{2}{\\beta}}{q^{2}})}\\sin{\\beta}. \\end{eqnarray} All the parameters of the gravitational potential will be fixed to the values of \\cite{Read:2005}: $M=1.2\\cdot 10^{-11}M_\\odot$, $h=4.5$ kpc, $v_0=175$ km/s, $R_c=13$ kpc and $q=1$. In Fig.~\\ref{fig:comb} the velocity shift caused by the Kuzmin disc and the CDM halo are plotted, as well as their sum, assuming a time span $\\Delta T=15$~yr. Due to its non-radial direction, the main effect of the Kuzmin potential is to shift the maximum and the zero-crossing of the function with respect to the spherically symmetric case. The signal itself is in principle completely independent on the distance of the galaxy; however we have expressed it as a function of observable quantities, like the viewing angle. In Fig.~\\ref{fig:comb} we assume a distance of the galaxy centre of $800$~kpc, namely the distance of Andromeda galaxy.  \\begin{figure} \\centering \\includegraphics[width=8.truecm]{fig1.eps} \\caption{ \\small \\small Kuzmin disc potential in cylindrical coordinates. The disc lies edge-on on the axis $R$. The contour shows an equipotential curve and $\\vec{a}_{K}$ represents the  force  field. Notice that $\\vec{a}_{K}$ does not point towards the origin. The viewing angle is  $\\hat{\\theta}$ (in this illustration we label angles with the $\\hat{}$ ). } \\label{fig:ill1} \\end{figure}  \\begin{figure} \\centering \\includegraphics[width=7.truecm]{fig2.eps} \\caption{ \\small \\small Velocity shift curves as a function of $\\beta$ for the Kuzmin disc in Newtonian configuration (dashed line), CDM halo (dotted line) and their combination (solid line) for $\\theta=0.03$. } \\label{fig:comb} \\end{figure}  \\subsection{Disc + MOND} \\label{suboutside}  The Modified Newtonian Dynamics (MOND) paradigm was  first proposed by Milgrom to reconcile discrepancies between general relativity and galaxy scale dynamics \\citep{1983ApJ...270..371M}. As presented in its  first version, MOND by itself violates conservation of momentum and energy. In this paper we concentrate on the Bekenstein-Milgrom formulation of MOND \\citep{1984ApJ...286....7B}, which leaves the Newtonian law of motion intact and modifies the standard Poisson equation for the Newtonian gravitational potential as follows: \\begin{equation} \\label{mond_pois} \\nabla \\Big[ \\mu \\Big( \\frac{|\\nabla \\psi|}{a_{0}}\\Big)\\nabla \\psi\\Big]=4\\pi G \\rho, \\end{equation} where $\\psi$  is the MOND gravitational potential, $\\nabla^{2}\\phi_{N}=4\\pi G \\rho$, and $a_{0}$ is a scale that was estimated by Milgrom to be $a_{0} = 1.2 \\cdot10^{-10}$m/s$^{2}$. The original shape of $\\mu$ that helps rendering the right profile of rotational velocities is $\\mu(|a|/a_{0}) = 1$ for $|a|\\gg |a_0|$ and $\\mu(|a|/a_{0}) = |a|/a_{0}$ for $|a| \\ll  |a_0|$.  Eq.~\\ref{mond_pois} is a non-linear equation which might be hard to solve analytically, except for a class of symmetric configurations. In \\cite{Brada:1995}, a class of disc-galaxy models has been introduced for which exact solutions of the MOND field equation exist. Subtracting the usual Poisson equation from the MOND equation (\\ref{mond_pois}) one has \\begin{equation} \\label{mond_pois2} \\nabla \\Big[ \\mu \\Big( \\frac{|\\nabla \\psi|}{a_{0}}\\Big)\\nabla \\psi - \\nabla \\phi_{N}\\Big]=0. \\end{equation} For configurations with spherical, cylindrical or plane symmetry the relation between the MOND  field and the Newtonian  field becomes \\begin{equation} \\label{mond_pois3} \\mu \\Big( \\frac{|\\nabla \\psi|}{a_{0}}\\Big)\\nabla \\psi = \\nabla \\phi_{N}. \\end{equation} This equation permits a straightforward relation between the two potentials, such that by assuming a matter density distribution, one can solve the Poisson equation for the Newtonian potential and then invert it to get the MOND one.  The function $\\mu$ defined in MOND has been reformulated during the years, moving from a step function as the aforementioned type to, e.g., $\\mu(x) = x/\\sqrt{1+x^{2}}$. As pointed out in \\cite{Brada:1995}, the function $I(x) = x\\mu(x)$ is therefore monotonic and invertible, and the inverse is related to the function  $\\nu(y) = I^{-1}(y)/y$. Eq.~\\ref{mond_pois3} now reads as \\begin{equation} \\label{out1} \\nabla \\psi=\\nu\\Big( \\frac{|\\nabla \\phi_{N}|}{a_{0}}\\Big)\\nabla \\phi_{N}. \\end{equation}   The exact MOND solution is then given by Eq.~(\\ref{out1}), with \\begin{equation} \\label{acc} \\vec{a}_M=-\\nabla\\psi=a_{0}I^{-1}\\Big(\\frac{a_{N}}{a_{0}}\\Big)\\frac{\\vec{a}_{N}}{a_{N}} \\end{equation} where $a_M$ and $a_N$ are the MOND and Newtonian acceleration, respectively.   For $\\mu(x) = x/\\sqrt{1+x^{2}}$, then  $\\nu(y) =[1/2 +\\sqrt{y^{-2} + 1/4}]$ and consequently \\begin{equation} \\label{accmond} a_{M}=a_{K}\\frac{\\Big(1+\\sqrt{1+\\frac{4a_{0}^{2}}{|a_{K}|^{2}}}\\Big)^{1/2}}{\\sqrt{2}}. \\end{equation}  The expression for the peculiar acceleration in MOND, outside the disc, is as follows: \\begin{eqnarray} \\label{accmond2} a_{s,M}= a_{M}\\cdot \\sin{(\\beta \\mp \\gamma)}, \\end{eqnarray} with \\begin{eqnarray} a_{M}= \\frac{MG\\Big(1+\\sqrt{1+\\frac{4a_{0}^{2}}{M^{2}G^{2}}R_{g}^{4}\\theta^{4}[1+(|\\tan{\\beta}|+\\frac{h}{R_{g}\\theta})^{2}]^{2}}\\Big)^{1/2}}{\\sqrt{2}R_{g}^{2}\\theta^{2}\\Big[1+\\Big(|\\tan{\\beta}|+\\frac{h}{R_{g}\\theta}\\Big)\\Big]}, \\end{eqnarray} and \\begin{equation} \\label{gam} \\gamma=\\arccos{\\frac{r^{2}+x^{2}-h^{2}}{2 r x}};\\qquad x^{2}=r^{2}+h^{2}-2rh\\sin{\\beta}. \\end{equation}  \\subsection{Comparison}   In order to compare peculiar acceleration predictions between the CDM and the MOND scenarios, it is necessary to match the MOND configuration to a Newtonian one leading to the same rotation velocity curve. Given Eq.~\\ref{accmond} for the acceleration, the related rotation velocity can be written as \\citep{Brada:1995} \\begin{equation} \\label{vel} v^{2}_{M}(R)=v_{\\infty}^{2}\\frac{u^{2}}{1+u^{2}}\\Big\\{ \\Big[1+\\frac{\\zeta^{2}}{4(1+u^{2})^{2}}\\Big]^{1/2}+\\frac{\\zeta}{2(1+u^{2})}\\Big\\}^{1/2} \\end{equation} with $u=R/h$ and  $\\zeta=MG/h^{2}a_{0}$ is a measure of how deep in the MOND regime we are. The asymptotic velocity  $v_{\\infty}=(MGa_{0})^{1/4}$ is set by the mass of the disc $M=1.2\\cdot10^{11} M_{\\odot}$  \\citep{Read:2005} and has to match the sum of the square velocities of the disc and CDM halo, which are \\begin{equation} \\label{velcdm} v^{2}_{K}(R)=\\frac{GMR^{2}}{(R^{2}+h^{2})^{3/2}},\\qquad v^{2}_{L}(R)=\\frac{v_{0}^{2}R^{2}}{R_{c}^{2}+R^{2}}, \\end{equation} respectively.  In Fig.~\\ref{fig:vel} the rotation curve of the galaxy is drawn both for the Newtonian and MOND configuration. Being the two curves almost undistinguishable, a further interesting point we want to address in this paper is the possibility of discerning a difference between the two gravity models by observing the velocity shift of a test particle outside the disc.  The different shape of the peculiar acceleration between Newtonian and MOND configurations are shown in Fig.~\\ref{fig:pec1} and \\ref{fig:pec2}. The predicted velocity shift reaches a maximum value along the line of sight,  $\\Delta v_{max}$, that depends on the gravitational configuration. The maximum also depends on the distance from the galactic centre, increasing with decreasing distance.  By estimating  $\\Delta v_{max}$ from a sample of test particles (e.g.\\ globular clusters) outside the disc of the galaxy along each line of sight, one might therefore hope to distinguish the Newtonian and the MOND gravitational potentials.  In our Galaxy, knowing precisely the coordinates of test particles one might trace out the distribution of the signal; in the following section we will focus on this analysis.  \\begin{figure} \\centering \\includegraphics[width=7.truecm]{fig3.eps} \\caption{ \\small \\small Rotation velocity curves for Newtonian configuration (long-short-dashed line) and MOND (solid line). The long-dashed line represents the CDM halo, while the dotted line represents the Kuzmin disc.} \\label{fig:vel} \\end{figure}  \\begin{figure} \\centering \\includegraphics[width=7.truecm]{fig4.eps} \\caption{ \\small \\small  Velocity shift curves caused by peculiar accelerations as a function of $\\beta$ for Newtonian configuration (thin lines, CDM halo+Kuzmin disc) and MOND (thick lines). For both configurations curves are drawn for three different values of $\\theta$, from top to bottom: $\\{0.05,0.06,0.07\\}$, corresponding to distances from the centre $\\{40,48,56\\}$ kpc. } \\label{fig:pec1} \\end{figure}  \\begin{figure} \\centering \\includegraphics[width=7.truecm]{fig5.eps} \\caption{ \\small \\small Velocity shift curves caused by peculiar accelerations as a function of $\\beta$ for Newtonian configuration (thin lines, CDM halo+Kuzmin disc) and MOND (thick lines). For both configurations curves are drawn for three different values of $\\theta$, from top to bottom: $\\{0.01,0.03,0.05\\}$, corresponding to distances from the centre $\\{8,24,40\\}$ kpc. } \\label{fig:pec2} \\end{figure}  \\section[]{Milky way} \\label{milkyway} Now that we have worked out the calculations for the generic spiral galaxy, it is interesting to focus our predictions to the specific case of test particles in the Milky Way. Our Galaxy is known to host almost 150 globular clusters, spread in the stellar halo, with distances up to $100$~kpc, which would constitute an ideal observational target. More importantly, knowing exactly the Galactic coordinates of the objects helps reconstructing and comparing the predicted and the observed peculiar accelerations.  \\begin{figure} \\centering \\includegraphics[width=9.truecm]{fig6.pdf} \\caption{ \\small \\small Kuzmin disc potential in cartesian coordinates in the Milky Way. The disc lies on plane $(x,y)$. Contours show equal steps of  $\\phi_{K}$ and $\\vec{a}_K$ represents the radial force  field. The grey star represents the sun and angles are labeled with the $\\hat{}$ . } \\label{fig:ill} \\end{figure}  Let us then write the contribution from peculiar acceleration to velocity shift in terms of Galactic coordinates, in particular the distance of the object from the sun $r_{s}$, the Galactic latitude $b$ and longitude $l$ (see Fig.~\\ref{fig:ill}). Firstly, we have to move from cylindrical coordinates $(z, R,\\varphi)$ to Galactic coordinates $(r_{s}, b, l)$, and neglecting $\\varphi$ due to the axisymmetry of the potential, the transformation reads \\begin{eqnarray} \\label{transf} z=r_{s}\\sin{b};\\qquad R=\\sqrt{s^{2}+r_{s}^{2}\\cos^{2}{b}-2s r_{s}\\cos{b}\\cos{l}}, \\end{eqnarray} where $s$ is the distance of the Sun from the Galactic centre.  The line of sight now corresponds to the direction of the distance from the object to us, i.e. $r_{s}$. The projection angle $\\theta'$ is related to the new coordinate system as follows: \\begin{equation} \\label{transf} \\cos{\\theta'}=\\frac{r_{s}^{2}+r^{2}-s^{2}}{2 r_{s}r}, \\end{equation} where $r^{2}=z^{2}+R^{2}$ (see Fig. \\ref{fig:ill}). In addition, as already mentioned in Sec.~\\ref{CDM}, for the Kuzmin potential the correction to the non-radial direction of the acceleration must again be taken into account, so that the projection angle changes to $\\cos\\left({\\theta'\\pm\\gamma}\\right)$, and $\\gamma$ is as in Eq.~(\\ref{gam}), where now \\begin{equation} \\label{beta} \\beta=\\arccos{\\frac{r^{2}+s^{2}-r_{s}^{2}}{2rs}}. \\end{equation}  We have calculated the expected signal for the 150 globular clusters of the \\cite{1996AJ....112.1487H} catalogue, excluding only three clusters whose position lies exactly on the disc (where the Kuzmin model breaks down). To be consistent with this catalogue assumption, in our calculations we fix the Sun distance from the Galactic centre to the value $s=8$ kpc. We map in Fig.~\\ref{fig:mapmond} and \\ref{fig:mapnewt} the 3D distribution of the velocity shift value for each cluster position, in cartesian coordinates. We find that, over  a time span of 15 years, signals as high as 22 cm/s in MOND and 21 cm/s in the Netwonian case can be expected. This is a rather large value when compared to an average cosmological signal of order $5$ cm/s found either for distant galaxies at $z\\sim$ 3--4 or for galaxies in clusters (see \\cite{2008PhLB..660...81A,2007MNRAS.382.1623B}). We will come back to this result later in Section \\ref{conclusions}.  In Fig.~\\ref{fig:mapdiff} we map the absolute value of the difference between the signals estimated in the two scenarios. Unfortunately, the maximum differences is only of order 1 cm/s, making it quite problematic to observe. However, supposing that such an accuracy will be achieved in the future, it is interesting to note that the two different scenarios result in a distinct morphology of the signal distribution, that could be used as a signature to identify the actual potential. This is shown in Fig.~\\ref{fig:array}, for three different distances from the Sun.  \\begin{figure} \\centering \\includegraphics[width=9.truecm]{fig7.eps} \\caption{ \\small \\small The velocity shift signal (for T=15 years) expected in the MOND configuration for the Milky Way globular clusters taken from Harris (1996), plotted in cartesian coordinates (kpc units).} \\label{fig:mapmond} \\end{figure} \\begin{figure} \\centering \\includegraphics[width=9.truecm]{fig8.eps} \\caption{ \\small \\small The velocity shift signal (for T=15 years)  expected in the CDM halo configuration for the Milky Way globular clusters taken from Harris (1996), plotted in cartesian coordinates (kpc units).} \\label{fig:mapnewt} \\end{figure} \\begin{figure} \\centering \\includegraphics[width=9.truecm]{fig9.eps} \\caption{ \\small \\small The difference in the velocity shift signal (for T=15 years) expected in the CDM halo and MOND configuration for the Milky Way globular clusters taken from Harris (1996), plotted in cartesian coordinates.} \\label{fig:mapdiff} \\end{figure} \\begin{figure*} \\centering \\includegraphics[width=14.truecm]{fig10.eps} \\caption{ \\small \\small Comparison in the velocity shift signal (for T=15 years) as seen in the sky from our position in the Milky Way, for the MOND, CDM halo configuration, and their difference (left to right) and for three distances from the Sun (3, 8 and 15 kpc, top to bottom). The sky signal is contour-plotted  as an equal-area Mollweide projection of the sky around us, where ideal meridians and parallels mark the Galactic longitude and latitude respectively. Regions with higher signal are lighter and the centre of the ellipse corresponds to $\\{b,l\\}=\\{0,0\\}$, i.e. to the direction of the Galactic centre.} \\label{fig:array} \\end{figure*} ",
        "conclusions": "\\label{conclusions}  The main purpose of this paper has been to investigate the possibility of reconstructing the gravitational potential of a galaxy by means of the so-called velocity shift signal. Assuming that either the  cosmological redshift of the galaxy is almost vanishing or it can be averaged out, the velocity shift produced by peculiar acceleration of test particles orbiting outside the galactic disc has been the key point of our analysis. In particular we focused on a generic spiral galaxy modelled as a baryonic disc component, e.g.\\ the Kuzmin disc, plus a CDM spherical halo. We also extended our calculation to a modified gravity configuration, in which the CDM halo is absent and the Poisson equation is modified following MOND \\citep{Brada:1995}: although the two scenarios are undistinguishable using rotation velocity curve, we addressed the question of whether peculiar acceleration might be employed to discriminate between them.  In Sec.~\\ref{pec} we have worked out the general treatment for peculiar acceleration in a non-specific galaxy seen by an external observer. The velocity shift signal  has been derived as a function of  the viewing angle $\\theta$ and the angle $\\beta$ between the radial position and the plane of the disc. The signal reaches a maximum value that depends on the gravitational potential and could in principle be observed along   each line of sight (Fig.~\\ref{fig:pec1} and \\ref{fig:pec2}, assuming a time span of 15 yr). As expected, the closer is the test particle to the galactic center the higher is the signal, reaching values comparable to the cosmological ones predicted by simulations and possibly detectable from high-resolution and ultra-stable spectrograph coupled to new generation telescope \\citep{2008MNRAS.386.1192L}.  In Sec.~\\ref{milkyway} we have focused our predictions to our Galaxy and used as test particles the 150  known globular clusters orbiting outside the disc of the Milky Way taken from \\cite{1996AJ....112.1487H}. As in this case we clearly are internal observers, the peculiar acceleration has been re-expressed in terms of Galactic coordinates, i.e. the distance from the Sun, the longitude and the latitude. The projected acceleration of the Sun on the disc, which has its maximun value on the order of few cm/s, is assumed to be substracted, and the cluster-centred acceleration of single stars in globular clusters is assumed to be averaged out.   We have mapped out the contours of the predicted signal of both configurations in a Mollweide projection of the sky at three different distances from us (Fig.~\\ref{fig:array}). The effect of being off-centred observers  combined with the non-radial direction of the acceleration in Kuzmin disc makes the pattern of the contours non trivial. However, while again the signal reaches the maximum value close to the Galactic centre for both scenarios, the difference signal is stronger at high Galactic longitude (close to $\\pi$) where the spherical CDM halo is more influent. This can be better seen by eye in Fig.~\\ref{fig:mapmond},\\ref{fig:mapnewt} and \\ref{fig:mapdiff} where we have mapped the distribution of the peculiar acceleration in Sun-centred cartesian coordinates.  In Table \\ref{topten_mond} and  \\ref{topten_newt} the first 10 globular clusters exhibiting the highest velocity shift signal are listed. They are the same for CDM halo configuration as for MOND, except the last one.   Their distance is very similar to the distance of the Sun from the Galactic centre, e.g.\\ $\\sim$8kpc,  and the absolute value of their latitude never exceeds 12$^{o}$: considering also their longitude, it turns out that they are all fairly close to the centre of the Milky Way. As already emphasised, this is not the case for the difference of the signal between the two scenarios. The first 10 globular clusters with the highest difference signal between the CDM halo configuration and MOND are shown in Table \\ref{topten_diff}. In view of possible future detections due to the tremendously improved technology \\citep{pasquini06,2008MNRAS.386.1192L} this list might indeed give indications on which particular clusters to look at. Although the feasibility of such observations is still under debate, this new technique would open a new window to the ability of testing the gravitational potential and of constraining the parameters involved in its expression, such as the mass of the galaxy and the scale lengths $R_{c}$ and $h$.  Finally, we point out that the advantage of adopting globular cluster as test particles is related  to the fact that they not only are bright, but more importantly composed of thousands of stars and their position is fairly well known; neverthless, in principle one might consider  using other probes, like high velocity clouds or HI regions. Moreover, one could also use the acceleration in the pulsar timing as another probe of the acceleration field in the Galaxy. However this requires further investigations that are beyond the purpose of this paper.  \\begin{table*} \\centering \\begin{minipage}{140mm} \\caption{The ten globular clusters having the highest velocity shift from peculiar acceleration in the MOND configuration. Astronomical data are from Harris (1996)  and $R_{gc}$ is the distance from the Galactic centre.} \\label{topten_mond} \\begin{tabular}{@{}llllllll@{}}  \\hline Name    &  $v$ (cm/s)\\footnote{Assuming $T=$15 years} &     RA  (hours)\\footnote{J2000 epoch}  & dec (degrees)   &   $l $ (degrees)  &   $b$ (degrees)   &  $r_s$ (kpc) & $R_{gc}$ (kpc)\\footnote{Assumes $R_0=8$ kpc} \\\\  \\hline   NGC 6712      &  21.5 &      18 53 04.3  & -08 42 22 &  25.35  & -4.32 &   6.9  & 3.5 \\\\ Lynga 7       &20.5      &  16 11 03.0  &-55 18 52  &328.77  & -2.79 &   7.2  & 4.2 \\\\ NGC 6535    &18.0       &   18 03 50.7 & -00 17 49 &  27.18 &  10.44  &  6.8  & 3.9 \\\\ NGC 5927    &17.9          &15 28 00.5 & -50 40 22 & 326.60 &   4.86  &  7.6 & 4.5 \\\\ NGC 6749     &17.3        & 19 05 15.3  & +01 54 03  & 36.20  & -2.20 &   7.9 &  5.0 \\\\ NGC 6760    &17.0         & 19 11 12.1  &+01 01 50  & 36.11 &  -3.92   & 7.4 &  4.8\\\\ NGC 6352    &16.3        &  17 25 29.2 & -48 25 22&  341.42 &   -7.17 &   5.7 &  3.3 \\\\ NGC 6717 (Pal 9)  & 14.9&     18 55 06.2 & -22 42 03  & 12.88 & -10.90   & 7.1  & 2.4 \\\\ NGC 6171 (M 107)   &  14.8 &   16 32 31.9  &-13 03 13  &  3.37 &   23.01 &   6.4 &  3.3 \\\\ NGC 6541     &     14.0 &    18 08 02.2  &-43 30 00 & 349.48 & -11.09 &   7.0 &  2.2  \\\\  \\hline  \\end{tabular} \\end{minipage} \\end{table*}%  \\begin{table*} \\centering \\begin{minipage}{140mm} \\caption{The ten globular clusters having the highest velocity shift from peculiar acceleration in the CDM halo configuration. Astronomical data are from Harris (1996) and $R_{gc}$ is the distance from the Galactic centre.}\\label{topten_newt} \\begin{tabular}{@{}llllllll@{}}  \\hline Name    &  $v$ (cm/s)\\footnote{Assuming $T=$15 years} &    RA  (hours)\\footnote{J2000 epoch}  & dec (degrees)   &   $l $ (degrees)  &   $b$ (degrees)   &  $r_s$ (kpc) & $R_{gc}$ (kpc)\\footnote{Assumes $R_0=8$ kpc}  \\\\  \\hline   NGC 6712      &    20.8 &    18 53 04.3  & -08 42 22 &  25.35  & -4.32 &   6.9  & 3.5 \\\\ Lynga 7      &19.7      &  16 11 03.0  &-55 18 52  &328.77  & -2.79 &   7.2  & 4.2 \\\\ NGC 6535   &17.2        &   18 03 50.7 & -00 17 49 &  27.18 &  10.44  &  6.8  & 3.9 \\\\ NGC 5927   & 17.1          &15 28 00.5 & -50 40 22 & 326.60 &   4.86  &  7.6 & 4.5 \\\\ NGC 6749  &16.6           & 19 05 15.3  & +01 54 03  & 36.20  & -2.20 &   7.9 &  5.0 \\\\ NGC 6760   &16.3          & 19 11 12.1  &+01 01 50  & 36.11 &  -3.92   & 7.4 &  4.8\\\\ NGC 6352    &16.0        &  17 25 29.2 & -48 25 22&  341.42 &   -7.17 &   5.7 &  3.3 \\\\ NGC 6717 (Pal 9) & 14.5 &      18 55 06.2 & -22 42 03  & 12.88 & -10.90   & 7.1  & 2.4 \\\\ NGC 6171 (M 107) &14.0  &     16 32 31.9  &-13 03 13  &  3.37 &   23.01 &   6.4 &  3.3 \\\\ IC 1276  (Pal 7)   & 13.7&    18 10 44.2 & -07 12 27 &  21.83 &   5.67  &  5.4  & 3.7 \\\\  \\hline  \\end{tabular} \\end{minipage} \\end{table*}%   \\begin{table*} \\centering \\begin{minipage}{140mm} \\caption{The ten globular clusters having the highest difference in signal between the CDM halo configuration and MOND. Astronomical data are from Harris (1996)  and $R_{gc}$ is the distance from the Galactic centre.}\\label{topten_diff} \\begin{tabular}{@{}llllllll@{}}  \\hline Name    &   $v$ (cm/s)\\footnote{Assuming $T=$15 years} &    RA  (hours)\\footnote{J2000 epoch}  & dec (degrees)   &   $l $ (degrees)  &   $b$ (degrees)   &  $r_s$ (kpc) & $R_{gc}$ (kpc)\\footnote{Assumes $R_0=8$ kpc} \\\\  \\hline  Pal 1         &  1.18 &      03 33 23.0  & +79 34 50 & 130.07  & 19.03  & 10.9 & 17.0 \\\\ NGC 2298  &1.18       &     06 48 59.2 & -36 00 19 &  245.63 & -16.01 &  10.7 & 15.7 \\\\ NGC 1851    &1.17    &      05 14 06.3  & -40 02 50 & 244.51 & -35.04 &  12.1 &  16.7 \\\\ NGC 1904 (M 79)  &1.16  &      05 24 10.6 & -24 31 27  & 227.23 &  -29.35  & 12.9  & 18.8 \\\\ NGC 288   &1.15   &         00 52 47.5 & -26 35 24 & 152.28 & -89.38 &    8.8 &  12.0  \\\\ NGC 5272 (M 3) &1.13   &       13 42 11.2  & +28 22 32 &  42.21 &  78.71 &  10.4  & 12.2 \\\\ NGC 5904 (M 5) &1.11  &       15 18 33.8 & +02 04 58  &  3.86  & 46.80 &   7.5  & 6.2 \\\\ NGC 1261    &   1.09&       03 12 15.3  & -55 13 01  & 270.54 &  -52.13 &   16.4  & 18.2 \\\\ NGC 362   &   1.08&         01 03 14.3 & -70 50 54 & 301.53 &  -46.25  &   8.5  & 9.4 \\\\ NGC 7099 (M 30) &1.07   &     21 40 22.0  & -23 10 45 &  27.18 & -46.83 &   8.0  & 7.1\\\\  \\hline   \\end{tabular} \\end{minipage} \\end{table*}%"
    },
    "0807/0807.4966.txt": {
        "abstract": "{The exact nature of the quiet solar chromosphere, and especially its temporal variations, are still subject of intense debate. One of the contentious issues is the possible role of magnetic field in structuring the quieter solar regions.} {We characterize the dynamics of the quiet inter-network chromosphere by studying the occurrence of acoustic shocks and their  relation with the concomitant photospheric structure and dynamics, including small scale magnetic structures.} {We analyze a comprehensive data set that includes high resolution chromospheric (\\CaII 854.2 nm) and photospheric (\\FeI 709.0 nm) spectra obtained with the IBIS imaging spectrometer in two quiet-Sun regions. This is complemented by high-resolution sequences of MDI magnetograms of the same targets. From the chromospheric spectra we identify the spatio- temporal occurrence of the acoustic shocks. We compare it with the photospheric dynamics by means of both Fourier and wavelet analysis, and study the influence of magnetic structures on the phenomenon.} {Mid-chromospheric shocks occur within the general chromospheric dynamics pattern of acoustic waves propagating from the photosphere. In particular, they appear as a response to underlying powerful photospheric motions at periodicities nearing the acoustic cut-off, consistent with 1-D hydrodynamical modeling. However, their spatial distribution within the supergranular cells is highly dependent on the local magnetic topology, both at the network and internetwork scale. We find that large portions of the internetwork regions undergo very few shocks, as ``shadowed'' by the horizontal component of the magnetic field. The latter is betrayed by the presence of chromospheric fibrils, observed in the core of the \\CaII line as slanted structures with distinct dynamical properties. The shadow mechanism appears to operate also on the very small scales of inter-network magnetic elements, and provides for a very pervasive influence of the magnetic field even in the quietest region analyzed. } {The magnetic field might play a larger role in structuring the quiet solar chromosphere than normally assumed. The presence of fibrils highlights a clear disconnection between the photospheric dynamics and the response of the geometrically overlaying chromosphere.  As these results hold for  a mid-chromospheric indicator such as the \\CaII 854.2 line, it is expected that diagnostics formed in higher layers, such as UV lines and continua, will be affected to a larger extent by the presence of magnetic fields, even in quiet regions. This is of relevance for the  chromospheric models that make use of such diagnostics.} ",
        "introduction": "Observations have long shown that the outer solar atmosphere is not in radiative equilibrium, with the net radiative loss of the quiet chromosphere estimated at between $\\approx$ 4 kWm$^{-2}$ \\citep{avrett81} and 14 kWm$^{-2}$ \\citep{anderson_89}. The identification of the source and exact mechanism of deposition of the energy necessary to maintain a stationary situation is often referred to as the problem of chromospheric heating. It has been a vexing problem in solar physics for decades.  In stellar atmospheres, strong density stratification makes it relatively easy for propagating acoustic waves, excited by the turbulent convection, to develop into shocks. Indeed this mechanism was proposed early on as a plausible means to provide the necessary energy input to the solar chromosphere \\citep{biermann_48,schwarzschild_48}. Much theoretical and numerical work has since been conducted to assess the viability of the acoustic heating as the basic heating mechanism for nonmagnetic chromospheres of slowly rotating stars \\citep[e.g.][]{narain_96,ulmrew_03}.  A landmark work on this topic was the 1-D radiative-hydrodynamical modeling of \\citet{carlsson_95,carlsson_97,carlsson_02}, who derived the chromospheric response to acoustic waves in the absence of magnetic fields. Using an observed photospheric piston, they could reproduce the temporal evolution of the chromospheric \\CaII H spectral line over long intervals with remarkable accuracy. In particular, they reproduced the occurrence of the H$_{2V}$  (K$_{2V}$) grains observed in the spectral profile of the resonance H and K lines of of Ca\\,II. The grains consist of quasi-periodic  brightenings in the violet wings of the lines, with a $\\sim$150--200 s repetition rate, and are associated with strong redshifts of the line core position \\citep[see][hereafter RU91, and Section \\ref{s_evid_shocks}]{rutten_91}. \\citet {carlsson_97}, hereinafter CS97, showed how the grains are due to weak acoustic shocks developing at heights of about 1 Mm in the non magnetic atmosphere, driven primarily by photospheric acoustic waves with periodicities close to the cutoff value of $\\sim$5.5 mHz. A controversial by-product of their model is that the intermittent acoustic shocks greatly enhance both the emissivity and the temperature but only locally, without causing a net outward increase in the {\\it average} chromospheric temperature \\citep{carlsson_95}.  If the CS97 shocks correctly represent the acoustic heating mechanism, the non-magnetic chromosphere would be essentially cold. This is contrary to the common view, derived from semi-empirical static models, of a chromosphere characterized everywhere by a steady temperature rise \\citep {avrett81,fontenla07,avrett_08}. The failure of the CS97 model to reproduce spatially and temporally resolved quiet Sun UV observations obtained with SUMER in chromospheric lines \\citep{carlssonetal_97} %Carlsson, Judge & Wilhelm 1997 ApJL 486 63 has further sharpened the controversy.  A strong debate has subsequently ensued about the possible shortcomings of the CS97 model and the general viability of the acoustic heating mechanism in shaping the quiet solar chromosphere. Much discussion has been devoted to the role of  high frequency acoustic waves \\citep[among many others,][]{fossum_05,wedemeyeretal_07,cuntz_07,carlssonetal_07}, %Fossum & Carlsson 2005 Nature 435 919 % Carlsson etc 2007, PASJ (Hinode data) % Wedemeyer et al Coimbra %Cuntz et al 2007 ApJL 657 57 the need to address the problem in 3-D \\citep[e.g.][]{wedemeyer_04,ulmschneider_05}, %Wedemeyer et al 2004, A&A and the role of magnetic fields in the quiet Sun \\citep[e.g.][] {judge_03,mcintosh_03,carlsson_07,wedemeyeretal_07}. % Judge Carlsson & stein 2003 ApJ 597 1158 At the bottom of the debate lies the fundamental question of whether the  internetwork chromosphere is wholly dynamic in nature \\citep{carlsson_07,martinez_07,wedemeyer_07} %Martinez-Sykora et al 2007 arXiv:0712.3854v1 or whether the dynamic variations represent only minor perturbations on a semi-static state of the kind described in semi-empirical models \\citep{kalkofen_99,avrett_08}. %Avrett & Loeser, 2008, ApJSS ...    \\begin{figure*} \\centerline{ \\includegraphics[scale=.6]{fig1.eps} } \\caption{Top row: data set 1; bottom row: data set 2. The images refer to data acquired around the middle of the observing sequences. Axes are given in arcsec. Panels {\\it a, e}: Broadband continuum at 710 nm. The image in data set 2 has been speckle reconstructed with the technique of \\cite{woeger_06}. Panels {\\it b, f}: co-temporal HR MDI maps, saturated at $\\pm$ 500 G. Data set 2 attains to an enhanced, bi-polar network region. Panels {\\it c, g}: Intensity in the red wing of the CaII 854.2 nm at about 0.1 nm from line core. Bright points clearly correspond to small scale magnetic features. Panels {\\it d, h}: Line core intensity of CaII 854.2 nm. Note the large extension of fibrils originating in the magnetic elements. } \\label{fig_fov} \\end{figure*}  From the observational point of view, gaining further insight into this issue requires new data that can simultaneously achieve high spatial and temporal resolution, and spectral information with different diagnostics, over extended fields of view (FOV). Most of the debate to this point has been based on (older) spectrographic observations obtained in the \\CaII H and K lines, and on UV spectral data obtained by Skylab or, more recently, by SUMER on board SOHO \\citep{wilhelm_95}. % Solar Physics di SOHO Spectrographic data suffer however from a limited FOV in fixed slit observations, or from a low cadence when scanning an extended region.  In the UV, the limited flux and instrumental response further lower the achievable spatial and temporal resolution, an important limitation if the dynamic picture just described for the chromosphere is the correct one \\citep{wedemeyer_07}.  Historically, imaging with broad-band filters has been widely used, both in the UV and visible range. Such data, however, often mix signals arising from vastly different regions of the solar atmosphere, as well as miss signatures of strong velocities that might shift the lines outside the filter band. For a general review on the issues of chromospheric observations we refer to \\citet{rutten_07}, only adding here that recently the POLIS spectrograph \\citep{beck_05} %Beck et al 2005 A&A 437 1159 has become operational, producing new high quality \\CaII H spectra that have been employed in studies of the quiet chromosphere \\citep{rezaei_07,beck_08,rezaei_08}. %Rezaei, et al 2007 A&A 466 1131 %Beck, Schmidt et al 2008 A&A 479 213  In this paper  we use an entirely different type of instrument, the Interferometric BIdimensional Spectrometer \\citep[IBIS,][]{cavallini_06}, to derive novel results about the dynamics of the quiet solar chromosphere and in particular the presence (or absence) of acoustic shocks. IBIS is an imaging spectrometer installed at the Dunn Solar Telescope (DST) of the US National Solar Observatory, and provides observations that combine high spatial and temporal resolution over an extended two-dimensional FOV, with the full spectral information in both photospheric and chromospheric lines. In this sense, it is an ideal instrument to overcome many of the observational shortcomings outlined above.  We analyze the behavior of the chromospheric \\CaII 854.2 nm line which, as shown in \\citet{cauzzi_08} (hereafter Paper I), is one of the most promising diagnostics for high resolution chromospheric studies. Equally important, the availability of spectral information over a 2-D FOV is crucial to understanding the very dynamic and non-local chromosphere. We further combine the \\CaII 854.2 observations with simultaneous and cospatial photospheric data also obtained with IBIS and magnetic data from  MDI \\citep{scherrer_95}. %Solar Physics SOHO issue This allows us to address the relationships between the chromospheric and photospheric dynamics, and the influence of magnetic field. We analyze two completely analogous data sets obtained on separate days, both times in quiet regions near disk center, and in so doing uncover the fundamental role played by the local magnetic topology.  The paper is organized as follows. In Section \\ref{s_obs} we describe the observations. Sections \\ref{s_evid_shocks} and \\ref{s_findshocks} provide evidence that mid-chromospheric acoustic shocks are clearly observed in the \\CaII 854.2 nm line, and outline the methods we use to identify them. In Sect. \\ref{s_properties} we report on the derived shock properties, both  in relation to earlier findings obtained primarily  with \\CaII K observations, and shock15.gc.as new results regarding their spatial distribution. Sections \\ref{s_mag} and \\ref{s_phot}  address the occurrence of the shocks in relation to the magnetic field structure and evolution, and the photospheric dynamics. Finally, in Sections \\ref{s_discussion} and \\ref {s_concl} we discuss our findings and provide conclusions.  % \\begin{figure} \\centerline{ \\includegraphics[scale=0.75]{fig2.eps} } \\caption{Full disk EIT images acquired in Fe\\,XII 19.5 nm on 2004 May 31 (left) and 2004 June 02 (right). The cross-hairs indicate the center of the IBIS FOV for each day. The IBIS FOV covers roughly the small coronal bright point intersected by the cross-hair on the right panel.} \\label{fig_foveit} \\end{figure} % ",
        "conclusions": "In our analysis, we reach two main conclusions, both of which are represented in Fig. \\ref{fig_three_chr}.  The first is that temporally resolved spectra of the \\CaII 854.2 nm provide a comprehensive picture of the occurrence of acoustic shocks in the quiet solar atmosphere. We determine that these shocks are a fundamental part of the general chromospheric dynamics, as they essentially make up the bulk of the three-minute oscillations in the middle chromosphere. We find that the shocks  occur in direct response (after a 120 second delay) to powerful photospheric oscillations near the acoustic cutoff frequency, with a primarily vertical propagation. No major role is found for waves with periods below 120 seconds or with a slanted propagation, consistent with the 1-D purely hydrodynamical simulations of \\citet{carlsson_97}. Portions of the quiet chromosphere are then the site of brief episodes when the acoustic shocks dramatically enhance the temperature (Fig. \\ref{fig_three_chr} {\\it a} and {\\it b}) mixed with extended periods of lower temperatures  (Fig. \\ref{fig_three_chr} {\\it c} and {\\it d}), as opposed to the steady-state mean temperature rise typically assumed in semi-empirical models \\citep[see the discussion in][]{carlsson_95,carlsson_07}. %Carlsson & Stein 1995 ApJL %Carlsson review in Coimbra  The second, more important conclusion is that these acoustic processes are very often significantly disturbed by the presence of magnetic field, even in the quiet Sun. Both at the network and (surprisingly) at internetwork scales, the signature of acoustic shocks is suppressed both in and around the footpoints of magnetic structures. Hence the distribution of shocks is considerably reduced compared to what would be inferred from the amplitude of the photospheric oscillations alone. The main player in this respect appears to be the magnetic topology,  and in particular  whether the field has an open or closed configuration. The areas of modified acoustic behavior coincide with the presence of chromospheric fibrils, betraying an atmosphere that is stratified differently with respect to the classical 1-D view.  These results provide an important hint concerning the longstanding controversy about why the dynamical model of CS97 fails to predict the ubiquitous emission observed in UV lines \\citep [e.g.][]{kalkofen_99,kalkofen_01,carlsson_07,avrett_08}. %Avrett & Loeser 2008 ApJS 175, 229 Our analysis is based on observations in the \\CaII 854.2 nm line, a spectral signature formed relatively low in the atmosphere. Our data provide uncontroversial evidence that even at these heights  the quiet atmosphere is highly structured by pervasive magnetic fields that connect different locations. The  presence of fibrils is a clear testimony of these inter-connections\\footnote{We note that this magnetic structuring  has been observed for decades in H$\\alpha$, a diagnostics that has figured little in the discussion on chromospheric heating.}. It is then natural to assume that other spectral diagnostics forming in higher, less dense parts of the atmosphere,  experience such a magnetic structuring to an even larger extent, much as advocated by \\citet{ayres_02}. %Ayres 2002, ApJ 575 1104  Far UV lines and continua are one such diagnostic, and they figure prominently in the derivation of the temperature structure in semi-empirical models, often as spatial and temporal averages \\citep[e.g.][]{avrett_08}. It follows that a substantial fraction of the UV flux could originate in magnetic structures, even in the {\\it quiet} internetwork chromosphere. This could happen either because of the volume filling of the magnetic field at these heights,  or because of the dominant weight of magnetic-related brightness in the averages: Fig. \\ref {fig_three_chr} illustrates how magnetic concentrations introduce a brightness component with a different spatial distribution and temporal evolution than for the purely acoustic case. This happens on spatial scales at the limit of the resolution of many modern instruments, including SUMER.  Any averaged or marginally resolved observation of the chromospheric flux might then contain a sizable magnetic component, in ways not immediately predictable by using only photospheric indicators. A dominance of the magnetic topology in the quiet upper chromosphere would also explain why UV observations obtained with SUMER have produced so many disparate results. These range from the clear signature of upward propagating acoustic waves reported by \\citet{wikstol_00} %Wikstol et al, 2000, ApJ 531 1150 all the way from the photosphere up to the base of the corona, as could happen in the case of a coronal hole, to instances of complete disconnection between the observed lower and upper chromospheric dynamics \\citep{judge_01,judge_03}, as would be the case for regions with a locally closed magnetic topology. In this scenario, obviously there is no need to reconcile the predictions of CS97, a purely non-magnetic model, with the UV observations, as they pertain to completely different physical regimes.   We thus conclude that the  radiative losses in the ``non-magnetic'' (quotes are necessary at this point) solar chromosphere are strongly influenced, and possibly dominated, by processes related to magnetic fields. We remark that this same conclusion has been reached by Judge and collaborators \\citep{judge_98,judge_03,judge_04}, in a precise analysis of UV emission lines  for both the case of the Sun and other stars with convective envelopes. % Judge & Carpenter 1998 ApJ 494 828 % Judge Carlsson & stein 2003 ApJ 597 1158 %Judge et al 2004, ApJ 609 392 The presence of fibrils in large fractions of the quiet atmosphere provides a plausible agent for such dominance. On the one hand, they disrupt the normal vertical propagation of acoustic waves from lower layers, strongly reducing the effects of a purely acoustic mechanism in shaping the chromosphere. On the other, they provide the means to transport energy of magnetic origin, or mediated by the magnetic elements, from the network footpoints towards the center of the supergranular cells. We note that this idea is consistent  with several recent works that have called into question whether the photospheric acoustic flux at high frequencies is sufficient to compensate the chromospheric radiative losses \\citep {fossum_05,fossum_06} or, even more pertinent,  whether different energy sources such as {\\it low-} frequencies magneto-acoustic waves, or atmospheric gravity waves, might play a more important role  \\citep{jefferies_06,straus_08}. % Jefferies et al 2006 portal paper % Straus et al 2008 ApJL 681 125"
    },
    "0807/0807.2007_arXiv.txt": {
        "abstract": "The NeXT (New exploration X-ray Telescope), the new  Japanese X-ray Astronomy Satellite following Suzaku,  is an international X-ray mission which is currently planed for launch in 2013. NeXT is a combination of wide band X-ray spectroscopy (3--80~keV) provided by multi-layer coating, focusing hard X-ray mirrors and hard X-ray imaging detectors, and high energy-resolution soft X-ray spectroscopy (0.3--10~keV) provided by thin-foil X-ray optics and a micro-calorimeter array.  The mission will also carry an X-ray CCD camera as a focal plane detector for a soft X-ray telescope and a  non-focusing soft gamma-ray detector. With these instruments, NeXT covers very wide energy range from 0.3~keV to 600~keV\\@.  The micro-calorimeter system will be developed by international collaboration lead by ISAS/JAXA and NASA\\@. The simultaneous broad bandpass, coupled with high spectral resolution of $\\Delta E $ $\\sim$7~eV by the micro-calorimeter will enable a wide variety of important science themes to be pursued. ",
        "introduction": "The NeXT (New exploration X-ray Telescope) mission has been studied as the next key X-ray astronomy mission of Japan, which will be developed under international collaboration (see Fig.~\\ref{Fig:NeXT}).  It has completed the pre-phase A study in 2006 and entered into Phase A since June 2007. A Request For Proposal (RFP) for the spacecraft was released by ISAS/JAXA to Japanese industry in November 2007 and NEC was selected. In May, 2008, System Definition Review, which is required before entering Phase B, was completed. We  anticipate that NeXT will enter Phase B in the middle of 2008 for the planned launch year of 2013. More recently, NASA has selected the US participation in NeXT as a mission of opportunity.   Under this program, the NASA/Goddard Space Flight Center will collaborate with ISAS/JAXA on the implementation of an X-ray microcalorimeter spectrometer\\cite{Ref:ProposaNASA}.  The NeXT mission objectives are to study the evolution of yet-unknown obscured super massive Black Holes (SMBHs) in Active Galactic Nuclei (AGN); trace the growth history of the largest structures in the Universe; provide insights into the behavior of material in extreme gravitational fields; determine the spin of black holes and the equation of state of neutron stars; trace particle acceleration structures in clusters of galaxies and SNRs; and investigate the detailed physics of jets.  In this paper, we will summarize the scientific requirements, the mission concept and the current baseline configuration of instruments of NeXT\\@.        \\begin{figure} \\centerline{\\includegraphics[width=12cm,angle=0]{Figure/figure1.eps}} \\caption{Artist's drawing of the NeXT satellite. The focal  length of the Hard X-ray Telescope (HXT) is 12m, whereas the Soft X-ray Spectrometer (SXS) and Soft X-ray Imager (SXI) will have a focal length of 6~meters.} \\label{Fig:NeXT} \\end{figure} ",
        "conclusions": ""
    },
    "0807/0807.2970_arXiv.txt": {
        "abstract": "{Beating the Earth's day-night cycle is mandatory for long and continuous time-series photometry and had been achieved with either large ground-based networks of observatories at different geographic longitudes or when conducted from space. A third possibility is offered by a polar location with astronomically-qualified site characteristics.} {In this paper, we present the first scientific stellar time-series optical photometry from Dome~C in Antarctica and analyze approximately 13,000 CCD frames taken in July 2007.} {The optical pilot telescope of the ``International Robotic Antarctic Infrared Telescope'', named ``small IRAIT'' (sIRAIT), and its $UBVRI$ CCD photometer were used in $BVR$ for a continuous 243 hours (10.15 days) with a duty cycle of 98\\% and a cadence of 155~sec. The prime targets were the chromospherically active, spotted binary star V841~Cen and the non-radially pulsating $\\delta$-Scuti star V1034~Cen.}{We confirm the known 0.2-day fundamental period of V1034~Cen and found a total of 23 further periods between 2.2 hours and 3.5 days. V841~Cen's $V$ amplitude due to spots appeared to be at a record high of 0\\fm4 in $V$ in July 2007. We present a spot-model analysis with a light-curve inversion technique and found the star with a spot filling factor of 44\\% of the visible hemisphere, among the highest ever measured values for active stars, and a temperature difference photosphere minus spot of 750$\\pm$100~K. Its odd-numbered (for a single site) rotation period was determined with a higher precision than before (5.8854$\\pm$0.0026 days) despite the comparably short data set. The rms scatter from a 2.4-hour data subset was 3~mmag in $V$ and 4.2~mmag in $R$. The differential data quality is 3--4 times better than with the 25cm Fairborn Automatic Photoelectric Telescope in southern Arizona and is likely due to the exceptionally low scintillation noise at Dome~C.}{We conclude that high-precision CCD photometry with exceptional time coverage and cadence can be obtained at Dome~C in Antarctica and be successfully used for time-series astrophysics.} ",
        "introduction": "Time-series photometry is a powerful tool to understand cosmic variabilities and their many underlying physical mechanisms, from Gamma Ray Bursts to the non-radial oscillations of our Sun. World-wide networks around the globe were organized to bypass the Earth's day-night cycle, e.g. by the ``Global Oscillation Network Group'' (Harvey et al. \\cite{gong}), the ``Whole Earth Telescope'' (Nather et al. \\cite{wet}) or the ``Multi Site Continuous Spectroscopy'' campaigns (Catala et al. \\cite{musicos}), just to name a few. Despite their tremendous success such networks must cope with the many different site characteristics, the individual weather patterns and, most importantly, with the different instrument/detector combinations and the respective calibration issues. An alternative to such networks is a single polar site.  The French-Italian Antarctic station \\emph{Concordia} at Dome~C at a height of 3200m on the east Antarctic plateau is currently, besides the U.S. South Pole station, the only plateau station populated over winter (=night). In principle, it thus enables astronomical observations comparable to regular observatories at temperate sites. Dome~C had received world-wide attention from the night-time astronomical community when it became known that the seeing conditions on the east Antarctic plateau are likely the best on the entire planet with a median seeing of 0.3\\arcsec\\ and occasionally even below 0.1\\arcsec\\ at a height of approximately 30m above ground (Lawrence et al. \\cite{law:nature}, Agabi et al. \\cite{agabi06}).  \\begin{figure*}[ht] \\centering \\begin{minipage}[t]{12cm} \\vspace{0pt} \\centering \\includegraphics[width=11cm]{F1.eps} \\end{minipage}% \\begin{minipage}[t]{5.8cm} \\vspace{0pt} \\centering \\caption{Identification of stars in the CCD FOV (8$\\arcmin\\times$5.3\\arcmin ). The three bright stars are our primary targets V841~Cen (right, $V$=8\\fm5) and V1034~Cen (left) and the comparison star CD-59\\degr 5309 (top). The comparison star and stars \\#4 and \\#7 were identified as variable (see Table~\\ref{T1}). The image shown is a composite of 20 individual R-band frames, each a 40-sec integration. The coordinates are for equinox 2000.0. The faintest stars are 16th magnitude.} \\label{F1} \\end{minipage} \\end{figure*}  Science cases that require continuous high-precision photometry uniquely benefit from the absent 24-hour day-night cycle and the consequently stable atmosphere in general. For example, $\\delta$~Scuti stars are known to have a complex surface oscillation spectrum involving many modes and frequencies. The detection of patterns of closely-spaced peaks in different modes enables the determination of the internal stellar structure. The essential observational restriction is the frequency resolution, which is proportional to the length of the photometric time series. It sets limits to the mode identification and thus their unique interpretation. Another example are spotted stars. Magnetic spots, like those on the Sun, are tracers of the internal dynamo activity. Spot size and temperature are related to the magnetic flux that the emerging flux tube transported up from the interior. However, surface velocity fields like differential rotation affect the magnetic field (and vice versa) so that spots are likely only indirectly linked to the dynamo. Observing the migration behavior of starspots from one stellar rotation to the other, however, may constrains some global properties of the dynamo.  Additionally, optical photometry of bright stars may be brought down to the photon-noise level because atmospheric scintillation noise appears to be a factor 3.6 smaller (Kenyon et al. \\cite{photom}) than at the best temperate sites. The detailed issues for optical photometry from Dome~C were highlighted by several authors, most recently by Vernin et al. (\\cite{vernin07}), Strassmeier (\\cite{kgs:roscoff}) and Kenyon et al. (\\cite{photom}). We also note the early attempts and experiences from the South Pole's ``Vulcan-South'' experiment\\footnote{http:/\\slash{}www.polartransits.org} that, unfortunately, suffered from the comparable harsh winds at the South Pole.  A continuous 1500-hour night opens up a new window for science cases like the search for extra-solar planets (e.g. Pont \\& Bouchy \\cite{pon:bou}, Deeg et al. \\cite{deeg}), for asteroseismology (e.g. Fossat \\cite{foss:astro}) and for stellar rotation and activity studies (e.g. Strassmeier \\& Ol\\'ah \\cite{eddi}). Several astronomical experiments are now planned for Dome~C (for a compilation see Strassmeier et al. \\cite{tenerife}). Among these pilot experiments is a small 25-cm optical precursor telescope for the 80-cm infrared telescope IRAIT (Tosti et al.~\\cite{irait}) and the 2$\\times$60cm Schmidt-telescope ICE-T (Strassmeier et al.~\\cite{icet}). sIRAIT was designed to experience the expected same difficulties in the same operative conditions foreseen for the larger projects with the goal of performing a period of full tests on a smaller scale.  In the current paper, we present and analyze the first data from sIRAIT from one CCD field obtained in July during winter 2007. Our primary target, V841~Cen, is a spotted, very active RS~CVn binary with occasional flares and with an orbital period so close to a multiple of a day (5.998 days) that time-series data from a single, non-polar site severely undersamples its light curve. Our secondary target is the $\\delta$-Scuti star V1034~Cen, a non-radially pulsating star in the lower part of the instability strip with a (likely) fundamental period of $\\approx$0.2 days.  \\begin{figure*}[tbh] \\includegraphics[angle=-90,width=\\textwidth,clip]{F2.eps} \\caption{Ten-day continuous differential $VR$ photometry of V841~Cen minus V1034~Cen. Note that the long-period variation with an amplitude of 0\\fm4 in $V$ is due to spots rotating in and out of view on V841~Cen while the short-period variations with an amplitude of $\\approx$0\\fm02 are due to non-radial pulsations of V1034~Cen. The insert shows a fraction of $V$ data that is nearly free of intrinsic short-term stellar variations. The residuals from a simple linear fit to a 2.4-hour subset suggest a rms scatter of just 3~mmag in $V$ (4.2~mmag in $R$). For such a long duration, this is 3--4 times better than an equally sized telescope at a temperate site. }\\label{F2} \\end{figure*} ",
        "conclusions": "We presented 243 continuous hours of optical photometry from Antarctica with a duty cycle of 98\\% and a cadence of 155 seconds. A 3~mmag rms precision in $V$ over 2.4 hours with the 25-cm sIRAIT telescope was achieved for the two bright 8\\fm5 target stars. This is a factor 3--4 better than what we had obtained with the 25-cm T1 automatic photoelectric telescope (APT) at Fairborn Observatory in southern Arizona (Strassmeier et al.~\\cite{str:bar}, Henry \\cite{h95a}). Most likely this is attributed to scintillation noise smaller by a factor 3--4 compared to temperate observing sites, as reported by Kenyon et al. (\\cite{photom}). We conclude that high-precision continuous photometry within the turbulent ground layer just one meter above ground is feasible at Dome~C, even with low-cost, partly commercial components. The main problems we had encountered with sIRAIT in its first winterover were all due to the quality of the equipment rather than with the harsh Antarctic environment. This makes us strongly believe that our proposed 2$\\times$60-cm, more optimized and robust, photometric facility ICE-T (Strassmeier et al.~\\cite{icet}) as well as the intermediate-step 40cm \\emph{a-step} experiment (Fressin et al. \\cite{astep}) are well suited for a site like Dome~C and could, for some favorable cases, even challenge photometry from space."
    },
    "0807/0807.0961_arXiv.txt": {
        "abstract": "It is shown, that primary proton spectrum, reconstructed from sea-level and underground data on muon spectrum with the use of QGSJET~01, QGSJET~II, NEXUS~3.97 and SIBYLL~2.1 interaction models, demonstrates not only model-dependent intensity, but also model-dependent form. For correct reproduction of muon spectrum shape primary proton flux should have non-constant power index for all considered models, except SIBYLL~2.1, with break at energies around 10--15~TeV and value of exponent before break close to that obtained in ATIC-2 experiment. To validate presence of this break understanding of inclusive spectra behavior in fragmentation region in p-air collisions should be improved, but we show, that it is impossible to do on the basis of the existing experimental data on primary nuclei, atmospheric muon and hadron fluxes. ",
        "introduction": "At present information on the characteristics of hadronic interactions in fragmentation region is still scarce or missing and experiments with `roman pots' are anticipated to improve the situation. Some of this information, in principle, could be obtained with the use of the data on cosmic ray (CR) muon and hadron spectra, provided primary spectra are known with high precision, but that is not the case. The obvious obstacle here is that at high energies primary cosmic ray (PCR) fluxes, measured in direct experiments, themselves are functionals of various interaction parameters plus their accuracy is appreciably affected by additional systematic effects~\\cite{grigorov1973_eng,jacee1986,msuhadrons2000_eng,runjob_eng}. In the series of our papers~\\cite{ya2004,ya_yaf2006_deficit_eng} we underlined, that these effects can lead to underestimation of light nuclei fluxes and thus can explain discrepancy between measured and calculated muon fluxes for $E_\\mu>100$~GeV. Preliminary data of ATIC-2~\\cite{atic2_RCRC2007}, covering the gap between magnetic spectrometer and  emulsion chamber experiments, seems to be in concordance with our conclusions, but situation is more complicated in fact, as further consideration will show. ATIC-2 slope of proton spectrum $\\gamma_p=2.63$~\\cite{zatsepin_3cmodel} for primary energies below 10~TeV is in remarkable contradiction with previously measured values $\\gamma_p=2.74$ and $\\gamma_p=2.80$ by RUNJOB~\\cite{runjob_apj2005} and JACEE~\\cite{jacee} experiments correspondingly, but this discrepancy is removed by steepening in ATIC-2 proton spectrum at energies above 10~TeV. These new data were already exploited in extensive calculations of muon and hadron fluxes with number of interaction models at different atmospheric depths and zenith angles in~\\cite{kochanov_2008} where it was shown, that their use allows to get reasonable agreement with the most of the data under appropriate choice of hadronic interaction parameters. In fact, it is not possible to obtain concordant conclusions on primary spectra or hadronic interactions parameters coming even from much smaller subset of the experimental data. In this paper we demonstrate this on the basis of analysis of the data on muon flux at vertical direction and only one set of the data on hadron flux of EAS-TOP~\\cite{eas_top_p}. First, we use the data of sea-level and underground experiments to obtain conclusions on behavior of muon spectrum at sea-level in the energy range 40~GeV--10~TeV. Further we analyze influence of uncertainties in the muon data and interaction parameters on properties of reconstructed primary proton fluxes. And finally we show, that such different in approach and characteristics interaction models as SIBYLL~2.1~\\cite{sibyll,sibyll2.1} and NEXUS~3.97~\\cite{nexus} can bring to hardly distinguishable predictions on muon and hadron fluxes. ",
        "conclusions": "The progress in CR and high-energy physics, achieved during last 10--15 years allowed to turn from statements about satisfactory (qualitative) concordance between different kinds of data to investigation of more fine effects. As an example, we managed to show that reconstructed from the data on vertical muon flux primary proton spectra have not only expected interaction model dependent intensities, but also model-dependent shapes. It is demonstrated, that application of QGSJET~01, QGSJET~II and NEXUS~3.97 models brings to proton spectrum with break at 10--15~TeV and power index $\\gamma_p$ before break close to that, measured in ATIC-2 experiment. Nevertheless one can see, that absolute proton flux for QGSJET~01 is hardly compatible with any data of direct experiments, and the break for all these three models is more moderate, compared to what can be inferred from ATIC-2 data, which though become less definite right in the break region. On the other hand SIBYLL~2.1 allows to reproduce shape of the muon spectrum with single power law proton spectrum, which is in reasonable agreement with both emulsion chamber and ATIC-2 data within experimental errors. Further improvement of our understanding of the situation, which is of primary astrophysical interest, can be achieved via experimental study as of muon CR component characteristics with water and ice neutrino telescopes and so of inclusive spectra $p+A\\to \\pi^\\pm, K^\\pm$ in fragmentation region. Reduction of uncertainties in the latter component with the use of the data on primary spectra, hadron and muon components, does not look possible, because of 1) poor correlation between muon and hadron production mechanisms, 2) ambiguity of existing CR experimental data and 3) possibility to realize self-consistent description of the data on muons and hadrons with the models, having remarkably differing inclusive spectra and underlying philosophies."
    },
    "0807/0807.0447.txt": {
        "abstract": " ",
        "introduction": "Since the classic work of 't Hooft and Veltman on quantum gravity on flat space-time~\\cite{'tHooft:1974bx}, many authors have studied the quantum behavior of gravitons. In particular, because of the potential relevance for inflationary cosmology, the quantum behavior of gravitons on a (locally) de Sitter background has been a widely studied subject over the past years~\\cite{Allen:1986ta,Antoniadis:1986sb,Allen:1986tt,Tsamis:1996qk,Tsamis:1992xa,Tsamis:1996qq,Tsamis:1996qm,Higuchi:2000ge,Higuchi:2000ye,Finelli:2004bm,Christensen:1979iy,Ford:1984hs,Abramo:1996gd,Abramo:1997hu,Abramo:1999wd,Mukhanov:1996ak,Losic:2006ht,Cognola:2005de,Janssen:2007ht,Abramo:1998hj,Abramo:2001dc,Abramo:2001dd,Iliopoulos:1998wq,JanssenProkopecMiao2008,Tsamis:2005je,Parker:1969au}. Of course de Sitter space is interesting for its relevance for inflationary cosmology. Therefore there are several works that study the potential influence of quantum behavior on inflationary observables~\\cite{Janssen:2009pb,Bilandzic:2007nb,Sloth:2006az,Sloth:2006nu}. Another line of research deals with the back-reaction of gravitational waves on the background space-time~\\cite{Abramo:1996gd,Abramo:1997hu,Mukhanov:1996ak,Losic:2006ht}. However of more interest for the present work is the one loop back-reaction by virtual gravitons on a de Sitter background, which has been calculated by several authors, using different techniques~\\cite{Finelli:2004bm,Tsamis:2005je,Parker:1969au}. Since it is not clear whether in these works exactly the same quantity is calculated and the renormalization schemes differ, the numerical coefficients differ. However the main result is the same: one loop graviton contributions to the expectation value of the energy momentum tensor result in a finite, time independent shift of the effective cosmological constant. Since the contribution can always be absorbed in a counterterm~\\cite{Tsamis:2005je}, the exact numerical coefficient coming from such a calculation has no physical meaning.\\\\  The goal of this paper is to go beyond the works mentioned above and calculate the one loop effective action induced by gravitons in a more general background space-time using dimensional regularization. The geometry we consider is a Friedmann-Lema\\^{\\i}tre-Roberston-Walker (FLRW) geometry with Hubble parameter $H=\\frac{\\dot{a}}{a}$ and the additional constraint that \\begin{equation}\\label{eps} \\epsilon\\equiv-\\frac{\\dot{H}}{H^2} \\end{equation} is a constant. If the universe contains only one type of matter, e.g. dust, radiation or a cosmological constant, this constraint is satisfied. In particular in matter era we have that $\\epsilon=3/2$, in radiation era $\\epsilon=2$ and de Sitter space corresponds to the limit $\\epsilon\\rightarrow 0$. One immediate problem with working in such a space-time, instead of de Sitter, is that for consistency of the Einstein equations the addition of matter fields is unavoidable. Whereas in de Sitter space, the only relevant metric fluctuations are the tensor modes (gravitational waves), in a more general setting one has to take the mixing of gravitational and matter degrees of freedom into account~\\cite{Brandenberger:2003vk,Mukhanov:1990me,Abramo:1998hj,Abramo:2001dc,Abramo:2001dd,Iliopoulos:1998wq}.\\\\  Naively one might think that any back-reaction will be insignificant, since it will in a cosmological setting typically be suppressed by $H^2/m_p^2$, with $m_p$ the Planck mass. The reason that this is not necessarily the case has to do with a second complication. This complication is that gravitons on a cosmological background space-time behave, apart from the tensorial structure, somewhat similar to massless scalars, with a certain amount of coupling to the Ricci scalar. Such a scalar field however is long known to posses problems in the infrared ~\\cite{Ford:1977in,Vilenkin:1982wt}, for a wide class of cosmological backgrounds. What happens is that due to cosmological particle production correlations of the Bunch-Davies vacuum grow too fast, causing the propagator to diverge in the infrared. Such a divergence should of course be regulated and in a recent work ~\\cite{Janssen:2008px} the massless scalar propagator on any constant $\\epsilon$ cosmological space-time was constructed, by assuming that the universe is described by a spatially compact manifold. Effectively this simply implies that any momentum integral has an infrared cut-off at some scale $k_0$. In this work, we shall use the propagators constructed in ~\\cite{Janssen:2008px} to obtain an infrared finite answer.\\\\ While the answers obtained in this way are infrared finite, particle production is of course a physical phenomenon. Long range correlations will therefore still be enhanced as time goes on. Therefore initially small quantum fluctuations, might -- at least in principle -- grow significantly in time. The contribution of these fluctuations to the stress energy tensor might therefore also grow in time, leading to a potentially significant back-reaction via the Einstein equations on the background space-time.  It is precisely this effect, that has led to several papers concerning quantum contributions to the stress energy tensor in the case of de Sitter space-time \\cite{Tsamis:1996qk,Tsamis:1996qm,Tsamis:1996qq,Abramo:1998hj,Abramo:2001dc,Abramo:2001dd,Linde:1982uu}. In these works it is found that at two loop order, the backreaction due to gravitons might become significantly, while for scalars this might happen at three loop order.  Such behavior, if indeed physically viable, has profound implications for the cosmological constant problem (see {\\it e.g.} Ref.~\\cite{Nobbenhuis:2004wn} for a review). In this paper we will not consider these higher loop effects, but we shall perform a one loop calculation, but in a more general space-time.\\\\   In section~\\ref{sgeo} we briefly review our background geometry. In section~\\ref{sprop} we briefly summarize the construction of the propagator as given in~\\cite{Janssen:2008px}. In section~\\ref{sprop2} we apply these results to calculate the relevant kinetic operators and propagators in a model with both matter and gravitational degrees of freedom. In section~\\ref{soneloop} we calculate the one-loop effective action contribution to the quantum corrected Friedmann equations and renormalize the theory. We discuss our results in section \\ref{s_disc} and we conclude in section \\ref{s_con}. ",
        "conclusions": "\\label{s_con} In this paper we have calculated the one loop contribution to the effective Friedmann equations due to graviton and matter quantum fluctuations. The background space-time we use is FLRW with the additional constraint that $\\epsilon$ is constant. In such a calculation many delicate issues have to be taken into account. First of all there is mixing between the gravitational and matter degrees of freedom. This mixing can (at least on-shell) be removed by field redefinitions (\\ref{fieldredef}) and (\\ref{Rotation}). Secondly, there is the issue of gauge freedom, which we can deal with in the standard way by adding a gauge fixing term (\\ref{gauge}) and the associated Fadeev-Popov ghost lagrangian (\\ref{ghost_lag}). Finally, to calculate the one loop contribution, one needs to know the propagators. The propagators needed will however typically be infrared divergent. To regulate this infrared divergence we assumed that the spatial sections of the universe are compact. This effectively implies that the propagators should be calculated using an infrared cut-off $k_0$. Using this approach the scalar propagator was constructed in \\cite{Janssen:2008px}. In the present work we have constructed the graviton and matter propagators by making use of the same method. All these technical issues can thus be resolved, leading to our final answer for the one loop Friedmann trace equation (\\ref{renormalized2}). Although the final answer might look intimidating, we can essentially write it as two types of contributions. One is  $H^4$ times a constant and the other is $H^4$ times some power of $z_0=|\\eta| k_0$. Since the tree level contribution is of the order of $H^2 m_p^2$, we see that -- as expected -- quantum corrections to the Friedmann equation are suppressed as $H^2/m_p^2$. There are thus only two possibilities for the quantum corrections to have a significant effect. Either the coefficient in front of $H^4$ is large, or any of the powers of $z_0$ grows fast enough to compensate the suppression after a sufficient amount of time. Looking at equation (\\ref{renormalized2}) we see that for typical values of the parameters, the coefficient will be of order one. One might think that the digamma functions might grow large. Indeed, these functions have a simple pole when the argument is a negative integer. However, as is shown in ~\\cite{Janssen:2008px}, these poles are due to a logarithmic infrared divergence. These poles should therefore be subtracted by the correction terms, and in ~\\cite{Janssen:2008px} it is shown that this indeed is what happens. Thus the only possibility for a significant backreaction of quantum fluctuations on the background space-time is if the correction terms grow fast enough in time. In section \\ref{s_disc} we showed however that this will not happen. The only possibility where the powers of $z_0$ grow fast enough is the special case of de Sitter ($\\epsilon=0$). However in this case, the growing power of $z_0$ drops out due to the vanishing of its prefactor.\\\\  It is however not clear whether these conclusions will also hold in more general cases. Generalizations could include for example studying a more general class of backgrounds, using different regularization schemes for the infrared. For example in~\\cite{JanssenProk} a scheme using mode matching is introduced. This scheme gives similar results for accelerating space-times as the scheme used here, but the results differ for decelerating space-times. Another generalization would be to include two or higher loop contributions into account. Indeed studies in de Sitter space show, for example, that at two loop order the backreaction due to gravitons might indeed grow in time~\\cite{Tsamis:1996qk,Tsamis:1996qm,Tsamis:1996qq}. This paper should be seen as a first step in trying to understand backreaction on a more general background space-time."
    },
    "0807/0807.1226_arXiv.txt": {
        "abstract": "{\\ls\\ and \\lsi\\ are two binaries that have been detected in  the TeV energy domain. These binaries are composed of a massive star and a compact object, possibly a young pulsar. The gamma-ray emission would be due to particle acceleration at the collision site between the relativistic pulsar wind and the stellar wind of the massive star. Part of the emission may also originate from inverse Compton scattering of stellar photons on the unshocked (free) pulsar wind.} {The purpose of this work is to constrain the bulk Lorentz factor of the pulsar wind and the shock geometry in the compact pulsar wind nebula scenario for \\ls\\ and \\lsi\\ by computing the unshocked wind emission and comparing it to observations.} {Anisotropic inverse Compton losses equations are derived and applied to the free pulsar wind in binaries. The unshocked wind spectra seen by the observer are calculated taking into account the $\\gamma-\\gamma$ absorption and the shock geometry.} {A pulsar wind composed of monoenergetic pairs produces a typical sharp peak at an energy which depends on the bulk Lorentz factor and whose amplitude depends on the size of the emitting region. This emission from the free pulsar wind is found to be strong and difficult to avoid in \\ls\\ and \\lsi. } {If the particles in the pulsar are monoenergetic then the observations constrain their energy to roughly 10-100 GeV. For more complex particle distributions, the free pulsar wind emission will be difficult to distinguish from the shocked pulsar wind emission.} ",
        "introduction": "Pulsars are fast rotating neutron stars that contain a large amount of rotational energy. A significant fraction of this energy is carried away by an ultra-relativistic wind of electrons/positrons pairs and possibly ions (see \\citealt{2007astro.ph..3116K} for a recent review). In the classical model of the Crab nebula \\citep{1974MNRAS.167....1R,1984ApJ...283..694K}, the pulsar wind is isotropic, radial and monoenergetic with a bulk Lorentz factor $\\gamma_0\\sim 10^6$ far from the light cylinder where the wind is kinetic energy-dominated ($\\sigma\\ll1$). The cold relativistic wind expands freely until the ram pressure is balanced by the surrounding medium at the standoff distance $\\rm{R_s}$. In the termination shock region, the pairs are accelerated and their pitch angle to the magnetic field are randomized, producing an intense synchrotron source. Moreover, the inverse Compton scattering of the relativistic electrons on soft photons produces high energy (HE, GeV domain) and very high energy (VHE, TeV domain) gamma-rays.  The shocked pulsar wind is thought to be responsible for most of the emitted radiation and gives clues about the properties of this region. However, our knowledge of the unshocked pulsar wind region is limited and based on theoretical statements. If the magnetic field is frozen into the pair plasma as it is usually assumed, there is no synchrotron radiation from the unshocked wind. Nevertheless, nothing prevents inverse Compton scattering of soft photons onto the cold ultra-relativistic pairs from occuring. The pulsar wind nebula (PWN) emission has two components: radiation from the shocked and the unshocked regions.  \\citet{2000MNRAS.313..504B} investigated the inverse Compton emission from the region upstream the termination shock of the Crab pulsar. Comparisons between calculated and measured fluxes put limits on the parameters of the wind, in particular the size of the kinetic energy dominated region. \\citet{2000APh....12..335B} investigated emission from an unshocked freely expanding wind with no termination shock in compact binaries. They computed spectra and light curves in the gamma-ray binary \\psrb, a system with a 48 ms pulsar and a Be star in a highly eccentric orbit. The resulting gamma-ray emission is a line-like spectrum.  In addition to \\psrb, two other binaries have been firmly confirmed as gamma-ray sources: \\ls\\ \\citep{2005Sci...309..746A} and \\lsi\\ \\citep{2006Sci...312.1771A}. They are composed of a massive O or Be star and a compact object in an eccentric orbit. The presence of a young pulsar was detected only in \\psrb\\ \\citep{1992ApJ...387L..37J}. Radio pulses are detectable but vanish near periastron, probably due to free-free absorption and interaction with the Be disk wind. The compact PWN scenario is most probably at work in this system and investigations were carried out to model high and very high energy radiation \\citep{1999APh....10...31K,2005MNRAS.356..711S,2007MNRAS.380..320K,2008MNRAS.tmp..347S}. In \\ls\\ and \\lsi\\, the nature of the compact object is still controversial but spectral and temporal similarities with \\psrb\\ argue in favor of the compact pulsar wind nebula scenario \\citep{2006A&A...456..801D}. The VHE radiation would therefore be produced by the interaction between the pulsar wind and the stellar companion wind. The massive star provides a huge density of seed photons for inverse Compton scattering with the ultra-relativistic pairs from the pulsar wind. Because of the relative position of the compact object, the companion star and the observer, the Compton emission is modulated on the orbital period. The vicinity of a massive star is an opportunity to probe the pulsar wind at small scales.  The component of the shocked pulsar wind was computed in \\citet{2008A&A...477..691D} for \\ls\\ and limits on the electron distribution, the pulsar luminosity and the magnetic field at the termination shock were derived. \\citet{2008arXiv0801.3427S} calculated the VHE emission in \\ls\\ as well, assuming a power law injection spectrum for the pairs in the unshocked pulsar wind and pair cascading. In this paper, we investigate the anisotropic inverse Compton scattering of stellar photons on the unshocked pulsar wind within the compact PWN scenario for \\ls\\ and \\lsi. Because of their tight orbits, the photon density is higher than in the Crab pulsar and \\psrb. A more intense gamma-ray signal from the unshocked pulsar wind is expected. The main purpose of this work is to constrain the bulk Lorentz factor $\\gamma_0$ of the pairs and the shock geometry. The next section presents the method and the main equations used in order to compute spectra in gamma-ray binaries. Section 3 describes and shows the expected spectra for \\ls\\ and \\lsi\\ with different parameters. Section 4 discusses the spectral signature from the unshocked pulsar wind. ",
        "conclusions": "Gamma-ray binaries are of particular interest in the study of pulsar wind nebula at very small scales. The massive star radiates a large amount of soft seeds photons for inverse Compton scattering on relativistic electrons from the pulsar. One expects two contributions in the gamma-ray spectral energy distribution: one from the shocked pulsar wind and another from the unshocked pulsar wind. The spectral signature from the unshocked region is strong and depends on the shock geometry and the initial energy of the pairs in the pulsar wind. A significant fraction of the pulsar wind power can be lost to inverse Compton scattering before a shock forms with the stellar wind \\citep{2005MNRAS.356..711S}. The shock location will thus be slightly closer in to the neutron star than calculated without taking into account the losses in the pulsar wind, assuming the wind is composed only of e$^+$e$^-$ pairs.  Significant emission from the free pulsar wind seems unavoidable. Inverse Compton losses in the free wind may be reduced if the shock occurs very close to the neutron star. This is unlikely as it would require a very strong stellar wind or a pulsar wind power that would be too low to produce the observed emission. If the pulsar wind is anisotropic then the orientation of the pulsar with respect to the observer can make the unshocked emission less conspicuous. This comes at the price of a peculiar orientation. If the pulsar wind is monoenergetic, then the line-like expected spectrum exceeds the observed very high energy power-laws for all geometries unless the pair energy is around 10 GeV or above 10 TeV. This pushes the direct emission from the wind in ranges where it may be constrained by future GLAST, HESS-2 or CTA measurements. The absence of any line emission will show that the assumption of a Crab-like monoenergetic, low $\\sigma$ pulsar wind was simplistic. If the pairs in the pulsar wind have a power-law distribution, then the unshocked emission is essentially indistinguishable from the shocked emission \\citep{2008arXiv0801.3427S}. A promising alternative is the {\\em striped wind} model in which the wind remains highly magnetised up to the termination shock, where the alternating field could be dissipated and accelerate particles (see \\citealt{2007astro.ph..3116K} and references therein). Future theoretical studies on the generation of pulsar relativistic winds in the context of a strong source of photons may be able to yield the particle distribution to expect and lead to more accurate predictions."
    },
    "0807/0807.4259_arXiv.txt": {
        "abstract": "Stationary solutions to the problem of particle acceleration at shock waves in the non-linear regime, when the dynamical reaction of the accelerated particles on the shock cannot be neglected, are known to show a prominent energy flux escaping from the shock towards upstream infinity. On physical grounds, the escape of particles from the upstream region of a shock has to be expected in all those situations in which the maximum momentum of accelerated particles, $p_{max}$, decreases with time, as is the case for the Sedov-Taylor phase of expansion of a shell Supernova Remnant, when both the shock velocity and the cosmic ray induced magnetization decrease. In this situation, at each time $t$, particles with momenta larger than $p_{max}(t)$ leave the system from upstream, carrying away a large fraction of the energy if the shock is strongly modified by the presence of cosmic rays. This phenomenon is of crucial importance for explaining the cosmic ray spectrum detected at Earth. In this paper we discuss how this escape flux appears in the different approaches to non-linear diffusive shock acceleration, and especially in the quasi-stationary semi-analytical kinetic ones. We apply our calculations to the Sedov-Taylor phase of a typical supernova remnant, including in a self-consistent way particle acceleration, magnetic field amplification and the dynamical reaction on the shock structure of both particles and fields. Within this framework we calculate the temporal evolution of the maximum energy reached by the accelerated particles and of the escape flux towards upstream infinity. The latter quantity is directly related to the cosmic ray spectrum detected at Earth. ",
        "introduction": "Supernova Remnants (SNRs) have long been suspected to be the main sources of galactic cosmic rays with energies up to the so called ``knee'' ($E\\sim 3 \\times 10^{15}$ eV). The existence of a flux of energetic particles escaping the accelerator from the upstream region is essential if this paradigm is indeed realized in Nature. The commonly accepted view of a typical SNR evolution is as follows. During the initial stage after the supernova explosion, the ejecta are in free-expansion with a high velocity $\\sim 10^4$ km/s. Non-linear acceleration is expected to be at work: efficient particle acceleration is associated with shock modification and magnetic field amplification in a complex chain of causality. During this stage, the maximum momentum achieved by the accelerated particles increases with time up to $p_{max}\\sim 10^{16}$ eV/c (see e.g. \\cite{bac07} and references therein). When the mass of the swept-up material becomes comparable with the mass of the ejecta, the Sedov-Taylor phase begins: the shock slows down and magnetic field damping is expected to become more effective (\\cite{ptuzira05}). As a consequence, the maximum momentum of particles that can return to the shock is expected to decrease. At each time $t$, particles with momentum exceeding the current $p_{max}(t)$ do not make it back to the shock and leave the system from upstream: this is the so called ``escape flux''.  The escape flux plays an essential role in the formation of the cosmic ray spectrum detected at Earth. If there were no escape from upstream during the Sedov-Taylor phase, all particles accelerated in a SNR would be advected downstream and undergo severe adiabatic energy losses before being injected into the ISM: in this case SNRs could not be responsible for accelerating cosmic rays with energies around the knee (\\cite{crspec}, \\emph{in preparation}, and \\cite{ptuzira05}).  While being of crucial importance during the Sedov-Taylor phase, it is clear that escape cannot occur during the free expansion phase of a SNR: the only particles that can leave the system from upstream are the ones at the highest energies, but, since $p_{max}$ is now increasing with time, no appreciable energy flux can escape the accelerator during this phase.  Kinetic approaches to non-linear particle acceleration (\\cite{malkov97,maldrury,blasi1,blasi2,ab05}) allow us to calculate the spectrum and the spatial distribution (including the absolute normalization) of the particles accelerated at the shock, even in the case when the diffusion coefficient is the result of magnetic field amplification by streaming instability induced by the accelerated particles themselves (\\cite{ab06}). These approaches, as well as others that have appeared in the literature (for instance \\cite{simple}), are based on the assumption that the acceleration process (and the shock modification) may be assumed to reach some sort of stationarity. In all these cases the calculations show that the shock is strongly modified by the presence of cosmic rays, and that the spectra are concave, with a slope at momenta close to the maximum momentum $p_{max}$ which is flatter than $p^{-4}$. All these calculations, independently of the techniques used to solve the equations, predict an escape flux of particles (and energy) towards upstream infinity: the shock becomes radiative, which is one of the very reasons why the shock modification becomes effective (namely the total compression factor becomes larger than 4).  It is worth recalling that the assumption of stationarity was widely used also in the framework of two-fluid models (\\cite{drury81p,drury81}), and its relevance in the case of modified shocks, leading to particle spectra harder than $p^{-4}$ at high energies, was discussed by \\cite{mv96} in the context of the so-called renormalized two-fluid models. It should be stressed, indeed, that the escape flux can arise both in the linear and non-linear regime of particle acceleration. The important difference between the two regimes, however, is that within the former the escaping particles carry away a negligible energy flux, while for modified shocks the bulk of the energy in the form of cosmic rays is stored in the highest energy particles and leaves the system with them.  The main weak point of the assumption of quasi-stationarity is that it leads to artificial escape fluxes in situations, such as the free expansion phase of a SNR, where escape would be unphysical (at least in the idealized cases of plane and perfectly spherical shocks). In this case, the appearance of an escape flux signals for the need for a fully time dependent approach. In other words, what appears as an escape flux is the energy which is channeled into particles with ever increasing momentum. During the Sedov-Taylor phase, as emphasized earlier, escape is in fact possible if the magnetic field in the shock region decreases. In this case, the leakage of particles towards upstream allows the system to relax to a quasi-stationary situation.  A number of time-dependent studies of non-linear shock dynamics have been presented in the literature (e.g. \\cite{fg87}), in particular for SNRs (e.g. \\cite{bv97,kj06} and references therein). All these calculations aim at solving the transport equations and the equation of fluid dynamics numerically. Most of these papers have however concentrated on the spectrum of accelerated particles rather than on the escape flux. In fact, in many of these papers the magnetic field is introduced by hand and is not a function of time, thereby not leading to a decrease with time of the maximum momentum of accelerated particles. In other cases the boundary conditions are imposed at upstream infinity, therefore the escape flux, even if it is present, would appear in the form of concentration of high energy particles at large distances from the shock. In these cases a recipe is needed to label these particles as {\\it escaping} the system (see for instance \\cite{kamae}).  The identification of the escape flux is all but a mathematical detail: as discussed by \\cite{ptuzira05}, the escape flux is crucial for establishing a connection between SNRs and the origin of cosmic rays, in that it allows the escaping particles to avoid the adiabatic losses related to the expansion of the supernova shell, and therefore become cosmic rays at the knee. In the absence of escape towards upstream infinity at the beginning of the Sedov-Taylor phase, SNRs cannot be the sources of cosmic rays up to the knee.  The contribution of the escape flux to the galactic cosmic ray spectrum was first estimated by \\cite{ptuzira05}. These authors included the magnetic field damping and the consequent decrease of $p_{max}$ during the Sedov-Taylor expansion, but in their study the shock structure and particle spectrum are fixed rather than self-consistently calculated.  In this paper we present the first attempt at carrying out a calculation of the escape flux within the framework of a kinetic approach to non-linear shock acceleration, self-consistently including particle acceleration, magnetic field amplification and the dynamical reaction of both on the shock. We focus our attention on shocks in shell type SNRs, discussing the physical meaning of the escape flux during the different phases of their evolution. In particular the flux of energetic particles leaving the remnant during the Sedov-Taylor phase is explicitely computed and its phenomenological implications for the origin of cosmic rays are discussed.  The paper is organized as follows: in \\S \\ref{sec:stat} we discuss the implications of the assumption of stationarity of the acceleration process. In \\S \\ref{sec:cons} we discuss the escape flux based on the most general version of the conservation equations. In \\S \\ref{sec:escape} we discuss how the escape flux is connected to the existence of a maximum momentum in the distribution of accelerated particles. In \\S \\ref{sec:snr} we apply our calculations to the different stages of evolution of a supernova remnant. We conclude in \\S \\ref{sec:concl}. ",
        "conclusions": "\\label{sec:concl}  The escape of accelerated particles from the shock region towards upstream appears in basically all approaches to non-linear particle acceleration at shock fronts. In kinetic approaches that solve the relevant equations assuming quasi-stationarity at any given time the escape flux appears either as a corrective term in the equation of energy conservation between downstream and upstream infinity, or as a result of imposing as a boundary condition that the distribution function of accelerated particles vanishes at a finite distance $x_0$ upstream of the shock. In this second case the escape flux is described by the non-vanishing spatial derivative of the distribution function at $x_0$. Such a quasi-stationary description of the acceleration process is meaningful only whenever the shock velocity decreases in time and the magnetization at the shock, as due to streaming of cosmic rays, also decreases in time. These two conditions lead to escape of the particles with momentum close to the highest achievable at any given time. This would be the case during the Sedov-Taylor phase of the expansion of a supernova shock. These approaches are therefore unfit to describe the phase of free expansion, where the maximum momentum increases with time and no escape towards upstream takes place.  Fully time-dependent, numerical calculations apply to arbitrary situations and should be able to describe the increase of momentum during the free expansion as well as the escape of particles during the Sedov phase. However, in most cases these calculations do not include the time dependence of the magnetic field at the shock. This leads to an increases of the maximum momentum even during the Sedov phase. Probably for this reason there is no detailed discussion in the literature of the escape flux from a SNR in the context of time dependent calculations. The appearance of the escape flux in numerical time dependent approaches would depend on the boundary conditions adopted in the specific case. If the distribution function is assumed to vanish at upstream infinity, the escaping particles should appear as a concentration of particles with large momenta (close to $p_{max}(t)$) at large distances from the shock: since the acceleration box extends to infinity, there is no real escape. On the other hand, if the boundary condition $f(x_0)=0$ is imposed at a finite distance from the shock, the escape flux (as a function of time) should appear automatically as a consequence of a non-vanishing $D\\partial f(x_0)/\\partial x$. It would be interesting to have confirmation of these expectations by running time dependent calculations with decreasing magnetization.  The existence of the escape flux is not simply a mathematical nuisance: as discussed by \\cite{ptuzira05} it is the very reason why supernova remnants can potentially be the sources of cosmic rays up to the knee. If the particles we observe as cosmic rays were the ones advected towards downstream of the shock, the adiabatic losses suffered during the expansion of the shell would drive these particles towards lower energies. It is therefore of the highest importance to go beyond the test particle calculation of \\cite{ptuzira05} and achieve an understanding of the escape flux in the context of the fully non-linear theory of diffusive shock acceleration.  In this paper we made a first step in this direction, by calculating the temporal evolution of the escape flux during the Sedov-Taylor phase of the expansion of a supernova in a medium with constant density. Our calculation included a recipe to describe the cosmic ray induced magnetic field, the dynamical reaction of the accelerated particles and the field itself. In addition to this we also discussed in some detail the limitations of the conservation equations as they are usually written, in that they do not allow to describe the magnetic field perturbations when they are not Alfv\\'en waves.  We showed that the escape flux may involve between few and 10-30 \\% of the shock ram pressure, while the particle acceleration efficiency at the same time reaches 50-60 \\%. This means that the energetics inferred from observations of secondary radiation from a remnant, for instance in the form of gamma rays, may not be the relevant one for the origin of cosmic rays, since it is not related in a trivial way to the energy flux escaping the accelerator.  The maximum energy up to which particles may get accelerated is reached at the beginning of the Sedov phase and is of order $10^{15}$ eV if only resonant amplification of the field is included. $E_{max}$ might be somewhat larger for some SNRs that at the very beginning of the Sedov-Taylor phase may experience the effect of non-resonant streaming instability \\cite[]{bell04}. This mechanism provides extremely efficient field amplification as long as the shock velocity is high, and hence is expected to play a very important role in the early Sedov-Taylor phase (and possibly during the free expansion phase). At present, non-resonant modes cannot be formally accounted for in the conservation equations, since a detailed description of the energy transfer between particles and waves is not available yet. We carried out all calculations in the simpler case of Alfv\\'en waves interacting with particles in a resonant way.  The amplification of the magnetic field, by either resonant or non-resonant streaming instability, has profound implications on the escape flux of particles towards upstream of a shock, and therefore on the spectrum of cosmic rays we observe at the Earth. The most obvious consequence of the magnetic field amplification is that of allowing for higher values of the maximum energy of accelerated particles, as shown by our Eqs.~\\ref{eq:emax1t} and \\ref{eq:emax2t}. However large magnetic fields exert a dynamical reaction on the plasma leading to a reduction of the compression in the precursor. This happens whenever the magnetic pressure exceeds the pressure of the background gas \\cite[]{apjlett}. As a result, the concavity of the spectrum of accelerated particles \\cite[]{long} is reduced and at the same time the escape flux at $p\\sim p_{max}$ decreases. It follows that larger field implies larger $p_{max}$ but not necessarily larger escape flux, as shown in Figs.~\\ref{fig:B1n003}, \\ref{fig:B1n01} and \\ref{fig:B5n01} (see the behaviour of the curves at early times).  The escape of accelerated particles from a cosmic ray modified shock has profound implications for the origin of cosmic rays, which will be discussed in detail in a forthcoming paper. Here we want to emphasize some general points:  1) the escape from upstream is the natural solution to the well known problem of explaining how the highest energy particles (say, close to the knee energy) could escape the system without suffering substantial adiabatic energy losses;  2) the magnetic field amplification is expected to switch from mainly non-resonant to mainly resonant at the beginning of the Sedov-Taylor phase. It can be easily understood that this may lead to peculiar changes in the spectrum of cosmic rays detected at the Earth, reflecting this transition;  3) the flux of escaping particles, once integrated in time during the SNR evolution may be very different from the concave instantaneous spectrum which can potentially be observed in a SNR, for instance by looking at its multifrequency emission. This point is certainly relevant for the purpose of addressing the commonly raised point of how a concave spectrum  of accelerated particles can reflect in an almost perfect power law over many orders of magnitude;  4) there is a further complication of all the picture, due to the acceleration of nuclei at energies that may be expected to scale as the charge of the nucleus (in the case of Bohm diffusion). Any calculation of the flux of single chemical species observed at the Earth must take these complex effects into account."
    },
    "0807/0807.3686_arXiv.txt": {
        "abstract": "{} {Determination and analysis of electron densities and filling factors of the diffuse ionized gas (DIG) in the solar neighbourhood.} {We have combined dispersion measures and emission measures towards 38 pulsars at distances known to better than 50\\%, from which we derived the mean density in clouds, $\\Nec$, and their volume filling factor, $\\Filfac$, averaged along the line of sight. The emission measures were corrected for absorption by dust and contributions from beyond the pulsar distance. } {The scale height of the electron layer for our sample is $0.93\\pm 0.13\\kpc$ and the midplane electron density is $0.023\\pm 0.004\\cmcube$, in agreement with earlier results. The average density along the line of sight is $\\avnel = 0.018\\pm 0.002\\cmcube$ and nearly constant. Since $\\avnel = \\Filfac \\Nec$, an inverse relationship between $\\Filfac$ and $\\Nec$ is expected. We find $\\Filfac (\\Nec) = (0.011\\pm 0.003)\\ \\Nec^{-1.20\\pm 0.13}$, which holds for the ranges $\\Nec = 0.05-1\\cmcube$ and $\\Filfac = 0.4-0.01$. Near the Galactic plane the dependence of $\\Filfac$ on $\\Nec$ is significantly stronger than away from the plane. $\\Filfac$ does not systematically change along or perpendicular to the Galactic plane, but the spread about the mean value of $0.08\\pm 0.02$ is considerable. The total pathlength through the ionized regions increases linearly to about $80\\pc$ towards $|z| = 1\\kpc$. } % {Our study of $\\Filfac$ and $\\Nec$ of the DIG is the first one based on a sample of pulsars with known distances. We confirm the existence of a tight, nearly inverse correlation between $\\Filfac$ and $\\Nec$ in the DIG. The exact form of this relation depends  on the regions in the Galaxy probed by the pulsar sample. The inverse $\\Filfac$--$\\Nec$ relation is consistent with a hierarchical, fractal density distribution in the DIG caused by turbulence. The observed near constancy of $\\avnel$ then is a signature of fractal structure in the ionized medium, which is most pronounced outside the thin disk. } % ",
        "introduction": "The full-sky H$\\alpha$ maps of the Milky Way that have recently become available (Dickinson et al.\\ \\cite{dickinson+03}; Finkbeiner\\ \\cite{finkbeiner03}; Bennett et al.\\ \\cite{bennett+03}; Hinshaw et al.\\ \\cite{hinshaw+07}) show a spectacular variation in intensity and structure. The classical \\ion{H}{II} regions in the disc are surrounded by extended, diffuse ionized gas (DIG) visible up to high latitudes, and supernova explosions, shocks and turbulence in the interstellar medium (ISM) have produced a wealth of cloud shapes, filaments, shells and voids.  Madsen et al. (\\cite{madsen+06}) found the variations in structure to be accompanied by variations in electron temperature. They derived a mean temprature of the DIG of about $8000\\,$K with strong differences between lines of sight.  A number of authors derived a scale height of the DIG in the solar neighbourhood of about $1\\kpc$ from the increase of the dispersion measures of pulsars with increasing distance to the Galactic plane (Reynolds\\ \\cite{reynolds91b}; Bhattacharia \\& Verbunt\\ \\cite{bhatta+91}; Nordgren et al.\\ \\cite{nordgren+92}; Gomez et al.\\ \\cite{gomez+01}; Cordes \\& Lazio\\ \\cite{cordes+lazio03}). Using dispersion measures towards 4 pulsars in globular clusters, Reynolds (\\cite{reynolds91a}) obtained a filling fraction of the DIG of $\\la 0.2$ through the full layer and a mean density in clouds of about $0.08\\cmcube$. This was recently confirmed by Berkhuijsen et al. (\\cite{elly+06}) from a much larger pulsar sample. These authors also found an increase of the filling factor with distance from the plane. Already the scanty data that were available to Kulkarni \\& Heiles (\\cite{kulkarni+88}) had indicated such an increase.  Turbulence causing hierarchical, fractal structure leads to an inverse relationship between the mean filling factor along the line of sight, $\\Filfac$, and the mean density in ionized regions, $\\Nec$ (Fleck\\ \\cite{fleck96}; Elmegreen\\ \\cite{elmegreen98}, \\cite{elmegreen99}). Pynzar (\\cite{pynzar93}), who was the first to investigate this relationship, derived $\\Filfac \\propto \\Nec^{-0.7}$ for a combination of DIG and \\ion{H}{II} regions, which was confirmed by Berkhuijsen (\\cite{elly98}). In an extensive study, Berkhuijsen et al. (\\cite{elly+06}, hereafter called BMM) obtained $\\Filfac \\propto \\Nec^{-1.07\\pm 0.03}$ for the DIG from a sample of 157 pulsars at latitudes $|b|> 5\\degr$. They used dispersion measures from the catalogue of Hobbs \\& Manchester (\\cite{hobbs+03}; see also Manchester et al.\\ \\cite{manchester+05}), emission measures from the WHAM survey (Haffner et al.\\ \\cite{haffner+03}), corrected for absorption (Diplas \\& Savage\\ \\cite{diplas+savage94}) as well as for contributions from beyond the pulsars, and pulsar distances from the model of the electron distribution in the Galaxy of Cordes \\& Lazio (\\cite{cordes+lazio02}). The statistical error in pulsar distances $<3\\kpc$, where 75\\% of the pulsars in the sample are located, is about 25\\% which is much smaller than the intrinsic spread in the dispersion measures and emission measures. Therefore, random errors in the model distances will not have influenced the statistical results. However, the distances from the model may also contain a systematic error. Based on these distances, the radial distribution of pulsars has a maximum at $R=3.5\\kpc$ from the Galactic centre (Lorimer\\ \\cite{lorimer04}; Yusifov \\& K\\\"u\\c{c}\\\"uk\\ \\cite{yusifov+04}), whereas other population~I objects (H2, \\ion{H}{I}, \\ion{H}{II} regions) peak near $R=5\\kpc$. The smoothness of the electron density model may lead to pulsar distances that are too large.  So it is important to check the results of BMM for pulsars with distances derived from observations. Since BMM did their work, the sample of pulsars with measured distances has increased considerably. Our new analysis is based on 38 pulsars, the distances of which are known to better than 50\\%. A further decrease of the constraint on the distance errror would make the sample too small for a statistical analysis. We also used an improved absorption correction to the emission measure of each pulsar instead of the statistical correction applied by BMM. ",
        "conclusions": "\\subsection{Comparison with BMM 2006}  It is interesting to compare our results with those of BMM who took the distances to the pulsars from the NE2001 model of Cordes \\& Lazio (\\cite{cordes+lazio02}) and used an average correction for the absorption of the H$\\alpha$ emission that only depends on latitude. Applying their method to our sample, we recalculated the main relationships using a scale height of $0.47\\kpc$ to correct for H$\\alpha$ emission from behind the pulsar. Table~\\ref{tab:4} shows that the resulting relations are nearly identical to those derived in Sect.~4, but with larger errors. This means that the less accurate distances and absorption corrections used by BMM have \\emph{not} influenced their statistical results. Therefore, we may compare results from our small sample with those from the much larger sample analyzed by BMM with confidence.  \\begin{table*} \\caption{Comparison with BMM for our sample of 34 pulsars} \\label{tab:4} \\centering \\begin{tabular}{l@{ $=$ }l r@{$\\pm$}l r@{$\\pm$}l r@{$\\pm$}l r@{$\\pm$}l} \\hline\\hline \\multicolumn{2}{c}{Relation}  &\\multicolumn{4}{c}{Our data} &\\multicolumn{4}{c}{BMM} \\\\ \\omit&\\omit &\\multicolumn{2}{c}{$a$}  &\\multicolumn{2}{c}{$b$} &\\multicolumn{2}{c}{$a$} &\\multicolumn{2}{c}{$b$} \\\\ \\hline \\noalign{\\smallskip} $\\langle\\nel^2\\rangle$ &$a\\, e^{-|z|/b}$  &0.0045&0.0007  &\\multicolumn{2}{c}{$1.4^{+0.6}_{-0.3}$} &0.0053&0.0011 &\\multicolumn{2}{c}{$1.4^{+1.5}_{-0.5}$} \\\\ \\noalign{\\smallskip} $\\Nec (z)$ &$a\\, e^{-|z|/b}$  &0.24&0.04  &\\multicolumn{2}{c}{$2.6^{+4.1}_{1.0}$} &0.26&0.06 &\\multicolumn{2}{c}{$2^{+10}_{-1}$} \\\\ \\noalign{\\smallskip} $\\Filfac (z)$ &$a\\, e^{|z|/b}$  &0.077&0.018  &\\multicolumn{2}{c}{$[23]^1$}  &0.077&0.023 &\\multicolumn{2}{c}{$[5]^1$} \\\\ \\noalign{\\smallskip} $\\Filfac (\\Nec)$   &$a\\, \\Nec^b$  &0.011&0.003 &$-$1.20&0.13  &0.014&0.003 &$-$1.18&0.11 \\\\ \\hline \\noalign{\\smallskip} \\multicolumn{10}{l}{1) Undetermined: $1/b$ smaller than errors} \\\\ \\end{tabular} \\end{table*}   A striking difference between the two samples is the behaviour of $\\Filfac(z)$. Where BMM found a significant increase of the filling factor to about 0.2 towards $|z| = 1\\kpc$, $\\Filfac(z)$ derived from our sample remains constant within the errors at a value of about 0.08. This difference must be due to the different locations in the Galaxy of the pulsars in the two samples. For example, 17 out of the 34 pulsars (50\\%) in our sample are within $1\\kpc$ from the Sun, whereas only 44 out of 157 pulsars (28\\%) of the BMM sample are within this distance. Since the ISM is a highly variable medium, differences between samples are to be expected. Hydromagnetic simulations of the evolution of the ISM by de Avillez \\& Breitschwerdt (\\cite{avillez+breit05}) may illustrate this point. Their Fig.~2a shows the density variation in a cut of $1\\kpc$ along and $10\\kpc$ perpendicular to the Galactic plane. Although the filling fraction of the DIG generally increases with height, there are also regions where it remains constant or even decreases away from the plane.   \\subsection{Extent of ionized regions}  The total pathlength through ionized regions perpendicular to the Galactic plane increases nearly linearly from about $7\\pc$ towards $|z| = 100\\pc$ to about $40\\pc$ towards $|z| = 500\\pc$ and about $80\\pc$ towards $|z|= 1\\kpc$ (Fig.~\\ref{fig:5}).  In the ISM dense regions are usually smaller than less dense regions, which also applies to the DIG (see Fig.~\\ref{fig:8}). Although the distribution of $L_\\mathrm{e} \\sin|b|$ with $\\Nec$ shows considerable spread, it is well fitted by a power law (see Table~\\ref{tab:2}). Regions of density $\\Nec = 1\\cmcube$  together occupy about $2\\pc$ in the $|z|$-direction and those of $\\Nec = 0.05\\cmcube$ about $300\\pc$.  \\begin{figure} % \\includegraphics[bb = 103 72 559 663,angle=270,width=8.8cm]{09675f8.eps} \\caption{Dependence of the total pathlength through ionized regions perpendicular to the Galactic plane, $L_\\mathrm{e} \\sin|b|$, on the mean density in clouds, $\\Nec$. Circles: pulsars with $\\Delta D/D < 0.2$; triangles: pulsars with $0.2 < \\Delta D/D < 0.5$. Full line: powerlaw fit given in Table~\\ref{tab:2}; dashed line: powerlaw fit to the sample of 157 pulsars of BMM, $L_\\mathrm{e} \\sin|b|(\\Nec) = (5.5\\pm 0.8) \\Nec^{-1.32\\pm 0.07}$. } \\label{fig:8} \\end{figure}  We calculated the same relationship for the BMM sample which gives a somewhat larger extent for the higher densities (see Fig.~\\ref{fig:8}). This power law is close to the one indicated by the distribution of 192 pulsars presented by Hill et al. (\\cite{hill+07}). However, these authors did not apply any correction to the observed emission measures, which leads to a flattening of the distribution and a shift to higher values of $\\Nec$ for the same pathlength (see Fig.~5 in BMM).   \\subsection{The $\\Filfac$--$\\Nec$ relation and density structure}  In spite of our small sample of pulsars, the nearly inverse correlation between $\\Filfac$ and $\\Nec$ shown in Fig.~\\ref{fig:6} is very tight, reaching a correlation coefficient of $0.88\\pm 0.09$ (see Table~\\ref{tab:2}). It does not change with distance along the plane and, as BMM explained, is insensitive to errors in emission measure. The near constancy of $\\avnel = \\Filfac \\Nec$ along the line of sight (see $D- \\avnel$ fit in Table~\\ref{tab:2}), first noted by Weisberg et al. (\\cite{weisberg+80}), was an early indication of an inverse relationship between $\\Filfac$ and $\\Nec$.  \\begin{figure} % \\includegraphics[bb = 98 56 554 669,angle=270,width=8.8cm]{09675f9.eps} \\caption{Dependence of the volume filling factor, $\\Filfac$, on mean density in clouds, $\\Nec$, as derived by different authors. Full line: this work; dashed line: dependence for the DIG within about $3\\kpc$ from the Sun (BMM); dotted line: dependence for the DIG in the bright emission ring in M31 (Berkhuijsen\\ \\cite{elly04}); dash-dot line: dependence for clouds of diffuse dust within $400\\pc$ from the Sun (Gaustad \\& Van Buren\\ \\cite{gaustad+93}). The parameters of the lines are given in Table~\\ref{tab:5}. } \\label{fig:9} \\end{figure}   We first compare our $\\Filfac$--$\\Nec$ relation with earlier determinations available in the literature, and then discuss what it tells us about the density structure in the DIG.  \\subsubsection{Comparison with other data}  Figure~\\ref{fig:9} and Table~\\ref{tab:5} show our $\\Filfac$--$\\Nec$ relation together with earlier determinations. For the same mean density in clouds, we find somewhat smaller filling factors than BMM which agree with the lower points in their distribution. The smaller errors in the BMM relation are not only due to the larger sample, but also to the smooth distribution of $\\avnel = \\DM /D$ of the electron density model of Cordes \\& Lazio (\\cite{cordes+lazio02}) used by BMM. Berkhuijsen (\\cite{elly04}) obtained the $\\Filfac$--$\\Nec$ relation for the DIG in the galaxy M\\,31 from a comparison of rotation measures of the polarized continuum emission (giving $\\avnel$ with magnetic field strengths of Fletcher et al.\\ \\cite{fletcher+04}) and thermal radio emission (giving $\\langle\\nel^2\\rangle$). These data refer to the DIG in the bright emission ring in M\\,31. The good agreement between the M\\,31 and the MW results suggests that the inverse correlation between $\\Filfac$ and $\\Nec$ is generally valid in the DIG in galaxies.  \\begin{table*} \\caption{Comparison of $\\Filfac -\\Nec$ relations $\\Filfac (\\Nec) = a\\, \\Nec^b$} \\label{tab:5} \\centering \\begin{tabular}{l r@{$\\pm$}l r@{$\\pm$}l c c c} \\hline\\hline Sample  &\\multicolumn{2}{c}{$a$} &\\multicolumn{2}{c}{$b$} &Range $\\Nec$  &Range $\\Filfac$ &Ref. \\\\ &\\omit&\\omit &\\omit&\\omit &$[\\cmcube]$ \\\\ \\hline DIG \\\\ MW\\quad N= 34 &0.0114&0.0025  &$-$1.20&0.13  &$0.05-1$  &$0.4\\pheins -0.011$ &1 \\\\ MW\\quad N=157 &0.0184&0.0011  &$-$1.07&0.03  &$0.03-2$  &$0.8\\pheins -0.009$ &2 \\\\ M31           &0.016 &0.004   &$-$1.24&0.30  &$0.19-1$  &$0.13-0.016$        &3 \\\\ \\noalign{\\smallskip} DIFFUSE DUST \\\\ $D < 400\\pc$  &0.056 &0.020   &$-$1.25&0.40  &$1-10$    &$0.06-0.003$        &4 \\\\ \\hline \\noalign{\\smallskip} \\multicolumn{8}{l}{References. (1) This work ; (2) Berkhuijsen, Mitra, \\& M\\\"uller\\ \\cite{elly+06};} \\\\ \\multicolumn{8}{l}{(3) Berkhuijsen\\ \\cite{elly04}; (4) Gaustad \\& Van Buren\\ \\cite{gaustad+93}} \\\\ \\end{tabular} \\end{table*}   Roshi \\& Anantharamaiah (\\cite{roshi+01}) observed radio recombination lines from the inner Galaxy at low latitudes. They derived filling factors of $\\le 0.01$ for extended regions of diffuse ionized gas with densities of $1-10\\cmcube$, which agree well with the extension of the curves for the DIG in Fig.~\\ref{fig:9}.  Thus the inverse correlation between $\\Filfac$ and $\\Nec$ at least holds for the density range $0.03-10\\cmcube$. Cordes et al. (\\cite{cordes+85}) estimated a filling factor of $10^{-4.0\\pm0.3}$ for clumps of about $1\\pc$ size causing scattering of pulsar signals. Our data predict this filling factor for clumps of density $\\Nec = 50\\cmcube$ and size $L_\\mathrm{e}/m = 0.9\\pc$, where $m$ is the number of clumps on the line of sight. The relation of BMM gives clumps of $\\Nec = 100\\cmcube$ for the same filling factor and sizes. So these clump properties also seem to follow the $\\Filfac$--$\\Nec$ relations derived for the DIG.  Furthermore, Gaustad \\& Van Buren (\\cite{gaustad+93}) obtained a similar relationship for clouds of diffuse dust within $400\\pc$ from the Sun, with mean densities of $1-10\\cmcube$. Figure~\\ref{fig:9} shows that these clouds have about 3 times higher filling factors than ionized clouds of the same mean density. Interestingly, their data agree very well with those of Pynzar (\\cite{pynzar93}, also shown in Berkhuijsen\\ \\cite{elly98}) derived from recombination lines.  \\subsubsection{Density structure of the DIG } The general validity of the inverse $\\Filfac$--$\\Nec$ relation indicates that it describes a basic property of the DIG, possibly even of the entire ISM (Berkhuijsen\\ \\cite{elly99}). At least two mechanisms could be responsible: thermal pressure equilibrium and turbulence causing a fractal density structure.  Thermal pressure equilibrium leads to an inverse $\\Filfac$--$\\Nec$ relation if it were widespread, but it seems only locally valid. In the MW large fractions of gas are observed in unstable regimes and have also been found in simulations of a turbulent ISM. Thermal pressure equilibrium of clouds appears to be of minor importance in the presence of turbulence (see Elmegreen \\& Scalo\\ \\cite{elmegreen+04}, and references therein).  Elmegreen (\\cite{elmegreen98}, \\cite{elmegreen99}) discussed the properties of diffuse ionized gas in a pervasive fractal ISM. One of these properties is an inverse correlation between volume filling factor and gas density, which is expected on theoretical grounds (Fleck\\ \\cite{fleck96}) as well as from simulations (Elmegreen\\ \\cite{elmegreen97}; Kowal \\& Lazarian\\ \\cite{kowal+07}). As this leads to a constant average density along the line of sight, the observed near constancy of $\\avnel$ with $\\Nec$ and $\\Filfac$ inversely varying over 1.5 decade may indicate that the density distribution of the DIG is fractal. Please note that the near constancy of $\\avnel$ alone  should only be regarded as a first indication, because the fits in Table~\\ref{tab:2} show that the statistical significance of the near constancy of $\\avnel$ is much lower than that of the nearly inverse $\\Filfac$--$\\Nec$ relation.  The fractal medium is characterized by filamentary, clumpy structures that tend to cluster together, with holes inside and large voids between them. The filling factor along the line of sight is determined by the outer regions of the filaments, i.e. by the largest scales, because the small, dense clumps in the inner parts hardly contribute. The voids occupy more space than the filaments, thus the filling factor of the filaments is small. The values of $\\Filfac \\simeq 0.1$ (Sect.~4.1) and $\\Filfac \\simeq 0.2$ (BMM) obtained when looking perpendicular to the Galactic plane towards $|z|=1\\kpc$ are consistent with this picture.  The tendency of filaments to cluster into larger complexes reduces the number of fractal clouds, $N_\\mathrm{f}$, along the line of sight. Elmegreen (\\cite{elmegreen98}) estimated $N_\\mathrm{f}=3/{\\rm kpc}$ locally. For $\\Filfac=0.1$ the total pathlength through the ionized regions is 100~pc/kpc, giving a line of sight through one fractal complex of about $30\\pc$ and a mean density $\\Nec$ of $0.15\\cmcube$ (from Fig.~\\ref{fig:6}). Using the results of BMM gives $\\Filfac=0.2$, a fractal complex of about $70\\pc$ along the line of sight and $\\Nec = 0.10\\cmcube$. Interestingly, these pathlengths are in the range of 10--100~pc that Ohno \\& Shibata (\\cite{ohno+93}) estimated for the cells of ionized gas causing $\\DM$ and rotation measure $\\RM$ towards pairs of pulsars close on the sky.  According to Elmegreen (\\cite{elmegreen98}), pervasive fractal structure only develops in regions outside the influence of SN shells, sites of star formation, chimneys etc. that are concentrated near the Galactic plane. This may explain our result of Sect.~4.2 that the exponent of the $\\Filfac$--$\\Nec$ relation is significantly smaller in the thin disc ($|z|<300\\pc$) than further from the plane where it is close to $-1$ (see Table~\\ref{tab:3} and Fig.~\\ref{fig:7}).  We conclude that our results on the mean electron densities and volume filling factors of the DIG outside the thin disc are consistent with a fractal ionized medium caused by turbulence, while the structure in the thin disc is dominated by the effects of SN shocks, stellar winds, star formation and other forces. The difference in density structure between these two regimes could be further analyzed when more pulsars with measured distance become available.  Recently, further observational evidence for turbulent structure of the DIG and of the diffuse atomic gas has been presented. Hill et al. (\\cite{hill+08}) found that the probability distribution function (PDF) of the emission measures observed perpendicular to the Galactic plane at $|b| > 10\\degr$ is lognormal as is expected for a turbulent medium from MHD simulations of the ISM (Elmegreen \\& Scalo\\ \\cite{elmegreen+04}). Furthermore, Berkhuijsen \\& Fletcher (\\cite{elly+08}) showed that the PDFs of the quantities $\\avnel$, $\\langle\\nel^2 \\rangle$, $\\Nec$ and $\\Filfac$ derived above as well as that of the mean \\ion{H}{I} density along the line of sight towards stars, $\\langle n_\\mathrm{HI}\\rangle$, are lognormal, consistent with a turbulent origin of density structure in the diffuse gas."
    },
    "0807/0807.3365_arXiv.txt": {
        "abstract": "Studies of the halo gas in the Milky Way and in nearby spiral galaxies show the presence of gas complexes that cannot be reconciled with an internal (galactic fountain) origin and are direct evidence of gas accretion. Estimating gas accretion rates from these features consistently gives values, which are one order of magnitude lower than what is needed to feed the star formation. I show that this problem can be overcome if most of the accretion is in fact \u201chidden\u201d as it mixes with the galactic fountain material coming from the disk. This model not only provides an explanation for the missing gas accretion but also reproduces the peculiar kinematics of the halo gas in particular the vertical rotation gradient. In this view this gradient becomes indirect evidence for gas accretion. ",
        "introduction": "A large amount of fresh gas accretion onto the Milky Way has been advocated for decades since the work of \\cite{twarog} showed the relative constancy of the Star Formation Rate (SFR) in the Galactic disk over the last 10 Gyrs. Roughly constant SFR requires constant accretion and this translates today into a need for gas accretion at a rate of the order $\\sim 1 \\moyr$.  Since the 1960s, the presence of a large number of neutral hydrogen (\\hi) clouds filling the sky around us and having, on average, negative velocities with respect to the Milky Way disk has been recognized (e.g.\\ \\cite[Hulsbosch 1968]{huls68}). Such clouds, named High Velocity Clouds (HVCs), given their largely anomalous radial velocities (\\cite[Wakker \\& van Woerden 1997]{wakker97}), were immediately regarded as possible evidence for gas accretion from intergalactic space onto our Galaxy (\\cite[Oort 1970]{oort70}). Alternative interpretations have been proposed, for example that the clouds are produced by the cooling of material expelled from the disk via a ``galactic fountain'' (\\cite[Bregman 1980]{bregman80}) or that they are a much more distant Local Group population (e.g.\\ \\cite[Blitz 1999]{blitz99}).  From the time of Oort's suggestion until very recently there were two major unknowns about the HVCs: their distances and their metallicities. It is thanks to a tenacious observational campaign carried out especially over the last few years that we now know both distances and metallicities for all the major HVCs (\\cite[Wakker et al.\\ 2007, 2008]{wakker07, wakker08}). The results leave no doubts: the HVCs are located in the Milky Way halo and have metallicities of about an order of magnitude lower than the average disk ISM metallicity. These properties make it most likely that this is accreting material falling onto the Milky Way for the first time.  The HVCs do not comprise the entire halo population, rather they are an extreme population at large heights and having particularly anomalous velocities. In the lower halo large amounts of cold gas are observed (e.g.\\ \\cite[Lockman 1984]{Lockman}), some of which go under the name of Intermediate Velocity Clouds (IVCs) and are regarded as a likely galactic fountain component (e.g.\\ \\cite[Wakker 2001]{Wakker01}). More recently, \\cite{kalberla08} showed that up to 10\\% of the Milky Way \\hi\\ gas is in fact extra-planar and highly turbulent. ",
        "conclusions": ""
    },
    "0807/0807.0293_arXiv.txt": {
        "abstract": "Massive stars deeply influence their surroundings by their luminosity and the injection of kinetic energy. So far, they have mostly been studied with spatially {\\it unresolved} observations, although evidence of geometrical complexity of their wind are numerous. Interferometry can provide spatially {\\it resolved} observations of massive stars and their immediate vicinity. Specific geometries (disks, jets, latitude-dependent winds) can be probed by this technique.  The first observation of a Wolf-Rayet (WR) star ($\\gamma^2$ Vel) with the AMBER/VLTI instrument yielded to a re-evaluation of its distance and an improved characterization of the stellar components, from a very limited data-set. This motivated our team to increase the number of WR targets observed with AMBER. We present here new preliminary results that encompass several spectral types, ranging from early WN to evolved dusty WC.  We present unpublished data on WR79a, a massive star probably at the boundary between the O and Wolf-Rayet type, evidencing some Wolf-Rayet broad emission lines from an optically thin wind. We also present new data obtained on $\\gamma^2$ Vel that can be compared to the up-to-date interferometry-based orbital parameters from North et al. (2007). We discuss the presence of a wind-wind collision zone in the system and provide preliminary analysis suggesting the presence of such a structure in the data. Then, we present data obtained on 2 dusty Wolf-Rayet stars: WR48a-b and WR118, the latter exhibiting some clues of a pinwheel-like structure from the visibility variations. ",
        "introduction": "\\label{section:intro} According to the classical MK spectral classification system, the A, B and O stars represent the group of hot or early type stars. The most permanent luminous stellar objects are the {\\bf O stars}. They exhibit strong, fast stellar winds of relatively low opacity, driven by their strong radiation field. Luminosities are found to be up to $10^6 L_{\\odot}$ (for O super-giants) with mass loss rates $\\dot{M}$ of up to 5 $\\times 10^{-6} M_{\\odot}$. Terminal wind velocities ($V_{\\infty})$ of up to 2000 km/s are found. Estimated surface temperatures are found to lie up to $\\sim$ 50\\,000 K. In an evolutionary context, O stars are the probable progenitors of Wolf-Rayet stars. Intense emission lines of various ionized elements are the spectral characteristics of Wolf-Rayet stars, originating from an extended, and rapidly expanding atmosphere. Among all stable stars, WR stars reveal the strongest mass loss via (radiative) wind mechanisms. Typical mass loss rates range from $(2-10)\\times 10^{-5} M_{\\odot}\\,yr^{-1}$, with wind velocities of 1000 - 2500 km/s. The Wolf-Rayet winds are mostly optically thick, so the surface is not visible. This means that the connection of spectral classification to photospheric core temperature is not possible as for normal stars via MK classification, and the WR classification is purely spectroscopic. WR stars are highly luminous, evolved, hot, massive stars, in the final state of their nuclear He burning. They create significant quantities of heavy elements and enrich the interstellar medium through mass-loss, and are also candidate as progenitor of Gamma-ray bursts.  It is of importance to probe the wind of these star, as the wind is the spectral reference for the classification and the modeling of these extreme stars. Long baseline interferometry can help constraining the spatial scale of the continuum and line forming regions, only weakly constrained for these objects by classical photometric and spectroscopic methods. % AMBER (Astronomical Multi BEam Recombiner) is the VLTI (Very Large Telescope Interferometer) beam combiner operating in the near-infrared\\cite{2007A&A...464....1P}. The instrument uses spatial filtering with fibers\\cite{2000SPIE.4006..299M}. The interferometric beam passes through anamorphic optics compressing the beam perpendicularly to the fringe coding in order to be injected into the slit of a spectrograph. The instrument can operate at spectral resolutions of 35, 1500 and 10,000, and efficiently deliver spectrally dispersed visibilities, closure phases and differential phases.  The paper is organized following increasing evolutionary stage of the different Wolf-Rayet stars observed so far. In Sect.~\\ref{section:WR79a}, we describe observations made on the star WR79a, at the boundary between O and WR types. In Sect.~\\ref{section:gammavel}, we present new AMBER observations made on the WC-type $\\gamma^2$ Velorum where we probe for the presence of a wind-wind collision zone. Then, in the last section, we present data acquired on two dusty Wolf-Rayet stars, where the observed data points are reminiscent of a ``pinwheel'' nebula. ",
        "conclusions": "We presented here a collection of AMBER observations on several Wolf-Rayet stars, and the different phenomena that can be characterized using the VLTI today: \\begin{itemize} \\item the wavelength-dependent spatial extent of the optically thick wind of single Wolf-Rayet stars (WR79a), \\item the presence of wind-collision zones by means of the free-free emission of the compressed gas, \\item the detection of ``pinwheel'' nebulae in dusty Wolf-Rayet stars which are too far to be resolved by single-dish telescopes. \\end{itemize}  The corresponding data processing and interpretations are still under progress, but they all show the great potential of AMBER to resolve the many different geometric features that evolved stars such as Wolf-Rayet can show to us."
    },
    "0807/0807.2891_arXiv.txt": {
        "abstract": "Observations of distances to Type-Ia supernovae can be explained by cosmological models that include either a gigaparsec-scale void, or a cosmic flow, without the need for Dark Energy. Instead of invoking dark energy, these inhomogeneous models instead violate the Copernican Principle. we show that current cosmological observations (Supernovae, Baryon Acoustic Oscillations and estimates of the Hubble parameters based on the age of the oldest stars) are not able to rule out inhomogeneous anti-Copernican models. The next generation of surveys for baryonic acoustic oscillations will be sufficiently precise to either validate the Copernican Principle or determine the existence of a local Gpc scale inhomogeneity. ",
        "introduction": "The Copernican Principle states that we do not occupy any special place in the Universe. The Friedmann--Lema\\^itre--Robertson--Walker (FLRW) models are built on this principle, and provide a remarkably precise description of cosmological observations \\cite{WMAP5,sdss,2df}. One could therefore think of this as an indirect demonstration of the Copernican Principle. However, homogeneous and isotropic FLRW models are not the only solutions of Einsteins equations which are able to fit cosmological observations. In particular, a number of inhomogeneous models have been  proposed, each of which are able to describe the evolution of distance with redshift as measured via type-Ia supernova (SnIa) without need of a cosmological constant \\cite{sn,sncmb,GH08a,GH08b,CFL08} (see \\cite{C07} for a review). Moreover, these models may be constructed in such a way that they describe the details of the cosmic microwave background power spectrum (CMB)~\\cite{sncmb,GH08a}. By eshuing the constraint of inhomogeneity, these inhomogeneous  models violate the Copernican Principle and suggest that we live near the center of a large (of several gigaparsecs diameter) highly isotropic void.  Inhomogeneous models are able to fit a variety of sets of cosmological observations without containing a cosmological constant because the lack of homogeneity offers a great degree of flexibility \\cite{MHE97}. For example, since the last scattering surface is separated from regions where supernova are observed by great distances, the property of inhomogeneity allows a model to be constructed which provides different physical densities in the regions from which these two sets of observational data are drawn. Thus, what is required to constrain inhomogeneous solutions to the apparent acceleration of the Universe includes several sets of data that measure a range of observables at comparable  redshifts. In this paper we confront two general classes of inhomogeneous, dark energy free models with observations of SnIa, with measurements of the baryon acoustic oscillations (BAO), and with the variation of the Hubble parameter with redshift. All these observational data are drawn from $z<1.8$. We begin by describing the Lama\\^itre--Tolman model (the simplest spherically symmetric inhomogeneous generalization of FLRW model), and then construct a range of 2-parameter models to describe a local large scale inhomogeneity. We then discuss constraints on the scale of these inhomogeneities based on existing and forthcoming data. ",
        "conclusions": "Homogeneous cosmological models which include a dark energy component and which are built on the Copernican Principle have been spectacularly successful in describing a suite of cosmological observations. However, alternatives which replace dark energy with an gigaparsec-scale inhomogeneity have also been postulated as explanations for the apparent acceleration of the Universe. In this paper we have shown that most inhomogeneous anti-Copernican models can already be ruled out by combining current cosmological observations (SNIa, H(z), and BAO). Our model were based on a simple parametrization of an inhomogeneity. We assumed that the inhomogeneity is modeled with a Gaussian-like profile, although other profiles are also possible. For example a similar analysis was recently proposed by Garc\\'ia-Bellido and Haugb\\o lle \\cite{GH08c} were the inhomogeneity was modeled using a hyperbolic tangent function. Their results also support our conclusion that there is still a range of parameter space that remains consistent with current constraints. As a result the success of Concordance models cannot be considered as a validation of the Copernican Principle. Forthcoming experiments will be capable of constraining all possible alternatives and will therefore either validate the Copernican Principle or determine the existence of a local Gpc-scale inhomogeneity.  \\ack This research was supported by the Peter and Patricia Gruber Foundation and the International Astronomical Union (KB) and by the Australian Research Council (JSBW). We thank Aleksandra Kurek for helpful discussions.    \\appendix"
    },
    "0807/0807.0636_arXiv.txt": {
        "abstract": "We explore how the slopes and scatters of the scaling relations of disk galaxies (\\vm-$L$[-$M$], $R$-$L$[-$M$], and \\vm-$R$) do change when moving from $B$ to $K$ bands and to stellar and baryonic quantities.  For our compiled sample of 76 normal, non--interacting high and low surface brightness (SB) disk galaxies, we find important changes, which evidence evolution effects, mainly related to the gas infall and star formation (SF) processes. We also explore correlations among the $(B-K)$ color, stellar mass fraction \\fs, mass $M$ (luminosity $L$), and surface density (SB), as well as correlations among the residuals of the scaling relations, and among these residuals and those of the other relations studied here. Some of our findings are: (i) the scale length \\hb~ is a third parameter in the baryonic Tully--Fisher relation (TFR) and the residuals of this relation follow a trend (slope $\\approx -0.15$) with the residuals of the \\hb-\\md~ relation; for the stellar and $K$ band cases, the scale length is not anymore a third parameter and the mentioned trend disappears; (ii) among the TFRs, the $B-$band TFR is the most scattered; in this case, the color is a third parameter, in agreement with previous works; (iii) the low SB galaxies break some observed trends in diagrams that include surface density, color, and \\fs, suggesting then a threshold in the gas surface density $\\Sigma_g$, below which the SF becomes independent of the gas infall rate and $\\Sigma_g$. Our results are interpreted and discussed in the light of $\\Lambda$ Cold Dark Matter--based models of disk--galaxy formation and evolution. These models are able to explain not only the baryonic scaling correlations, but also most of the processes responsible for the observed changes in the slopes, scatters, and correlations among the residuals when changing to stellar and luminous quantities.  The galaxy baryon fraction, \\fb, is required to be smaller than 0.05 on average. We detect some potential difficulties for the models: the observed color-$M$ and surface density-$M$ correlations are steeper, and the intrinsic scatter in the baryonic TFR is smaller than those predicted. ",
        "introduction": "Disk galaxies are the main population of galaxies in the local Universe. The study of their global scaling relations is of paramount relevance for understanding the formation and evolution of galaxies in general.  The main disk--galaxy observed global scaling relations are those between maximum circular velocity, \\vm, and luminosity, $L$; between disk scale length, $R$, and $L$; and between \\vm\\ and $R$. Each one of these relations can be established for different optical/near--IR (NIR) pass--bands. Under some assumptions, analogous relations can be calculated for the corresponding stellar and baryonic quantities. The changes in the zero points, slopes and scatters of these relations, when changing from one color band to another, and to stellar and baryonic quantities, are expected to be intimately related to the stellar--population properties and radial distributions of disk galaxies. In turn, these properties and distributions depend mainly on the disk assembly and star--formation (hereafter SF) histories. Therefore, the study of the scale relations for different passbands and for the stellar and baryonic cases offers valuable information on the structural, dynamical and SF properties of disk galaxies, as well as important constraints on models of their formation and evolution. The present paper is focused on this study.  \\subsection{Background on the disk scaling relations}  The most robust and studied of the disk--galaxy scaling relations is the one between $L$ and \\vm, the so--called Tully--Fisher relation (TFR). Originally discovered in the $B$ passband (Tully \\& Fisher 1977), it was afterward shown that it also applies to infrared bands. More recently, after applying a velocity (luminosity)--dependent extinction correction, Tully \\& Pierce (2000) have confirmed that the slope of the $L$-$\\vm$ correlation gets steeper systematically from $B$ ($\\approx 2.9$) to $K$ ($\\approx 3.5$) bands. This result suggests some dependence of galaxy colors on \\vm\\ for a given galaxy luminosity or mass. The amplitude of the scatter of the TFRs changes also with wavelength. These dependencies of the slope and scatter of the TFR on wavelength are related mainly to stellar population effects, which on their own, are related to the assembly and SF histories of galaxies.  Regarding the intrinsic TFR scatters, they are found to be small, posing a challenge for scenarios aimed to explain the origin of disk galaxies (Eisenstein \\& Loeb 1996; Avila-Reese et al. 1998; Mo et al. 1998; Firmani \\& Avila-Reese 2000).  The TFR has widely been studied for estimating extragalactic distances. In this case the independent variable (predictor) is the circular velocity, and the galaxy samples are carefully pruned to minimize the scatter. The latter implies a selection of samples for a narrow range of morphological types and surface brightnesses (hereafter SBs). However, for studies that aim to understand the origin of the TFR and use it to constrain models of galaxy formation rather than circular velocity, the luminosity should be the predictor, and the sample should be as wide as possible in morphological types, SBs, and colors (Courteau et al. 2007). Hereafter, we will refer to the TFR as the \\vm-$L$(-$M$) correlation.   By using NIR magnitudes --which are good tracers of stellar masses-- and gas contents when available, some authors have calculated the stellar and baryonic TFRs (e.g., McGaugh et al. 2000; Bell \\& de Jong 2001; Mayer \\& Moore 2004; Gurovich et al. 2004; Pizagno et al. 2005; de Rijcke et al. 2007). As discussed in Firmani \\& Avila-Reese (2000; hereafter FA-R), on one hand the baryonic TFR appears to be closely related to the cosmological halo maximum circular velocity--mass relation, and on the other hand, when compared to the stellar TFR, it reflects some aspects of the SF efficiency for different types of galaxies. The behavior of the scatter, when one changes from $B-$ to $K-$band TFRs and to the stellar and baryonic TFRs, certainly contains relevant information on the processes of disk galaxy evolution such as gas infall, SF, and stellar population histories.   The radius for a given luminosity (or the disk SB) has been considered as a potential third parameter that correlates with the scatter of the TFR (e.g., Kodaira 1989; Willick 1999). The scale length $R$ is indeed another structural parameter of disk galaxies that span a wide range of values. Some authors have reported that the scatter of the TFR is reduced up to 50\\% when $R$ is introduced as a third parameter (e.g., Koda, Sofue \\& Wada 2000; Han et al. 2001).  However, if the scatter of a given relation decreases upon introducing a third variable, it does not necessarily mean that such a variable is a statistically significant parameter. Indeed, some authors have shown that $R$ (or central SB) is not a physical third parameter in the TFR (e.g., Han 1991; Tully \\& Verheijen 1997; Willick et al. 1997; Courteau \\& Rix 1999; FA-R), but the radius could be a third parameter in the baryonic TFR (FA-R; Avila-Reese, Firmani \\& Zavala 2002; Zavala 2003; Dutton et al. 2007).  It is well known that the radius correlate with $L$ (Freeman 1970) and \\vm\\ (Tully \\& Fisher 1977), although these correlations are much more scattered than the TFR. Disk galaxies in the Log$L$--Log$\\vm$--Log$R$ parameter--space indeed lie in a thin oblique plane whose edge--on projection is close to the Log$L-$Log\\vm\\ (TFR) projection (Koda, Sofue \\& Wada 2000; Han et al. 2001; Shen, Mo \\& Shu 2002; Gnedin et al. 2007). This is the so called fundamental plane of disk galaxies. Here we construct from observations and analyze separately each one of the three projections of these planes.   \\subsection{Construction and interpretation of the scaling relations: how do they change from luminous to stellar and baryonic quantities?}  In most of the previous works, the disk--galaxy scaling correlations were constructed for samples biased in morphological types and SBs (pruned for reducing the scatter), and without selecting the environment. However, {\\it for theoretical and interpretative studies, the  observational galaxy sample should be as wide as possible in the range of morphologies, SBs, and colors} in order to include the whole population of disk galaxies, later to be compared with model predictions. Besides, since galaxy models are commonly used for isolated galaxies, the observational sample should also be selected to contain galaxies {\\it not affected by the environment}. On the other hand, the limited information of the samples employed in previous works, did not allow one to explore how the scaling correlations and their scatters change from luminous to stellar and baryonic quantities.   In this paper we shall use a sample of $76$ disk galaxies compiled from the literature and homogenized in Zavala et al. (2003a; hereafter ZAHF) and Zavala (2003).  This sample is not in any sense complete, an aspect that is crucial for volume--averaged population studies (e.g., the luminosity function), but is less relevant for studies on the correlations of physical properties. The sample includes high and low SB (hereafter HSB and LSB, respectively) disk {\\it normal}\\footnote{Here we understand by normal galaxies: non--dwarf, non--interacting, nearly isolated local galaxies.} galaxies of all morphological types, and it provides photometric information, at least in $B$ and $K$ bands, as well as information on the rotation curve and the HI integrated flux. None of the previous related works have presented a sample with these characteristics, which are crucial for the analysis proposed here. By using this sample, we will explore the differences in the slope and scatter of each scaling relation defined in the $B$ and $K$ bands, and for stellar and baryonic quantities. The results of this exploration will be interpreted in the light of disk--galaxy formation and evolution models.  The current theory of galaxy formation and evolution is based on the hierarchical $\\Lambda$ Cold Dark Matter (\\LCDM) cosmological model (for recent reviews, see Baugh 2006; Avila-Reese 2007). Several  authors have attempted to predict the scaling relations of disk galaxies and their scatters within this model (e.g., Avila-Reese et al. 1998,2002; Mo et al. 1998; Steinmetz \\& Navarro 1999; Navarro \\& Steinmetz 2000; FA-R; Koda et al. 2000; Shen et al. 2002; Zavala 2003; Courteau et al. 2003,2007; Dutton et al. 2005,2007; Gnedin et al. 2007; Courteau et al. 2007). The results are encouraging. We shall use some of these results, mainly those of FA-R, to interpret and discuss the new inferences from observations presented here.  Note that it was a common practice in the literature to compare model results with the observed {\\it luminous} properties and correlations, while most of the galaxy models did not actually include or self--consistently treat the process of gas transformation into stars, as well as did not follow the evolutionary path of halos and galaxies. As FA-R showed (see also Avila-Reese et al. 2002; Dutton et al. 2005), on one hand, there are systematical differences in the predicted scaling relations and their scatters among the luminous, stellar, and baryonic quantities; differences that we will explore here. On the other hand, properties such as the stellar and baryonic $M$, $R$, \\vm, gas fraction, and integral colors depend not only on the initial parameters, but also on the evolutionary path of the halo--disk system (Stringer \\& Benson 2008). Contrary to the ``static population'' models, such as those of Mo et al. (1998) and extensions of them, the FA-R models are evolutionary. With some simplifications, they follow self--consistently (see Appendix B): the halo and disk assembly, taking into account the halo adiabatic contraction; the self--regulated disk SF and feedback; the luminosity evolution; and the secular bulge formation.  During the completion of the present work, several of the papers mentioned above, with results and a philosophy similar to ours, appeared in the literature. Although some overlapping is present, our contribution is original in the sense that we focus the research {\\it on the changes in the scaling relations and their scatters when moving from optical to NIR bands and to stellar and baryonic quantities}.  Our sample is similar in number to others used in recent works (e.g., Gnedin et al. 2007) and smaller with respect to Dutton et al. (2007) and Courteau et al. (2007), who use $\\sim 1300$ objects with $I-$ band photometry and $H\\alpha$ kinematics, but only $\\sim 360$ out of them, mainly HSB galaxies, have available $K-$ band photometry. These large samples are undoubtedly useful and we will compare, when possible, with their results. However, we can not use them because they lack several observables (e.g., gas content, optical photometry) and requirements (e.g., inclination limits, LSB galaxies in both optical and $K$ bands) that we have imposed for our study. We hope that our study will serve as a benchmark for similar analyzes based on large multiwavelength galaxy surveys, which are currently being completed.  This paper is organized as follows. In \\S 2 we describe the disk galaxy sample to be used. The three main scaling correlations and their scatters are constructed and discussed in \\S 3 for bands $B$ and $K$ and for the inferred stellar and baryonic quantities. Some other correlations among global galaxy quantities are also reported. In \\S 4, in the light of previous \\LCDM--based disk galaxy formation and evolution models, we interpret the changes in the slopes and scatters of the scaling relations when changing from $B$ to $K$ bands and to stellar and baryonic quantities. A summary of the results and the conclusions of this work are presented in \\S 6. ",
        "conclusions": "  $\\bullet$ The slopes of the baryonic \\vm-\\md, \\hb-\\md, and \\vm-\\hb~ correlations can be interpreted as a direct imprint of the cosmological \\LCDM~ halo \\vmh-\\mh, \\rh-\\mh, and \\vmh-\\rh~ relations, especially the former (baryonic TFR), which is the tightest one. The models show that this relation is robust to systematical or statistical variations in the disk baryonic mass fraction, \\fb, for $\\fb<0.08$. Then, the main sources of scatter are the dispersions in the halo spin parameter $\\lambda$ and halo MAH (or concentration). The FA-R models provide a lower limit to the scatter around the relation due to these two dispersions, namely \\sigintrV$\\approx 0.053$ (\\fb=0.05=const.), which is already slightly larger than our observational inference.  $\\bullet$ In combination, the slopes and scatters of the baryonic \\vm-\\md, \\hb-\\md, and \\vm-\\hb~ correlations inferred here from observations, as well as the correlations among their residuals, imply that the average value of \\fb~ in normal disk galaxies could not be larger than $\\sim 0.05$ and that \\fb~ can not depend strongly on $M$ (or $L$); if any, a moderate anti--correlation with $M$ is permitted. A trend of \\fb~ nearly independent on mass, with average values of $\\fb<0.05$, has been reported recently by other authors, who extended the analysis to disk dwarf galaxies and inferred halo masses directly from the observational data (Blanton et al. 2007; see also Baldry et al. 2008; Begum et al. 2008).   $\\bullet$ The models are able to explain qualitatively the observed (weak) anti--correlation among the residuals of the {\\it baryonic} \\vm-\\md~ and \\hb-\\md~ correlations and the related finding that the radius (or surface density) is a third parameter in the former correlation. The physical parameter beyond these behaviors is $\\lambda$. Model disks with low $\\lambda$ have high surface density and contribute significantly to \\vm~, in such a way, that for a given mass, \\vm~ correlates with \\hb~ or \\sigd.  Instead, disks with high $\\lambda$ values are of low surface density and have a negligible gravitational contribution to \\vm; then, the dependence of \\vm~ on \\sigd~ (or on \\hb~ for a fixed mass) tends to disappear. The models can also explain why the anti--correlation among the residuals becomes negligible in the stellar or $K-$band cases: for a given mass, galaxies not only shift to the high-\\vm~ side in the TFR diagram as the radius is smaller (\\sigd~ is higher), but also shift to the high-\\ms~ (or $L_K$) side due to a higher efficiency in transforming gas into stars. As a result, the scatter around the stellar TFR becomes independent of \\hs~ or \\sigs, and smaller than in the baryonic case.  $\\bullet$ The (self--regulated) SF efficiency effect just mentioned produces a dependence of the stellar mass fraction \\fs~ on surface density (or brightness). Such a non--trivial prediction roughly agrees with the one found for our processed galaxy catalog (Fig. \\ref{stfrac}), and it is associated mainly with the $\\lambda$ parameter: the smaller the $\\lambda$, the higher the disk surface density and the more efficient the transformation of gas into stars. Although less relevant, \\fs~ and \\sigs~ are also affected by the gas infall history, which is connected to the cosmological halo MAH. Since the MAH plays a role in the scatter of the baryonic TFR, trends among the residuals of this relation and those of the color-\\md~ and \\fs-\\md~ relations are predicted. These trends were found for our observational catalog, mainly for the HSB galaxies (Fig. \\ref{residuos}). For LSB galaxies, it could be that the SF--rate timescale becomes larger than the that of the gas infall rate, the connection between MAH and color/\\fs~ then disappears.   $\\bullet$ {\\it Potential difficulties.-} Our observational results show that both (i) the color--$M$ and (ii) surface density--$M$ correlations are steeper than the models would predict. We speculate that these shortcomings suggest that late minor mergers and interaction--induced SF, which are more common in massive galaxies, should play some role and/or that both, the galaxy spin parameter and \\fb~ anti--correlate with mass; at least qualitatively, both anti--correlations, in combination are not expected to affect the scaling correlations presented here. (iii) The fact that the intrinsic scatter around the predicted baryonic TFR appears to be larger than that inferred here, implies that there is no room for other physical or systematical sources of scatter around the TFR such as disk ellipticity. On the other hand, the intrinsic scatter around the stellar TFR becomes smaller in the models than in the observations. However, external effects, such as the interaction--driven SF, work in the direction of adding some scatter to the stellar TFR. (iv) The FA-R models (rescaled to $\\sigma_8<1$) agree with the observed baryonic and stellar TFR zero points as well as with the \\mh--to--$L$ ratios inferred directly from observations. The potential difficulty that the SAM/HOM find in matching both the TFR zero point and the luminosity function could be solved by taking into account the shortcomings we have discussed in \\S 5.   The results obtained in this work pose relevant questions for future studies. We showed that the determination of the scaling relations for disk galaxies in different bands and for stellar and baryonic quantities, as well as for the correlations among their residuals, bring valuable information on the nature and evolution of galaxies. The completion of large galaxy samples, unbiased and broad in properties, with detailed photometric, dynamic, and gas--content information is crucial for this kind of study. The preferred approaches for such an undertaking are: follow--up observations in $H_\\alpha$ and/or HI emission lines for galaxies from the largest optical/NIR band surveys as SDSS (Pizagno et al. 2005), and synergy of existing wide--area homogeneous surveys conducted at optical/NIR bands and HI emission line, for example SDSS with HI surveys such as ALFALFA (Giovanelli et al. 2005).  According to the philosophy followed in this paper, substantial advance comes from the adequate comparison of observations with theoretical models. Thus, in order to confirm, rule out, or perhaps expand, some of the conclusions attained here, a deep model exploration of the formation and evolution of the disk galaxy population should be carried out."
    },
    "0807/0807.2633_arXiv.txt": {
        "abstract": "Extending previous results [Phys. Rev. Lett. 97, 081301 (2006)], we explore the cosmological implications of a new quintessence scenario driven by a slow rolling homogeneous scalar field whose equation of state behaved as freezing over the entire cosmic evolution, is approaching $-1$ today, but will become thawing in the near future, thereby driving the Universe to an eternal deceleration. We argue that such a mixed behavior, named \\emph{hybrid}, may reconcile the slight preference of current observational data for freezing potentials with the impossibility of defining observables in the String/M-theory context due to the existence of a cosmological event horizon in asymptotically de Sitter universes as, e.g., pure freezing scenarios. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0807/0807.5013_arXiv.txt": {
        "abstract": "High time resolution observations of \\psr\\ (or PSR J0908$-$4913) over a wide range of frequencies have enabled us to determine the geometry and beam shape of the pulsar.  We have used the position angle traverse to determine highly-constrained solutions to the rotating vector model which show conclusively that \\psr\\ is an orthogonal rotator. The accuracy obtained in measuring the geometry is unprecedented. This may allow tests of high-energy emission models, should the pulsar be detected with GLAST. Although the impact parameter, $\\beta$, appears to be frequency dependent, we have shown that this is due to the effect of interstellar scattering.  As a result, this pulsar provides some of the strongest evidence yet that the position angle swing is indeed related to a geometrical origin, at least for non-recycled pulsars. We show that the beam structures of the main pulse and interpulse in \\psr\\ are remarkably similar. The emission comes from a height of $\\sim$230~km and is consistent with originating in a patchy cone located about half way to the last open field lines.  The rotation axis and direction of motion of the pulsar appear to be aligned. ",
        "introduction": "Pulsars are highly magnetized, rotating neutron stars which emit beams of radio emission at a low height above the star's magnetic poles.  The emission is confined to the open field lines and the beam only sweeps out a small fraction of the celestial sphere. For this reason we see emission from only one pole in the vast majority of pulsars.  In fortunate geometrical circumstances, however, the second pole should also be observable as a radio emitting site, if the spin axis is almost orthogonally orientated to the line-of-sight to Earth. Such pulsars can be identified by the detection of an ``interpulse'' separated by $\\sim 180^\\circ$ in pulse longitude from the main pulse.  Less than 30 pulsars have been identified to show interpulses \\cite{rib08,wj08a} from the known population of $\\sim$1800 pulsars.  Not all interpulses may arise from the second magnetic pole, but some of them can also be interpreted within a geometry of an aligned rotator where the observer's line-of-sight never, or hardly, leaves the emission cone during the pulse period (e.g.~Manchester \\& Lyne 1977).\\nocite{ml77}  The geometry of a pulsar may be determined from observations of their polarization which has been a major industry since shortly after their discovery. In the more than 40 years since, a few key points are well established (see e.g.~Lorimer \\& Kramer 2005).\\nocite{lk05} Pulsars are essentially rotating dipoles and the super-strong magnetic field naturally implies a high degree of linear polarization with the position angle (PA) of the radiation parallel (or perpendicular) to the magnetic field direction. Therefore, as the dipole field sweeps past the observer once per rotation of the pulsar, the observer sees a characteristic S shaped sweep of the position angle. Known as the rotating vector model (RVM), it provides a simple mapping between the relevant geometrical angles and the position angle variations \\cite{rc69a}.  The PA as a function of pulse longitude, $\\phi$, can be expressed as \\begin{equation} {\\rm PA} = {\\rm PA}_{0} + {\\rm arctan} \\left( \\frac{{\\rm sin}\\alpha \\, {\\rm sin}(\\phi - \\phi_0)}{{\\rm sin}\\zeta \\, {\\rm cos}\\alpha - {\\rm cos}\\zeta \\, {\\rm sin}\\alpha \\, {\\rm cos}(\\phi - \\phi_0)} \\right) \\end{equation} Here, $\\alpha$ is the angle between the rotation axis and the magnetic axis, and $\\zeta=\\alpha+\\beta$, with $\\beta$ being the angle at closest approach of the line of sight to the magnetic axis.  $\\phi_0$ is the corresponding pulse longitude at which the PA is then PA$_{0}$.  Note that we follow the observer's sign convention for the PA, which differs from that introduced by Damour \\& Taylor (1992\\nocite{dt92}, see also Everett \\& Weisberg 2001\\nocite{ew01}).  The RVM has been used extensively to derive geometrical angles in a number of pulsars \\cite{ran83,lm88}. Unfortunately there are many factors which mean that $\\alpha$ and $\\beta$ are often subject to large uncertainties. These include a covariance between $\\alpha$ and $\\beta$ which cannot be broken over the usually small longitude range over which pulsars emit (e.g.~von Hoensbroech \\& Xilouris 1997\\nocite{hx97a}, Everett \\& Weisberg 2001\\nocite{ew01}), but also orthogonal modes and propagation effects in the pulsar magnetosphere \\cite{ew01}. For these reasons multi-frequency studies of a given pulsar rarely give similar values for $\\alpha$ and $\\beta$ as the uncertainties are large (e.g. Mitra \\& Li 2004\\nocite{ml04}).  Some information about the geometry can also be obtained from the evolution of pulse profiles with frequency.  Observational data show that low frequency pulse profiles tend to be wider in extent than high frequency profiles \\cite{tho91a} and the explanation for this is that radio emission originates at different heights in diverging dipolar magnetic field lines.  In this model, high frequency radio emission originates from closer to the stellar surface than low frequency emission with emission heights ranging from several tens to several hundreds of km \\cite{mr02a}. Recent evidence tends to suggest that, at a given frequency, the emission height varies across the polar cap with lower emission heights over the magnetic axis \\cite{gg03}.  Also, observations indicate that low-frequency profiles are often dominated by strong central components whereas high-frequency profiles show strong outrider components in many cases. This has been interpreted as indicating a different spectral index for core and cone emission \\cite{ran83}, while geometrical factors may be at least partly responsible \\cite{kwj+94}.  When considering emission heights, it is intriguing to also consider the observational consequences of relativistic effects in the magnetosphere. Blaskiewicz, Cordes \\& Wasserman (1991)\\nocite{bcw91} pointed out that, to first order, the main effect is that the PA swing is delayed with respect to the total intensity profile and the magnitude of this effect is directly propotional to the emission height. This effect should therefore be frequency dependent if different frequencies originate at different heights. The studies by Blaskiewicz et al.~(1991), von Hoensbroech \\& Xilouris (1997) and Mitra \\& Li (2004)\\nocite{ml04} showed evidence for this effect in a sample of pulsars.  In this paper, we present observations and their analysis of \\psr\\, a 100-ms pulsar discovered by D'Amico et al.~(1988)\\nocite{dmd+88} which allows us, as it turns out, to study our described understanding of pulsar emission. Initial observations showed that the profile consisted of two widely separated components.  Polarization observations at 0.66 and 1.6~GHz were carried out by Qiao et al.~(1995)\\nocite{qmlg95} and Wu et al.~(1993)\\nocite{wmlq93} respectively. They showed that the profile is virtually 100 per cent linearly polarized with an indication of a small fraction of circular polarization. The time resolution of the observations was low and each component was largely featureless. Wu et al. (1993)\\nocite{wmlq93} attempted an RVM fit to the PA swing and determined that $\\alpha=58\\degr$ and $\\beta=30\\degr$.  This fit indicated that the emission arises from a single pole and supported the idea that young pulsars had wide profiles \\cite{man96}.  In a survey looking for pulsar wind nebulae, Gaensler et al.~(1998)\\nocite{gsfj98} discovered an unusual looking nebula around \\psr. The head of the nebula is clearly resolved and a bow-shock indicates that the direction of motion of the pulsar is at an angle of $\\sim$315\\degr\\ (measured as north through east).  Polarization properties of the pulse profile can give an indication of the position angle of the pulsar's rotation axis and we will investigate if the correlation between the velocity and rotation axis seen by Johnston et al. (2005)\\nocite{jhv+05} is also evident in this pulsar, and whether it confirms the direction of motion suggested by Gaensler et al.~(1998).  The organisation of the paper is as follows. In Section 2 we briefly discuss multi-frequency observations of \\psr\\ as part of a bigger campaign to measure polarization in southern pulsars generally.  In Section 3 we show how the profile of \\psr\\ evolves with frequency and present our RVM fits, and re-examine in particular the question of whether this pulsar shows emission from a single, wide pole or from two poles. We discuss the implications of the RVM fits in Section 4 and derive a model for the beam of \\psr.  \\begin{figure} \\centerline{\\psfig{figure=fig1.ps,angle=0,width=9.5cm}} \\caption{\\label{fig:profiles} Polarization profiles of \\psr\\ obtained at three frequencies. While the interpulse shows hardly any frequency evolution, the main pulse exhibits an unusual behaviour (see text). The circular polarisation changes sign between the low and high frequencies. The insets show a Gaussian decomposition of the main and interpulse. } \\end{figure} ",
        "conclusions": "High time resolution observations over a range of frequencies for \\psr\\ have enabled us to determine the geometry and beam shape of the pulsar.  Uniquely amongst pulsars, even amongst those with interpulses, we have been able to determine highly constrained RVM fits which show conclusively that \\psr\\ is an orthogonal rotator. We have shown that the value of $\\beta$ appears to be frequency dependent; however careful modelling shows that this appears to be entirely due to the effect of interstellar scattering of the profiles. While other factors usually dominate the uncertainties in the geometry derived from RVM fits, it should be considered in particular for low-frequency observations or observations of short-period pulsars. It will also be relevant for future high-sensitivity detections and studies of interpulse sources with the Square-Kilometre-Array \\cite{ckl+04}.  Interestingly, the relativlely high ratio of spin-down luminosity $\\dot{E}$ to distance $d$ as measured by $\\dot{E}/d^2 = 1.1\\times 10^{34}$ erg s$^{-1}$ kpc$^{-2}$ makes this pulsar potentially detectable with the Gamma-Ray Large Area Space Telescope (GLAST). This would be an exciting opportunity to test competing $\\gamma$-ray emission models, given the very well determined geometry of this pulsar. For instance, outer gap models (e.g.~ Chiang \\& Romani 1994)\\nocite{cr94} predict that orthogonal rotators should be $\\gamma$-ray bright as we look directly into the null-charge surface, i.e.~the bottom part of the outer gap (Yadigaroglu \\& Romani 1995).  \\nocite{yr95}   We have shown that the total intensity structure of the MP and the IP in \\psr\\ have similarities. While the degree of linear polarisation is very large in both MP and IP, the circular polarisation changes handedness with increasing frequency with an intriguing mirror-symmetry between the MP and IP. The emission comes from a height of $\\sim$230~km in both and is consistent with an origin in a cone with a radius of only one half of the beam opening angle. The cone is patchy as we do not detect the leading emission from the MP.  In summary, this pulsar represents a show-case for many features of pulsar radio emission, combined in a single, remarkable source. But most of all, this source provides conclusive evidence for the geometrical interpretation of the PA swing within the RVM. Many pulsars, however, in particular recycled pulsars, show large deviations from a simple RVM swing in which cases the geometrical interpretation has to be questioned. In this case, however, the evidence is incontrovertible."
    },
    "0807/0807.3918_arXiv.txt": {
        "abstract": "s{ Gamma-ray bursts are violent events occurring randomly in the sky. In this review, I will present the fireball model, proposed to explain the phenomenon of gamma-ray bursts. This model has important consequences for the production and observation at Earth of gravitational waves, high energy neutrinos, cosmic rays and high energy photons, and the second part of this review will be focused on these aspects. A last section will briefly discuss the topic of the use of gamma-ray bursts as standard candles and possible cosmological studies.} ",
        "introduction": " ",
        "conclusions": "The science of GRBs is still very active. Open questions are related to the nature of the central engine, the acceleration mechanism at play during the first stages of the burst, the properties of the progenitors that lead to a GRB rather than a simple supernova, ...  France, in collaboration with China and Italy is deeply involved in a future satellite, {\\it SVOM} \\cite{svom0}. SVOM will be devoted to the study of the prompt and early afterglow phases, thus observing directly the last part of the acceleration stage of the fireball. Used with non-photonic instruments still in construction at the moment (such as Antares, IceCUBE or the Advanced VIRGO/LIGO) and very high energy experiments (AUGER, GLAST), this will provide very interesting constraints on these problems, ... and should of course raise new unexpected questions."
    },
    "0807/0807.4824_arXiv.txt": {
        "abstract": "s{Detection of the air-fluorescence radiation induced by the charged particles of extensive air showers is a well-established technique for the study of ultra-high energy cosmic rays. Fluorescence telescopes provide a nearly calorimetric measure of the primary energy. Presently the main source of systematic uncertainties comes from our limited accuracy in the fluorescence yield, that is, the number of fluorescence photons emitted per unit of energy deposited in the atmosphere by the shower particles. In this paper the current status of our knowledge on the fluorescence yield both experimental an theoretical will be discussed.} ",
        "introduction": "Fluorescence telescopes have been successfully used for the detection of ultra-high energy cosmic rays ($> 10^{18}$ eV) since the pioneering Fly's Eye experiment~\\cite{flys_eye}. In this technique the fluorescence radiation induced by the charged particles of the extensive air shower generated by a primary cosmic ray is registered at ground by wide-angle telescopes. Assuming that the intensity of the air-fluorescence light is proportional to the energy deposited in the atmosphere by the shower, this technique provides a nearly calorimetric measure of the energy of the primary cosmic ray. Therefore it has the advantage, as compared with methods relying on simulations (e.g. surface arrays working in standalone mode), of being nearly model independent. In spite of this advantage, fluorescence telescopes are presently limited by the uncertainty in the fluorescence yield, that is, the calibration parameter which converts number of fluorescence photons into absolute energy units. For instance in the Pierre Auger Observatory~\\cite{auger} the uncertainty in the fluorescence yield contributes a 14\\% to the total systematic error in the energy calibration which is presently 22\\%.  In order to improve the accuracy of this parameter, dedicated laboratory experiments~\\cite{experiments} are carrying out precise measurements of the air-fluorescence emission. In these experiments an electron beam excites air at certain pressure and temperature conditions. A large set of experimental parameters are measured, not only the absolute value of the fluorescence yield but also the spectral features of the fluorescence radiation and the dependence with atmospheric parameters (pressure, temperature, humidity, etc.). On the other hand, progress on the theoretical understanding of the various processes leading to the air-fluorescence emission is being carried out~\\cite{blanco}. \\begin{figure}[t] \\centering \\epsfig{width=1\\textwidth, file=figure1.eps} \\caption{a) Molecular levels of N$_2$ and N$_2^+$ involved in the generation of air-fluorescence and cross section versus electron energy for the excitation of the corresponding upper levels. b) At high electron energy most of the fluorescence light is generated by secondary electrons.} \\label{mol-levels} \\end{figure} ",
        "conclusions": "Our understanding on the processes leading to generation of air fluorescence has increased significantly in the last years~\\cite{5th_FW_SP}. The world-wide campaign for the experimental determination of the fluorescence yield has achieved remarkable results, in particular in the measurement of the various dependencies with atmospheric parameters. The fundamental assumption of proportionality between fluorescence intensity and deposited energy has been verified both theoretically and experimentally.  In regard with the determination of the absolute value of the fluorescence yield new data are available. However the interpretation of the results is not straightforward. A comparison using the procedure discussed here shows a general agreement with typical differences of about 15\\%. For a real improvement in the accuracy of fluorescence telescopes an uncertainty better than 10\\% in the fluorescence yield is necessary. Several experiments claim high accuracy, for instance, the reported uncertainty of the FLASH experiment is of about 8\\%. In addition the AIRFLY collaboration will publish soon a final absolute value with an error below 10\\%. A discussion on these and other high accuracy measurements have been presented elsewhere~\\cite{5th_FW_SP}. Discrepancies between these experiments go beyond the reported accuracies and therefore some experimental effort is still necessary to clarify the situation."
    },
    "0807/0807.4479_arXiv.txt": {
        "abstract": "This paper describes a new algorithm for Monte Carlo integration, based on the Field Estimator for Arbitrary Spaces (\\F). The algorithm is discussed in detail, and its performance is evaluated in the context of Bayesian analysis, with emphasis on multimodal distributions with strong parameter degeneracies. Source code is available upon request. ",
        "introduction": "\\label{secIntro}  Many problems in Physics, as well as in other branches of science, involve the computation of an integral \\be I=\\int_V f(\\xx)\\ \\dd\\xx \\label{eqI} \\ee over a given multidimensional domain $V$. In most circumstances, the value of $I$ has to be estimated numerically, evaluating the function $f$ at a series of $N$ sample points $\\xx_i$.  Numerical quadrature algorithms (see e.g. \\cite{Press+92}) sample $f$ on a regular grid or at a carefully selected set of unequally-spaced points (Gaussian quadrature). This approach is very efficient for smooth functions, but it is certainly far from optimal when the integrand is strongly peaked. Adaptive techniques provide a solution to this problem by refining those sub-intervals of the integration domain where $f$ is found to vary on smaller scales.  In multiple dimensions, Monte Carlo methods usually yield better accuracy for a given $N$ than repeated integrations using one-dimensional deterministic algorithms. The essence of the Monte Carlo method is that the sample points $\\xx_i$ are chosen (to some extent) at random. The most basic version uses a uniform probability distribution over the integration region, and the estimator for the integral $I$ is given by \\be \\IMC = \\frac{V}{N}\\sum_{i=1}^{N} f_i \\label{eqI0} \\ee where $V$ denotes the integration volume and $f_i\\equiv f(\\xx_i)$. The error associated to $\\IMC$ can be estimated as \\be \\DMC^2 = \\frac{V^2}{N^2}\\sum_{i=1}^{N} \\left( f_i - \\frac{\\IMC}{V} \\right)^2 . \\label{eqD0} \\ee  Although the error decreases as $N^{-1/2}$ (this is, in fact, why Monte Carlo methods are preferable in multidimensional problems), the basic Monte Carlo is still computationally inefficient when the function $f$ has a very narrow peak compared to the total integration volume. On the one hand, a large number of sample points would be required to obtain good estimates $\\IMC$ and $\\DMC$. On the other hand, and more important, it is indeed quite possible that the algorithm achieves convergence, reporting a wrong value of the integral with an extremely small estimate of the error, before finding the maximum (or maxima) of $f$.  A refinement of this method, known as importance sampling, consists in drawing the points from a distribution similar in form to the integrand, so that they are more likely to come from the regions that contribute most to $I$. Rather than sampling the whole integration volume uniformly, one chooses a probability distribution $g(\\xx)$ that adapts to the shape of the integrand. For any $g(\\xx)$, the integral~(\\ref{eqI}) can be re-written as \\be I=\\int_V \\frac{f(\\xx)}{g(\\xx)}~g(\\xx)~\\dd\\xx \\ee and the estimators of $I$ and its error become \\be \\IS{I} = \\frac{1}{N}\\sum_{i=1}^{N} \\frac{f_i}{g_i} \\label{eqI_IS} \\ee and \\be \\IS{\\Delta}^2 = \\frac{1}{N^2}\\sum_{i=1}^{N} \\left( \\frac{f_i}{g_i} - \\IS{I} \\right)^{\\!2} \\label{eqD_IS} \\ee respectively. The key issue is, of course, the choice of the distribution $g(\\xx)$. These two expressions reduce to~(\\ref{eqI0}) and~(\\ref{eqD0}) for the particular case of uniform sampling, i.e. all $g_i\\equiv g(\\xx_i)=1/V$. In the ideal case, when $g(\\xx)=|f(\\xx)|/I$, all points contribute exactly the same amount to the sum in~(\\ref{eqI_IS}) and zero in~(\\ref{eqD_IS}), and the algorithm takes only a few $N$ to achieve the desired accuracy. Unfortunately, choosing $g(\\xx)$ is not straightforward when the shape of $f$ is unknown \\emph{a priori}, and evaluating it is actually part of our goal.  A perfect example of such a situation is Bayesian inference, where one tries to select the model $M$ that best describes the data $\\vv{D}$ and estimate the values of the model parameters $\\vv{\\Theta}$. For a given model, Bayes' Theorem states that \\be p(\\vv{\\Theta}|M,\\vv{D}) = \\frac{ p(\\vv{\\Theta}|M)~p(\\vv{D}|M,\\vv{\\Theta}) }{ p(\\vv{D}|M) } \\label{eqP} \\ee where $p(\\vv{\\Theta}|M)$ is the prior probability distribution of the model parameters, $p(\\vv{D}|M,\\vv{\\Theta})$ is the likelihood of the data, $p(\\vv{\\Theta}|M,\\vv{D})$ is the posterior probability distribution of the parameters, and $p(\\vv{D}|M)$ is the evidence for model $M$, \\be p(\\vv{D}|M) = \\int p(\\vv{\\Theta}|M)~p(\\vv{D}|M,\\vv{\\Theta})~\\dd\\vv{\\Theta} \\label{eqE} \\ee The posterior probability distribution~(\\ref{eqP}) contains all the information needed for parameter estimation, and the evidence plays in this case the role of a mere normalization constant. Nevertheless, the value of the integral~(\\ref{eqE}) becomes critical in model selection. Applying again Bayes' Theorem, two models $M_1$ and $M_2$ can be compared by evaluating the Bayes' factor, \\be \\frac{ p(M_1|\\vv{D}) }{ p(M_2|\\vv{D}) } = \\frac{ p(M_1)~p(\\vv{D}|M_1) }{ p(M_2)~p(\\vv{D}|M_2) } \\ee where $p(M_i)$ denotes the prior probability of each model, $p(\\vv{D}|M_i)$ are their evidences, and $p(M_i|\\vv{D})$ are the posterior probabilities. Even if one is only interested in parameter estimation, it is good practice to quote the evidence (as well as the assumed priors) in order to make possible the comparison with future work.  The computational part of a Bayesian analysis consists thus in the evaluation of the integral $I=p(\\vv{D}|M)=\\int f(\\vv{\\Theta})~\\dd\\vv{\\Theta}$, where $f(\\vv{\\Theta})= p(\\vv{\\Theta}|M)~p(\\vv{D}|M,\\vv{\\Theta})$, over the region of the parameter space where $p(\\vv{\\Theta}|M)>0$. Since each evaluation of the likelihood may involve a significant computational cost, the algorithm should locate the the maxima of $f(\\vv{\\Theta})$ as efficiently as possible, even when this function has a complicated structure. In particular, it is not uncommon that $f(\\vv{\\Theta})$ presents several maxima, corresponding to different sets of parameter values that provide a good fit to the data. Which solution is to be preferred depends, within the Bayesian framework, on the value of the integral (the evidence) over the region associated to each peak. Even more often, there are strong degeneracies in parameter space. The maxima of the function $f$ can be extremely elongated, generally with a curved shape arising from a non-linear relationship between two or more parameters.  These two problems (multimodality and curved degeneracies) have proven difficult to overcome for many standard algorithms, especially as the number of dimensions increases. Markov Chain Monte Carlo methods (see e.g. \\cite{MacKay03}) typically require a relatively large number of evaluations. Nested sampling \\cite{Skilling04} provides a more efficient alternative, transforming the multidimensional integral~(\\ref{eqI}) into a single dimension by means of a re-parameterization in terms of the `prior mass' $p(\\vv{\\Theta}|M)~\\dd\\vv{\\Theta}$. This formalism has recently been extended to cope with the problems of multimodality and curved degeneracies by resorting to clustering algorithms \\cite{Shaw+07,FerozHobson_07}. On the other hand, traditional importance sampling methods (e.g. Vegas \\cite{Lepage78}) assume a separable weight function, i.e. $g(\\xx)=\\prod_{d=1}^Dg_d(x_d)$, which can give a good approximation to $f(\\xx)$ only as long as the characteristic features are well aligned with the coordinate axes. An adaptive refinement technique has been implemented in the algorithm Suave \\cite{Hahn05} to improve the performance in the general case.  Here I present a new method, based on the Field Estimator for Arbitrary Spaces (\\F) developed by \\cite{AscasibarBinney05} to estimate the continuous density field underlying a given point distribution. The algorithm is fully described in Section~\\ref{secMethod}, results are presented in Section~\\ref{secResults}, and the main conclusions are summarized in Section~\\ref{secConclus}. ",
        "conclusions": "\\label{secConclus}  This paper presents a new algorithm to carry out numerical integration in multiple dimensions or, in a Bayesian context, sample from the posterior distribution and compute the evidence for the assumed model. The method, dubbed `\\F\\ sampling', is a variant of importance sampling where the weight of each point is computed with the help of the Field Estimator for Arbitrary Spaces \\cite{AscasibarBinney05}.  Its performance has been tested for several toy problems with known analytical solution, specifically designed to contain multimodal distributions and significant degeneracies. The results suggest that \\F\\ sampling provides an interesting alternative to other methods for problems of low dimensionality. In particular, it is able to discover most isolated maxima of the integrand, and it can perfectly sample from distributions with pronounced degeneracies. As the number of dimensions increases, the ability of the algorithm to identify peaks decreases and the required computational resources (in terms of both time and memory) grow exponentially.    \\appendix"
    },
    "0807/0807.3545_arXiv.txt": {
        "abstract": "We present results from a pilot study of radial stellar population trends in early-type galaxies using the VLT VIMOS integral field unit (IFU). We observe twelve galaxies in the cluster Abell 3389 (z $\\approx$ 0.027). For each galaxy, we measure 22 line-strength indices in multiple radial bins out to at least the effective radius. We derive stellar population parameters using a grid inversion technique, and calculate the radial gradients in age, metallcity and $\\alpha$-abundance. Generally, the galaxies in our sample have flat radial trends in age and [$\\alpha$/Fe], but negative gradients in [Z/H] (--0.20 $\\pm$ 0.05 dex). Combining our targets with two similar, long-slit studies to increase sample size, we find that the gradients are not correlated with the central velocity dispersion or $K$-band luminosity (both proxies for galaxy mass). However, we find that the age and metallicity gradients are both anti-correlated with their respective central values (to $>$ 4$\\,\\sigma$), such that galaxies with young cores have steeper positive age gradients, and those with metal-rich centres have strong negative [Z/H] gradients. ",
        "introduction": "\\label{sec:intro}  Stellar population gradients offer a valuable probe of the star formation and chemical evolution of early-type galaxies. Simple models of classic monolithic collapse \\citep{lar74-585,car84-403} can be ruled out by the relative weakness of the observed metallicity gradients \\citep[e.g.][]{tam04-617}. Generally, age gradients are expected to be small for the short collapse times required for $\\alpha$-enhancement \\citep[e.g.][]{tho99-655}, while the $\\alpha$-enhancement itself could show positive or negative gradients \\citep*{pip07-0706.2932}. Within the hierarchical $\\Lambda$CDM models \\citep[e.g.][]{col00-168,del06-499,bow06-645}, merger gas content and feedback prescriptions can be investigated \\citep{kaw03-908,kob04-740}; mergers would tend to dilute or erase Z/H gradients \\citep[e.g.][]{whi80-1}, although dissipation could trigger star formation at the core, creating age gradients \\citep{bar91-65,mih94-47}.  Traditionally, observational studies of radial trends in early-type galaxies have employed long-slit spectroscopy in local cluster and field environments (e.g. \\citealt*{gor90-217,car93-553}; \\citealt{meh03-423}; \\citealt*{san06-823,nor06-815}; \\citealt{red07-1772,san07-759}). In general, most studies find flat age gradients, small $\\alpha$-abundance gradients (either positive or negative), and large negative radial trends in metallicity.  Long-slit spectrographs generally have a good optical throughput. However, the spatial data is immediately reduced to one-dimension and is dependent on the slit orientation, with most of the light from the galaxy undetected. The use of an integral field unit (IFU), which enables both spectral and spatial coverage, allows the analysis of the entire galaxy. With careful treatment, light from the entire circumference can be binned to give radial trends unbiased by a slit orientation.  One notable use of an IFU in this field is the Spectrographic Areal Unit for Research on Optical Nebulae (SAURON) project \\citep{bac01-23,dez02-513}. The instrument has good spatial resolution allowing the creation of maps for each measured parameter, particularly in kinematics. However, the spectral range of SAURON is limited to four useful line indices (H$\\beta$, Fe5015, Mgb5177, Fe5270$_{\\rm s}$). From spatial maps of the iron lines, \\citet{kun06-497} deduce that the metallicity gradients are generally negative and are roughly aligned with the light profiles. The H$\\beta$ maps are typically flat, with a strong central peak in a few galaxies, interpreted as recent star formation activity in the core. Stellar population maps for the SAURON early-type spiral sample are analysed in \\citet{pel07-445}.  \\begin{figure*} \\includegraphics[viewport=0mm 0mm 278mm 93mm,width=165mm,angle=0,clip]{figures/thumbs.eps} \\caption{Thumbnail images from the Digitized Sky Survey, showing a 100 arcsec square field around the target galaxies. The boxes correspond to the 27$\\times$27 arcsec VIMOS field of view, underlining how well matched our chosen galaxies are to the instrument.} \\label{fig:targets} \\end{figure*}  \\begin{table*} \\centering \\caption{Parameters of the target galaxies in Abell 3389. ID and morphology from \\citet[][D80]{dre80-565}. Positions from 2MASS. $R$-band magnitude from NED; $K$-band magnitude from 2MASS. Effective circular radius ($r_{\\rm e}$) calculated from 2MASS $J$-band images. Central velocity dispersion, $\\sigma$ (km s$^{-1}$), as measured in Section \\ref{sec:sigma} using the $r_{\\rm e}$/3 aperture.} \\label{tab:targets} \\begin{tabular}{@{}clccrrrlr} \\hline ID & \\multicolumn{1}{c}{NGC} & RA & Dec & \\multicolumn{1}{c}{m$_R$} & \\multicolumn{1}{c}{m$_K$} & \\multicolumn{1}{c}{$r_{\\rm e}$} & \\multicolumn{1}{c}{Morph.} & \\multicolumn{1}{c}{$\\sigma$} \\\\ (D80) & \\multicolumn{1}{c}{\\#} & (J2000) & (J2000) & \\multicolumn{1}{c}{mag} & \\multicolumn{1}{c}{mag} & arcsec & \\multicolumn{1}{c}{(D80)} & \\multicolumn{1}{c}{(km s$^{-1}$)} \\\\ \\hline D40 & & 06 21 59.7 & --64 59 19 & 14.2 & 11.16 & 4.2 & S0 & 178.8 $\\pm$ 6.1 \\\\ D41 & 2233 & 06 21 40.1 & --65 02 00 & 13.5 & 10.47 & 4.6 & S0 & 184.6 $\\pm$ 2.4 \\\\ D42 & 2230 & 06 21 27.5 & --64 59 34 & 12.6 & 9.66 & 10.5 & E & 253.9 $\\pm$ 3.1 \\\\ D43 & & 06 21 13.9 & --65 00 58 & 14.6 & 11.58 & 3.1 & S0 & 149.8 $\\pm$ 4.4 \\\\ D44 & & 06 20 54.4 & --65 01 08 & 14.0 & 11.13 & 3.8 & S0 & 221.5 $\\pm$ 4.5 \\\\ D48 & & 06 23 49.0 & --64 57 15 & 14.8 & 11.94 & 2.9 & E & 144.6 $\\pm$ 2.3 \\\\ D49 & & 06 23 07.4 & --64 55 50 & 14.7 & 11.64 & 2.3 & E & 180.9 $\\pm$ 3.9 \\\\ D52 & 2235 & 06 22 22.1 & --64 56 03 & 12.4 & 9.35 & 12.3 & D & 252.1 $\\pm$ 2.4 \\\\ D53 & & 06 22 04.8 & --64 57 36 & 15.1 & 12.18 & 2.8 & E & 113.5 $\\pm$ 4.4 \\\\ D60 & & 06 22 19.8 & --64 14 04 & 13.5 & 10.58 & 4.5 & E & 194.0 $\\pm$ 1.9 \\\\ D61 & & 06 21 58.3 & --64 53 44 & 15.3 & 12.37 & 2.1 & S0 & 146.6 $\\pm$ 6.7 \\\\ D89 & 2228 & 06 21 15.6 & --64 27 32 & 12.8 & 9.79 & 9.3 & SB0 & 247.7 $\\pm$ 2.8 \\\\ \\hline \\end{tabular} \\end{table*}  Here, we assess the feasibility of using the VLT VIsible MultiObject Spectrograph (VIMOS) IFU for studying radial trends of stellar populations by observing a small sample of early-type galaxies in the nearby cluster Abell 3389 (z $\\approx$ 0.027). With a large future study in mind, we develop the observing strategy and analysis techniques, comparing our results with those of previous work. The paper is organised as follows: Section \\ref{sec:data} describes the sample, observations and data reduction; Section \\ref{sec:analysis} details the measurement of the gradients in the absorption lines and in the stellar populations; Section \\ref{sec:sp_grads} analyses the stellar population gradients; Section \\ref{sec:conclusion} highlights our main conclusions. ",
        "conclusions": "\\label{sec:conclusion}  Using the VIMOS integral field unit, we observed a sample of early-type galaxies in the local cluster Abell 3389. From the resulting data-cube, we measured 22 line indices in radial bins out to the effective radius, and successfully derived stellar population gradients. Our key findings are that early-type galaxies have:  \\begin{itemize} \\item strong negative gradients in metallicity ($\\sim$ --0.2 dex per decade in radius); \\item radial trends in age and $\\alpha$-abundance that are consistent with zero; \\item stellar population gradients not significantly correlated with the central velocity dispersion or $K$-band luminosity (proxies for total galaxy and stellar mass respectively); \\item a strong anti-correlation between the central values and gradients of both age and [Z/H]. \\end{itemize}  In a future paper we will report on an additional sample of 19 early-type galaxies in the Shapley Supercluster, observed with the VIMOS IFU."
    },
    "0807/0807.3290_arXiv.txt": {
        "abstract": "{Observations with space-borne X-ray telescopes revealed the existence of soft, diffuse X-ray emission from the inner regions of planetary nebulae. Although the existing images support the idea that this emission arises from the hot shocked central-star wind which fills the inner cavity of a planetary nebula, existing models have difficulties to explain the observations consistently. } {We investigate how the inclusion of thermal conduction changes the physical parameters of the hot shocked wind gas and the amount of X-ray emission predicted by time-dependent hydrodynamical models of planetary nebulae {with central stars of normal, hydrogen-rich surface composition.} } {We upgraded our 1D hydrodynamics code NEBEL by to account for energy transfer due to heat conduction, which is of importance at the interface separating the hot shocked wind gas (`hot bubble') from the much cooler nebular material. With this new version of NEBEL we recomputed a selection of our already existing hydrodynamical sequences and obtained synthetic X-ray spectra for representative models along the evolutionary tracks by means of the freely available CHIANTI package. } {Heat conduction leads to lower temperatures and higher densities within a bubble and brings the physical properties of the X-ray emitting domain into close agreement with the values derived from observations. The amount of X-rays emitted during the course of evolution depends on the energy dumped into the bubble by the fast stellar wind, on the efficiency of `evaporating' cool nebular gas via heat conduction, and on the bubble's expansion rate. We find {from our models} that the X-ray luminosity of a planetary nebula increases during its evolution across the HR diagram until stellar luminosity and wind power decline. Depending on the central-star mass and the evolutionary phase, our models predict X-ray [0.45--2.5 keV] luminosities between $10^{-8}$ and $10^{-4}$ of the stellar bolometric luminosities, in good agreement with the observations. Less than 1\\% of the wind power is radiated away in this X-ray band. Although temperature, density, and also the mass of the hot bubble is significantly altered by heat conduction, the dynamics of the whole system remains practically the same. } {Heat conduction allows the construction of nebular models which predict the correct amount of X-ray emission and at the same time are fully consistent with the observed mass-loss rate \\emph{and} wind speed. Thermal conduction must be considered as a viable physical process for explaining the diffuse X-ray emission from planetary nebulae with closed inner cavities. Magnetic fields must then be absent or extremely weak. } ",
        "introduction": "The modern, very successful concept for the formation and evolution of the main structures of Planetary Nebulae (PNe) is based on the dynamical effects caused by (i) the interaction of a rapidly varying central-star stellar wind with the slow AGB wind ejected earlier, and (ii) the heating of this circumstellar material by photoionization.    The wind from the central star is very fast, exceeding in most cases 1000~\\kms, and passes through a shock before making contact with the dense slow AGB wind.  A tenuous but very hot ($\\ga\\! 10^7~\\mbox{K}$) `bubble' is formed since $9/16$ of the kinetic wind energy is converted into internal energy behind a strong adiabatic shock. The bubble is separated from the outer, much cooler nebular gas by a contact discontinuity (or contact surface). The bubble's thermal expansion accelerates the inner layers of the PN. For recent reviews see \\citet{SchSt.03} and \\citet{StSch.06}.  The shocked wind gas including the thin transition regime towards the PN proper are the likely sites of diffuse X-ray emission, mainly by thermal bremsstrahlung and line emission \\citep{VK.85}. Any positive detection of diffuse X-ray emission coming from the inner cavities enclosed by nebular shells is a direct confirmation of the wind interaction scenario as sketched above. Thus there was considerable interest to observe selected PNe with X-ray satellites once they became available.  Since the surface brightness in X-rays is expected to be quite small \\citep[see][]{VK.85}, only very few positive detections have been reported to date.   Summaries of X-ray observations of PNe conducted to date with the different satellites are presented in \\citet{Chuetal.03}, \\citet{G.06}, and \\citet{Kast.07}.  Although the observational confirmation of the existence of shocked wind material was very gratifying, the properties of the X-ray emitting gas were disturbing.  It turned out that the X-ray spectra are rather soft, indicating temperatures of only $\\approx\\! (1\\ldots 3)\\times10^6~\\mbox{K}$.  The electron densities of the emitting volumes vary between about 20 and 200~\\ccm.  These results are in sharp contrast with theoretical expectations since the \\emph{observed} wind velocities and mass-loss rates demand bubble temperatures of $\\approx\\!\\!10^7\\dots 10^8~\\mbox{K}$, and electron densities far \\emph{below} the observed ones.  \\citet{ASB.06} and \\citet{ASBB.07} proposed that the remains of the much slower wind blown during the early PN evolution (`early wind') when heat conduction is unimportant are responsible for the properties of the observed X-ray emission.  Also collimated outflows (jets) can play a role. So it appears that several physical mechanisms exist (heat conduction and/or mixing, early wind, jets) which all may contribute to the X-ray emission, as pointed out by \\citet{SoKa.03}. The jet-wind interaction has recently been studied in detail by \\citet{AMS.08}.  In two fundamental papers on the properties of interstellar bubbles by \\citet{CCW.75} and \\citet{Wetal.77} it was demonstrated that heat conduction by electrons across the contact surface which separates the hot shocked wind gas from the much cooler swept-up matter is a natural mechanism to account for the observed X-ray and EUV line emissions. Heat conduction enforces `evaporation' of cool gas into the hot bubble, leading to a shell of matter at the bubble's outer edge with properties just ideal for explaining the observed X-ray and EUV line emission.  Since planetary nebulae are virtually scaled down versions of the interstellar bubbles (or \\textsc{H\\,ii} regions), heat conduction should be important for them as well.  To our knowledge, \\citet{So.94} was the first to look into the consequences of heat conduction in PNe.  He concluded that the observed X-ray emission must come from the heat conduction front, and not from the hot bubble as a whole. Heat conduction is also capable of explaining the unusually strong UV lines (e.g. of \\ion{O}{vi}) seen in some PNe, which cannot be produced by regular photoionization. \\citet{So.94} noted also that weak tangential magnetic fields would suppress heat conduction very effectively, and hence also the X-ray emission. But even without magnetic fields X-rays from old, large PNe may escape detection because the X-ray surface brightness falls below the detection limit of existing X-ray satellites. We note in passing that the same considerations apply to wind-blown bubbles around massive stars as well \\citep[cf.][]{WWW.94}.  Despite of their potential importance, quantitative theoretical calculations of X-ray and EUV emissions from PNe based on the concept of heat conduction are rare. \\citet{ZP.96,ZP.98} succeeded in developing analytical solutions based on the concept of \\citet{Wetal.77} and were able to present quantitative predictions for a few specified cases.  Using  analytical solutions of heat conduction, \\cite{GChuG.04} modelled the transition layers between the hot bubble and the nebular regime for \\object{NGC~6543} in order to compute the \\textsc{O\\,vi} emission lines seen in the FUSE spectra.  Two hydrodynamical studies addressed recently the issue of X-ray emission from planetary nebulae \\citep{SS.06,ASBB.07}.  The basic philosophy of both studies is to avoid thermal conduction and to select the properties of the wind such that temperature and density of the X-ray emitting gas agree with the observations. No relation between stellar parameters and wind properties is considered, and the feedback of radiation to the hydrodynamics is neglected. To explain the observed X-ray luminosities, these authors came up with wind parameters which disagree strongly with both the predictions of the theory of radiation-driven winds \\citep[e.g.][]{Pauletal.88} \\emph{and} with the directly observed mass-loss rates and wind speeds of the objects in question. More precisely, the necessary mass-loss rates are too large and the corresponding wind speeds too low -- by as much as a factor 5.% \\footnote{The only exception is the central star of NGC~2392 whose wind speed is much too low for its position in the Hertzsprung-Russell diagram.}  The results of \\citet{ASBB.07} are particularly interesting.  Consistence with existing observations in terms of X-ray luminosity and temperature of the X-ray emitting gas could be achieved only for rapidly decaying central-star winds: while the wind speed increases linearly with time, the mass-loss rate  must decline quite rapidly to very small values \\citep[see][Figs. 6 and 8 therein]{ASBB.07}. For the best choice of \\citet{ASBB.07}, the central-star wind reaches $10^{-10}$ \\Mdot\\ already at a modest wind speed of only 750 \\kms.  We iterate that combinations of mass-loss rates and outflow speeds as they are found in \\citet{ASBB.07} are in severe conflict with both current theories of radiation-driven winds from hot stars (see Fig.~\\ref{mod.prop} in Sect. \\ref{bubb.struc} for details) \\emph{and} with the observations (see Table~\\ref{tab.xray} in Sect. \\ref{xray.lum}), with the apparent exception of NGC~2392. In our opinion the existing studies of the X-ray emission from planetary nebulae which neglect thermal conduction by electrons are therefore far from being convincing. After all, thermal conduction is a physical process inherent to all hydrodynamical systems and becomes important in rarefied plasmas wherever the mean free path of the electrons is large enough. Thermal conduction can \\emph{only} be modified or even suppressed by the existence of magnetic fields \\citep[cf.][]{BBF.90}.  Since no self-consistent radiation-hydrodynamics computations including heat conduction have been performed to date\\footnote{ \\cite{MF.95} computed the X-ray emission of their 2D hydrodynamics models, but heat conduction was not considered.}, and urged by the fact that the new observations which became available by the Chandra and XMM-Newton satellites lack a convincing interpretation by detailed modeling, we decided to update our 1D radiation-hydrodynamics code NEBEL by incorporating a thermal conduction module.  We recomputed then some of our sequences from \\citet[Paper~I hereinafter]{Peretal.04} with heat conduction self-consistently included.  Based on the new thermal structure of the bubble we modelled the X-ray and EUV emissions at selected positions along the evolutionary sequences by using the well-documented CHIANTI code.  We emphasize that in our study (i) wind and stellar evolution are consistently connected within the framework of the theory of radiation-driven winds, and (ii) no fit to the data by manipulating the models is done.  It is our belief that only with such an approach it is possible to interprete the data appropriately and to deduce meaningful constraints regarding the relevant physics. Obviously, our approach is only applicable to objects which do not depart too much from sphericity, have closed shells where a bubble of shocked, hot wind gas is able to exist, and have central stars of normal, hydrogen-rich surface composition.  We avoid to discuss in the present work objects with [WC] central stars for two reasons: (i) Their evolution is at present not known, and (ii) the physics of heat conduction in hydrogen-free and carbon-rich plasmas has still to be worked out.  We note that the X-ray emission from a hydrogen-free plasma is expected to be stronger, making it more likely to detect extended X-ray emission from PNe with [WC] central stars. The present number of positive detections is, however, too small as to make any definitive statements concerning the ratio of PNe X-ray sources with normal and [WC] central stars, respectively.  The computational details of how the thermal conduction is treated are described in the Sect.~\\ref{computate}, followed by a discussion of the resulting bubble structures in Sect.~\\ref{bubb.struc}. Section~\\ref{x-rays} is devoted to the X-ray emission emerging from the conduction front and how it develops with time. In Sect.~\\ref{comp.obs} we discuss extensively how our models compare with the existing observations. The paper concludes with Sect.~\\ref{diss.concl}. Part of the results presented here can be found in \\citet{SSW.06}. ",
        "conclusions": "\\label{diss.concl}  We presented a detailed numerical approach towards an understanding of the diffuse soft X-ray emission from planetary nebulae based on the concept of thermal conduction. Since thermal conduction is a physical process inherent to all hydrodynamical systems, becoming important wherever the mean free path of the electrons is sufficiently large, we included a thermal conduction module into our 1D radiation-hydrodynamics code. We were able to compute the thermal structure of the shocked wind gas inside the nebular cavity self-consistently with the hydrodynamics once the stellar AGB remnant, its initial circumstellar envelope, and the post-AGB wind model are specified. By doing so, there is some freedom in the treatment of heat conduction in cases where the mean free path of the electrons becomes comparable to the characteristic temperature scale length. Magnetic fields that may play a role for the shaping of PNe are not considered.  Thermal conduction has a twofold effect favoring thermal X-ray emission from the shocked wind gas: it lowers the temperature gradient across the bubble/nebula interface and heats the cool nebular matter, forcing it to `evaporate' inwards.  With time, heat conduction accumulates additional matter in the bubble with a characteristic temperature of some $10^{6}$ K, which quickly dominates the bubble's mass budget. The amount of added matter is controlled by heat transfer across the bubble from the inverse wind shock towards the nebula. The bubble mass may increase or decrease with time, depending on the evolution of the stellar wind power.  We selected several of the hydrodynamic sequences presented in \\citetalias{Peretal.04} and recomputed them, without changing any of the other parameters or boundary conditions, with our heat conduction treatment included. The X-ray emission was computed post-facto by employing the CHIANTI code, slightly adapted to match our purposes. We were able to find good agreement with existing observations, both with respect to X-ray luminosity and surface brightness distribution, if {method~2} is used for the treatment of heat conduction. At the same time, the stellar wind model employed here, based on the theory of radiation-driven winds, implies wind speeds in excess of 1000 \\kms\\ for typical nebular models (cf. Fig.~\\ref{xray.wind}). We conclude that our modeling is fully consistent with both the observed wind speeds of more than 1000 \\kms\\ \\emph{and} the observational evidence that the soft X-ray emission comes from regions with electron temperatures of about $2\\times 10^6$~K (see also Table~\\ref{tab.xray}). In contrast, our models without heat conduction fail to reproduce the characteristics of the observed X-ray emission by large amounts if realistic wind parameters are assumed.  The basic findings from our modeling can be summarized as follows: \\begin{itemize} \\item  The energy budget of the hot bubble is substantially altered by radiative losses at the bubble/nebula interface, both in models with and without thermal conduction. \\item  Heat conduction has a strong influence on the thermal structure of the hot bubble, but leaves the dynamics of the whole system is virtually unchanged, i.e.\\ the shaping of the nebular shells and their expansion properties are not affected. \\item  In the absence of heat conduction, the temperature of the hot bubble depends only on the velocity of the central star wind. The resulting X-ray temperature is much too high, and the X-ray luminosity much too low, to explain the existing observations. \\item  If heat conduction is substantial, the bubble temperature is a function of the stellar wind power and the bubble size. The compact nebulae around massive, short-lived central stars are therefore expected to have hotter bubbles than the  nebulae of low-mass, slowly evolving central stars. \\item  The X-ray luminosity is determined by the total emission measure, which decreases by expansion and increases by `evaporation' from the main nebula. The competition between expansion \\emph{and} `evaporation' rules the temporal evolution of the X-ray luminosity. \\item  According to our PN simulations with a time-dependent wind model, the X-ray luminosities increase with time during the main phase of evolution when the wind power increases with time, in contrast to the case of a constant wind power that would imply a decrease of the X-ray luminosity with time. \\item  For the energy range usually considered for planetary nebulae, the X-ray luminosity is below about 1\\% of the stellar wind power, and between $10^{-4}$ and $10^{-8}$ of the stellar bolometric luminosity. The exact numbers depend on the actual evolutionary state and the wind model used. \\item  For given effective temperature, the X-ray emission is largest for nebulae around massive central stars even though their small bubbles contain only little mass. The X-ray emission measure of their bubbles is nevertheless large, primarily because the high electron densities overcompensate the small amount of X-ray emitting gas that massive central stars can accumulate during their short lifetimes. \\item  In contrast to the X-ray luminosity, the emission of the UV \\ion{O}{vi} lines is not very sensitive to the effects of thermal conduction. \\item  Magnetic fields must be absent or extremely weak in all objects with diffuse X-ray emission since their presence would strongly depress thermal conduction and hence `evaporation'. The absence or weakness of magnetic fields implies also that they cannot be responsible for shaping these objects. \\end{itemize}  Our numerical treatment as described here leads to X-ray luminosities that are substantially below those found by \\citet{ZP.98, ZP.96} with their analytical approach, although their wind model and central-star evolution are very similar. We believe that their analytical approach overestimates the evaporation rate because radiative cooling of the gas at the conduction front by line emission is not considered (see also Sect.~\\ref{bubble.time}).  It is likely that the bubble/PN interface becomes dynamically unstable, leading to direct mixing between hot bubble and cool nebular matter \\citep[cf.][]{SS.06}. The net effect would be similar to heat conduction, i.e.\\ a reduction of temperature gradients allowing a more efficient X-rays emission, as already mentioned by \\citet{Chuetal.97}. The 2D simulations of a spherical bubble performed by \\citet{SS.06} suggest that the mixing region is confined to a rather thin shell at the surface of the bubble.  If so, the X-ray emission would be limb brightened, very similar to our models computed according to {method~1}. For the time being, the existing observations do not seem to be consistent with such a limb brightening.  \\citet{Georgetal.08} found recently that the wind of the central star of \\object{NGC~6543} is less depleted in iron compared to the plasma emitting the diffuse X-rays \\citep{Georgetal.06}. They concluded that the X-ray emitting plasma ``is derived from nebular gas rather than the stellar wind\".  This finding is fully consisting with the heat-conduction models. Clearly, more studies of this kind would be very rewarding.  Finally we want to emphasise that the models introduced in this work are entirely based on normal chemical compositions in the stellar and circumstellar envelopes. The results presented here should therefore \\emph{not} be used to interprete X-ray emission from objects with hydrogen-deficient central stars such as \\object{BD+30\\degr 3639}, \\object{NGC~40} because their evolution is different and not yet understood. Additionally, one has to deal with heat conduction and X-ray emission in a practically hydrogen-free plasma."
    },
    "0807/0807.3773_arXiv.txt": {
        "abstract": "H.E.S.S. observed TeV blazar PKS 2155--304 in a strong flare state in 2006 July. The TeV flux varied on timescale as short as a few minutes, which sets strong constraints on the properties of the emission region. By use of the synchrotron self-Compton model, we found that models with the bulk Lorentz factor $\\sim 100$, the size of the emission region $\\sim 10^{15}$ cm, and magnetic field $\\sim 0.1$ G explain the observed spectral energy distribution and the flare timescale $\\sim$ a few minutes. This model with a large value of $\\Gamma$ accounts for the emission spectrum not only in the TeV band but also in the X-ray band. The major cooling process of electrons/positrons in the jet is inverse Compton scattering off synchrotron photons. The energy content of the jet is highly dominated by particle kinetic energy over magnetic energy. ",
        "introduction": "In active galactic nuclei various physical processes such as particle acceleration, emission of high energy photons, and so on, take place. Accretion of matter onto the central black holes is thought to be the energy source of these processes. When a fraction of matter is ejected from the accretion disks, relativistic jets of plasmas are formed. The relativistic jets are thought to explain intense and variable emission from blazars. The beaming effect amplifies the radiation from jets depending on the opening angle of the jets, bulk Lorentz factor, $\\Gamma$, and the angle between the line of sight and the jet axis, $\\theta$. Some blazars are known to emit very high energy $\\gamma$-rays in the TeV energy band. They are, for example, PKS 2155--304, Mrk 421, and Mrk 501 \\citep[e.g.,][]{wag07}.  Recent observations of TeV $\\gamma$-rays have found remarkably short time variability of TeV blazars. In 2006 July PKS 2155--304 with redshift $z = 0.116$ showed an outburst of TeV $\\gamma$-rays. The average flux during the outburst was more than 10 times typical values observed from the object \\citep{aha-flare07}.  During this period X-rays were also monitored by {\\it Swift}. In the 0.3--10 keV energy band the X-ray flux increased by a factor of 5 \\citep{fos07}. The observations by the High Energy Stereoscopic System (H.E.S.S.) report that the timescale of variation is a few minutes \\citep{aha-flare07}. The well-resolved burst of the TeV $\\gamma$-ray flux from PKS 2155--304 varied on timescales only $\\sim 200$ s. Recently MAGIC also observed short time variation in Mrk 501 \\citep{alb07}. When the size of the emission region in the comoving frame of the jet is denoted by $R$, the observed timescale of the variability sets a limit on $R$, i.e., $R \\lesssim c t_\\mathrm{var} {\\cal D}/(1+z)$, where $c$ is the speed of light and ${\\cal D} = [\\Gamma (1- \\beta \\cos \\theta)]^{-1}$ is the beaming factor with $\\beta = (1-1/\\Gamma^2)^{1/2}$. \\cite{aha-flare07} argued that ${\\cal D} \\gtrsim$ 60 - 120 $R/R_\\mathrm{S}$ is required to explain $t_\\mathrm{var} \\sim 2$ min, where $R_\\mathrm{S} = 2 G M/c^2$ is the Schwarzschild radius with $G$ and $M \\sim 10^9 M_\\sun$ being the gravitational constant and the central black hole mass, respectively. Such large values of $\\cal D$ are also suggested by \\cite{bgl08} recently, based on the requirements of radiative cooling time and optical depth for $\\gamma$-rays due to electron-positron pair production.  The emission spectra of blazars are characterized by two peaks in the $\\nu$-$\\nu F_\\nu$ representation, where $F_\\nu$ is the differential energy flux. The lower energy peak is located in the optical -- X-ray bands and the higher energy peak is located in very high energy $\\gamma$-ray bands. The lower energy peak is most probably by synchrotron radiation of nonthermal electrons/positrons. Radiation mechanisms of high energy $\\gamma$-rays are thought to be inverse Compton scattering off soft photons in the leptonic models \\citep[e.g.,][]{mar92} and hadronic interaction of relativistic particles and photons in the hadronic models \\citep[e.g.,][]{mb92,mu03}. The soft photons of the leptonic models are supplied by synchrotron radiation by relativistic electrons/positrons in the jet (synchrotron self-Compton model) or by the external sources such as accretion disks \\citep{ds93} and disk radiation scattered around the jet \\citep{sikora94}. These emission models assume that the emission region is one zone.   Observations of TeV $\\gamma$-rays of blazars have revealed that large values of ${\\cal D} \\gtrsim 10$ are required to fit the emission spectra of very high energy $\\gamma$-rays by inverse Compton scattering \\citep[e.g.,][]{ktk02}. On the other hand, the observations of the apparent velocity of VLBI knots show that parsec-scale jets are subrelativistic or at most mildly relativistic \\citep[][]{pe04}. To reconcile the discrepancy between these values of $\\Gamma$, the deceleration of jets is considered. \\cite{gk03} assumed that TeV $\\gamma$-rays are produced by inverse Compton scattering off synchrotron photons that are emitted by decelerated jet components. \\cite{gtc05}, on the other hand, proposed a spine-layer model. In this model they assumed that a fast moving emission region is surrounded by a slow moving sheath. Gamma-rays are then produced by inverse Compton scattering off the photons emitted in the spine and the sheath. Although various emission models of jets with multiple radiation zones have been proposed as mentioned above, a simple synchrotron self-Compton (SSC) model is still worth use in obtaining physical parameters of the emission regions in jets.  In this paper we show that a simple SSC model can explain the TeV $\\gamma$-rays and X-rays of PKS 2155--304. PKS 2155--304 is an interesting source because of its strong TeV $\\gamma$-ray emission and the short time variability of TeV emission. Since the redshift of PKS 2155--304 is 0.116, the absorption of TeV $\\gamma$-rays by extragalactic background light (EBL) is effective \\citep[e.g.,][]{sms06}. This is an ideal object to test EBL models as well as emission models.  In \\S \\ref{sec:model} we describe the parameters of our numerical calculations and in \\S \\ref{sec:results} the values of the parameters are determined by fitting the observed data.  Finally discussion is given in \\S \\ref{sec:sum}. ",
        "conclusions": "\\label{sec:sum}    In this work we have shown that a large value of the bulk Lorentz factor, $\\Gamma \\sim 100$, is required to explain the emission from PKS 2155--304 in the TeV band as well as the X-ray band. With $R = 9.6 \\times 10^{14}$ cm and $\\Gamma = 90$, the timescale $t_\\mathrm{var} \\sim 400$ s is obtained, while $t_\\mathrm{var} \\sim 200$ s is obtained for $R \\sim 5 \\times 10^{14}$ cm and $\\Gamma \\sim  100$ - 150. The value of ${\\cal D} = \\Gamma$ as large as 100 requires the viewing angle as small as 0.01 rad, which makes the case an unlikely coincidence if the opening angle of the jet is the same order as $1/\\Gamma$. However, if the opening angle is larger than $1/\\Gamma$, we can observe a portion of the jet that aligns with the line of sight, thus avoiding an excessively small observability. Considering that Mrk 501 also exhibited a similar short timescale TeV flare, we think that the opening angle is larger than $1/ \\Gamma$ (by a factor $a$). If this is the case, $R$ represents the lateral size along the jet motion, while the transverse size is larger than $R$ by the same factor $a$. Thus the real size of the emission region is larger than $R$ and only a part of the emission regions is observed owing to a strong relativistic beaming effect. Then the kinetic power should be by a factor $a^2$ larger than a conventional estimate given by $\\pi c R^2 \\Gamma^2 u_e$. A further speculative  possibility is that the opening angle is as small as $1/\\Gamma$ but the jet axis wanders within a large angle and the flare is observed when its direction happens to fall in the line of sight.   Our results show that the jet plasma is highly particle energy dominated over magnetic energy, i.e., $u_e/u_B \\sim 10^4$ in the flare state in 2006. The ratio $u_e/u_B$ has been discussed in many papers. In \\citet{ktk02} we discussed that TeV blazars are generally particle dominated typically by a factor of 10. For a quiescent state of PKS 2155--304 we obtained the ratio of 3. Thus the ratio of $10^4$ seems unusually high, although there is no serious theoretical reason to expect a ratio of unity except that it leads to a minimum kinetic power for a given observed spectrum. The obtained high value of the ratio basically results from the fact that the SSC luminosity is much larger than the synchrotron luminosity. In this respect, we note that the X-ray and TeV observations are not strictly simultaneous; if the X-ray flux during the TeV flare is higher than the adopted value in this paper, the ratio decreases correspondingly. Also, some soft photon sources other than synchrotron radiation will help to decrease the ratio.  However, considering a good fit to the observed spectrum obtained here, such possibilities are not easy to realize and the reduction will not be significant. \\citet{lev07} and \\citet{bgl08} adopted the view that the TeV flare is caused by external Compton mechanism.  In this case, the model may avoid an extremely large value of $u_e/u_B$. However, considering the short timescale of radiative deceleration, the acceleration of the emitting plasma to the large bulk Lorentz factor itself becomes very difficult because of the large radiation drag effect. Also note that both papers did not make detailed multi-wavelength fitting.   As mentioned above, \\cite{bgl08} argued that inverse Compton scattering off external soft photons is favored as the dominant radiation process to produce TeV $\\gamma$-rays. The major differences between their model and ours are that they assumed the dominance of the Poynting flux over the kinetic energy flux as well as the synchrotron radiation as a major cooling process to estimate particle's cooling timescale, while in our model the energy flux is dominated by nonthermal particle's kinetic energy and the cooling is dominated by SSC.   In comparison between the flare state in 2006 July and the quiescent state in 2003 July, the powers in radiation and electrons increased by factors 2 and 4.2, respectively, during the flare in 2006, if Model A is adopted. The increase in the jet power is significant and the efficient acceleration of electrons is implied. The size of the blob for the 2006 July flare is $\\sim 10^{15}$ cm. This is about 3 $R_\\mathrm{S}$, if the mass of the central black hole is $10^9 M_\\sun$. A large energy injection into a very compact region occurred in the flare. The obtained parameter values for the flare event and various phases are not well ordered and we found no simple tendencies from quiescent states to flares. The TeV flare may be caused by a concentration of energy into a sub-horizon scale region, while during quiescent states we may be seeing larger scale energy dissipation on longer timescales.   In applying our model, we assumed a low flux model of EBL, i.e., LLL of \\cite{dk05}. With this EBL model the one-zone SSC model is found to be viable for this TeV blazar, although other models proposed by \\cite{gk03} and \\cite{gtc05} might explain the TeV spectrum with higher fluxes of EBL.  This should be studied in future work.   Radiation below $10^{15}$ Hz is not fitted with our model. We assumed that the emission comes from more extended regions far away from the central black hole \\citep[see also][]{aha-quiet-model05}. If the radiation below $10^{15}$ Hz is emitted by the jet plasma, inverse Compton scattering produces radiation in the MeV -- GeV band. Observation in this energy band is particularly important to set constraints on the emission mechanisms of blazars."
    },
    "0807/0807.4365_arXiv.txt": {
        "abstract": "{We investigate the possibility that the decrease in the relative luminosity of the 158~$\\mu$m [C{\\sc \\,ii}] line with the far-infrared luminosity in extragalactic sources stems from a stronger contribution from the heated dust emission in the more distant sources. Due to the flux limited nature of these surveys, the luminosity of the detected objects increases with distance. However, the \\carb\\ luminosity does not climb as steeply as that of the far-infrared, giving the decline in the $L_{\\rm [CII]}/L_{\\rm FIR}$ ratio with $L_{\\rm FIR}$. Investigating this further, we find that the \\carb\\ luminosity exhibits similar drops as measured against the carbon monoxide and radio continuum luminosities. The former may indicate that at higher luminosities a larger fraction of the carbon is locked up in the form of molecules and/or that the CO line radiation also contributes to the cooling, done mainly by the [C{\\sc \\,ii}] line at low luminosities. The latter hints at increased activity in these galaxies at greater distances, so we suggest that, in addition to an underlying heating of the dust by a stellar population, there is also heating of the embedded dusty torus by the ultra-violet emission from the active nucleus, resulting in an excess in the far-infrared emission from the more luminous objects. ",
        "introduction": "The $^2{\\rm P}_{3/2}\\rightarrow^2{\\rm P}_{1/2}$ fine-structure line of C$^+$, \\carb, is believed to be a cooling pathway for the diffuse gas in galaxies \\citep{dm72}. This transition traces photo-dissociation regions (PDRs), where the ultra-violet radiation from young stars dominates the heating of the gas.  From this process, the \\carb\\ luminosity may reach up to 1\\% the total luminosity of the galaxy \\citep{cgtw85,sgg+91,wmb+91}, thus being the most powerful emission line in many galaxies. From a study of 60 normal galaxies with the Long Wavelength Spectrometer (LWS) on-board the Infrared Space Observatory (ISO, \\citealt{caa+96}), \\citet{mkh+01} find that for far-infrared luminosities of $L_{\\rm FIR}\\lapp10^{11}$ \\Lo, the $L_{\\rm [CII]}/L_{\\rm FIR}$ ratio is roughly constant at $\\log_{10}(L_{\\rm [CII]}/L_{\\rm FIR})\\sim-2.5$, but drops rapidly above these values (Fig. \\ref{ratio-fir}). This normalised decrease in the brighter galaxies is attributed to an increased grain charge, lowering the kinetic energy of the liberated photo-electrons which deliver heat to the gas.  \\citet{nocr01} extended this sample to include starburst galaxies and active galactic nuclei (AGN), all of which are found to follow the same trend, suggesting that far-infrared emission in the active galaxies also arises primarily from star forming activity. Rather than an increase in the charge carried by the dust grains, \\citet{nocr01} suggest that the $L_{\\rm [CII]}/L_{\\rm FIR}$ decrease with $L_{\\rm FIR}$ is due to higher gas densities resulting in higher collision rates, thus de-exciting the ionised carbon through a non-radiative process.  \\begin{figure} \\centering \\includegraphics[angle=270,scale=0.75]{ratio-fir.ps} \\caption{The $L_{\\rm [CII]}/L_{\\rm FIR}$ ratio versus $L_{\\rm FIR}$ for the low redshift ($z\\leq0.1296$) galaxies and the high redshift quasar searches (cf. figure 2 of \\citealt{mcc+05}). Throughout this paper the symbols are colour coded according to the source reference -- red \\citep{mkh+01}, blue \\citep{nocr01}, green \\citep{lsf+03} and black (the two high redshift \\carb\\ detections of \\citealt{mcc+05,iye+06}). The downwards arrows show the $3\\sigma$ upper limits. As per \\citet{lsf+03,mcc+05}, the hollow symbols indicate where the $\\approx80\\arcsec$ LWS aperture (\\citealt{caa+96}) subtends $\\lapp10$ kpc and the filled symbols where the aperture subtends $\\gapp10$ kpc (at angular diameter distances of $\\gapp26$ Mpc or $z\\gapp0.006$ -- throughout this paper we use $H_{0}=71$~km~s$^{-1}$~Mpc$^{-1}$, $\\Omega_{\\rm matter}=0.27$ and $\\Omega_{\\Lambda}=0.73,$ \\citealt{svp+03}).} \\label{ratio-fir} \\end{figure} \\citet{lsf+98,lsf+03} confirm the decrease in $L_{\\rm [CII]}/L_{\\rm FIR}$ with $L_{\\rm FIR}$ to higher luminosities ($L_{\\rm FIR}\\gapp10^{12}$ \\Lo), by observing \\carb\\ in a sample of ultra-luminous infrared galaxies (ULIRGs). This trend is in part attributed to much of the far-infrared emission arising from dust-bounded photo-ionised gas which does not contribute to the \\carb\\ emission.  Most recently, the far-infrared luminosities have been taken up another notch by the two high redshift detections of \\carb, which again follow the same decline in $L_{\\rm [CII]}/L_{\\rm FIR}$ \\citep{mcc+05,iye+06}. Once more, this indicates different excitation conditions than in local galaxies \\citep{iye+06} and possibly extremely high star formation rates ($\\sim3000$ \\Mo\\ yr$^{-1}$, \\citealt{mcc+05}) in the early Universe. From all of these studies (summarised in Fig. \\ref{ratio-fir}), there is no doubt that the relative strength of the \\carb\\ line drops with far-infrared luminosity. However, such flux limited surveys are subject to a selection effect, where only the brighter sources are detected at large distances.  It is therefore possible that the \\carb\\ deficit is caused by changing demographics of the galaxies at larger distances, a possibility we investigate in this paper. ",
        "conclusions": "In addition to the well documented drop in the $L_{\\rm [CII]}/L_{\\rm FIR}$ ratio with far-infrared luminosity in extragalactic sources, we find similar decreases with the molecular gas and radio continuum luminosities. This indicates that there is a \\carb\\ deficit in relation to each of these properties, which due to the flux limited nature of the surveys, suggests a relative \\carb\\ decline with luminosity. If evolutionary in nature, the order of magnitude decrease in the mean $L_{\\rm [CII]}/L_{\\rm FIR}$ ratio over the past 12\\% of the history of the Universe, could be due to a decrease in the metallicities, although, as per some of the literature, we see no evidence of this. The decline in $L_{\\rm [CII]}/L_{\\rm FIR}$ is, however, dominated by the ULIRGs at redshifts of $z\\sim0.1$ as well as the two QSOs, at $z=4.69$ and $6.42$, and so rather than a detectable evolutionary effect, the decreasing $L_{\\rm [CII]}/L_{\\rm FIR}$ ratio is more likely the cause of a change in the demographics of the objects which can be detected at these distances.  We suggest that the excess FIR and radio luminosities arise from additional AGN activity, where the former is the result of dust in the embedded circumnuclear torus being heated by ultra-violet photons, in addition to the underlying ultra-violet emission from the stellar population.  Both \\citet{nocr01} and \\citet{lsf+03} also advocate non-PDR mechanisms as being responsible for some of the far-infrared emission.  Furthermore, if the radio emission was due to the same ionised gas as traced by the \\carb\\ emission, we would not expect a relative decrease in $L_{\\rm [CII]}$ with $L_{\\rm radio}$. This may also be indicative of an increasing AGN contribution to the luminosity of these objects and towards the high end (i.e. the ULIRGs of \\citealt{lsf+03}), these sources would be considered radio galaxies.  Whether caused by an embedded AGN or vigorous star forming activity, the heating of the dust is the most likely explanation for the decrease in the \\carb\\ luminosity in relation to that of the FIR. So although the gas is not expected to be heated by quite the same extent as the dust, the changing galaxy types, as traced by the increasing luminosities, may imply different cooling mechanisms. One possibility is an increase in the relative contribution of the cooling by the [O{\\sc \\,i}] line, which is known to become more dominant at higher ultra-violet fluxes \\citep{kwhl99}, thus tracing the warmer dust \\citep{mkh+01}. Furthermore, like the molecular gas and radio continuum luminosities, in comparison with the FIR, the [O{\\sc \\,i}] also exhibits an excess over the \\carb\\ luminosity, although this is also depleted for the ULIRGs. An additional coolant may therefore be the CO emission, in which the higher rotational transitions are found to rival the cooling by the \\carb\\ line in one of the ULIRGs of the sample \\citep{piv07}, with the two high redshift endpoints (for which these transitions have also been observed) exhibiting no FIR excess in relation to the CO.  We may therefore expect the warm molecular gas to also be located in the torus from which the additional FIR luminosity is arising, with the cooler gas, as traced by the low excitation rotational transitions, located at larger radii. Beyond the torus, in the main galactic disk, is also where most of the cool neutral gas, as traced by the \\HI\\ 21-cm absorption, is believed to reside (\\citealt{cww08l} and references therein).  Observations of the higher rotational CO transitions in the ULRIGs could determine whether these could contribute to the cooling budget in these extreme FIR environments (due to starburst/AGN activity), where the \\carb\\ emission is apparently lacking. As well as this, our suggestion that the relative decline in \\carb\\ with FIR is due to a changing demographic, and any associated evolutionary effects, could be further tested by:  \\begin{enumerate} \\item Radio flux measurements of the two high redshift sources. At $z=4.69$ and $6.42$ the 1.4 GHz continuum flux would be redshifted to $250$ and $191$ MHz, respectively. Both of these frequencies are accessible by the Giant Metrewave Radio Telescope, although such low frequencies could be subject to severe interference.  \\item $[$O{\\sc \\,i}$]$ observations of these quasars in order to verify that the excess in this line over the \\carb\\ line extends beyond the local ($z\\lapp0.13$) galaxies.  \\item Confirming the $L_{\\rm [CII]}/L_{\\rm FIR}\\lapp10^{-4}$ limit in PSS 2322+1944 at $z=4.12$, as referred to in \\citet{mcc+05} [Benford et al., in prep.], but not since published.  \\item Further observations of the \\carb\\ transition at $0.1296<z<4.69$, filling in the redshift gap in the $L_{\\rm [CII]}/L_{\\rm FIR}$--luminosity distance distribution (Fig. \\ref{ratio-lbt}). At $z\\gapp1$, sub-millimetre observations would also cover the redshift range where star formation is expected to be most prevalent \\citep{pf95,llhc96}. However, at $\\gapp400$~GHz these observations are difficult, requiring the very best atmospheric conditions.  \\end{enumerate}"
    },
    "0807/0807.0940_arXiv.txt": {
        "abstract": "We carried out a comprehensive far-ultraviolet (UV) survey of $^{12}$CO and H$_2$ column densities along diffuse molecular Galactic sight lines in order to explore in detail the relationship between CO and H$_2$. For this survey we measured new CO abundances from absorption bands detected in {\\it Hubble Space Telescope} spectra for 62 sight lines, and new H$_2$ abundances from absorption bands in {\\it Far Ultraviolet Spectroscopy Explorer} data for 58 sight lines. In addition, high-resolution optical data were obtained at the McDonald and European Southern Observatories, yielding new abundances for CH, CH$^+$, and CN along 42 sight lines to aid in interpreting the CO results. A plot of log $N$(CO) versus log $N$(H$_2$) shows that two power-law relationships are needed for a good fit of the entire sample, with a break located at log $N$(CO, cm$^{-2}$) = 14.1 and log $N$(H$_2$) = 20.4, corresponding to a change in production route for CO in higher-density gas. Similar logarithmic plots among all five diatomic molecules allow us to probe their relationships, revealing additional examples of dual slopes in the cases of CO versus CH (break at log $N$ = 14.1, 13.0), CH$^+$ versus H$_2$ (13.1, 20.3), and CH$^+$ versus CO (13.2, 14.1). These breaks are all in excellent agreement with each other, confirming the break in the CO versus H$_2$ relationship, as well as the one-to-one correspondence between CH and H$_2$ abundances. Our new sight lines were selected according to detectable amounts of CO in their spectra and they provide information on both lower-density ($\\leq$ 100 cm$^{-3}$) and higher-density diffuse clouds. The CO versus H$_2$ correlation and its intrinsic width are shown to be empirically related to the changing total gas density among the sight lines of the sample. We employ both analytical and numerical chemical schemes in order to derive details of the molecular environments. In the denser gas, where C$_2$ and CN molecules also reside, reactions involving C$^+$ and OH are the dominant factor leading to CO formation via equilibrium chemistry. In the low-density gas, where equilibrium-chemistry studies have failed to reproduce the abundance of CH$^+$, our numerical analysis shows that nonequilibrium chemistry must be employed for correctly predicting the abundances of both CH$^+$ and CO. ",
        "introduction": "The most abundant molecule in the Cosmos, H$_2$, has no permanent dipole moment and is thus lacking permitted pure rotation and vibration-rotation transitions. On the other hand, rotational transitions of CO, such as $J$~= 1$-$0 at 115 GHz, have been routinely and extensively observed in molecular clouds, which are also too cold for detection of excited levels of H$_2$ (\\citealt{Papad02} and references therein). Radio telescopes are thus used to map CO emission in the interstellar medium (ISM) in order to delineate global distributions of molecular clouds in our Galaxy and in other galaxies.  It is an empirical and theoretical foundation of radio mapping that the velocity-integrated emission intensity of CO ($W_{\\rm CO}$, in units of K km s$^{-1}$) from a molecular cloud is proportional to the total virialized mass of the cloud, and hence to its hydrogen content \\citep{Lars81, Youn91}. Radio astronomers thus utilize $W_{\\rm CO}$ as a proxy for H$_2$, by employing the ``$X$-factor'' $X_{\\rm CO}$ = $N$(H$_2$)/$W_{\\rm CO}(1-0)$, where $N$ is the observed column density and $X_{\\rm CO}$ will be assumed hereafter to be in units of 10$^{20}$ cm$^{-2}$ (K km s$^{-1}$)$^{-1}$.  An average value is ascribed to giant molecular clouds (GMCs) in the Milky Way, $X_{\\rm CO}$ $\\approx$ 4 \\citep{Youn82, Dick86}. \\citet{Polk88} found significant millimeter-wave emission from CO that was not associated with GMCs. This lowers the value of $X_{\\rm CO}$, having a mean value in the solar neighborhood of 1.8 $\\pm$ 0.3 \\citep{Dame01}. A dependence of $X_{\\rm CO}$ on the metallicity (primarily C/H) of the gas has been found in, e.g., studies of $\\gamma$-ray emission across the Galaxy. This intensity is the product of interactions between cosmic rays and ``stationary'' gas; thus the $\\gamma$-ray intensity is proportional to the amount of gas along the line of sight. \\citet{Stro04} found that $X_{\\rm CO}$ varies between 0.4 and 10.0 from the inner Galaxy to its outer regions, indicating lower gas metallicity in the outer Galaxy.  The GMCs, as well as the smaller dark clouds, are opaque enough that the CO in their cores is not dissociated by the interstellar far-ultraviolet (far-UV) radiation field. Consequently, CO is unobservable in the UV, necessitating its detection via millimeter-wave emission. On the other hand, diffuse molecular clouds (as well as envelopes of dark clouds) have visual extinctions lower than 5 mag, enabling a direct determination of column densities via line absorption in the UV. Both CO and H$_2$ are photodissociated by far-UV radiation, resulting in a variable $X_{\\rm CO}$ that depends on the efficiency of self shielding (as well as mutual shielding) of the two species \\citep{vanD88}, and thus on their column densities. One of our goals here is to study the behavior of $X_{\\rm CO}$ under diffuse-ISM conditions.  The \\citet{Fede80} study of diffuse interstellar clouds showed that there is an approximate quadratic relationship between CO and H$_2$, such that $N$(CO) $\\propto$ [$N$(H$_2$)]$^B$, with $B$ $\\approx$ 2. Furthermore, \\citet{Fede80} remarked that for a group of lower-$N$ sight lines, a shallower slope with $B$ $\\approx$ 1.5 was more appropriate, thus signaling the possibility of $B$ varying with $N$. Indeed, for sight lines with log $N$(CO) $\\leq$ 16, or values higher than those that were available to \\citet{Fede80}, an even steeper relationship with $B$ $\\approx$ 3 was found in the study of \\citet{Pan05}. Both \\citet{Burg07} and \\citet{Sonn07} confirmed the steepening of the slope near log $N$(CO) $\\approx$ 15, attributing it to self-shielding of CO. Thus a variable power law in CO versus H$_2$ shows that the abundance of CO relative to H$_2$, and hence $X_{\\rm CO}$, are increasing with $N$(H$_2$) along diffuse molecular sight lines. In this paper we explore the CO versus H$_2$ relationship in more detail, namely, the trend of CO (as well as of CH, CH$^+$, and CN) versus H$_2$ in diffuse molecular clouds and the dependence of these correlations on $N$(H$_2$) and on physical parameters, such as the total gas (hydrogen) density, $n_{\\rm H}$.  Our study differs from the recent work of \\citet{Burg07} and \\citet{Sonn07} in a number of ways. First, unlike \\citeauthor{Burg07} and like \\citeauthor{Sonn07}, our methodology is based on profile fitting of the data, with detailed decomposition into cloud component structures \\citep{Shef07}. We do not consider apparent optical depth nor curve of growth treatments, which are always less preferred to spectrum synthesis by profile fitting \\citep{Sonn07}, nor do we build a grid of models to look for a solution with a single effective $b$-value \\citep{Burg07}. Second, we follow up on these measurements of $N$ values with two methods of chemical analysis, analytical and numerical, in order to derive $n_{\\rm H}$ at the sites where CO is detected. Third, we are able to discern two regimes of CO formation in terms of $n_{\\rm H}$. CO is associated with the similarly heavy diatomic molecules C$_2$ and CN inside denser and colder clumps of gas \\citep{Fede94, Pan05, Sonn07}, whereas in low-density clouds CO is related to the formation and chemistry of CH$^+$ \\citep{Zsar03}. Remarkably, this transition in the photochemistry of CO will be shown ($\\S$ 3) to affect also the trends of other correlations among the diatomic molecules analyzed here.  In $\\S$ 2 we detail our sources for data and our methods of reduction and analysis. Next, in $\\S$ 3 observational results are presented in terms of derived component structures and correlations between molecular column densities. In $\\S$ 4, 5, and 6 we explore some of the physical conditions of the CO-harboring gas in terms of empirical relationships, analytical chemical analysis, and detailed numerical modeling with Cloudy, respectively. A discussion will be given in $\\S$ 7, followed by the conclusions in $\\S$ 8. ",
        "conclusions": "Our study of diffuse molecular clouds employed a new and extensive sample of sight lines with UV observations of CO and H$_2$ to explore in detail the power-law relationship between the two species. The slope of log $N$(CO) $\\propto$ $B$ $\\times$ log $N$(H$_2$) was shown to require two components, one with $B$ = 1.5 $\\pm$ 0.2 for log $N$(H$_2$) $\\le$ 20.4, and another $B$ = 3.1 $\\pm$ 0.7 for VUV sight lines with higher $N$(H$_2$). The break in slope arises from a change in CO production, with CH$^+$ + O important at low $N$(CO) and C$^+$ + OH at $N$(CO) $>$ 10$^{14}$ cm$^{-2}$. The ratio CO/H$_2$ has a dependence on $N$(H$_2$) that results in an increase by $\\sim$3.5 orders of magnitude over the range of log $N$(H$_2$) $\\approx$ 19.5--22.0. Causes for variation in $X_{\\rm CO}$ include (a) the $n_{\\rm H}/I_{\\rm UV}$ ratio, which affects production and destruction, including self shielding, and (b) the metallicity of gas.  Together with the CO and H$_2$ we also analyzed new data for the carbon-bearing diatomic molecules CH, CH$^+$, and CN (as well as C$_2$) that are accessible through ground-based spectroscopy. The linear relationship between $N$(CH) and $N$(H$_2$) was confirmed again, both directly using these two molecules, as well as indirectly by showing that the CO versus CH relationship follows that for CO versus H$_2$. After determining fitted relationships of both CO and CH versus H$_2$ we were able to employ fit parameters in the prediction of $N_{\\rm p}$(H$_2$) for three sight lines without H$_2$ data. Analyzing $N$(CH$^+$) versus H$_2$ and CO resulted in two more confirmations of the power-law break displayed by CO versus either H$_2$ or CH, showing that this break separates the regime of low-density photochemistry from that involving high density. As for $N$(CN), all our regression fits returned slopes with $B$ $\\leq$ 1.8, somewhat shallower than earlier reports. Since essentially all CN detections are along high-density sight lines, the absence of a detected break in slope for the (smaller) CN sample is not surprising.  Many of the sight lines here are helping us to explore molecular environments that are associated with low $n_{\\rm H}$. As such, these lines of sight probe regions where nonequilibrium CH$^+$ chemistry is dominating the production of CO, as confirmed by modeling with Cloudy. For those sight lines with higher $n_{\\rm H}$ it was possible to include (equilibrium) chemistry of CH, C$_2$, and CN to predict molecular abundances and gas density. Such predictions were found to have tighter correlations when $N$(C$_2$) is part of the input into the chemical model. For the entire range of densities we showed that the empirical ratio $N$(CN)/$N$(CH$^+$) is better suited than $N$(CN)/$N$(CH) as an indicator of the average $n_{\\rm H}$ along diffuse sight lines.  We also considered rotational (excitation) temperatures in our modeling of CO and H$_2$, showing that $T_{0J}$(CO) does not vary for $J$ = 1--3. On the other hand, $T_{0J}$(H$_2$) increases with $J$, with indistinguishable slopes between log $T_{0J}$ and log $T_{01}$ for $J$ = 2--4. Further analysis of the excitation of both molecules should help constrain the conditions in diffuse molecular clouds.  As related in $\\S$ 7.4, it is our understanding that the regime of low-$N$(CO), low-$n_{\\rm H}$ diffuse molecular clouds is also sampled by a variety of non-UV observational methods, which nonetheless result in a significant number of CO nondetections. Thus the true nature of diffuse molecular clouds is best revealed by synoptic knowledge extracted from studies spanning the electromagnetic spectrum from radio and FIR, through visible and UV, to gamma ray observations."
    },
    "0807/0807.0989_arXiv.txt": {
        "abstract": "We report on the results of the latest solarFLAG hare-and-hounds exercise, which was concerned with testing methods for extraction of frequencies of low-degree solar p modes from data collected by Sun-as-a-star observations.  We have used the new solarFLAG simulator, which includes the effects of correlated mode excitation and correlations with background noise, to make artificial timeseries data that mimic Doppler velocity observations of the Sun as a star. The correlations give rise to asymmetry of mode peaks in the frequency power spectrum. Ten members of the group (the hounds) applied their ``peak bagging'' codes to a 3456-day dataset, and the estimated mode frequencies were returned to the hare (who was WJC) for comparison. Analysis of the results reveals a systematic bias in the estimated frequencies of modes above $\\approx 1.8\\,\\rm mHz$. The bias is negative, meaning the estimated frequencies systematically underestimate the input frequencies.  We identify two sources that are the dominant contributions to the frequency bias. Both sources involve failure to model accurately subtle aspects of the observed power spectral density in the part (window) of the frequency power spectrum that is being fitted.  One source of bias arises from a failure to account for the power spectral density coming from all those modes whose frequencies lie outside the fitting windows. The other source arises from a failure to account for the power spectral density of the weak $l=4$ and 5 modes, which are often ignored in Sun-as-a-star analysis. The Sun-as-a-star peak-bagging codes need to allow for both sources, otherwise the frequencies are likely to be biased. ",
        "introduction": "The solar Fitting at Low-Angular degree Group (\\emph{solarFLAG})\\footnote{See http://bison.ph.bham.ac.uk/\\,\\~\\,wjc/Research/FLAG.html} has as its main aims the development and refinement of techniques for analysis of data from the low-degree (low-$l$) p modes, from observations made of the ``Sun as a star'' . These data play a crucial r\\^ole in studies of the deep radiative interior and core of the Sun.  The input data for probing the solar interior are the mode parameters, such as individual frequencies, frequency splittings, damping rates and powers. The mode frequencies may be used to infer the internal hydrostatic structure (sound speed, density); accurate and precise frequencies are a vital prerequisite for ensuring that robust inference is made on the structure.  Analysis of the Sun-as-a-star (and also the resolved-Sun) helioseismic data requires application of complicated algorithms to extract estimates of the mode parameters. This usually involves fitting multi-parameter models to the resonant peaks in the frequency power spectrum of the observations. An important aim of the solarFLAG program is to quantify levels of bias arising from, and precision achievable in, these \\emph{peak bagging} procedures. ``Hare and hounds'' exercises on realistic artificial data form the framework for this activity.  In a first study (Chaplin et al. 2006) we looked in detail at the accuracy and precision of rotational frequency splittings extracted from a 3456-d set of artificial Sun-as-a-star data, to which ten members of the solarFLAG applied their peak-bagging codes. The parameters we look at in this paper are the low-$l$ mode frequencies returned by the peak-bagging codes.  The sets of artificial timeseries data used in our first study did not include any asymmetry of the simulated p-mode peaks in the frequency power spectrum; this asymmetry is exhibited by the real solar p modes. The peak-bagging codes must be able to cope with the asymmetry to in principle allow accurate estimation of the mode frequencies. It was therefore clear to us that to expedite a meaningful hare-and-hounds study on the mode frequencies we would need to generate artificial data with asymmetry included. This we have now done, and this paper reports on results of a hare-and-hounds exercise conducted with the new asymmetric artificial data, to which ten members of the solarFLAG applied their peak-bagging codes.  We have used a simple, but very powerful method to introduce in the time domain the effects of asymmetry, which is based on the framework proposed by Toutain, Elsworth \\& Chaplin (2006). There are two main factors in the method that contribute to the asymmetry of the artificial mode peaks. First, background noise is correlated with the excitation functions of the modes. Second, overtones of the same angular degree and azimuthal order have excitation functions that are correlated in time (see Chaplin, Elsworth \\& Toutain 2008). In this framework, correlation of the excitation follows naturally from invoking correlations with the background noise. As we shall see in this paper, the power spectral density of the resulting asymmetric mode peaks must be modelled accurately, otherwise estimates of the mode frequencies returned by the peak-bagging codes will be biased. This is the main result of the paper.  The layout of the rest of the paper is as follows. Section~\\ref{sec:simu} gives a brief overview of the solarFLAG simulator, which was used to make the artificial timeseries data for the hare-and-hounds exercise.  We also discuss in this section the basic attributes of the artificial dataset analyzed by the ten hounds. A detailed description of the simulator, which pays particular attention to the impact of correlations in the data, is given in Chaplin et al. (in preparation).  Section~\\ref{sec:fit} summarizes the main elements of the fitting strategies that were adopted by the hounds. Section~\\ref{sec:res} then presents the main results of the hare-and-hounds exercise. We look in detail at how the estimated frequencies of the hounds compared, not only against the input frequencies (results which bear on the accuracy of the peak-bagging procedures), but also against one another (results which bear on the precision inherent in the estimated frequencies). In Section~\\ref{sec:disc} we identify the origins of a systematic frequency bias that is reported in Section~\\ref{sec:res}.  Finally, we conclude in Section~\\ref{sec:sum} with a summary of the main conclusions of the paper, where we also discuss implications of the frequency bias for the fitting strategies. ",
        "conclusions": "\\label{sec:sum}  We have used the new solarFLAG simulator, which includes the effects of correlated mode excitation and correlations with background noise, to make artificial timeseries data that mimic Doppler velocity observations of the Sun as a star. The correlations give rise to asymmetry of mode peaks in the frequency power spectrum.  A 3456-day dataset was used as the input data for the latest solarFLAG hare-and-hounds exercise. This paper reports on the results of that exercise, which was concerned with testing methods for extraction of p-mode frequencies of low-degree (low-$l$) modes. Ten hounds applied their peak-bagging codes to the hare-and-hounds dataset. Peak-bagging involves maximum-likelihood fitting of mode peaks in the frequency power spectrum to multi-parameter fitting models.  Each hound returned peak-bagging estimates of the frequencies of the artificial $l=0$ to 3 modes to the hare (who was WJC) for further scrutiny.  Analysis of the results showed clear evidence of a systematic bias in the estimated frequencies of modes above $\\approx 1.8\\,\\rm mHz$. The bias is negative, meaning the estimated frequencies systematically underestimate the input frequencies. A follow-up analysis on independent realizations of the hare-and-hounds dataset showed that in some fits the bias could be as much as three- to four-times as large as the frequency uncertainties. Over the affected range of mode frequencies, the average bias is typically one to two-times the frequency uncertainties.  We identified two sources that are the dominant contributions to the frequency bias. Both sources involve failure to model accurately subtle aspects of the observed power spectral density in the part (window) of the frequency power spectrum that is being fitted.  One source of bias arises from a failure to account for the power spectral density of the weak $l=4$ and 5 modes. The other source arises from a failure to account for the power spectral density coming from all those modes whose frequencies lie outside the fitting windows (``outlying'' power).  The main lesson to be drawn from this paper is that the Sun-as-a-star peak-bagging codes need to allow for both sources, otherwise the frequencies given by analysis of real Sun-as-a-star data will in all likelihood be biased. The identification, and measurement, of the bias from ``outlying'' power is the most important new finding of the paper. Can we afford to ignore its effects? The short answer must be no. Our analysis suggests the magnitude of its frequency bias may be up to three-times as large as the frequency uncertainties, depending on the impact of realization noise, and it could present problems for helioseismic inference on the solar structure from inversions of the mode frequencies.  The precise sizes of biases given on the real Sun-as-a-star data will clearly depend on how closely our artificial data resemble those real data. We are certainly now able to reproduce frequency power spectra that bear a close resemblance to the real spectra -- courtesy of the new solarFLAG simulator -- and this suggests our bias estimates do have quantitative merit where helioseismic predictions concerning the real data are concerned. However, we should bear in mind that there may be some aspects of the real Sun-as-a-star data that are not reproduced exactly in the artificial data.  One such detail concerns the exact shapes shown by the asymmetric mode peaks. In the real solar p-mode data, it is assumed that there is also a contribution to the asymmetry owing to the radial location, extent, and multipole properties, of the acoustic sources. The question then arises: Is the underlying form of the power spectral density due to these contributions the same as that from the correlated noise modelled in the solarFLAG simulator? Subtle differences would affect the sizes of the frequency biases.  We finish with a few comments on implications of the results of this paper for the peak-bagging codes.  The standard Sun-as-a-star approach is to go through the frequency power spectrum fitting a pair of modes at a time. For such an approach to be useful our results here stress the need for power from the outlying modes to be fully accounted for in the fitting windows (so-called ``pseudo whole-spectrum'' fitting). The other option is to fit all modes in the frequency power spectrum in one go (so-called ``whole spectrum'' fitting). Either way, it is not sufficient to have an approximate, or first-order, estimate of the outlying, or total, power spectral density: the estimate must be very accurate, otherwise the frequency bias will remain, or additional bias may be introduced.  In order to provide such an estimate, it is necessary to describe accurately the power spectral density of each mode peak a long way from its resonant frequency (i.e., in the decaying wings of the peaks). The fitting formalism most often used to model the asymmetric power spectral density of the mode peaks -- that due to Nigam \\& Kosovichev (1998) -- fails in such a description. This is because the Nigam \\& Kosovichev formalism is an approximation (low-order expansion) that is usable only at frequencies close to resonance. Far from resonance, the modelled power spectral density tends to a constant offset not shown by the real data\\footnote{The ``Fano profile'' formula in Gabriel et al. (2001) also fails in this regard, because it retains those terms that lead to the offset far from resonance.}  Clearly the requirements on the sought-for fitting model are that it should describe the asymmetric peaks both close to resonance, and far from resonance where we know the power falls off significantly in the real p-mode peaks.  There is, potentially, a very simple solution to this problem. If the high-order terms in the Nigam \\& Kosovichev formalism (those in the square of the asymmetry parameter) are disregarded, it turns out that the resulting, truncated formalism can satisfy both of the above requirements. Indeed, it has the correct form to model accurately the asymmetric shapes given by correlated noise (see Toutain, Elsworth \\& Chaplin 2006). The issue then arises as to whether this is sufficient to describe the shapes of the real p-mode peaks (see previous comments above). Another approach is to generate a model of resonant spectrum in one go, with all the overtone structure included, rather than model the spectrum as a superposition of individual Lorentzian-like peaks. This may be accomplished using a suitable form for the acoustic potential of the solar cavity, together with information on the acoustic source and the reflection and transmission properties of the upper cavity boundary (e.g., see Jefferies, Vorontsov \\& Giebink 2004).  Finally, we note that the detailed conclusions drawn in this paper are for the Sun-as-a-star observations. The Sun-as-a-star techniques are of course also directly applicable to stars that show Sun-like oscillations, and we are now starting to get \\emph{asteroseismic} peak-bagging results on other stars (e.g., see Fletcher et al. (2006), who analyzed data collected by WIRE on $\\alpha$\\,Cen~A; and Appourchaux et al. (2008), who have analyzed the first Sun-like oscillations data collected by CoRoT, on the star HD49933). A peak-bagging pipeline is being constructed by the asteroFLAG group\\footnote{http://www.issibern.ch/teams/Astflag} (Chaplin et al. 2008a) for application to the asteroseismic data that will be collected on hundreds of Sun-like stars by the NASA Kepler mission (Christensen-Dalsgaard 2007). Even though the intrinsic signal-to-noise ratios will be lower than for the Sun-as-a-star data some stars will be monitored continually for several years, meaning we should be able to constrain the p-mode parameters to high levels of precision.  With the p-mode parameters reflecting different intrinsic properties of the stars, we should expect to be confronted with many different potential bias problems in different parts of the color-magnitude diagram occupied by the Sun-like oscillators (Appourchaux et al. 2006a, b; Chaplin et al. 2008b). Similar studies to the one undertaken here for the Sun as a star are now being made by asteroFLAG to prepare the peak-bagging analysis for Kepler.   \\subsection*"
    },
    "0807/0807.2803_arXiv.txt": {
        "abstract": "EPOXI (EPOCh $+$ DIXI) is a NASA Discovery Program Mission of Opportunity using the Deep Impact flyby spacecraft. The EPOCh (Extrasolar Planet Observation and Characterization) Science Investigation will gather photometric time series of known transiting exoplanet systems from January through August 2008. Here we describe the steps in the photometric extraction of the time series and present preliminary results of the first four EPOCh targets. ",
        "introduction": "The high-precision time series produced from the EPOCh observations, with their substantial time and phase coverage, allow us to pursue three important science goals with respect to transiting exoplanets. First, we will search for evidence of reflected light from the giant transiting planets at secondary eclipse (Rowe et al. and Matthews et al., these proceedings). Second, EPOCh will be sensitive to terrestrial planets in these systems, using two detection methods. The first will be a search for direct transit events caused by approximately Earth-sized planets, particularly ``hot Earths'' that are predicted to be captured in low order mean motion resonances with the giant planets (\\cite[Croll et al. 2007a,b]{Croll2007a}). EPOCh will also perform transit timing measurements for the giant planets, and search for timing variations due to terrestrial planet perturbers (\\cite[Agol et al. 2005]{Agol05}; \\cite[Holman \\& Murray 2005]{Holman05}). The third science goal of the EPOCh mission is a precise determination of the system parameters, notably the planet radius, and a search for planetary satellites and circumplanetary rings through their perturbations to the transit light curve (\\cite[Brown et al., 2001]{Brown01}; \\cite[Barnes \\& Fortney, 2004]{Barnes04}). These proceedings will summarize the steps we take to produce the time series, as well as describe the preliminary search for additional transiting bodies.  Each time series consists of continuous 50-second integrations through a clear visible filter with the 30 cm aperture High Resolution Imager (HRI) instrument. These observations have a typical duration of 3 weeks per target. The spacecraft will then continue en route to its 2010 encounter with comet Hartley-2, the target of the DIXI Science Investigation. For a history of the EPOCh observations and more information about the EPOXI mission, please see Christiansen et al. (these proceedings). ",
        "conclusions": "We have observed seven transits of HAT-P-4, five transits of TrES-3, three partial transits of XO-2, and seven transits of GJ 436. Figure 2 shows the EPOCh observations, phase-folded and binned, for these four systems. Our strategy is to conduct the science analysis of the data while working in parallel to refine the photometric precision through improvements to our extraction pipeline. We are currently conducting a Markov Chain Monte Carlo modeling and reflected light analysis to refine the system parameters and constrain the planetary albedos. These analyses are explained in more detail in Christiansen et~al. (these proceedings). We are also conducting preliminary searches for shallow transits made by additional bodies. For GJ 436, this search is especially interesting, since the eccentricity of the known transiting planet points to the possible presence of a second planetary companion. We are just beginning to develop the software to search for these additional transits, which will take steps briefly outlined here.  The set of lightcurves for additional transiting bodies is a three parameter family if we make the simplifying assumption of no limb darkening of the host star and an inclination angle $i$ of 90$^{\\circ}$: radius of the planet $R_{p}$, period of the planet $P$, and phase $\\phi$. Using the Mandel and Agol (2002) routines for generating lightcurves given these parameters, we compute a grid of models corresponding to additional possible planets. For each of these models, we compute the $\\chi^{2}$, using the out-of-transit standard deviation to estimate the error in each point. At the positions of local minima, we will change to a routine that more finely samples the $\\chi^{2}$ space so that we can accurately gauge the significance levels of the parameters. The improvement in the $\\chi^{2}$ over the null hypothesis will also determine whether we can claim a positive detection.  The photometric precision of the GJ 436 time series (with S/N=44 for each transit event) would enable a detection of a 1.3 R$_{\\oplus}$ planet with S/N of 4 and a 2.1 R$_{\\oplus}$ with S/N of 10, even if the planet produced only a single transit event."
    },
    "0807/0807.0201_arXiv.txt": {
        "abstract": "{The \\begin{em}ESO-Spitzer Imaging extragalactic Survey\\end{em} (ESIS) is the optical follow up of the \\begin{em}Spitzer Wide-area Infra-Red Extragalactic\\end{em} survey (SWIRE) in the ELAIS-S1 region of the sky.} {In the era of observational cosmology, the main efforts are focused on the study of galaxy evolution and its environmental dependence.  Wide area, multiwavelength, extragalactic surveys are needed in order to probe sufficiently large volumes, minimize cosmic variance and find significant numbers of rare objects. } {We present VIMOS $I$ and $z$ band imaging belonging to the ESIS survey. A total of $\\sim4$ deg$^2$ were targeted in $I$ and $\\sim1$ deg$^2$ in $z$. Accurate data processing includes removal of fringing, and mosaicking of the complex observing pattern. Completeness levels and photometric uncertainties are estimated through simulations. The multi-wavelength data available in the area are exploited to identify high-redshift galaxies, using the IR-peak technique.} {More than 300000 galaxies have been detected in the $I$ band and $\\sim50000$ in the $z$ band. Object coordinates are defined within an uncertainty of $\\sim0.2$ arcsec r.m.s., with respect to GSC 2.2. We reach a 90\\% average completeness at 23.1 and 22.5 mag (Vega) in the $I$ and $z$ bands, respectively. On the basis of IRAC colors, we identified galaxies having the 1.6$\\mu$m stellar peak shifted to $z=1-3$. The new $I,\\ z$ band data provide reliable constraints to avoid low-redshift interlopers and reinforce this selection. Roughly 1000 galaxies between $z=2-3$ were identified over the ESIS $\\sim4$ deg$^2$, at the SWIRE 5.8$\\mu$m depth (25.8 $\\mu$Jy at 3$\\sigma$). These are the best galaxy candidates to dominate the massive tail ($M>10^{11}$ M$_\\odot$) of the $z>2$ mass function.} {} ",
        "introduction": "\\label{sect:intro}   In the recent years, the observational effort in cosmology has focused on the understanding of galaxy formation and tracing the processes that transformed the smooth and homogeneous Universe detected right after the Big Bang into the clumpy, clustered structures we see at $z=0$.  With this respect, the critical parameters to be probed turned out to be baryonic mass and the epoch of its major assembly \\citep[e.g.][]{pozzetti2007,rudnick2006,rudnick2003, dickinson2003b,brinchmann2000}. In this respect, particularly controversial is the origin of the recently discovered population of massive ($M>10^{11}$ M$_\\odot$) galaxies at redshift $z>1.5-2.0$ \\citep[e.g.][]{berta2007b,fontana2006,fontana2004,franceschini2006, bundy2005,daddi2004,cimatti2002c}.  This puzzle has been approached with two opposite ``philosophies'': very-deep, pencil-beam surveys, or wide-area, shallower surveys.  The first approach has mostly relied on space-based and 8m telescopes observations, that provided the deepest and clearest views of the distant Universe so far. Some examples, are the Hubble Deep Fields \\citep{williams1996,williams2000}, the GOODS survey \\citep{dickinson2003a} and the Hubble Ultra Deep Field \\citep{beckwith2003}, pushing the detection limit to the 30th magnitude.  Due to their small areas, these surveys are affected by non negligible cosmic variance, which can be a significant source of uncertainty in the study of galaxy statistical properties, even at high redshift \\citep[see for example][]{somerville2004}.  Moreover the search for rare objects, such as massive high-redshift galaxies, requires large volumes to be sampled, in order to detect them in representative numbers. For example, \\citet{berta2007b} find a number density of $\\sim 3\\times 10^{-5}$ $[$h$_{70}^3$ Mpc$^{-3}]$ for $M>1.6\\times10^{11}$ M$_\\odot$ galaxies at $z=2-3$, which translates in only a few hundreds sources per square degree.  In overcoming these effects, wide field surveys, such as COSMOS \\citep[2 deg$^2$][]{scoville2007a}, VVDS \\citep[16 deg$^2$][]{lefevre2004}, CFHTLS \\citep[4 to 410 deg$^2$ at different depths, e.g.][]{nuijten2005} play a crucial role.  In this framework, the {\\em Spitzer Wide-area Infra-Red Extragalactic} survey (SWIRE, 49 deg$^2$) is quite unique. The availability of Spitzer IRAC (3.6, 4.5, 5.8, 8.0 $\\mu$m) data samples the restframe near-IR emission of galaxies up to redshift $z=3$ at least, directly probing the stellar mass assembly of galaxies in the distant Universe. Despite SWIRE being a rather shallow survey (reaching 5$\\sigma$ depths of 3.7 and 43 $\\mu$Jy at 3.6 and 5.8 $\\mu$m, respectively), it detects $2\\times10^{11}$ M$_\\odot$ galaxies up to $z=3$ \\citep{berta2007b}. MIPS (mainly 24$\\mu$m) data secure the identification of star forming systems in the same redshift range.  An extensive multi-wavelength coverage is required. When dealing with large datasets over such a wide area, a spectroscopic characterization of the survey is beyond the multiplexing capabilities of current instrumentation, and photometric investigation is clearly the only viable approach to be carried out in a reasonable amount of time.  The {\\em ESO-Spitzer wide-area Imaging Survey} (ESIS) is an ESO Large Programme (P.I. Alberto Franceschini), targeting the SWIRE ELAIS-S1 field. This region includes the absolute minimum of the Galactic 100 $\\mu$m emission in the Southern sky (0.37 $[$MJy/sr$]$).  In paper-I \\citep{berta2006} we presented the ESIS WFI/2.2m observations in the $B$, $V$, $R$ bands, illustrating also the potential of panchromatic studies of galaxies. Here VIMOS $I$ and $z$ band observations are presented, completing the optical ESIS survey.  The paper is organized as follows. Section \\ref{sect:observations} presents the survey and the observing strategy. Section \\ref{sect:datared} and Appendix \\ref{sect:app_fringing} describe the main steps in the data reduction, including astrometric, photometric calibration and mosaicking. The quality of data products is tested in Sect. \\ref{sect:quality}, while the contents of the released catalogs are described in Sect. \\ref{sect:catalogs}. In Sect. \\ref{sect:selection} we show how VIMOS $I$ and $z$ band data help in the selection and in constraining the spectral energy distributions (SEDs) of high-redshift galaxies. Finally, Sect. \\ref{sect:summary} summarizes our findings. ",
        "conclusions": "\\label{sect:summary}   In the era of wide-area cosmological surveys, we have presented $I$ and $z$ band wide field imaging of ELAIS-S1, carried out with the VIMOS camera on VLT. These observations are part of the {\\em ESO-Spitzer Imaging extragalactic Survey} (ESIS) and cover $\\sim4$ deg$^2$ in the $I$ band and $\\sim1$ deg$^2$ in $z$. The nominal exposure time is 1800s and 3600s in $I$ and $z$ respectively, but the coverage is not uniform, due to the very complex observing pattern.  As a result of accurate data processing, nine independent mosaics for the $I$ band and one single, 1 deg$^2$ wide frame in the $z$ band were produced. The final catalogs include more than 300000 $I$ band sources and 50000 $z$ band objects. Equatorial coordinates of VIMOS objects are constrained within $\\sim0.2$ r.m.s. arcsec, when compared to the GSC 2.2 catalog. Relative astrometric accuracy between the nine $I$ band images is as good as $\\sim0.1$ arcsec r.m.s.  Photometric uncertainties were estimated through simulations, and turn out to be $\\sim 0.05,\\ 0.20$ mag at $I=21.5,\\ 23.5$ and slightly larger in the $z$ band. The survey reaches a 90\\% average completeness at 23.1 and 22.5 (Vega magnitudes) in the $I$ and $z$ band respectively.  The ESIS VIMOS survey provides optical $I$ counterparts for $>60$\\% (137206/219877) of SWIRE/Spitzer sources in the ELAIS-S1 area.  We have exploited the new $I$ and $z$ band data --- in combination to the other multi-wavelength observations available in ELAIS-S1 --- to search for high-redshift ($z=1-3$) galaxies. Our selection is based on the detection of the 1.6$\\mu$m stellar peak in the IRAC wavelength domain \\citep[IR-peakers][]{berta2007b}. The $I$ band data give a valuable constraint on IR-peakers: requiring $(I-3.6)_{AB}>2$ avoids low-redshift ($z\\lesssim0.6$) interlopers. The availability of VIMOS data over $\\sim4$ deg$^2$ allows the selection to be extended over the whole SWIRE-ESIS field. A total of 973 5.8$\\mu$m-peak galaxies at $z=2-3$ were found over $\\sim4$ deg$^2$ at the SWIRE 5.8$\\mu$m 3$\\sigma$ depth (25.8 $\\mu$Jy).  The $z$ band data were used to perform the $BzK$ \\citep{daddi2004} selection over $\\sim1$ deg$^2$ and verify the consistency of the IR-peakers class. Our high-$z$ objects lie in the locus of $z>1.4$ star forming galaxies. Apparently, no passive IR-peakers are detected at SWIRE-ESIS depth, over the central ESIS 1 deg$^2$ field.  At the SWIRE depth, these galaxies have stellar masses $M>10^{11}$ M$_\\odot$ at $z>2$. The identification of a significant number of these rare high-$z$ sources is possible only thanks to the wide area probed by SWIRE-ESIS. This survey is giving a significant contribution --- among others --- in studying the very massive tail of the stellar mass function at $z>2$, and constraining the physical properties of its constituents \\citep[e.g.][]{berta2007b,lonsdale2006}.  Updates on the ESIS project are available on the web page {\\tt http://www.astro.unipd.it/esis}."
    },
    "0807/0807.1946_arXiv.txt": {
        "abstract": "We investigate the total kinetic powers ($L_{\\rm j}$) and ages ($t_{\\rm age}$) of powerful jets in FR II radio galaxies by  comparison of the dynamical model of expanding cocoons with observations. We select four FR II radio sources (Cygnus A, 3C 223, 3C 284, and 3C 219), for which the mass-density profiles of intracluster medium (ICM) are known in the literature. It is found that large fractions $\\gtrsim 0.02 - 0.7$ of the Eddington luminosity ($L_{\\rm Edd}$) are carried away as a  kinetic power of jet. The upper limit of estimated $2 L_{\\rm j} / L_{\\rm Edd}$ are larger than unity ($\\lesssim 10$) for some sources,  suggesting a possibility of super-Eddington mass accretions. As a consequence of the large powers, we also find that the total energy stored in the cocoon ($E_{\\rm c}$)  exceeds the energy derived from the minimum energy condition for the energy of radiating non-thermal electrons and magnetic fields ($E_{\\rm min}$): $4< E_{\\rm c}/E_{\\rm min} <310$. % This implies that most of the energy in cocoon is carried by invisible components such as thermal leptons (electron and positron) and/or protons. ",
        "introduction": "Relativistic jets in active galactic nuclei (AGN) are a fundamental aspect of plasma accretion onto supermassive black holes (SMBHs). Although the formation mechanism of relativistic jets remains a longstanding problem, it is well established that they carry away some fractions of the available accretion power in the form of collimated beam (e.g., Begelman et al. 1984 for review). The total kinetic powers of AGN jets $L_{\\rm j}$ is one of the most basic physical quantities characterizing the jet. A lot of authors have investigated $L_{\\rm j}$ in various ways so far \\citep[e.g.,][]{RS91,  CF93, WRB99}. It is, however, difficult to estimate $L_{\\rm j}$, since most of the observed emissions from AGN jets are of non-thermal electron origin and it is hard to detect the electromagnetic signals from the thermal and/or proton components. Hence, the free parameter describing the amount of the invisible plasma components always lurks in the estimates of $L_{\\rm j}$ based on the non-thermal emissions.       The estimate of $L_{\\rm j}$ for low-power Fanaroff-Riley I (FRI) radio galaxies has been motivated by the observations of ``X-ray cavity\" % which is the region embedded in ICM with the suppressed X-ray surface brightness and coincides with the  radio lobe \\citep{BVF93}. The cavities (or cocoons) are supposed to be a direct evidence of the displacement of the ambient ICM by the shocked jet matter. Dynamical models of cavities are a good tool, since the invisible plasma components as well as non-thermal electrons play a role for expansions of cavities. For FR I sources, the total kinetic energy of the jet has been estimated as \\begin{eqnarray} L_{\\rm j}t_{\\rm age} \\sim \\frac{\\hat{\\gamma}_{\\rm c}}{\\hat{\\gamma}_{\\rm c}-1} P_{\\rm c} V_{\\rm c},  \\nonumber \\end{eqnarray} where $t_{\\rm age}$, $P_{\\rm c}$, $V_{\\rm c}$, $\\hat{\\gamma}_{\\rm c}$, are the source age, the pressure, the volume, the adiabatic index of the cavity, respectively \\citep{FCB02, ADF06}. In these studies, however, the thermal pressure of surrounding ICM ($P_{\\rm ICM}$) is substituted for the one in the cavity (i.e., $P_{\\rm c}\\sim P_{\\rm ICM}$). This assumption may be applied only to subsonic expansions.     On the other hand, the cocoon pressure of powerful Fanaroff-Riley class II (FRII) radio galaxies is expected to be larger than that of the surrounding ICM, which is expressed as $P_{\\rm c} > P_{\\rm ICM}$ (Begelman \\& Cioffi 1989 hereafter BC89), and the cocoon of FR II radio galaxies is likely to be expanding super-sonically. Then the substitution of ICM pressure for the cocoon pressure is not justified. A new estimate of $L_{\\rm j}$ for FR II radio galaxies by use of the dynamical model of cocoon expansions is proposed by \\citet{KK05} (hereafter KK05), in which $L_{\\rm j}$ and  $t_{\\rm age}$ are derived from the comparison of the cocoon model with the morphology of the cocoon obtained by radio observations. It should be stressed that  $P_{\\rm c}$ is not assumed but solved in this model. Hence it can be applied even to the cocoons with $P_{\\rm c}>P_{\\rm ICM}$. So far, however, this estimate of $L_{\\rm j}$  has been applied only to Cygnus A. The expansion of the number of samples is evidently crucially important for exploring general characteristics of the powerful AGN jets. For this purpose, we apply the method of KK05 to other bright FR II radio galaxies, for which the physical conditions of the associated ICM  have been estimated in the literature.   In the present work, we especially focus on the ratio of $L_{\\rm j}/L_{\\rm Edd}$, where $L_{\\rm Edd}$ is the Eddington luminosity of AGN, since $L_{\\rm j}/L_{\\rm Edd}$ is a more fundamental quantity than $L_{\\rm j}$ from the point of view of the jet formation physics. Another interesting quantity we examine in this work is the ratio of the internal energy deposited in the cocoon ($E_{\\rm c}$) to the  minimum energy ($E_{\\rm min}$) obtained by the minimum energy condition for radiating non-thermal electrons and magnetic fields (e.g., Miley 1980). Some of the previous works studying this quantity (e.g., Hardcastle \\& Worrall 2000; Leahy \\& Gizani 2001) reported that the cocoon pressure expected from the inferred $E_{\\rm min}$ is smaller than the pressure of ambient matter, suggesting the difference between $E_{\\rm c}$ and $E_{\\rm min}$. due to the lack of minimum pressure against the pressure of ambient medium. Although these studies obtained the lower limit of the ratio, it is the value of $E_{\\rm c}$ that is of greater importance.           According to a large sample of galaxies collected recently, the fraction of AGNs in all the galaxies is suggested to be $\\sim20 \\-- 40\\%$, larger than previously thought \\citep{KHT03, MNG03}. The interest in AGNs is gaining momentum in the context of the co-evolution of galaxies and their central black holes (e.g., Kawakatu et al. 2003, Granato et al. 2004; Di Matteo et al. 2005). AGN outflows in particular are likely to be a key ingredient in this context (Silk \\& Rees 1998; Fabian 1999; King 2003). The AGN feedback by the outflows may be also promising to explain the tight correlations between the  ratio of the mass of SMBH ($M_{\\rm BH}$) to that of galactic bulge (Kormendy \\& Richstone 1995; Magorrian et al. 1998) and the ratio of $M_{\\rm BH}$  and the stellar velocity-dispersion in the bulge \\citep{FM00, TGB02}. In this sense, a robust estimate of the basic quantities such as $L_{\\rm j}/L_{\\rm Edd}$ and $t_{\\rm age}$ of radio loud AGNs at low $z$ is an important first step for understanding the AGN feedback processes in the universe.                        The outline of the paper is as follows. In \\S2, the model of the expanding cocoon by KK05 is briefly reviewed. In \\S3, we explain how to extract the key quantities from the observations of four nearby FR II radio galaxies, Cygnus A, 3C 223, 3C 284, and 3C 219, which are required for the comparison with the model. We then estimate  the total kinetic power, $L_{\\rm j}$, and the dynamical ages, $t_{\\rm age}$, in  \\S4. Finally in \\S5, we summarize our results and discuss some implications on the physics of AGN jet. Throughout the paper, we adopt a cosmology with $H_0=71~{\\rm km~s^{-1}}$, $\\Omega_{\\rm M}=0.3$, and  $\\Omega_{\\rm \\Lambda}=0.7$ \\citep{SVP03}. ",
        "conclusions": "\\label{SD}       In this paper we have investigated the total kinetic power and age of  powerful FR II jets. We have selected four FR II radio galaxies (Cygnus A, 3C 223, 3C 284, and 3C 219), for which the surrounding ICM densities and pressures are known in the literature. Below we summarize our main results.   (I) {\\it  Large fractions of the Eddington power in the range of  $\\gtrsim 0.02 - 0.7$ are carried away as a kinetic power of  jet in the FR II sources.}  (II) {\\it The energy deposited in the cocoon, $E_{\\rm c}$, exceeds the minimum energy, $E_{\\rm min}$, by a factor of $4 - 310$.} %     Although our results  allow a wide range of $L_{\\rm j}$ and $E_{\\rm c}$, interesting implications can still be obtained and will be discussed below. In \\S 5.1., we address some issues concerning the ratio of $L_{\\rm j}$ to $L_{\\rm Edd}$ by referring to the studies of X-ray binaries. In \\S 5.2., the energetics in the cocoon is constrained from the obtained $\\eta_{\\rm c}$.     \\subsection{$L_{\\rm j} / L_{\\rm Edd}$}      Postulating that the relativistic jet is powered by the release  of gravitational energy ($L_{\\rm acc}$) of accreting matter (e.g., Marscher et al. 2002), the jet power can be expressed as $2L_{\\rm j} = \\epsilon_{\\rm j} L_{\\rm acc}$, where $\\epsilon_{\\rm j}$ is the efficiency of energy conversion ($0<\\epsilon_{\\rm j}<1$). Hence, $2 L_{\\rm j} / L_{\\rm Edd}$ gives the minimum mass accretion rate normalized by the Eddington mass accretion rate. Our results then suggest that quite high mass accretion rates, at least above $0.02L_{\\rm Edd}$, are required to produce FRII radio sources. Moreover, since $10 \\gtrsim 2 L_{\\rm j} /L_{\\rm Edd} \\gtrsim 0.65$ is predicted for 3C 219, some FRII radio sources may have super-Eddington mass accretion rates. The theory of accretion disk predicts that the accretion disk of these objects is optically-thick and called the {\\it slim-disk} (e.g., Abramowicz et al. 1988). From three distinctive X-ray properties (the large photon index $\\Gamma \\gtrsim 2$, rapid variability and soft X-ray excess), narrow line Seyfert 1 galaxies (NLS1s) are considered to be super-Eddington objects and, therefore, are inferred to have a slim-disk  (e.g., Pounds et al. 1995; Boller et al. 1996; Mineshige et al. 2000; Collin \\& Kawaguchi 2004; Shemmer et al. 2006). If  3C 219 is indeed super-Eddington, it is expected to show the above mentioned X-ray features like NLS1s. Note, however, that 3C 219 has a relatively hard X-ray spectrum with $\\Gamma =1.58 < 2 $ (Shi et al. 2005), and thus the physical state of the accretion disk in 3C 219 could be different from those in NLS1s.   In order to explore the nature of the accretion disk in 3C 219, we  compare the characteristics of AGNs with those of X-ray binaries (XRB), since both of them have common physical processes such as disk accretions, relativistic jets, and quenching of these jets (e.g., Heinz \\& Sunyaev 2003; Ho 2005; McHardy et al. 2006). Thanks to their much shorter dynamical timescales, XRBs in various states have been observed in great detail and are found to occupy particular X-ray spectral states (Fender et al. 2004; Remillard \\& McClintock 2006 for a review) as follows;   (i) $L_{\\rm acc}/L_{\\rm Edd}\\lesssim 0.01$ (low/hard state: LS),   (ii) $0.01 \\lesssim L_{\\rm acc}/L_{\\rm Edd}\\lesssim 0.3$ (high/soft state: HS),   (iii) $L_{\\rm acc}/L_{\\rm Edd}\\gtrsim 0.3$ (very high state: VHS).     State (i) is usually accompanied by a jet. For $L_{\\rm acc}/L_{\\rm Edd}>0.01$, the  radio emission is quenched in state (ii), while in state (iii), the soft VHS, which has an  X-ray spectrum dominated by a  steep power-law component (photon index  $\\Gamma >2$), is radio quiet and the hard VHS and/or the transition from the hard VHS to the soft VHS is accompanied by relativistic ejection events.   If the analogy between XRBs and AGNs holds, NLS1s may be in the soft VHS, since many NLS1s have a  steep power-law component and  high Eddington ratios $>$ 0.3 (e.g., Collin \\& Kawaguchi 2004), that is, higher than the upper limit set by the stability of the Shakura \\& Sunyaev (1973) disk (Shakura \\& Sunyaev 1973). On the other hand, 3C 219 (FR II) may correspond to the hard VHS and/or the transition state because it has hard X-ray spectra and high Eddington luminosities. In order to judge whether AGNs  are scaled-up XRBs, it is essential to confirm that radio loud AGNs actually have the states analogous to the spectral states (especially VHS) in XRBs\\footnote{It has been well established that low-luminosity AGNs are the high-mass counterpart of XRBs in LS's (Ho 2005).}.      The co-evolution of a central black hole (BH) and its host galaxy together with AGN feedbacks have been intensively studied in  various ways (e.g., Silk \\& Rees 1998; Di Matteo et al. 2003; Granato et al. 2004). The  intensive surveys of QSOs show that the number density of QSOs  is peaked at  $z\\approx 2$ \\citep{FNL01}. The existence of QSOs at $z\\gtrsim 2$  with a central BH of a mass smaller than predicted from the bulge BH-mass relation is naturally expected in the  build-up of SMBHs toward  $z\\approx 2$. For the exploration of the co-evolution processes, dusty spheroidal galaxies (e.g., galaxies emitting sub-millimeter radiations and  ultra-luminous infrared galaxies) are  the ones to be scrutinized, since the dusty-gas in them is one of the key quantities for the co-evolution. Kawakatu et al. (2003) pointed out the possibility of radio loud AGNs at high-$z$ being QSOs in the early evolution phase (we call them proto-QSOs) that contain a growing BH. It is, however, difficult to explore the nature of proto-QSOs at high-$z$ by  observations in optical and X-ray bands owing to severe dust-absorptions. In contrast, the estimate of $L_{\\rm j}$ and $t_{\\rm age}$ presented in this work is applicable even to high-$z$ sources (e.g., Schmidt et al. 2006) without being suffered from the dust extinction. Therefore, the estimate of $L_{\\rm j}/L_{\\rm Edd}$ based on the cocoon dynamics is a potential new powerful tool to discover proto-QSOs among high-$z$ radio loud AGNs.                     \\subsection{The energetics} \\label{content}    Lastly, we discuss the energetics in the cocoon. Summing up in advance, the total internal energy of invisible components such as thermal leptons and/or protons tends to be larger than those of radiating non-thermal electrons and magnetic fields. $E_{\\rm c}$ is expressed by components as $E_{\\rm c} = (U_{\\rm e} + U_{\\rm B} + U_{\\rm inv})V_{\\rm c}$, where $U_{\\rm e}$, $U_{\\rm B}$, and $U_{\\rm inv}$ are the energy densities of  non-thermal leptons (electrons and positrons), magnetic fields, and  invisible particles, respectively. The energy ratio $\\eta_{\\rm c}\\equiv E_{\\rm c}/E_{\\rm min}$ can be then expressed as $\\eta_{\\rm c} =  (U_{\\rm e} + U_{\\rm B} + U_{\\rm inv}) / U_{\\rm min}$, where  $U_{\\rm min}=E_{\\rm min}/V_{\\rm c}$ is the minimum energy density. From the obtained values of $\\eta_{\\rm c}$, we investigate here the contribution of $U_{\\rm inv}$ to the total energy by evaluating $U_{\\rm e}/U_{\\rm min}$ and $U_{\\rm B}/U_{\\rm min}$.    It is useful to express $U_{\\rm e}/U_{\\rm min}$ and $U_{\\rm B}/U_{\\rm min}$ in terms of $U_{\\rm e}/U_{\\rm B}$, since $U_{\\rm e}/U_{\\rm B}$ has been intensively investigated by a lot of authors (e.g., Isobe et al. 2002; Kataoka et al. 2003; Croston et al. 2004; Kataoka \\& Stawarz 2005; Croston et al. 2005). Since the synchrotron luminosity $L_{\\nu}$ is proportional to $U_{\\rm e}U_{\\rm B}^{3/4}V_{\\rm c}$, we obtain the relation $U_{\\rm e}\\propto (U_{\\rm e}/U_{\\rm B})^{3/7}$, or equivalently $U_{\\rm B}\\propto (U_{\\rm e}/U_{\\rm B})^{-4/7}$ for fixed values of $L_{\\nu}$ and $V_{\\rm c}$. From this relation and Eq. (\\ref{minimum}), we can derive the  following expressions: \\begin{eqnarray} \\frac{U_{\\rm e}}{U_{\\rm min}} \\simeq 0.5 \\left(\\frac{U_{\\rm e}}{U_{\\rm B}}\\right)^{3/7}, \\qquad \\frac{U_{\\rm B}}{U_{\\rm min}} \\simeq 0.5 \\left(\\frac{U_{\\rm e}}{U_{\\rm B}}\\right)^{-4/7}. \\label{ratio2} \\end{eqnarray} Hence,  $U_{\\rm inv}/U_{\\rm min}$ is given by \\begin{eqnarray} \\frac{U_{\\rm inv}}{U_{\\rm min}} \\simeq \\eta_{\\rm c} - 0.5 \\left\\{ \\left(\\frac{U_{\\rm e}}{U_{\\rm B}}\\right)^{3/7} + \\left( \\frac{U_{\\rm e}}{U_{\\rm B}}\\right)^{-4/7} \\right\\}. \\label{inv} \\end{eqnarray} Recent observations show that the ratio of $U_{\\rm e}$ to $U_{\\rm B}$ is $1 \\lesssim U_{\\rm e} / U_{\\rm B} \\lesssim 10$  on average. Substituting these values in  Eq. (\\ref{ratio2}), we  find $0.13\\lesssim U_{\\rm B}/U_{\\rm min} \\lesssim 0.5$ and $0.5\\lesssim U_{\\rm e}/U_{\\rm min} \\lesssim 1.3$. Since the obtained range of $\\eta_{\\rm c}$ is $\\sim 4-310$, $U_{\\rm inv}/U_{\\rm min}$ is evaluated as $3\\lesssim U_{\\rm inv}/U_{\\rm min} \\lesssim 310$ from Eq. (\\ref{inv}). Thus, we conclude that the internal energy  of invisible particles must be larger than  % those of radiating non-thermal leptons and magnetic fields ($U_{\\rm e}, U_{\\rm B} \\lesssim U_{\\rm inv}$) to explain the result obtained in this paper that $E_{\\rm c}$ is larger than $E_{\\rm min}$ by a factor of 4 at least."
    },
    "0807/0807.3458_arXiv.txt": {
        "abstract": "{ In this paper we report on the spectral analysis of seven X-ray transients, which were found to be active during a monitoring campaign of the Galactic center carried out in 2006 and 2007 using the X-ray telescope aboard the \\textit{Swift} satellite. This campaign detected new outbursts of five known X-ray transients and discovered two new systems. Their 2-10 keV peak luminosities range from $\\sim 10^{34}~\\mathrm{to}~6 \\times 10^{36}~\\mathrm{erg~s}^{-1}$, which implies that all seven X-ray transients are subluminous compared to the bright X-ray transients that have peak luminosities of $10^{37-39}$~erg~s$^{-1}$. Two of the sources discussed in this paper are confirmed neutron star systems (AX J1745.6-2901 and GRS 1741-2853), while the five others have an unknown nature. We discuss the characteristics of the observed outbursts and the duty cycles of the various systems. Several of the detected transients seem to undergo enhanced X-ray activity with levels intermediate between quiescence and full outburst. We discuss the possibility that the subluminous appearance of the eclipsing X-ray burster AX J1745.6-2901 is due to line-of-sight effects. We detected two type-I X-ray bursts with a duration of 50-60 seconds from AX J1745.6-2901, which we discuss in view of the bursting behavior of low-luminosity X-ray transients. Assuming that we are dealing with accreting neutron stars and black holes, we estimate the time-average accretion rate, $\\langle \\dot{M}_{\\mathrm{long}} \\rangle$, of the transients, which is an important input parameter for binary evolution models that attempt to explain the nature of subluminous X-ray transients. Our estimates lie in the range of $ 3 \\times 10^{-13}~\\mathrm{M}_{\\odot}~\\mathrm{yr}^{-1} \\lesssim \\langle \\dot{M}_{\\mathrm{long}} \\rangle \\lesssim 1 \\times 10^{-10}~\\mathrm{M}_{\\odot}~\\mathrm{yr}^{-1}$, if the systems are neutron star X-ray binaries and between $ 4 \\times 10^{-14}~\\mathrm{M}_{\\odot}~\\mathrm{yr}^{-1} \\lesssim \\langle \\dot{M}_{\\mathrm{long}} \\rangle \\lesssim 2 \\times 10^{-11}~\\mathrm{M}_{\\odot}~\\mathrm{yr}^{-1}$ for a scenario where the accreting object is a black hole. Some of the systems have such low estimated mass-accretion rates that they possibly pose a challenge for binary evolution models. } ",
        "introduction": "Our Galaxy harbors many X-ray transients that spend most of their time in a dim, quiescent state, but occasionally they experience bright X-ray outbursts (typically lasting weeks to months) during which their X-ray luminosity increases by more than a factor of 100. Many of these transient X-ray sources can be identified with compact objects (neutron stars or black holes) accreting matter from a companion star. In such systems, the X-ray outbursts are ascribed to a sudden strong increase in the accretion rate onto the compact object. X-ray transients can be classified based on their 2-10 keV peak luminosity\\footnote{All fluxes and luminosities quoted in this paper are for the 2-10 keV energy band, unless otherwise stated.}, $L_{X}^{\\mathrm{peak}}$. The \\textit{bright} X-ray transients ($L_{X}^{\\mathrm{peak}}=10^{37-39}$~erg~s$^{-1}$) have been known and extensively studied since the early days of X-ray astronomy. However, in the past decade it became clear that a group of subluminous X-ray transients ($L_{X}^{\\mathrm{peak}}<10^{37}$~erg~s$^{-1}$) also exists, where the distinction is made between \\textit{faint} \\citep[$L_{X}^{\\mathrm{peak}}=10^{36-37}$~erg~s$^{-1}$, e.g.,][]{heise99,zand01} and \\textit{very-faint} \\citep[$L_{X}^{\\mathrm{peak}}=10^{34-36}$~erg~s$^{-1}$, e.g.,][]{sidoli99, porquet05,muno05_apj622,wijn06_monit} systems. Although the faint to very faint X-ray transients exhibit qualitatively different behavior than the brighter systems \\citep[e.g.,][]{cornelisse02, okazaki01, king00}, this classification based on peak luminosities is not strict and hybrid systems are known to exist \\citep[e.g.,][]{wijn02}.  In particular the study of very-faint X-ray transients (VFXTs) is hampered by the sensitivity limitations of X-ray instruments, and consequently their nature is not understood well. To date, about 30 members are known, most of which are found very close to $\\mathrm{Sgr~A}^{\\ast}$ \\citep[within $\\sim$ 10 arcminutes; ][]{muno05_apj622}, but this might be a selection effect due to all the high-resolution X-ray observations in this region. Several VFXTs were found at larger distances from $\\mathrm{Sgr~A}^{\\ast}$ as well \\citep[e.g., ][]{hands04,heinke2008}. A significant fraction ($\\sim 1/3$) of the VFXTs have exhibited type-I X-ray bursts \\citep[e.g., ][]{cornelisse02} and can thus be identified with neutron stars accreting matter from, most likely, a low-mass (i.e., $M \\lesssim 1 \\mathrm{M_{\\odot}}$) companion. The low outburst luminosities characteristic of VFXTs combined with what is known about their duty cycles, imply that these low-mass X-ray binaries (LMXBs) have very low time-averaged mass accretion rates, which could challenge our understanding of their evolution \\citep{king_wijn06}.  There might also be other types of sources that can produce subluminous X-ray outbursts. It is conceivable that some systems are compact objects that are transiently accreting at a very low level from the strong stellar wind of a high-mass star or the circumstellar matter around a Be star \\citep[e.g.,][]{okazaki01}. In addition, some strongly magnetized neutron stars ($B \\sim10^{14}-10^{15}$~G, magnetars) are observed to experience occasional X-ray outbursts with peak luminosities of $\\sim10^{35}$~erg~s$^{-1}$ \\citep{ibrahim04,muno07_magnetar} and can thus be classified as VFXTs. The cause of their outbursts is unknown, but is likely related to magnetic field decay \\citep[e.g.,][]{ibrahim04}. Furthermore, \\citet{mukai08} recently pointed out that classical novae can be visible as 2-10 keV X-ray sources with luminosities in the range of a few times $10^{34-35}~\\mathrm{erg~s}^{-1}$ for weeks to months \\citep[see Fig.~1 of][]{mukai08}. The X-ray emission is thought to emerge from shocks within the matter that is ejected during the nova.  Here we present the analysis of seven X-ray transients, that were found active during a monitoring campaign of the Galactic center (GC) by the X-ray telescope (XRT) aboard the \\textit{Swift} satellite \\citep{kennea_monit}, carried out in 2006 and 2007. ",
        "conclusions": "\\label{sec:discuss}  \\begin{table*} \\begin{threeparttable}[t] \\begin{center} \\caption[]{{Overview of the outburst duration, fluency and estimated long-term averaged accretion rates of the seven transients detected during the \\textit{Swift}/GC monitoring campaign of 2006 and 2007.}} \\begin{tabular}{l l l l l l} \\hline \\hline Source & Year &  $t_{\\mathrm{ob}}$\\tnote{a}& \\~{F }\\tnote{b} & $\\langle \\dot{M}\\rangle _{\\mathrm{NS}}$\\tnote{c} & $\\langle \\dot{M} \\rangle _{\\mathrm{BH}}$\\tnote{d} \\\\ \\hline AX J1745.6-2901\\tnote{e} & 2006 & $>16$ &$\\gtrsim 1.8 \\times 10^{-4}$ &  &  \\\\ & 2007 & $>78$ &$\\gtrsim 6.5 \\times 10^{-3}$ &  $\\sim (5-15) \\times 10^{-11}$ &   \\\\ CXOGC J174553.5-290124 & 2006 & $>12$ &$\\gtrsim 1.6 \\times 10^{-5}$ & $\\sim (5-13) \\times 10^{-13}$ & $\\sim (7-18) \\times 10^{-14}$ \\\\ CXOGC J174540.0-290005  & 2006 & 2 & $1.3\\times 10^{-5}$ & $ \\sim (3-15) \\times 10^{-13}$ & $ \\sim (4-21) \\times 10^{-14}$ \\\\ Swift J174553.7-290347 & 2006 & 2 & $8.0\\times 10^{-6}$ &$\\lesssim 1 \\times 10^{-12}$ & $\\lesssim 2 \\times 10^{-13}$ \\\\ Swift J174622.1-290634 & 2006 &  5 & $5.0\\times 10^{-6}$ & $\\lesssim 4 \\times 10^{-13}$ &  $\\lesssim 6 \\times 10^{-14}$\\\\ GRS 1741.9-2853\\tnote{e} & 2006 & 1 & $2.6 \\times 10^{-6}$ &  &  \\\\ & 2007 & $>13$ & $\\gtrsim 1.4 \\times 10^{-3}$ & $\\sim (2-8) \\times 10^{-11}$ &  \\\\ XMM J174457-2850.3 & 2007 & $<12$ & $\\lesssim 7.5 \\times 10^{-4}$ &  $\\sim (1-10) \\times 10^{-11}$ &  $\\sim (2-20) \\times 10^{-12}$\\\\ \\hline \\end{tabular} \\label{tab:mdot} \\begin{tablenotes} \\item[a]Outburst duration in weeks. \\item[b]Outburst fluency in units of $\\mathrm{erg~cm}^{-2}$ in the 2-10 keV energy band, which was calculated by multiplying the mean unabsorbed outburst flux by the outburst duration. \\item[c]Estimated long-term averaged accretion rate ($\\mathrm{M}_{\\odot}~\\mathrm{yr}^{-1}$) for a neutron star with $M=1.4~\\mathrm{M_{\\odot}}$ and $R=10~\\mathrm{km}$ \\item[d]Estimated long-term averaged accretion rate ($\\mathrm{M}_{\\odot}~\\mathrm{yr}^{-1}$) for a black hole with $M=10~\\mathrm{M_{\\odot}}$ and $R=30~\\mathrm{km}$. \\item[e]AX J1745.6-2901 and GRS 1741.9-2853 both display type-I X-ray bursts and are thus confirmed neutron star systems. \\end{tablenotes} \\end{center} \\label{tab:obdurations} \\end{threeparttable} \\end{table*}  We have presented the spectral analysis of seven X-ray transients that were found to be active during a monitoring campaign of the field around $\\mathrm{Sgr~A}^{\\ast}$ using \\textit{Swift}/XRT, carried out in 2006-2007. Two new transients were discovered (Swift J174622.1-290634 and Swift J174553.7-290347) and renewed activity from five known sources was observed (AX J1745.6-2901, CXOGC J174553.5-290124, CXOGC J174540.0-290005, GRS 1741.9-2853 and XMM J174457-2850.3). Adopting source distances of 8 kpc, we can infer peak luminosities in the range of $\\sim 1 \\times 10^{34} - 6 \\times 10^{36}~\\mathrm{erg~s}^{-1}$ in the 2-10 keV energy band. The two transients AX J1745.6-2901 and GRS 1741.9-2853 are hybrid systems, that display very-faint outbursts with 2-10 keV peak luminosities of $L_{\\mathrm{X}}<10^{36}~\\mathrm{erg~s}^{-1}$, as well as outbursts with luminosities in the range of $10^{36-37}~\\mathrm{erg~s}^{-1}$, which are classified as faint. The other five systems display 2-10 keV peak luminosities of $10^{34-36}~\\mathrm{erg~s}^{-1}$, i.e., in the very-faint regime. We have observed a large variation in spectral properties, outburst luminosities and outburst durations (see Tables~\\ref{tab:spectra}~and~\\ref{tab:mdot}). In that respect, the subluminous transients are not different from the well-known bright systems.   \\subsection{The nature of the detected transients} AX J1745.6-2901 and GRS 1741.9-2853 are both known X-ray bursters, which makes it very likely that these are neutron stars in LMXBs, since type-I X-ray bursts have never been detected from high-mass X-ray binaries. For AX J1745.6-2901 a LMXB nature is confirmed by its orbital period of 8.4 hours. The nature of the remaining five transients is unknown. However, XMM J174457-2850.3, CXOGC J174553.5-290124 and CXOGC J174540.0-290005, all have been in outburst more than once in the past decade. This likely rules out a white dwarf accretor, since recurrent novae display outburst cycles of decades rather than years. This suggests a neutron star or black hole nature for XMM J174457-2850.3, CXOGC J174553.5-290124 and CXOGC J174540.0-290005. For CXOGC J174540.0-290005, observations reported by \\citet{wang06_atel935} could not detect a near-IR counterpart, while the observations would have detected a main sequence star down to spectral type B5. This suggests that if CXOGC J174540.0-290005 is an X-ray binary, it is likely a LMXB.  The two new transients Swift J174553.7-290347 and Swift J174622.1-290634 were observed at 2-10 keV peak luminosities of $2.0 \\times 10^{35}~\\mathrm{erg~s}^{-1}$ and $7.0 \\times 10^{34}~\\mathrm{erg~s}^{-1}$. Although such luminosities are quite uncommon for white dwarf systems, \\citet[][]{mukai08} showed a few examples of classical novae that reach peak values of several times  $10^{34-35}~\\mathrm{erg~s}^{-1}$. Thus, in absence of other outbursts from Swift J174553.7-290347 and Swift J174622.1-290634, we cannot exclude the possibility that these two systems harbor accreting white dwarfs.   \\subsection{subluminous X-ray transients in quiescence}\\label{quiescence} The quiescent luminosity of X-ray transients sometimes holds clues to the nature of the system. The \\textit{ASCA} burster AX J1745.6-290 is very likely associated with CXOGC J174535.6-290133, which was detected several times with \\textit{Chandra} at a level of a few times $10^{32}~\\mathrm{erg~s}^{-1}$ (see Sect.~\\ref{subsec:bron1}). This is consistent with the neutron star nature of AX J1745.6-290, since black hole systems with an orbital period of $\\sim$~8~hours are significantly fainter \\citep[e.g.,][]{narayan97,menou99,lasota07}. GRS 1741.9-2853 is also a confirmed neutron star system and displays a similar quiescent level of $\\sim 10^{32}~\\mathrm{erg~s}^{-1}$ \\citep[2-8 keV,][]{muno03_grs}.  The possible quiescent counterpart of the new subluminous X-ray transient Swift J174553.7-290347, the \\textit{Chandra}-detected X-ray source CXOGC J174553.8-290346, displays a 2-10 keV X-ray luminosity of $\\sim 7 \\times 10^{31}-4 \\times 10^{32}~\\mathrm{erg~s}^{-1}$, depending on the assumed spectral model (see Sect.~\\ref{subsec:bron4}). The quiescent luminosity of XMM J174457-2850.3 is also in this regime; $\\sim 10^{32}~\\mathrm{erg~s}^{-1}$ \\citep[2-10 keV;][]{sakano05}. If Swift J174553.7-290347 and XMM J174457-2850.3 are X-ray binaries, their quiescent luminosities are relatively high and might point towards a neutron star nature \\citep[e.g.,][]{lasota07}, athough the orbital period of both these systems is unknown. We note that the absorption towards our transients is very high ($> 6 \\times 10^{22}~\\mathrm{cm}^{-2}$). Therefore, any thermal emission from the neutron star surface cannot be observed and we can only detect contributions from a powerlaw component, which is frequently observed for neutron stars at similarly low quiescent luminosities \\citep[e.g.,][]{jonker07_eos}.  Two other transients, CXOGC J174553.5-290124 and CXOGC J174540.0-290005, were not detected in quiescence, but have upper limits on their luminosities of $< 9 \\times 10^{30}~\\mathrm{erg~s}^{-1}$ and $< 4 \\times 10^{31}~\\mathrm{erg~s}^{-1}$ respectively \\citep[2-8 keV;][]{muno05_apj622}. Such low quiescent luminosities are more common for black hole X-ray binaries than for neutron star systems \\citetext{e.g., \\citealp{garcia01,lasota07}, but see \\citealp{jonker06,jonker07}}.   \\subsection{The outbursts of subluminous X-ray transients}\\label{subsec:outbursts} The disk instability model \\citep[e.g.,][]{king98, dubus99, lasota01} provides a framework to describe the outburst cycles of transient LMXBs. However, it is unclear why some X-ray transients, such as the ones discussed here, undergo outbursts with very low peak luminosities. AX J1745.6-2901 has a known orbital period of 8.4 hours, which allows for a maximum luminosity of $\\sim 2 \\times 10^{38}~\\mathrm{erg~s}^{-1}$ \\citep[for a hydrogen-dominated disk and a neutron star mass of $1.4~\\mathrm{M_{\\odot}}$;][]{lasota07}. Yet, its observed peak luminosity is over an order of magnitude lower (see Table~\\ref{tab:spectra}). Since AX J1745.6-2901 displays eclipses, we must look at the system at high inclination. For several eclipsing X-ray binaries observations suggest that these are intrinsically bright but appear faint because the bright center of the system is blocked by the outer edge of the disk and the corona \\citep[e.g.,][]{parmar00, kallman03,muno05_apj633}. This may also be the case for AX J1745.6-2901, for which \\citet{maeda1996} derived an inclination angle of $i \\sim 70^{\\circ}$. To include inclination effects, the observed X-ray luminosity should be corrected by a factor $\\xi_{p}$, which relates to the inclination, $i$, as $\\xi_{p}^{-1}= 2 |\\cos i|$ \\citep{fujimoto88,lapidus85}. In 2007, AX J1745.6-2901 displayed a 2-10 keV peak luminosity of $6.1 \\times 10^{36}~\\mathrm{erg~s}^{-1}$, which corrects to $9.2 \\times 10^{36}~\\mathrm{erg~s}^{-1}$ for the suggested inclination of $i \\sim 70^{\\circ}$ ($\\xi_{p} \\sim 1.5$).  It is thus conceivable that AX J1745.6-2901 is a bright X-ray transient that is obscured due to line of sight effects, although it would still seem to be at the lower end of the luminosity range for bright systems (peak luminosities of $\\sim 10^{37-39}~\\mathrm{erg~s}^{-1}$ in the 2-10 keV energy band). For comparison, the quasi-persistent neutron star system MXB 1659-29 has an orbital period of 7.1 hours \\citep{mxb1659_eclipses}, which is close that of AX J1745.6-2901. However, MXB 1659-29 displays an average 2-10 keV outburst luminosity of $7 \\times 10^{36} ~(d/10 ~\\mathrm{kpc})~\\mathrm{erg~s}^{-1}$ \\citep[][]{oosterbroek01_mxb,sidoli01_mxb}, which is about a factor of 4 higher than the average 2-10 keV outburst luminosity observed for AX J1745.6-2901 in 2007; $1.6 \\times 10^{36} ~\\mathrm{erg~s}^{-1}$. Possibly, the inclination of AX J1745.6-2901 is somewhat higher than the $i \\sim 70^{\\circ}$ suggested by \\citet{maeda1996}. AX J1745.6-2901 might have a subluminous appearance due to line of sight effects, but it is important to note that statistical arguments show that such effects cannot account for the entire population of subluminous X-ray transients, and that most systems must have low intrinsic luminosities \\citep[see the discussion of][]{wijn06_monit}.   Although taking into account inclination effects potentially pushes AX J1745.6-2901 into the regime of bright X-ray transients, this does not provide an explanation for the peculiar outburst behavior of the source. As discussed in Sect.~\\ref{subsec:bron1}, the system was likely in quiescence for several years before it was seen active in 2006 for more than 4 months. At that time, the source reached a peak luminosity of $9.2 \\times 10^{35}~\\mathrm{erg~s}^{-1}$, which would classify the system as very-faint. However, after several months of quiescence (see Fig.~\\ref{fig:lc}), the source reappeared displaying a peak luminosity of $6.1 \\times 10^{36}~\\mathrm{erg~s}^{-1}$ (i.e., in the faint regime) and remained active for over 1.5 years (see Sect.~\\ref{subsec:bron1}). Thus, the outburst observed in 2006 was subluminous by about a factor of 6 compared to the 2007 outburst, yet the system maintained this low luminosity for months. It is unclear if this behavior can be explained in terms of a disk instability model. In 1993 and 1994, different outburst luminosities of $2 \\times 10^{35}~\\mathrm{erg~s}^{-1}$ and $9 \\times 10^{35}~\\mathrm{erg~s}^{-1}$ were reported for AX J1745.6-2901 \\citep[3-10 keV,][]{maeda1996}. This is on the same time scale as the discussed \\textit{Swift} detections, suggesting that the behavior observed in 2006 and 2007 could be typical for the source.  GRS 1741.9-2853 also displayed two separate outbursts with very different characteristics in terms of peak luminosity and outburst duration during the \\textit{Swift}/XRT monitoring observations. A short, $\\sim1-$week outburst was detected in 2006, which had a 2-10 keV peak luminosity of $9.2 \\times 10^{34}~\\mathrm{erg~s}^{-1}$. A few months later, the source exhibited a much longer ($\\gtrsim 13$~weeks) outburst, that reached a peak luminosity of $2.0 \\times 10^{36}~\\mathrm{erg~s}^{-1}$ (2-10 keV). Possibly, the short 2006 outburst of GRS 1741.9-2853 was an X-ray precursor for the 2007 outburst. Such behavior is observed for several bright X-ray transients \\citep[both neutron star and black hole systems, see][and references therein]{chen97}. Both Swift J174553.7-290347 and CXOGC J174540.0-290005 displayed short, $\\sim$ 2-week outbursts that had an average luminosity of a few times $10^{34}~\\mathrm{erg~s}^{-1}$. This kind of activity resembles the small accretion outburst of GRS 1741.9-2853 in 2006 (see Fig.~\\ref{fig:lc}), but for these two systems no longer outbursts have been observed. XMM J174457-2850.3 seems to undergo X-ray activity at different luminosity levels as well (see Sect.~\\ref{subsec:bron7}). It is unclear what causes the varying accretion luminosities. However, this phenomenon is also observed for bright X-ray transients and is thus not restricted to the subluminous systems discussed here.  Current disk instability models do not provide an obvious explanation for accretion outbursts that last several years, rather than the usual weeks to months, such as observed for AX J1745.6-2901. A few bright systems are known to undergo quasi-persistent outbursts \\citep[see e.g.,][]{wijnands04_quasip}. There are also two X-ray transients that exhibit prolonged outbursts at low luminosities. XMMU J174716.1-281048 has likely been continuously active since its initial discovery in 2003, displaying a typical  2-10 keV luminosity of a few times $10^{34}~\\mathrm{erg~s}^{-1}$  \\citep[e.g.,][]{delsanto07,degenaar07_xmmsource}. Furthermore, AX J1754.2-2754 recently made a transition to quiescence \\citep{bassa08}, after exhibiting an accretion outburst with a 2-10 keV luminosity of several times $10^{34-35}~\\mathrm{erg~s}^{-1}$, which likely lasted for 7-8 years \\citep[][]{sakano02,delsanto07_ascabron,chelovekov07_ascabron}. This source was again found active in July 2008 \\citep{jonker08}. The detection of type-I X-ray bursts identifies both these systems as neutron star LMXBs, just like AX J1745.6-2901.   \\subsection{X-ray bursts from subluminous X-ray transients} The properties of type-I X-ray bursts are set by the conditions in the flash layer such as the temperature, thickness, hydrogen abundance and the fraction of carbon-nitrogen-oxygen (CNO) elements in the layer \\citep[e.g.,][]{fujimoto81,bildsten98,peng2007}. These conditions can vary drastically as the mass-accretion rate onto the neutron star ($\\dot{M}$) varies, which results in flashes with different characteristics for different $\\dot{M}$ regimes \\citep[e.g., ][]{fujimoto81,peng2007,cooper07}.   The \\textit{Swift}/XRT monitoring observations of 2006 caught two type-I X-ray bursts from AX J1745.6-2901 (see Sect.~\\ref{subsec:bron1}). The average 2-10 keV luminosity of the 2006 outburst was $3.9 \\times 10^{35}~\\mathrm{erg~s}^{-1}$, from which we can estimate an instantaneous mass accretion rate onto the neutron star of $\\sim1 \\times 10^{-10}~\\mathrm{M_{\\odot}~yr}^{-1}$. If we include a correction factor to account for inclination effects, as discussed in Sect.~\\ref{subsec:outbursts}, this value increases to $\\sim1.5 \\times 10^{-10}~\\mathrm{M_{\\odot}~yr}^{-1}$. The bursts had a duration of $50-60~\\mathrm{seconds}$ (see Fig.~4), which suggests triggering in a mixed hydrogen/helium environment. This is in line with the classical predictions for the estimated mass-accretion rate \\citep[e.g.,][]{fujimoto81}.  We discussed in Sect.~\\ref{subsec:bron1}, that the type-I X-ray bursts observed from AX J1745.6-2901 have 0.01-100 keV peak luminosities of $\\sim 10^{38}~\\mathrm{erg~s}^{-1}$, close to the Eddington limit of a neutron star. These peak values should also be corrected for the inclination effects discussed in Sect.~\\ref{subsec:outbursts}. However, the X-ray burst (originating form the neutron star surface) and the outburst emission (emerging from the accretion disk) are attributed to geometrically different regions and may therefore have different degrees of isotropy \\citep{fujimoto88,lapidus85}. Although, the X-ray burst emission will be partly intercepted and re-radiated by the accretion disk, it was shown that the degree of anisotropy is less than for the emission coming from the accretion disk \\citep{fujimoto88,lapidus85}. Inclination effects are expected to reduce the X-ray burst emission by a factor $\\xi_{b}^{-1}= 0.5 + |\\cos i |$ \\citep{fujimoto88,lapidus85}. For the suggested inclination of $i=70^{\\circ}$ \\citep{maeda1996}, we thus obtain a correction factor of $\\xi_{b}=1.2$. This implies peak luminosities for the type-I X-ray bursts observed from AX J1745.6-2901 on June 3 and June 14, 2006, of  $(1.2-1.6) \\times 10^{38}~\\mathrm{erg~s}^{-1}$ and $(0.62-1.3) \\times 10^{38}~\\mathrm{erg~s}^{-1}$ (0.01-100 keV) respectively.  This is below, but close, to the Eddington luminosity for a neutron star \\citep[$2.0 \\times 10^{38}\\mathrm{~erg~s}^{-1}$ for a hydrogen-rich and $3.8 \\times 10^{38}\\mathrm{~erg~s}^{-1}$ for a hydrogen-poor photosphere; e.g.,][]{kuulkers03_xrb}.    \\subsection{Long-term average accretion rates}\\label{mdot_estimate} Presuming that the detected transients are accreting systems, we attempted to estimate their long-term time-averaged accretion rates using the method and described in Sect.~\\ref{subsec:accrates}. We explored the scenarios of both neutron star and black hole accretors (except for AX J1745.6-2901 and GRS 1741.9-2853, since these are confirmed neutron star systems), which resulted in the estimated long-term mass-accretion rates\\footnote{Note the caveat mentioned for the black hole cases in Sect.~\\ref{subsec:accrates}.} listed in Table~\\ref{tab:obdurations}. The two confirmed neutron star systems, AX J1745.6-290 and GRS 1741.9-2853 have the highest estimated accretion rates of the seven discussed transients ($\\sim 10^{-11}-10^{-10}~\\mathrm{M_{\\odot}~yr}^{-1}$). This arises from the fact that GRS 1741.9-2853 is active quite regularly and AX J1745.6-290 can be in outburst for a very long time (over 1.5~years). The regime estimated for these two sources can be well explained within current LMXB evolution models. The same is likely true for XMM J174457-2850.3, which was active several times since its discovery in 2001 and has an estimated long-term mass-accretion rate of $\\gtrsim 10^{-11}~\\mathrm{M_{\\odot}~yr}^{-1}$.   The estimates for the remaining four systems, CXOGC J174553.5-290124, CXOGC J174540.0-290005, Swift J174553.7-290347 and Swift J174622.1-290634 are much lower; $\\lesssim 10^{-12}~\\mathrm{M_{\\odot}~yr}^{-1}$ for accreting neutron stars and even an order of magnitude lower for black hole X-ray binaries, $\\lesssim 10^{-13}~\\mathrm{M_{\\odot}~yr}^{-1}$ (see Table~\\ref{tab:obdurations}). Comparing our results with a theoretical toy-model of \\citet{king_wijn06}, who explored the mechanism of Roche-lobe overflow at low accretion rates, suggests that if these transients are LMXBs, their low time-averaged mass-accretion rates might pose difficulties explaining their existence, without invoking exotic scenarios such as accretion from a planetary donor or an intermediate mass black hole as the accreting primary \\citep{king_wijn06}. These are thus interesting systems to track and monitor in the future.  Apart from evolutionary scenarios and line-of-sight effects, there other possible explanations for the subluminous X-ray appearance of these transients. For example, in particular for the systems containing a black hole, the liberated accretion power may not be primarily dissipated as X-rays but rather via radiatively inefficient flows \\citep[e.g.,][]{blandford99,fender03,narayan}. Furthermore, in neutron star systems the \"propeller mechanism\" can possibly operate, so that only a small fraction of the mass transferred from the donor can be accreted onto the neutron star~\\citep[e.g.,][]{illarionov1975,alpar2001,romanova2005}.  The discussed examples of AX J1745.6-2901, GRS 1741.9-2853 and XMM J174457-2850.3 illustrate that X-ray transients can display different behavior in terms of peak luminosity, outburst duration and recurrence time from year to year. It is currently not understood whether these variations should be interpreted as, e.g., being due to changes in the mass-transfer rate from the donor star or as the result of instabilities in the accretion disk. Such issues need to be resolved before we can fully comprehend the nature of subluminous X-ray transients."
    },
    "0807/0807.2562_arXiv.txt": {
        "abstract": "Sources of gravitational waves (GW) and emitters of high energy (HE) neutrinos both involve compact objects and matter moving at relativistic speeds. GW emission requires a departure from spherical symmetry, which is the case if clumps of matter are accreted around black holes or neutron stars, and ejected in relativistic jets, where neutrinos are believed to be produced. Both messengers interact weakly with the surrounding matter, hence point directly to the heart of the engines that power these emissions. Coincidences between GW interferometers ({\\it e.g.} \\vo) and HE $\\nu$ telescopes ({\\it e.g.} \\anta) would then give a unique insight on the physics of the most powerful objects in the Universe. The possibility, observability and detectability for such GW/HE $\\nu$ coincidences are analysed. ",
        "introduction": "\\label{sec:intro}  The forthcoming years should be very exciting both in gravitational wave (GW) astronomy and high energy neutrino (HE $\\nu$) astronomy. The \\vo~interferometer~\\cite{virgo}, currently closed down for upgrade, should be taking data with an improved injection system in 2009, whereas the \\anta~collaboration has completed the deployment and connections of its 12 lines~\\cite{antares} early in June 2008, starting the operation of the first {\\it undersea} neutrino telescope. In the mean time, the two \\lo~interferometers (ITF) are in operation~\\cite{ligo}, and {\\sc IceCube} is deploying its 1~km$^3$ neutrino telescope in the ice of the South Pole, already having 40 lines in data taking mode~\\cite{icecube}.  Both GW sources and HE $\\nu$ emitters involve compact objects and matter moving at relativistic speeds. As a result, coincidences between GW interferometers and neutrino telescopes can be envisaged~\\cite{coinc}, and should then be feasible by 2009, with the joint operation of \\vo, \\lo, \\anta~and {\\sc IceCube}. Together, they would give a unique insight on the physics of the most powerful objects in the Universe. Some classes of astrophysical object, invisible in electromagnetic channels, may be observable only {\\it via} their  gravitational waves and high energy neutrino emissions. Finally, in many quantum gravity (QG) models~\\cite{qg}, the propagation velocity of a particle depends on the energy:  \\begin{equation} c^2 p^2 = E^2 \\left[1 + \\xi \\left( \\frac{E}{E_{\\textrm{{\\tiny QG}}}} \\right) + {\\mathcal O}\\left( \\frac{E^2}{E_{\\textrm{{\\tiny QG}}}^2} \\right)+ \\ldots \\right] \\label{eq:intro_qg1} \\end{equation}  \\noindent where $\\vert \\xi \\vert \\simeq 1$, and $E_{\\textrm{{\\tiny QG}}}$ is the energy scale at which QG effects arise. Hence, measuring a non-zero time delay between gravitational wave bursts and high energy neutrino transients would then allow to probe Quantum Gravity effects at the Planck energy level.  Such time coincidences require GW bursts (localized in time), and would allow GW antennae to confirm a burst detection, and neutrino telescopes to sign the cosmic origin of the signal. It must be first demonstrated that such a coincidence process is:  \\begin{itemize} \\item[$\\bullet$] possible: sources of both GW and HE $\\nu$, able to emit signals possibly coincident in time, exist in the Galaxy. This is discussed in section~\\ref{sec:sources}; \\item[$\\bullet$] observable: the visibility sky maps of \\vo~and \\anta~are not orthogonal. This point is addressed in section~\\ref{sec:obs} ; \\item[$\\bullet$] detectable: the resulting coincidence detection probability, for a fixed accidental coincidence rate, which constrains the authorized background level for each individual detector, must be estimated. This point is developed in section~\\ref{sec:detec}. \\end{itemize} ",
        "conclusions": "Time coincidences between GW interferometers, such as \\vo, currently down for upgrade, and high energy neutrino telescopes, such as \\anta, now fully operational, are thus feasible: common sources able to emit coincident signals in both channels exist in our own galaxy. Microquasars, during major outbursts, and flares from SGRs are possible targets. \\anta~and \\vo~visibility sky maps are not othogonal, allowing for coincident detections. Taking into account the relationship between detection efficiencies and false alarm rates in \\anta~and \\vo, coincident observations appear possible.  Such coincidences can be performed using \\anta~5-lines and {\\it VSR} data (end of 2007) as a test that could be performed before the upgrade of \\vo~next year, which should improve the sensitivity at low and high frequency to reach the design sensitivity. This {\\sc Virgo+} upgrade will correspond with the {\\it routine} operation of the full \\anta~detector, provinding the opportunity to perform the time coincidences presented in this paper.  Finally, {\\it circa} 2015, a km$^3$ neutrino telescope should be operating in the Mediterranean Sea~\\cite{km3}, along with an {\\sc Advanced Virgo} interferometer~\\cite{adv_virgo}, with enhanced sensitity at low frequency. Figure~\\ref{fig:adv_virgo}~\\cite{mypaper} shows the ejected mass needed in a microquasar ejection event to obtain a SNR = 5 in {\\sc Virgo+} and {\\sc Advanced Virgo}~\\cite{adv_virgo} as a function of the acceleration time (see section~\\ref{sec:sources}): less extremes ejection scenarios could be probed, and interesting constraints on accretion/ejection models could be set by this novel multimessenger approach.  \\begin{figure}[h!] \\centerline{\\includegraphics[width=\\linewidth]{figures/snr5_adv_virgo.png}} \\caption{Ejected mass needed for SNR=5 for a microquasar at 1 kpc for {\\sc Virgo+} and {\\sc Advanced Virgo}. \\label{fig:adv_virgo}} \\end{figure}  \\noindent \\textbf{Aknowledgements : } \\textit{I would like to thank the Organizing Committee, especially John Carr, for trusting me for this plenary talk.}"
    },
    "0807/0807.3614_arXiv.txt": {
        "abstract": "{ Among the tools available for the study of the dark energy driving the expansion of the Universe, Baryon Acoustic Oscillations (BAO) and their effects on the matter power spectrum are particularly attractive. It was recently proposed to study these oscillations by mapping the  21cm emission of the neutral hydrogen in the redshift range $0.5<z<3$. We discuss here the precision of such measurements using radio-interferometers consisting of arrays of dishes or north-south oriented cylinders. We then discuss the resulting uncertainties on the BAO scales and the sensitivity to the parameters of the Dark Energy equation of state. } ",
        "introduction": "One of the goals of current cosmological research is to fully characterize the ``dark energy'' that drives the apparent acceleration of the expansion of the Universe. Among the available tools \\citep{DETF}, much interest has been generated by the features in the matter power spectrum that result from Baryon Acoustic Oscillations (BAO) in the pre-recombination Universe.  These oscillations cause a peak in the matter correlation function at a comoving distance, $s\\sim105h^{-1}Mpc$, equal to the acoustic horizon at recombination. Equivalently, the BAO induce ``wiggles'' in the matter power spectrum with peaks at comoving wave numbers $k=(n+1/2)\\pi/s$ for $n=2,4,....$. The theoretical  positions of these peaks  are known to a precision of order 1\\% so they can be used as  reliable standard rulers to study the expansion history.  After first being seen in the CMB anisotropy spectrum \\citep{boomerang,wmap}, BAO effects were subsequently observed in the low-redshift ($z<1$) galaxy correlation function and power spectrum by the SDSS survey \\citep{BAO2,SDSSPk} and the 2dGFRS survey \\citep{BAO1}. With future data in the redshift range $0.5<z<3$, it is hoped to provide precise constraints on dark-energy parameters. It has been proposed to do this with redshift surveys using galaxy optical spectra \\citep{wfmos} or HI emission \\citep{SKA} or by mapping in three dimensions the pattern of Lyman-$\\alpha$ absorption of distant quasars \\citep{lyalpha}.  An alternative elegant approach was recently proposed by \\citet{HSHSMoriond} and further developed by \\citet{HSHSPRL}. They propose to use neutral hydrogen (HI) as a tracer of matter and simply map out the 21cm emission with an angular resolution that is insufficient to detect individual galaxies (where most HI is concentrated) but sufficient to observe the BAO wiggles. At redshift $z=1.5$, the acoustic horizon subtends 0.0334rad on the sky which can be resolved with telescopes of size of order 100m. These would be considerably smaller and cheaper than the 1km elements necessary to detect individual high redshift galaxies \\citep{SKA}.   Mapping the matter distribution using HI 21 cm emission as a tracer has been extensively discussed in literature (e.g. \\citet{Furlanetto,fftt}). Several projects, such as LOFAR \\citep{Rottgering} or MWA \\citep{Bowman} aim at detecting the reionization epic ($z\\sim10$). Detecting BAO features around $z\\sim1$ using HI radio emission has also been discussed by \\citet{Wyithe}.   In Section \\ref{pssec} we review the expectations for the HI power spectrum $P_{HI}(k)$ and its relation to the large-scale structure matter power spectrum $P_{LSS}(k)$. Section \\ref{datasec} describes the interferometric observations that are assumed for this study and Section \\ref{recsec} presents a simplified procedure to use these observations to reconstruct the HI power spectrum in a small volume near the zenith. While this procedure is not entirely realistic, it allows us to give estimates of the reconstructed noise power spectrum, $P_{noi}(k)$ (Fig. \\ref {powerfig}). Section \\ref{radsourcesec} discusses the problem of subtraction of foreground and background radio sources. Finally, in Section \\ref{cosmosec} we describe the constraints on the cosmological parameters that can be derived from the determination of the BAO peaks in the power spectrum. ",
        "conclusions": "In this paper we have discussed the measurement of the HI power spectrum with interferometric surveys. By presenting a simplified procedure for reconstructing the Fourier space map of the GHz sky, we derived expressions for the noise power spectrum.  Adding this spectrum to the expected HI power spectrum, we determined with what precision the positions of the BAO peaks can be measured. This led to  a sensitivity to cosmological parameters that is competitive with other BAO projects.  In calculating the noise power spectrum, we considered three interferometer arrays using cylinders or dishes. Packed-cylinder and packed-dish arrays have comparable noises so the choice for one or the other must be based on cost and technical considerations, with cylinder arrays having simpler mechanical design but requiring more receivers."
    },
    "0807/0807.3752_arXiv.txt": {
        "abstract": "The metal-poor post-AGB star HD~46703 is shown to be a single-line spectroscopic binary with a period of 600 days, a high velocity of $-$94 km~s$^{-1}$, and an orbital eccentricity of 0.3.  Light curve studies show that it also pulsates with a period of 29 days. High$-$resolution, high signal-to-noise spectra were used for a new abundance study.  The atmospheric model determined is T$_{eff}$ = 6250 K, log $\\it g$ = 1.0, V$_t$ = 3.0 km~s$^{-1}$, and a metal abundance of [M/H] = $-$1.5.  A low carbon abundance and lack of s-process element enhancement indicate that the star has not experienced third dredge-up on the AGB.  The sulfur and zinc abundances are high compared with iron, and the chemical abundances show a clear anti$-$correlation with condensation temperature.  The abundance depletion pattern is similar to that seen in other post$-$AGB binaries, and, like them, is attributed to the chemical fractionation of refractory elements onto dust stored in a circumbinary disk and the re-accretion of volatiles in the stellar atmosphere. The infrared excess is small but the excess energy distribution is very similar to what can expected from a disk. HD~46703 joins the growing list of depleted, post-AGB stars which are likely surrounded by a dusty and stable circumbinary disk. ",
        "introduction": "During the past decennia, it has been realized that post-AGB stars are chemically much more diverse than anticipated. In the canonical picture, the photospheric content of a post-AGB star is the end-product of the chemical enrichment induced by the thermal pulses of AGB stars, and as such post-AGB stars are ideal tracers for the 3rd dredge-up nucleosynthesis. Some objects are the most s-process enriched objects known to date \\citep[e.g.,][and references therein]{vanwin03}, while others are not enriched at all. This dichotomy is very strict, in the sense that mildly enhanced objects are not often observed  \\citep[and references therein]{vanwin03}. For the moment it is not clear whether this dichotomy is caused by a selection effect, or if it is the result of two different evolutionary channels.  In this study, we focus on HD~46703. More than 20 years ago, Bond and Luck drew particular attention to HD~46703 with their abundance studies \\citep{luc84,bon87}.  They found the object to be iron-poor, [Fe/H] = $-$1.6, but relatively overabundant in carbon ([C/Fe]=+1.0), nitrogen ([N/Fe]=+1.8), and oxygen ([O/Fe]=+1.1) and underabundant in s-process elements (Y,Zr,Ba; ([s-process/Fe]=$-$0.6 to $-$0.8).  This combination is hard to explain.  The atmosphere model they determined was T$_{\\rm eff}$ = 6000$\\pm$150 K, log ${\\it g}$ = 0.4$\\pm$0.4, and V$_{\\rm turb}$ = 3.5$\\pm$0.5 km~s$^{-1}$.  They also found that it was a high velocity star, with V$_{\\rm r}$ = $-$106 km~s$^{-1}$.  On this basis, they concluded that HD~46703 was a low-mass, post-AGB star. Surprisingly, they also found a strong overabundance of sulfur, [S/Fe] = +1.2, even though the other iron-peak elements had similar abundances to iron \\citep{bon87}.  HD~46703 was also shown to be a light variable, with variations peak-to-peak of 0.22 mag over a 12 year period \\citep[1966$-$1979;][]{luc84} and 0.1 mag over three 1$-$2 week observing intervals 1981$-$1982 \\citep{bon84}.  The star was bluer when brighter. On this basis, it was classified as a low-amplitude, semi-regular variable and has been assigned the name V382 Aur.  \\citet{wat92} included HD~46703 among a small group of post-AGB objects for which they proposed that the selective depletion in the stellar atmosphere arose from the formation of a circumstellar dust disk, in which refractory elements are tied up in the dust and then the volatiles are selectively re-accreted in the atmosphere of the star. Since then, atmospheric depletion patterns have been detected in many more objects (see the recent papers by \\citet{gir05, maa07, rey07b} and references therein). Although the process itself is still not well explored theoretically, it seems to be a very common phenomenon among Galactic post-AGB stars as well as in their peers in the LMC.  The original sample of four extremely depleted objects turned out to be binaries \\citep{vanwin95} in which there was strong evidence for the presence of a circumbinary disk. Many of the objects subsequently determined to have the chemical depletion anomaly were found to be binaries as well \\citep[and references therein]{vanwin03}.  Using N-band interferometry, \\citet{der06} and \\citet{der07} were able to resolve the disks around some depleted objects, and the main conclusion of their analysis was that the circumstellar disks must be very compact and stable.  It is as yet still unclear whether all such depleted stars are binaries surrounded by a circumstellar disk \\citep{ruy06}, or whether there are other circumstances in which depleted photospheres can be formed \\citep{gir05,maa07}.  For large-amplitude pulsators, standard chemical studies based on static atmospheres fall short, and \\citet{rao05} and \\citet{gir05} link the detection of chemical anomalies to a not-well-documented effect in which an anticorrelation between [X/H] and the first ionization potential is created. For pulsational stars with a significant photometric amplitude, any detection of a chemical anomaly should be studied in more detail and preferably using spectra over the whole pulsation cycle.  In this study, we therefore present a study of HD\\,46703 which combines a detailed chemical study using high-resolution optical spectroscopy with a critical discussion of our photometric and radial velocity monitoring results. This object is an ideal one in which to study in detail the chemical anomalies and to confront those with the pulsational analysis, the nature of the circumstellar material, and the binary connection of the depletion process. The paper is organized as follows: we first analyze the radial velocity to investigate binarity, then the light curves to investigate pulsation, and then the high-resolution spectra to determine the chemical abundance. Finally we discuss the evolutionary nature of the object. ",
        "conclusions": "HD~46703 has a high velocity (at galactic latitude +20$\\arcdeg$), a low gravity, and a low carbon abundance, all of which point to it being an evolved, low-mass, post-AGB star. The chemical depletion pattern indicates the presence of a circumbinary disk, for which there is also evidence from the observed intrinsic polarization \\citep{tra94}. The depletion of Fe and other Fe-peak elements make the estimate of the initial metallicity difficult. Assuming the photospheric Zn abundance was not affected by photospheric depletion, the initial metallicity was about $-$1.0.  \\citet{ruy06} present a systematic spectral energy distribution study of post-AGB stars around which disks are suspected.  This whole sample is  a significant fraction of all known post-AGB stars in the Galaxy \\citep[e.g.,][]{szc07}. Recently we found similar objects in the Large Magellanic Cloud \\citep[e.g.,][]{rey07b}. HD~46703 is among the minority of this disk sample that show no clear signature of a near-IR excess \\citep[Figure A.1.]{ruy06}. In most cases, dust is present at the sublimation temperature, which is not the case in HD~46703. This suggests a larger disk radius. Nonetheless, the disk is likely only resolvable by interferometric means \\citep{der07}.  HD~46703 is by no means the only disk post-AGB star for which the binary orbit is determined \\citep[e.g.,][and references therein]{vanwin07}. One of the surprising results of all the orbital studies so far is the fact that many objects survived their phase of strong binary interaction in rather wide orbits. Moreover, many systems, like HD~46703, are eccentric, despite the fact that tidal effects as well as orbital-energy dissipation (during a likely common envelope phase) are expected to rapidly circularize the orbits. Episodic mass loss caused by periastron mass transfer in an elliptic orbit has been proposed as a main mechanism \\citep{sok00,bon08}, but also disk-binary interactions may be strong enough to pump the eccentricity \\citep{art91}. The latter mechanism is only efficient when the dust extends outward to a distance of only a few times the major axis of the orbit. For HD 46703, the lack of a near-IR excess points to the lack of dust at the sublimation radius and thus a larger disk radius. Orbital pumping by the actual disk is not likely the main eccentricity pumping mechanism for HD~46703.  The global picture that emerges of all those objects is that a binary star evolved in a system that is (presently) too small to accommodate a full grown AGB star. During a badly understood phase of strong interaction, a circumbinary dusty disk was formed, but the binary system did not suffer a dramatic spiral in. What we observe now is an F$-$K % supergiant in a binary system that is surrounded by a circumbinary dusty disk in a bound orbit. Understanding the evolution of the disks, as well as their impact on the evolution of the central stars, become important challenges in the study of the final evolution of a significant fraction of binary stars."
    },
    "0807/0807.1561_arXiv.txt": {
        "abstract": "We present radiation transfer models that demonstrate that reflected light levels from three dimensional (3D) exoplanetary atmospheres can be more than 50\\% lower than those predicted by models of homogeneous or smooth atmospheres. Compared to smooth models, 3D atmospheres enable starlight to penetrate to larger depths resulting in a decreased probability for the photons to scatter back out of the atmosphere before being absorbed. The increased depth of penetration of starlight in a 3D medium is a well known result from theoretical studies of molecular clouds and planetary atmospheres. For the first time we study the reflectivity of 3D atmospheres as a possible explanation for the apparent low geometric albedos inferred for extrasolar planetary atmospheres. Our models indicate that 3D atmospheric structure may be an important contributing factor to the non-detections of scattered light from exoplanetary atmospheres. We investigate the self-shadowing radiation transfer effects of patchy cloud cover in 3D scattered light simulations of the atmosphere of HD209458b. We find that, for a generic planet, geometric albedos can be as high as 0.45 in some limited situations, but that in general the geometric albedo is much lower. We conclude with some explanations on why extrasolar planets are likely dark at optical wavelengths. ",
        "introduction": "Since the discovery of the planet orbiting 51~Peg over a decade ago \\citep{may95}, the detection and characterization of extrasolar planets has become an area of intense study. To date, over two hundred and seventy extrasolar planets have been discovered primarily by the measurement of the doppler shifting of stellar spectral lines due to the motion of the star around the centre of mass of the exoplanetary system. Other indirect techniques that have succesfully detected extrasolar planets include transit searches and microlensing (e.g. \\cite{pon06,bea06}). The first detection of direct light from an extrasolar planet was the infrared emission due to reprocessed starlight in the warm atmospheres of the planets orbiting TrES-1 and HD 209458 \\citep{cha05,dem05}. Additionally, a handful of planets may have been detected using direct imaging \\citep{neu05}  Many extrasolar planets orbit within a few astronomical units of their parent star. The direct detection of reflected starlight from exoplanetary atmospheres is therefore a realistic goal (e.g., \\citet{sea00}). However, despite intense effort observers have been unable to detect stellar light reflected from extrasolar planetary atmospheres. \\citet{cam99} and \\citet{cha99} developed techniques that exploit the doppler effect to separate the spectral lines of an extrasolar planet from the lines of the parent star, but both have been unable to definitively detect the faint signature. A detection of reflected light from the planet orbiting Tau Bootis by \\citet{cam99} was unverifiable in the following observing season \\citep{cam04}. Subsequent work has set an upper limit of 0.39 for the planet's geometric albedo, ($A_g$), assuming a grey albedo and Venus-like phase function \\citep{lei03}. Using similar methods the group found the geometric albedo of the companion to HD 75289 has an upper limit of $A_g < 0.16$ \\citep{lei03a}. Analysis of data from the {\\it MOST} satellite has set an optical (400-700 nm) limit of $A_g < 0.17$ on HD~209458b \\citep{row07}.  These non-detections of reflected light and upper limits on the geometric albedo imply that compared to solar system planets, the extrasolar planets are very dark at optical and near-infrared wavelengths. In comparison, the geometric albedos of solar system gas giant planets range from 0.41 to 0.52 \\citep{cox00}. Particles in the exoplanetary atmospheres could have very low scattering albedos or the particles could have scattering phase functions that are very forward throwing so that starlight is directed deep into the atmosphere from where it has a small probability of scattering back out. \\cite{sud00} explored parameters which can influence the geometric albedo of an extrasolar planet, with a one-dimensional radiation transfer model. The focus of our paper is the effect of three-dimensional atmospheric structure on the radiation transfer of starlight and hence the overall reflectivity of exoplanetary atmospheres.  Previous studies have shown the importance of three-dimensional radiation transfer effects on the penetration of starlight into dark clouds \\citep{boi90}, the analysis of reflection nebulae \\citep{mat02}, and the infrared spectrum from dusty ultracompact H~{\\sc ii} regions \\citep{ind06}.  In the planetary atmosphere community much effort has been devoted to developing three-dimensional radiation transfer techniques to study the penetration of Solar radiation through clouds in the Earth's atmosphere \\citep{cah05}. Studies of clouds in the Earth's atmosphere have shown a fractal structure \\cite{lov82}. Additionally, \\citet{kuc04} showed that a three-dimesional Monte Carlo radiation transfer simulation in fractal clouds results in a more accurate representation of radiation penetration in real clouds than a plane-parallel homogenous simulation. Though we lack the large amount of data of atmospheric scientists studying the Earth, this research is very relevant to extrasolar planetary atmospheric simulations.  The work we present in this paper is motivated by the non-detections of reflected light from exoplanets and the clear three-dimensional nature of clouds in the atmospheres of the Earth and other planets in the Solar System. We present three-dimensional\\footnote{We use 3D to mean that we can vary the density and characteristics of our model for each point in our x,y, and z dimensional grid, and introduce structure that varies in every direction.} radiation transfer simulations of the reflected light from inhomogeneous exoplanetary atmospheres. We have not computed the 3D temperature and pressure structure of an atmosphere, rather to investigate 3D radiation transfer effects on reflected light we take 1D atmospheres and convert them to 3D via a hierarchical clumping algorithm. This approach allows us to identify and characterize the different reflectivities of uniform and 3D atmospheres.  We then extend our analysis to study a more detailed atmospheric model incorporating multiple species of scatterers, cloud decks, and vertical density gradients within the atmosphere. Although we do not solve for the 3D temperature, pressure, and density structure, we attempt to build a more realistic simulation and will focus specifically on current atmospheric models for the extrasolar planet HD 209458b. We have simulated the atmosphere and found that in all cases the geometric albedos for the simulations of extrasolar planet HD 209458b are lower, and in many cases much lower, than that of Jupiter, suggesting radiation transfer effects in a 3D atmosphere might help explain why to date, reflected light has not been detected from an extrasolar planet.  In \\S\\ref{sec:code} we describe our numerical radiation transfer code, the algorithm we use for generating a hierarchically clumped 3D atmosphere, and the adopted scattering properties of particles within the atmosphere. Results of our radiation transfer simulations for a wide range of atmospheric optical depth, porosity, and particle scattering properties are presented in \\S\\ref{sec:simres} and \\S\\ref{sec:dis} summarizes our generalized results with respect to the scattered light levels. \\S\\ref{sec:209458} presents 3D radiation transfer models for the scattered light from the atmosphere of  HD 209458b and in \\S\\ref{sec:209sum} we summarize our findings. ",
        "conclusions": "\\label{sec:209sum}  The primary goal of this paper was to investigate the effects of 3D radiation transfer on the optical reflected light levels from extrasolar planetary atmospheres. We find that compared to smooth models, 3D atmospheres reflect less light because starlight can penetrate deep into the atmosphere, increasing its probability of being absorbed. As found in other 1D and 2D scattered light simulations, other atmospheric parameters serve to reduce the geometric albedos of extrasolar planets, most notably the scattering albedo of individual particles in the planetary atmosphere \\citep{dlu74}. Indeed, we suspect that it is the albedo of particles in the atmosphere that results in the very low geometric albedo observed from HD209458b.  We realize that our model is not a perfect representation of an extrasolar atmosphere, and that other effects not explored here might contribute to a lower geometric albedo. Specifically, strong absorption of alkali metals sodium and potassium could lower the albedo throughout the visible spectrum (\\cite{sud00} and \\cite{sea00}). However, we have tried to isolate some effects, specifically from the 3d structure of the atmosphere, that might contribute to the overall albedo of extrasolar planets.  HD 209458b is a dark planet, with a geometric albedo $A_g < 0.17$. One possibility is that clouds are absent and the low albedo is due to atomic or molecular absorption. To study the albedos in the presence of clouds, we investigated a fiducial 3D scattered light model which considered the densities, opacities, and albedos of hydrogen, iron, and enstatite, and their respective placements in the atmosphere. This fiducial model gives a geometric albedo $A_g > 0.4$, so we know that it does not accurately represent the atmosphere of HD 209458b. Upon varying parameters in our model we found significant drops in the geometric albedo. Parameters that strongly affected the geometric albedo include making the clouds more porous, lowering the condensation rate of enstatite, and assuming amorphous enstatite instead of glassy. Indeed, we were surprised by how easily the geometric albedo could be reduced with a relatively small change to the constituents of the atmosphere. However, we feel the most important parameter influencing the reflected light levels is the purity of enstatite and the assumption that it sits alone and above all the other condensates in the atmosphere. Models have suggested that HD 209458b has a turbulent atmosphere, and if clouds are indeed present we expect that mixing of enstatite with other, lower albedo condensates is what keeps the geometric albedo so low."
    },
    "0807/0807.3746_arXiv.txt": {
        "abstract": "We describe the discovery of a 0.68+0.52~$M_{\\sun}$ eclipsing binary (EB) with an 8.4-day orbital period, found through a systematic search of ten fields of the Trans-atlantic Exoplanet Survey (TrES). Such long-period low-mass EBs constitute critical test cases for resolving the long standing discrepancy between the theoretical and observational mass-radius relations at the bottom of the main sequence. It has been suggested that this discrepancy may be related to strong stellar magnetic fields, which are not properly accounted for in current theoretical models. All previously well-characterized low-mass main sequence EBs have periods of a few days or less, and their components are therefore expected to be rotating rapidly as a result of tidal synchronization, thus generating strong magnetic fields. In contrast, the binary system described here has a period that is over three times longer than previously characterized low-mass main sequence EBs, and its components rotate relatively slowly. It is therefore expected to have a weaker magnetic field and to better match the assumptions of theoretical stellar models. Our follow-up observations of this EB yield preliminary stellar properties that suggest it is indeed consistent with current models. If further observations confirm a low level of activity in this system, these determinations would provide support for the hypothesis that the mass-radius discrepancy is at least partly due to magnetic activity. ",
        "introduction": "\\label{secIntro}  Despite a great deal of work that has been done to understand the structure of low-mass ($<$ 0.8 $M_{\\sun}$) main sequence stars \\citep[e.g.,][]{Chabrier00}, models continue to underestimate their radii by as much as 15\\% \\citep{Lacy77rad, Torres02, Ribas06}.  This is a significant discrepancy, considering that for solar-type stars the agreement with the observations is typically within 1--2\\% \\citep{Andersen91, Andersen98}. In recent years an intriguing hypothesis has been put forward, suggesting that strong magnetic fields may have bloated these stars, either through chromospheric activity \\citep[e.g.,][]{Ribas06, Torres06, LopezMorales07, Chabrier07} or through magnetically induced convective disruption \\citep{Torres06}. Such strong magnetic fields are expected to be formed by the dynamo mechanism of rapidly rotating stars.\\footnote{Dynamo theory predicts that this mechanism operates only in partially convective stars. However, the strong magnetic activity observed in fully convective low-mass stars indicates that they also possess a mechanism for generating strong magnetic fields \\citep[see][and references therein]{Browning07}.} To test this hypothesis, one needs to measure both the masses and radii of low-mass stars, which thus far can be done most accurately with eclipsing binary (EB) systems. However, all well characterized low-mass main sequence EBs have orbital periods shorter than three days (see Table~\\ref{tablePreviousEBs}) and are therefore expected to have synchronization timescales shorter than $\\sim$100 Myr \\citep[][see Figure~\\ref{figTsync} and further description in \\S\\,\\ref{secPhysical}]{Zahn77, Zahn94}. As a result of these short periods and short synchronization timescales, the rotations of these binary components are expected to have accelerated to the point that they now match the rapid angular velocity of their orbits. With such rapid rotations, these binary components could have a wide range of dynamo-induced magnetic field strengths. To better constrain current stellar models, we set out to find systems with slowly rotating components. Such systems would presumably have comparably weak magnetic fields, thus being more consistent with the model assumptions. Furthermore, by comparing the mass-radius relations of binary components with well determined levels of magnetic activity, one could test various magnetic disruption models.  We note here that in addition to EB analysis, long-baseline optical interferometry has also been used recently to measure the radii of nearby low-mass stars \\citep{Lane01, Segransan03, Berger06}.  While these stars are single and are therefore expected to rotate slowly, their masses can only be estimated through empirical mass-luminosity relations or other indirect methods. Those determinations are thus less fundamental, in a sense, and arguably of lesser value for accurately constraining stellar models and testing the magnetic disruption hypothesis.  \\placefigure{figTsync} ",
        "conclusions": "\\label{secConclusions}  A comparison with solar-metallicity models by \\cite{Baraffe98} for a mixing length parameter of $\\alpha_{\\rm ML} = 1.0$ is presented in Figure~\\ref{figIsochrones}. Our mass and radius determinations for T-Lyr1-17236 (see Table~\\ref{tableSystemParam}) are shown along with those of the low-mass systems listed in Table~\\ref{tablePreviousEBs}. The location of the models in this diagram depends only slightly on age because these stars evolve very slowly. The age of T-Lyr1-17236 is difficult to establish independently. The space motions derived in \\S\\,\\ref{secPhysical} do not associate the system with any known moving group, and are quite typical of the thin disk. Thus, all we can say is that it is not likely to be very old. We display in Figure~\\ref{figIsochrones} two models for ages of $1$~Gyr and $10$~Gyr, which likely bracket the true age of T-Lyr1-17236. Within the errors, our measurements for the two components are consistent with the models, which would in principle support the magnetic disruption hypothesis. Unfortunately, however, the uncertainties in the radius measurements ($\\sim$7\\% and $\\sim$10\\%) are still large enough that our statement cannot be made more conclusive. Further follow-up observations, especially rapid-cadence and precise photometric measurements during multiple eclipses, should significantly reduce the uncertainties in the radii and thus provide far stronger constraints on the theoretical models of low-mass stars. Additionally, higher-quality spectroscopic observations than ours are needed to confirm that the level of chromospheric activity in the system is relatively low.  If after such observations, the masses and radii of the T-Lyr1-17236 components remain consistent with the stellar models, then the magnetic disruption hypothesis will be strengthened. However, if further observations find that the components of T-Lyr1-17236 are larger than predicted by current stellar models, as is the case with most other similar systems investigated in sufficient detail, then this will provide evidence that additional mechanisms need to be included in the models of the structure of low-mass main sequence stars \\citep[see, e.g.,][]{Chabrier07}.  \\placefigure{figIsochrones}  It is important to note here that T-Lyr1-17236 falls within the field of view of the upcoming NASA Kepler Mission \\citep{Borucki03}. The Kepler Mission will not return data for all stars within its field of view; rather, the targets will be selected by the Kepler team. We see at least two reasons why such monitoring of T-Lyr1-17236 would be of significant value. First, the data would greatly refine the estimates of the physical parameters of the component stars and may permit a search for their asteroseismological modes. Second, the data would enable a search for transits of exoplanets, which are expected to orbit in the same plane as that defined by the stellar orbits.  Finally, we note that our findings in this paper confirm the accuracy of the MECI algorithm (see Figure~\\ref{figMECI}), which can be further used to find additional long-period low-mass EBs, and indeed a variety of other interesting targets. We have shown in a recent paper \\citep{Devor08} how this can be done with comparable ease by systematically searching the ever-growing body of LC survey datasets. We hope that this new approach for locating rare EBs will motivate additional studies of these vast, largely untapped datasets, which likely harbor a wealth of information on the formation, structure, dynamics, and evolution of stars."
    },
    "0807/0807.1269_arXiv.txt": {
        "abstract": "{The metallicity dependence of the Cepheid $PL$-relation is of importance in establishing the extra-galactic distance scale.} {The aim of this paper is to investigate the metallicity dependence of the $PL$-relation in $V$ and $K$ based on a sample of 68 Galactic Cepheids with individual Baade-Wesselink distances (some of the stars also have an HST-based parallax) and individually determined metallicities from high-resolution spectroscopy.} {Literature values of the $V$-band, $K$-band and radial velocity data have been collected for a sample of 68 classical cepheids that have their metallicity determined in the literature from high-resolution spectroscopy. Based on a $(V-K)$ surface-brightness relation and a projection factor derived in a previous paper, distances have been derived from a Baade-Wesselink analysis. $PL$- and $PLZ$-relations in $V$ and $K$ are derived.} {The effect of the adopted dependence of the projection factor on period is investigated. The change from a constant $p$-factor to one recently suggested in the literature with a mild dependence on $\\log P$ results in a less steep slope by 0.1 unit, which is about the 1-sigma error bar in the slope itself. The observed slope in the $PL$-relation in $V$ in the LMC agrees with both hypotheses. In $K$ the difference between the Galactic and LMC slope is larger and would favour a mild period dependence of the $p$-factor. The dependence on metallicity in $V$ and $K$ is found to be marginal, and independent of the choice of $p$-factor on period. This result is severely limited by the small range in metallicity covered by the Galactic Cepheids. } {} ",
        "introduction": "Obtaining accurate distances to stars is a non-trivial matter. Cepheids are considered an important standard candle as they are bright and are thus the link between the distance scale in the nearby universe and that further out via those galaxies that contain both Cepheids and SNIa.  Distances to local cepheids may be obtained via main-sequence fitting for those Cepheids in clusters (e.g. Feast 1999) or via determination of the parallax.  Until recently only Polaris had an accurate parallax determination via Hipparcos.  Benedict et al. (2007) recently published absolute trigonometric parallaxes for nine Galactic Cepheids using the {\\it Fine Guidance Sensor} on board the {\\it Hubble Space Telescope}, and revised, more accurate, Hipparcos parallaxes have also become available (van Leeuwen et al. 2007, van Leeuwen 2007).  In addition, distances to Cepheids can be obtained from the Baade-Wesselink method. This method relies on the availability of surface-brightness (SB) relations to link variations in colour to variations in angular diameters and understanding of the projection ($p$-) factor that links radial velocity  to pulsational velocity variations.  The SB relations can be obtained using Cepheids that have accurate interferometrically determined angular diameters over the pulsation phase and accurate multi-wavelength light curves.  This has allowed Kervella et al. (2004) and Groenewegen (2007; hereafter G07) to derive such relations, in particular using the $(V-K)$ colour, which gives the highest precision in the derived angular diameter.   The $p$-factor can also be obtained from Cepheids with interferometrically determined angular diameters over the pulsation phase when in addition accurate RV curves are available (see e.g. M\\'erand et al. 2005). G07 assembled all currently available data to find that a constant $p$-factor represents the best fit to the available data. A strong dependence on period, like $p \\sim - 0.15 \\log P$ as proposed by Gieren et al. (2005), could be ruled out. A moderate dependence of $p \\sim - 0.03 \\log P$ as used in the Gieren et al. papers (1993, 1997, 1998, Storm et al. 2004, Barnes et al. 2003), or $p \\sim - 0.064 \\log P$ as advocated by Nardetto et al. (2007) are also consistent with the available data.   Another aspect of the Cepheid $PL$-relation that is still a matter of debate is the metallicity dependence of the Cepheid $PL$-relation, and its dependence on wavelength. Observations seem to consistently indicate that metal-rich Cepheids are brighter, and various estimates have been given in the literature, $-0.88 \\pm 0.16$ mag/dex ($BRI$ bands, Gould 1994), $-0.44^{+0.1}_{-0.2}$ mag/dex ($VR$ bands, Sasselov et al. 1997), $-0.24 \\pm 0.16$ mag/dex ($VI$ bands, Kochanek 1997), $-0.14 \\pm 0.14$ mag/dex ($VI$ bands, Kennicutt et al. 1998), $-0.21 \\pm 0.19$ in $V$, $-0.29 \\pm 0.19$ in $W$, $-0.23 \\pm 0.19$ in $I$, $-0.21 \\pm 0.19$ mag/dex in $K$ (Storm et al. 2004), and $-0.29 \\pm 0.10$ mag/dex ($BVI$ bands, Macri et al. 2006).   The potential caveat is that no individual abundance determinations of individual Cepheids are being used in these studies but rather abundances of nearby H{\\sc ii} regions, or even a mean abundance of the entire galaxy.  Groenewegen et al. (2004) tried to attack the problem by using Cepheids with individually determined metallicities, 37 in our Galaxy, 10 in the LMC and 6 in the SMC. Considering the Galactic sample, a metallicity effect was found in the zero point in the $VIWK$ $PL$-relation of $-0.6 \\pm 0.4$ or $-0.8 \\pm 0.3$ mag/dex, depending on the in- or exclusion of one object. The distances to the galactic Cepheids mostly came from BW-distances available at that time from Fouqu\\'e et al. (2003).  In the last few years significant progress has been made in several areas. On the one hand more Cepheids have interferometrically determined angular diameters. Combined with improved theoretical studies (Nardetto et al. 2007, and references therein) this has led to an improved understanding of the $p$-factor and SB relations, as was studied in G07. On the other hand, more basic photometric light curves and radial velocity data have become available and new studies have presented metallicity determinations for significant samples (see references below).   In this paper, Baade-Wesselink distances to Cepheids with metallicity determinations are revisited, using the $p$-factor and the SB $(V-K)$ colour relation derived in G07, updated as described in Sect.~3. Section~2 describes the selection of the photometric and radial velocity data.  Section~3 outlines how the data are modelled. Section~4 describes how the binary Cepheids are treated, and new and updated orbital elements are presented. Section~5 describes the results. Section~6 presents conclusions. ",
        "conclusions": "The period-luminosity in the $V$ and $K$-band and the dependence on metallicity are investigated for a sample of 68 Galactic Cepheids with an individual metallicity determination from high-resolution spectroscopy. Distances are derived using the Baade-Wesselink technique using the most recent surface-brightness relation and estimates for the projection factor.  The difference in slope between a constant $p$-factor (consistent with available interferometric data) and a mild dependence on period (following theoretical work by Nardetto et al. 2007) results in a difference in slope in the $PL$-relation which is of the same magnitude as the error in slope.  The slope of the LMC $PL$-relation in $V$ and $K$ is derived taking into account the difference in period distribution between the Galactic sample and the LMC sample. In $V$ the slope for the LMC sample agrees with the galactic one for both assumptions regarding the dependence of $p$ on period. In $K$ the comparison of the slopes favours the mild period dependence of the $p$-factor on period.  The metallicity dependence of the $PL$-relation is investigated and no significant dependence is found. A firm result is not possible as the range in metallicity spanned by the current sample of galactic Cepheids is 0.3-0.4 dex, while previous work suggested a small dependence on metallicity only (typically $-0.2$ mag/dex, see the introduction).  It would be extremely valuable to analyse a sample of LMC and SMC Cepheids in a similar way. Unfortunately, the sample of Cepheids with adequate optical, NIR and RV light curves (see e.g.  Storm et al. 2004, 2005) for which a Baade-Wesselink distance has been derived (Gieren et al. 2005) and the sample of Magellanic Cloud Cepheids for which a metallicity has been derived from high-resolution spectroscopy (Mottini et al. 2008) have no overlap. High-resolution spectroscopy of MC Cepheids with Baade-Wesselink distances would therefore be extremely valuable."
    },
    "0807/0807.1119_arXiv.txt": {
        "abstract": "{ Accretion and outflow processes are of fundamental importance for our understanding of the formation of stars and planetary systems. To trace these processes, diagnostic spectral lines such as the {\\BrG} $2.166~\\mu$m line are widely used, although due to a lack of spatial resolution, the origin of the line emission is still unclear. } { Employing the AU-scale spatial resolution which can be achieved with infrared long-baseline interferometry, we aim to distinguish between theoretical models which associate the {\\BrG} line emission with mass infall (magnetospheric accretion, gaseous inner disks) or mass outflow processes (stellar winds, X-winds, or disk winds). } { Using the VLTI/AMBER instrument, we spatially and spectrally ($\\lambda/\\Delta\\lambda=1500$) resolved the inner ($\\lesssim 5$~AU) environment of five Herbig~Ae/Be stars (HD\\,163296, HD\\,104237, HD\\,98922, MWC\\,297, V921\\,Sco) in the {\\BrG} emission line as well as in the adjacent continuum. From the measured wavelength-dependent visibilities, we derive the characteristic size of the continuum and {\\BrG} line-emitting region. Additional information is provided by the closure phase, which we could measure both in the continuum wavelength regime (for four objects) as well as in the spectrally resolved {\\BrG} emission line (for one object). The spectro-interferometric data is supplemented by archival and new VLT/ISAAC spectroscopy. } { For all objects (except MWC\\,297), we measure an increase of visibility within the {\\BrG} emission line, indicating that the {\\BrG}-emitting region in these objects is more compact than the dust sublimation radius.  For HD\\,98922, our quantitative analysis reveals that the line-emitting region is compact enough to be consistent with the magnetospheric accretion scenario. For HD\\,163296, HD\\,104237, MWC\\,297, and V921\\,Sco we identify an extended stellar wind or a disk wind as the most likely line-emitting mechanism. Since the stars in our sample cover a wide range of stellar parameters, we also search for general trends and find that the size of the {\\BrG}-emitting region does not seem to depend on the basic stellar parameters (such as the stellar luminosity), but correlates with spectroscopic properties, in particular with the {\\Ha} line profile shape. } { By performing the first high-resolution spectro-interferometric survey on Herbig~Ae/Be stars, we find evidence for at least two distinct {\\BrG} line-formation mechanisms. Most significant, stars with a P-Cygni {\\Ha}~line profile and a high mass-accretion rate seem to show particularly compact {\\BrG}-emitting regions ($R_{\\mathrm{Br}\\gamma}/R_{\\mathrm{cont}} <0.2$), while stars with a double-peaked or single-peaked {\\Ha}-line profile show a significantly more extended {\\BrG}-emitting region ($0.6 \\lesssim R_{\\mathrm{Br}\\gamma}/R_{\\mathrm{cont}} \\lesssim 1.4$), possibly tracing a stellar wind or a disk wind. } ",
        "introduction": "\\label{sec:intro}  Accretion disks around young stellar objects (YSOs) are at the focus of astronomical research, not only because they play an essential role in the star-formation process, but also because they provide the stage where planet formation takes place.  Historically, these disks were discovered due to their characteristic infrared excess emission, which is believed to trace dust grains, providing (for the expected temperatures and densities) the dominant source of opacity. In recent years, there has been substantial progress in constraining the detailed three-dimensional geometry of the dust disk using, for example, combined modeling of the spectral energy distribution (SED) and spatially resolved infrared interferometric observations. While the thermal emission from the dust disk is likely to be the dominant contributor to the infrared excess emission observed towards YSOs, it is believed that the dust content makes up only a small fraction of the total mass of the system. 99\\% of the mass is likely contributed by gas, in particular hydrogen, and can mainly be traced by the spectral lines, which for pre-main-sequence stars are often found in emission.  While some information about the kinematics of the gas can already be extracted from the line profile, the spatial origin of the gas emission and the physical processes they trace are still strongly debated.  In this context, the Brackett-$\\gamma$ ({\\BrG}) 2.1661~$\\mu$m line is of special importance. It was found that the luminosity $L($\\BrG$)$ of this line (as determined from the circumstellar component of the line equivalent width) seems to correlate with the mass accretion luminosity $L_{\\mathrm{acc}}$, as determined from UV veiling. This empirical $L($\\BrG$)$--$L_{\\mathrm{acc}}$ relationship has been established for pre-main-sequence stars of various masses, including brown dwarfs \\citep{nat04}, T~Tauri~stars \\citep{muz98a}, as well as intermediate-mass Herbig~Ae/Be stars \\citep{cal04,van05}. Based on this empirical correlation, observers started to use the {\\BrG} luminosity as an estimator for the mass accretion rate \\citep[e.g.][]{gar06}, which follows from $L_{\\mathrm{acc}}$ using $\\dot{M}_{\\mathrm{acc}} = L_{\\mathrm{acc}} R_{\\star} / G M_{\\star}$.  Therefore, it is now essential to identify the process(es) involved in the formation of the {\\BrG} emission line in YSOs. The main scenarios which have been proposed up to now include mass infall, as well as mass outflow mechanisms:  \\begin{figure*}[htbp] \\centering \\includegraphics[angle=0,width=18cm]{9946fig1.eps}\\\\ \\caption{Illustration of the regions which have been proposed as the origin of the permitted hydrogen recombination line emission observed towards HAeBe stars (this sketch is not to scale; read Sect.~\\ref{sec:intro} for details about the individual mechanisms).} \\label{fig:windorigin} \\end{figure*}  \\noindent {\\bf (a)~Magnetospheric accretion:} The line emission might emerge from matter which is accreted onto the star through magnetospheric accretion columns \\citep{van05}.  This infall is supposed to happen very close to the star, inside the co-rotation radius, where the Keplerian angular velocity matches the stellar angular velocity.  \\noindent {\\bf (b)~Gaseous inner disk:} Inside of the dust destruction radius, the gas continues to accrete towards the star, forming a gaseous inner accretion disk.  The recombination line emission from ionized hydrogen in this disk might contribute to the line emission observed towards HAeBe stars. \\citet{muz04} estimated that especially for high mass-accretion rates $\\dot{M}_{\\mathrm{acc}} > 10^{-7}~M_{\\sun} $yr$^{-1}$, the flux contribution of the inner gaseous disk might be of importance.  \\noindent {\\bf (c)~Stellar wind:} The P~Cygni line profile observed in the hydrogen lines of several HAeBe~stars might also indicate mass loss through stellar winds \\citep[e.g.][]{mes68,cat87,str98}. While magnetically accelerated stellar winds seem plausible for massive stars rotating close to their break-up velocity, this scenario does not seem to work for lower-mass YSOs \\citep{fer06}.  \\noindent {\\bf (d)~Stellar-field driven wind (X-wind):} In the X-wind model \\citep{shu94}, the outflows from YSOs are launched from a narrow interaction region where the stellar magnetosphere truncates the accretion disk.  Keplerian rotation leads to a winding up of the field lines and to the formation of magnetic surfaces. Charged particles from the stellar wind will then get trapped in these surfaces, accelerated, and collimated into a beam.  \\noindent {\\bf (e)~Disk-field driven wind (disk wind):} Another mechanism which has been proposed for the launching and collimation of outflows and jets observed towards YSOs are magnetocentrifugally driven disk winds \\citep{bla82,pud83,fer97}. Material from the disk surface of the rotating disk is centrifugally accelerated along the open magnetic field lines of the accretion disk. In contrast to the X-wind model, where the wind is launched from a narrow disk annulus around the inner truncation radius of the accretion disk, disk winds originate from a wide range of radii, likely extending from the co-rotation radius out to several AUs. \\newline  Since most of these processes are believed to take place on (sub-)AU scales and were therefore not accessible with direct imaging techniques, most earlier studies tried to constrain the spatial distribution and kinematics of the line-emitting gas from the shape of line profiles. However, since these line profile fitting techniques are known to be highly ambiguous \\citep[e.g.][]{cat99}, spatially resolved observations are urgently required to constrain the physical mechanism and spatial origin of important tracer lines like {\\BrG}. Since the different emitting mechanisms noted above can be associated to distinct spatial regions (see Fig.~\\ref{fig:windorigin}), spatially resolved observations should allow us to identify and physically characterize the true underlying emitting mechanism.  The VLTI/AMBER instrument combines, for the first time, the milli-arcsecond spatial resolution achievable with IR long-baseline interferometry with a good spectroscopic resolution ($R=\\lambda/\\Delta\\lambda=1500$ or 12\\,000) in the near-infrared $K$-band. In this study, we use the unique capabilities of this instrument in order to measure the geometry and position of the {\\BrG} line-emitting region relative to the continuum-emitting region, providing direct information about the processes involved.  A particularly well-studied class of YSOs are the Herbig Ae/Be stars (HAeBes). These are intermediate-mass, pre-main sequence stars.  Over the last few years, the continuum emission from a rather large number of HAeBes could already be studied with broad-band infrared interferometry. These studies revealed a correlation between the size of the NIR continuum emitting region and the stellar luminosity \\citep{mon02}, suggesting that for most Herbig~Ae stars, the NIR emission likely traces dust at the dust sublimation radius.  Only very little is known about the spatial origin of the hydrogen recombination line emission. First AMBER observations of the Herbig~Be star MWC\\,297 have shown that the {\\BrG}-emitting region around this star is more extended than the continuum region \\citep{mal07}. Surprisingly, observations on the less luminous Herbig~Ae star HD\\,104237 did not show any change in visibility along the {\\BrG} line \\citep{tat07a}. \\citet{eis07b} found a visibility increase within the {\\BrG} line for the Herbig~Ae star MWC\\,480. This might suggest that the {\\BrG} line traces fundamentally different mechanisms for Herbig~Ae and Be stars. In this study, we present observations on three HAeBes, which are resolved for the first time with spectro-interferometry (HD\\,163296, HD\\,98922, V921\\,Sco). For HD\\,104237, we present new observations which were obtained at longer baseline lengths, adding substantial new information to the initial AMBER results presented by \\citet{tat07a}.  Finally, we re-reduced the archival data on MWC\\,297 and HD\\,104237 in order to provide a uniformly calibrated sample of five stars for our interpretation.  The stellar parameters which we adopted for the stars in our sample are shown in Tab.~\\ref{tab:targetstars}.  This paper is structured as follows:  In Sect.~\\ref{sec:observations} we discuss the spectro-interferometric and spectroscopic data and the applied data reduction procedures, followed by an outline of the obtained results (Sect.~\\ref{sec:results}). Then, we present the models which were fitted to the data in order to constrain the geometry of the continuum- (Sect.~\\ref{sec:modelingcont}) and {\\BrG}-emitting region (Sect.~\\ref{sec:modelingbrg}). After interpreting the individual objects in detail (Sect.~\\ref{sec:interpretation}), we discuss general trends which we find in our small sample of HAeBe stars (Sect.~\\ref{sec:discussion}). Finally, we conclude with a summary of our results and discuss the potential of future observations in Sect.~\\ref{sec:conclusions}. ",
        "conclusions": "\\label{sec:conclusions}  We summarize the results from our investigation using VLTI/AMBER data of five HAeBe stars as follows: \\begin{enumerate} \\item The $K$-band continuum-emitting region was spatially resolved for all objects. For the Herbig~Ae and late Herbig~Be stars the characteristic size of this region scales roughly with the square-root of the stellar luminosity, supporting the scenario that for these sources, the continuum is dominated by thermal emission from the inner rim of the dust disk. For the early-type Herbig~Be stars MWC\\,297 and V921\\,Sco, on the other hand, the continuum-emitting region is significantly more compact than predicted by the $R_{\\mathrm{subl}} \\propto L_{\\star}^{1/2}$ relation, maybe suggesting that for these objects the NIR continuum is dominated by hot gas emission. Fits of inclined ring geometries to the visibilities measured on HD\\,163296 indicate object elongation, consistent with a notably inclined disk geometry, while for HD\\,98922 and V921\\,Sco we find no indications for a large disk inclination. \\item Closure phases were measured for four objects in the $K$-band continuum regime.  For V921\\,Sco we measured the CP also within the spectrally resolved {\\BrG} line, suggesting a nearly centro-symmetric line-emitting region (in agreement with the low inclination angle). \\item In the spectrally resolved {\\BrG}~lines of HD\\,163296, HD\\,104237, HD\\,98922, and V921\\,Sco, we measure an increase of visibility, indicating that the line-emitting region is more compact than the continuum-emitting region. MWC\\,297 shows a decrease of visibility within the {\\BrG}-line, although this is likely related to the compactness of the continuum-emitting region and not to fundamental differences in the line formation mechanism. \\item Judging on the currently available limited statistics, we do not find a trend which relates the size of the {\\BrG} line-emitting region with the stellar luminosity, as was proposed by \\citet{eis07b}. Instead, the continuum-normalized size of the line-emitting region $R_{\\mathrm{Br}\\gamma}/R_{\\mathrm{cont}}$ might correlate with the stellar effective temperature, although more observations will be required to confirm this trend. Maybe the most remarkable correlation we found was that stars with a P-Cygni {\\Ha}~line profile and a high mass-accretion rate seem to show particularly compact {\\BrG}-emitting regions ($R_{\\mathrm{Br}\\gamma}/R_{\\mathrm{cont}} <0.2$), while stars with a double-peaked or single-peaked {\\Ha}-line profile show a significantly more extended {\\BrG}-emitting region ($0.6 \\lesssim R_{\\mathrm{Br}\\gamma}/R_{\\mathrm{cont}} \\lesssim 1.4$). \\item  Only for HD\\,98922, the emitting-region is compact enough to support the hypothesis that most of the {\\BrG}-emission emerges from magnetospheric accretion columns.  In the other cases, the {\\BrG}-emitting region is only slightly more compact than the dust sublimation radius, supporting the idea that the emission mainly emerges from an extended stellar wind or a disk-wind.  This makes it unlikely that {\\BrG} primarily traces magnetospheric accretion in HAeBe stars, suggesting that even for Herbig Ae stars, this line is only an indirect tracer of the mass accretion rate. \\end{enumerate}"
    },
    "0807/0807.1928_arXiv.txt": {
        "abstract": "The Canadian \\textit{MOST} satellite is a unique platform for observations of bright transiting exoplanetary systems.  Providing nearly continuous photometric observations for up to 8 weeks, \\textit{MOST} can produce important observational data to help us learn about the properties of exosolar planets.  We review our current observations of HD 209458, HD 189733 with implications towards the albedo and our progress towards detecting reflected light from an exoplanet. ",
        "introduction": "We train our detectors to monitor the periodic cycle of a planet transiting a star, being eclipsed by the star and the change of the planet's illuminated phase.  Our primary goal: to determine what fraction of photons come from the star and what fraction come from the planet.  The easiest feature to detect and measure is when the planet transits the star from the observers perspective where the depth of the transit is determined by the ratio of the projected surface area of a planet and star.\\footnote{Ignoring emission and transmission of photons from the planet and assuming uniformity on the surface of the star.}  When a planet is eclipsed by the star only photons emitted by the star are detected.  The ratio of the observed flux before first contact to the observed flux during total eclipse determines the ratio of the planet's reflected flux to the star's incident flux at the substellar point and determine the geometric albedo.  With the \\textit{MOST} ({\\it Microvariablity and Oscillations of STars}) satellite we have observed two exoplanetary systems with a goal of measuring the reflected light signature from the planet.  For our first candidate, HD 209458 (G0V, V=7.7), observations spanning 14, 45 and 29 days have been obtained in 2004, 2005 and 2007 respectively.  From the 2004 and 2005 observation runs we set a $1\\, \\sigma$ upper limit on the planet-to-star flux ratio of 16 ppm.  Details of this study may be found in \\cite[Rowe \\etal\\ (2006)]{row06} and \\cite[Rowe \\etal\\ (2007)]{row07}.  We have also observed the chromospherically active star HD 189733 (K2V, V=7.7) in an attempt to measure the albedo of the transiting Jupiter-mass planet.  Our observations span 22 and 31 days for our 2006 and 2007 campaigns.  The presence of large spots and large scatter seen in optical photometry has hindered our progress.  \\subsection{The Albedo}  Understanding what an albedo tells us about the nature of a planet requires a quick review of the definition of the Bond albedo, which in turn defines the geometric albedo and the phase function.  Bond's definition of an albedo, from 1861, is: ``Let a sphere S be exposed to parallel light.  Then its albedo ($A_B$) is the ratio of the whole amount of light reflected from S to the whole amount incident on it\" \\cite[(Russell 1916)]{rus16}.  The incident flux ($L_{in}$) is the stellar flux density at the distance of the planet from the star ($a$) times the cross section of the planet ($\\pi R_p^2$) which, by definition of the Bond albedo states, that the emergent flux is $L_{out}=A_B\\ L_{in}$.  For an anisotropic scatter the planetary flux density observed at a distance $D$ is \\begin{equation}\\label{eq:Fp} F_p=C\\ \\phi (\\alpha)\\ \\frac{L_{out}}{4 \\pi D^2}, \\end{equation} where $C$ is a constant and $\\phi$ is the phase function describing the brightness of the planet at phase angle $\\alpha$ normalized such that $\\phi(\\alpha = 0) = 1$.  Integrating over the surface of a sphere with radius $D$, \\begin{equation} L_{out} = \\int_S C\\ \\phi (\\alpha)\\ \\frac{L_{out}}{4 \\pi D^2}\\ dS, \\end{equation} we can solve for the normalization constant $C$ as the flux density integrated over the surface of the planet much equal total flux received.  We find, \\begin{equation}\\label{eq:C} C=\\frac{2}{\\int_0^{\\pi} \\phi (\\alpha)\\ {\\rm sin} \\alpha\\ d\\alpha} = \\frac{4}{q}, \\end{equation} where we have defined the phase integral, \\begin{equation} q= 2 \\int_0^{\\pi} \\phi(\\alpha)\\ {\\rm sin} \\alpha\\ d\\alpha, \\end{equation} which quantifies the scattering properties of the reflective surface over all phase angles and is proportional to the total amount of light reflected by the planet. The geometrical albedo is defined at phase $\\alpha=0$.  Using the definition of the Bond albedo with Equation \\ref{eq:Fp} and replacing $L_{in}$ with the flux density of the star ($F_*$) at the planet-to-star distance ($a$) we obtain, \\begin{equation}\\label{eq:fp} F_p = C\\ A_B \\frac{F_* R_p^2}{4 a^2}, \\end{equation} which we rearrange to solve for the Bond albedo and substitute Equation \\ref{eq:C} for $C$ and obtain, \\begin{equation} A_B = \\frac{4}{C} \\frac{F_p}{F_*} \\frac{a^2}{R_p^2} = \\frac{F_p}{F_*} \\frac{a^2}{R_p^2} q = A_g\\ q, \\end{equation} which gives us the geometric albedo ($A_g$).  The geometric albedo can be physically interpreted as the ratio of the planetary flux density at full phase ($\\alpha=0$) to a Lambertian surface of the same cross section.  The ratio of the planetary to stellar flux can be written as an approximation given by, \\begin{equation}\\label{ag_approx} \\frac{F_p}{F_*} \\approx A_g \\left( \\frac{R_p}{R_J} \\right)^2 \\left( \\frac{a}{0.05\\ {\\rm AU}} \\right)^{-2}  100\\ \\rm{ppm}, \\end{equation} where $R_J$ is the radius of Jupiter.  Equation \\ref{ag_approx} demonstrates that ideal exoplanet candidates for measuring low amplitude reflected light signatures have large radii and orbit at small distances from its host star. For close-in planets with radii similar to Jupiter and have a geometric albedo of 10\\% the photometric signature of planet being eclipsed by the star has an amplitude of approximately 10 ppm. ",
        "conclusions": "In the absence of clouds, all hot Jupiter models do predict extremely low visible-wavelength geometric albedos, due to strong, broad absorption lines of neutral atomic Na and K.  Some model atmospheres include self-consistent cloud formation in a 1D, complete cloud cover scenario (\\cite[Ackerman \\& Marley 2001]{ack01}, \\cite[Cooper \\etal\\ 2003]{coo03}). Nevertheless, the physics of cloud formation is a process that is not well understood or constrained both for particle size and densities which are governed by the competing effects of condensation and coagulation versus sedimentation (Marley et al. 1999). The low reflected fraction of incident radiation, as measured for HD 209458b, readily rules out reflective clouds.  New Spitzer observations at 3.6, 4.5, 5.8 and $8.0\\, \\mu$m \\cite[(Knutson \\etal\\ 2007)]{knu07} indicate a temperature inversion which requires an extra, unknown absorber at low pressures \\cite[(Burrows \\etal\\ 2007)]{bur07}.  Our low albedo limit means that, if the absorber is a cloud, then it must not be highly reflective.  With the 2007 photometric reductions for the HD 209458 and HD 189733 datasets nearing completion, we plan to refine our albedo estimates for both planets.  With the addition of onboard image stacking we gain a factor better than 6 for number of detected photons and the new datasets have a potential to push the albedo limits to a few percent at high significance. Both HD 189733 and HD 209458 are scheduled to be reobserved in 2008.  \\textit{MOST} will also observe the highly eccentric (e=0.93) planetary system HD 80606 (G5V, V=8.9).  The large eccentricity means that the planet distance ranges from 0.03 to 0.84 AU over a 111 day period.  The large eccentricity also means that the dayside temperature of the planet is expected to change from $\\sim$ 750 K to 1250 K on the timescale of a week.  This offers the potential to study the reflective properties of an exoplanet as a function of temperature.  Rayleigh scattering, Mie scattering and molecular absorption are the dominant mechanisms that determine the reflected and emitted spectra of an extrasolar giant planet \\cite[(Marley \\etal\\ 2006)]{mar06}.  In the bluer portion of the optical spectra (wavelengths shorter than 600nm) Rayleigh scattering in a clear atmosphere will reflect a large fraction of the stellar flux outwards. In the red portion (wavelengths greater than 600nm) photons will be absorbed deep in the atmosphere that will make the reflected spectrum relatively dark.  Towards the infrared, the spectrum will begin to be dominated by thermal radiation from the planet. Figure \\ref{fig1} shows the bandpass for the \\textit{Kepler} and \\textit{CoRoT} (whitelight) spacecraft.  Since the bandpasses of these missions, including \\textit{MOST}, extends towards the red (and beyond) portion of the spectrum, the reflected light signature for planets that orbit close to the host stars can be contaminated by thermal radiation.  As a final note, when attempting to measure low albedos around giant close-in exoplanets precaution should be taken ensure that reflected light signatures are not confused by similar lightcurves produced by thermal radiation."
    },
    "0807/0807.4216_arXiv.txt": {
        "abstract": "\\begin{center} \\bf{Abstract} \\end{center}  We discuss corrections to the Yukawa  matrices of the Standard Model (SM) fermions in intersecting D-brane models due to four-point interactions. Recently, an intersecting D-brane model has been found where it is possible to obtain correct masses and mixings for all quarks as well as the tau lepton.  However, the masses for the first two charged leptons come close to the right values but are not quite correct. Since the electron and muon are quite light, it is likely that there are additional corrections to their masses which cannot be neglected. With this in mind, we consider contributions to the SM fermion mass matrices from four-point interactions. In an explicit model, we show that it is indeed possible to obtain the SM fermion masses and mixings which are a better match to those resulting from experimental data extrapolated at the unification scale when these corrections are included. These corrections may have broader application to other models. ",
        "introduction": "In recent years, intersecting D-brane models, where the chiral fermions arise at the intersections between D6-branes (Type IIA) in the internal space~\\cite{bdl} with the T-dual Type IIB description in terms of magnetized D-branes~\\cite{bachas} have provided an exciting approach towards constructing semi-realistic string vacua (for reviews, see~\\cite{Blumenhagen:2005mu, Blumenhagen:2006ci}). Indeed, such models provide promising setups which may accommodate semi-realistic features of low-energy physics. Given this, it is an interesting question to see how far one can get from a particular string compactification to reproducing the finer details of the Standard Model (SM) as a low-energy effective field theory.  The Standard Model has an intricate structure, with three-generations of chiral fermions which transform as bifundamental representations of $SU(3)_C \\times SU(2)_L \\times U(1)_Y$. In addition to the fact that the SM fermions are replicated into three distinct generations, the different generations exhibit a distinct pattern of mass hierarchies and mixings. Interestingly, intersecting D-brane models may naturally generate the SM fermion mass hierarchies and mixings, as well as an explanation for the replication of chirality. In short, D6-branes (in Type IIA) fill four-dimensional Minkowski space-time and wrap 3-cycles in the compact manifold, with a stack of $N$ D6-branes having a gauge group $U(N)$ (or $U(N/2)$ in the case of $\\mathbf{T^6/(\\Z_2 \\times \\Z_2)}$) in its world volume. The 3-cycles wrapped by the D-branes will in general intersect multiple times in the internal space, resulting in chiral fermions in the bifundamental representation localized at the intersections between different stacks. The multiplicity of such fermions is then given by the number of times the 3-cycles intersect.  The Yukawa couplings in intersecting D6-brane models arise from open string world-sheet instantons that connect three D6-brane intersections \\cite{Aldazabal:2000cn}. For a given triplet of intersections, the minimal world-sheet action which contributes to the trilinear Yukawa couplings is weighted by a factor $\\exp(-A_{abc})$, where $A_{abc}$ is the world-sheet area of the triangle bounded by the branes $a, b,$ and $c$. Since there are several possible triangles with different areas, mass hierarchies may inherently arise. The Yukawa couplings depend on both the D-brane positions in the internal space as well as on the geometry of the underlying compact manifold. Effectively, these quantities are parameterized by the vacuum expectation values (VEVs) of open and closed-string moduli.   Despite substantial progress in constructing semi-realistic vacua with intersecting D-branes, there are many phenomenological challenges remaining, besides the usual moduli stabilization problem.   Typically, some or all of the Yukawa couplings are typically forbidden by selection rules which arise from global $U(1)$s which become massive via a generalized Green-Schwarz mechanism.  It has been found that only Pati-Salam models can have all the SM fermion Yukawa couplings present at the stringy tree level, although some couplings which are perturbatively forbidden may be generated in principle via D-brane instanton effects~\\cite{Ibanez:2006da, Blumenhagen:2006xt, Florea:2006si, Abel:2006yk, Billo:2007py}. Also, there has generally been a rank one problem in the SM fermion Yukawa matrices, preventing the generation of masses and mixings for the first two families of quarks and leptons. For the case of toroidal orientifold compactifications, this can be traced to the fact that not all of the SM fermions are localized at intersections on the same torus~\\cite{Cremades:2003qj, Chamoun:2003pf, Kitazawa:2004nf, Dutta:2005bb}. However, one example of an intersecting D6-brane model in Type IIA on the $\\mathbf{T^6/(\\Z_2\\times \\Z_2)}$ orientifold has recently been discovered in which these problems may be solved~\\cite{Cvetic:2004ui,Chen:2006gd}.  Thus, this particular model may be a step forward to obtaining realistic phenomenology from string theory. Indeed, as we have recently shown~\\cite{Chen:2007px}, it is possible within the moduli space of this model to obtain the correct SM quark masses and mixings, the tau lepton mass, and to generate naturally small neutrino masses via the seesaw mechanism. In addition to these features, the model exhibits automatic gauge coupling unification, and it is possible to generate realistic low-energy supersymmetric particle spectra, a subset of which may produce the observed dark matter density.   In spite of the successes of this model, the electron and muon masses come close to the right values, but they are still not quite correct. Since the electron and muon masses are very light, it is likely that there are additional corrections that must be considered, as first suggested a long time ago in \\cite{Ellis:Nano}.  This idea was applied later in string theory, for the case of the free-fermionic-formulation, where the general rules for calculating higher order corrections, from multipoint-functions, to the Yukawa couplings were first given in \\cite{Kalara:1990fb}. With this in mind, in this paper we consider contributions to the Yukawa couplings from four-point functions. Four-point string interactions in intersecting D6-brane scenarios can be regarded as the scattering amplitudes of four string matter fields \\cite{Lust:2004cx}, and are consistent to the K\\\"ahler metrics of twisted matter \\cite{Bertolini:2005qh}. Both the quantum and classical contributions have been studied in two special cases where one has only an independent angle and the other has two independent angles in a closed quadrilateral \\cite{Cvetic:Abel}. This was then promoted to the generalized calculation of $N$-point amplitudes and four-point functions without angle constraints \\cite{Abel:2003yx}. This analysis is based on the discussion of the twisted closed string interactions on orbifolds \\cite{Hamidi:1986vh, Dixon:1986qv, Burwick:1990tu}, where the analogy between open strings at brane intersections and closed strings on orbifolds is known. The complete amplitudes for Yukawa couplings, as a three-point limit of the four-point amplitude from conformal field theory, are also discussed in \\cite{Lust:2004cx, Abel:2003yx, Cvetic:Abel}.   It is natural to ask whether the four-point function corrections to the Yukawa couplings can generate the correct electron and muon masses in the model discussed in \\cite{Chen:2007px}; indeed the correct electron and muon masses can be obtained by including four-point function corrections in principle. However, the D6-branes in this model can not form a closed quadrilateral for the required four-point couplings on the relevant two torus. Two stacks of D-branes overlap so the ``angle'' at their ``intersection'' turns out as a straight line. Instead of using the conventional calculational procedure in this model, we can still apply the idea of four-point function corrections by taking the vector-like fields at the ``intersection'' on the overlapped D6-branes in a asymptotic limit of the ``angle'' to $\\pi$, a physical mechanism which has not yet been clarified. Therefore, we shall consider another model where the standard four-point function corrections can be calculated. As we shall see later in this model, if the SM fermion mass matrices receive contributions from only three-point functions, one can only obtain the correct quark masses and tau lepton mass, but it is not possible to explain the CKM quark mixings, and the electron and muon masses. Introducing corrections from four-point functions, one can indeed obtain the correct quark and the tau lepton masses, and the CKM quark  mixings, and as well as the electron mass. However, the muon mass is still about 36\\% smaller than the desired value.  In this paper, we review the construction of four-point functions, and include both trilinear Yukawa couplings and four-point interactions to obtain the SM fermion mass matrices in an explicit intersecting D-brane model.  We show that it is indeed possible to better match the masses and mixings for all fermions when four-point corrections are included.  In addition, these corrections may have broader application to other models. ",
        "conclusions": "We have discussed corrections to the SM fermion Yukawa couplings in intersecting D6-brane models due to four-point interactions and presented a working example, and demonstrated that these corrections can improve the best fits for the SM fermion masses and mixings. In a concrete model, we first calculated the SM fermion masses and mixings from three-point functions. Considering only these contributions, we can obtain the correct quark masses and tau lepton mass, but the CKM quark mixings are not large enough, and the electron and muon masses are far from the desired values. After including the corrections to the SM fermion Yukawa couplings from four-point functions, we can obtain the correct quark masses and CKM mixings, and the correct electron and tau lepton mass scales. However, the muon mass is still around 36\\% smaller than the desired value.  In this work, we have not considered the moduli stabilization problem, and have essentially treated the moduli VEVs as free parameters. As is obvious, there exists fine-tuning in our discussion of the D-brane positions and the Higgs VEVs. The four-point interactions have nine dependent worldsheet instanton parameters as well as six additional string-scale Higgs particles and their VEVs.  To stabilize these undetermined variables, one may consider a $\\mathbf{\\Z_2\\times \\Z_2'}$ orientifold model with discrete torsion where the D-branes wrap rigid cycles, thus stabilizing the open-string moduli.  In this case, one would also expect corrections to the Yukawa couplings from D-brane instantons. Indeed, E2-branes required for this construction must also wrap rigid cycles, so there is additional motivation to consider this background.  We plan to pursue these possibilities in our future research."
    },
    "0807/0807.0603_arXiv.txt": {
        "abstract": "Debris disks around main-sequence stars are believed to derive from planetesimal populations that have accreted at early epochs and survived possible planet formation processes. While debris disks must contain solids in a broad range of sizes~--- from big planetesimals down to tiny dust grains~--- debris disk observations are only sensitive to the dust end of the size distribution. Collisional models of debris disks are needed to ``climb up'' the ladder of the collisional cascade, from dust towards parent bodies, representing the main mass reservoir of the disks. We have used our collisional code to generate five disks around a sun-like star, assuming planetesimal belts at 3, 10, 30, 100, and $200\\AU$ with 10 times the Edgeworth-Kuiper-belt mass density, and to evolve them for 10~Gyr. Along with an appropriate scaling rule, this effectively yields a three-parametric set of reference disks (initial mass, location of planetesimal belt, age). For all the disks, we have generated spectral energy distributions (SEDs), assuming homogeneous spherical astrosilicate dust grains. A comparison between generated and actually observed SEDs yields estimates of planetesimal properties (location, total mass etc.). As a test and a first application of this approach, we have selected five disks around sun-like stars with well-known SEDs. In four cases, we have reproduced the data with a linear combination of two disks from the grid (an ``asteroid belt'' at $3\\AU$ and an outer ``Kuiper belt''); in one case a single, outer component was sufficient. The outer components are compatible with ``large Kuiper belts'' of 0.2--50 earth masses (in the bodies up to $100\\km$ in size) with radii of $100$--$200\\AU$. ",
        "introduction": "Since the IRAS discovery of the excess infrared emission around Vega by \\citet{aumann-et-al-1984}, infrared surveys with IRAS, ISO, Spitzer, and other space-based and ground-based telescopes have shown the Vega phenomenon to be common for main-sequence stars \\cite[e.g.][]{meyer-et-al-2004,beichman-et-al-2005,najita-williams-2005,rieke-et-al-2005, bryden-et-al-2006,siegler-et-al-2006, su-et-al-2006, trilling-et-al-2007,hillenbrand-et-al-2008,trilling-et-al-2008}. The observed excesses are attributed to circumstellar disks of second-generation dust, sustained by numerous planetesimals in orbit around the stars. Jostling collisions between planetesimals grind them all the way down to smallest dust grains which are then blown away by stellar radiation. While the bulk of such a debris disk's mass is hidden in invisible parent bodies, the observed luminosity is dominated by small particles at dust sizes. Hence the studies of dust emission have a potential to shed light onto the properties of parent planetesimal populations as well as planets that may shape them and, ultimately, onto the evolutionary history of circumstellar planetary systems.  However, there is no direct way to infer the properties of invisible planetesimal populations from the observed dust emission. Dust and planetesimals can only be linked through models. First, dynamical models can be used to predict, for a given planetesimal family (mass, location, age, etc.), the distribution of dust. Such models have become available in recent years \\citep[e.g.][]{thebault-et-al-2003,krivov-et-al-2006,thebault-augereau-2007,wyatt-et-al-2007,loehne-et-al-2007}. After that, standard thermal emission models will describe the resulting dust emission. Comparison of that emission to the one actually observed would then reveal the probable properties of underlying, dust-producing planetesimal families.  In this paper, we follow this approach and generate a set of hypothetical debris disks around G2 dwarfs with different ages (10~Myr -- 10~Gyr), assuming debris dust to stem from planetesimal belts with different initial masses at different distances from the central star. For every set of these parameters, we simulate steady-state dust distributions with our collisional code \\citep{krivov-et-al-2005,krivov-et-al-2006,loehne-et-al-2007}. This is different from a traditional, ``empirical'' approach, in which dust distributions are postulated, usually in form of power laws, parameterized by ranges and exponents that play a role of fitting parameters \\citep[e.g.][]{wolf-hillenbrand-2003}. Interestingly, replacing formal dust distributions with those coming out of dynamical modeling does not increase the number of fitting parameters. Just the opposite: the number of parameters reduces and those parameters that we keep free all have clear astrophysical meaning. The most important are location of a parent planetesimal belt and its current mass \\citep{wyatt-et-al-2007}.  Having produced a set of model debris disks,  we compute thermal emission fluxes in a wide range of wavelengths from mid-infrared to millimeter. In so doing, we completely abandon simple blackbody or modified blackbody calculations and solve a thermal balance equation instead. At this stage, we assume compact spherical grains composed of astronomical silicate \\corrected{\\citep{laor-draine-1993}} and employ standard Mie calculations to compute dust opacities. Although this is still a noticeable simplification, it represents a natural step towards considering realistic materials and using more involved methods of light scattering theory that we leave for subsequent papers.  As a test and a first application of the results, we re-interprete available observational data on a selection of disks around sun-like stars with well known spectral energy distributions (SEDs).  This paper is organized as follows. Section 2 describes the dynamical and thermal emission models. In section~3, a set of reference disks is introduced and the model parameters are specified. Section~4 presents the modeling results for this set of disks: size and spatial distribution of dust, dust temperatures, and the generated SEDs. Application to selected observed disks is made in section 5. Section~6 summarizes the paper. ",
        "conclusions": "Debris disks around main-sequence stars may serve as tracers of planetesimal populations that have accumulated at earlier, protoplanetary and transitional, phases of systems' evolution, and have not been used up to form planets. However, observations of debris disks are only sensitive to the lowest end of the size distribution. Using dynamical and collisional models of debris disks is the only way to ``climb up'' the ladder of the collisional cascade, past the ubiquitous $\\mu$m-sized grains towards parent bodies and towards the main mass reservoir of the disks.  The main idea of this paper has been to take a grid of planetesimal families (with different initial  masses, distances from a central star etc.), to collisionally ``generate'' debris disks from these families and evolve them with the aid of an elaborated  collisional code, and finally, to calculate SEDs for these disks. A comparison/fit of the observed SEDs with the pre-generated SEDs is meant to allow quick conclusions about the properties of the planetesimal belt(s) that maintain one or another observed disk.  Our specific results are as follows:  1. We have produced five reference disks around a G2V star from planetesimal belts at 3, 10, 30, 100, and $200\\AU$ with 10 times the EKB mass density and evolved them for 10~Gyr. With an appropriate scaling rule (Eq.~\\ref{scaling1}), we can translate these results to an arbitrary initial disk mass and any age between 10~Myr and 10~Gyr. Thus, effectively we have a three-parametric set of reference disks (initial mass, location of planetesimal belt, age). For all the disks, we have generated SEDs, assuming astrosilicate (with tests made also for blackbody and amorphous carbon).  2. We have selected five G2V stars with good data (IRAS, ISO/ISOPHOT, Spitzer/IRAC, /IRS, /MIPS, Keck II/LWS, and JCMT/SCUBA) and tested our grid against these data. For all five systems, we have reproduced the data points within the error bars with a linear combination of two disks from the grid (an ``asteroid belt'' at $3\\AU$ and an outer ``Kuiper belt''). This automatically gives us the desired estimates of planetesimals (location, total mass etc.).  3. A comparison of the observational data on the five stars with the grid of models leads us to a conclusion that the cold emission (with a maximum at the far-IR) is compatible with ``large Kuiper belts'', with masses in the range 3--50 earth masses and radii of $100$--$200\\AU$. These large sizes trace back to the facts that the collisional model predicts the observed emission to stem from micron-sized dust grains, whose temperatures are well in excess of a blackbody temperature at a given distance from the star \\citep[as discussed, e.g., in][]{hillenbrand-et-al-2008}. This conclusion is rather robust against variation in parameters of collisional and thermal emission models, and is roughly consistent with disk radii revealed in scattered light images (e.g. HD~107146). Still, quantitative conclusions about the mass and location of the planetesimal belts would significantly depend on (i) the adopted model of collision outcomes \\corrected{(which, in turn, depend on the dynamical excitation of the belts, i.e. on orbital eccentricities and inclinations of planetesimals)} and (ii) the assumed grains' absorption and emission efficiencies. For example, a less efficient cratering (retaining more grains with radii $\\sim10\\mum$ in the disk) and/or more ``transparent'' materials (making dust grains of the same sizes at the same locations colder) would result in ``shifting'' the parent belts closer to the star.  In future, we plan to extend this study in two directions. First, we will investigate more systematically the influence of the dust composition by trying relevant materials with available optical data rather than astrosilicate; this should be done consistently in the dynamical/collisional and thermal emission models. Second, it is planned to extend this study to stars with a range of spectral classes. This will result in a catalog of disk colors that should be helpful for interpretation of data expected to come, most notably from the Herschel Space Observatory."
    },
    "0807/0807.0573_arXiv.txt": {
        "abstract": "{Though still beyond our observational abilities, Population III stars are interesting objects in many perspectives. They are responsible for the re-ionisation of the inter-galactic medium. They also left their chemical imprint in the early Universe, imprint which can be deciphered in the most metal-poor stars in the halo of our Galaxy.} {Rotation has been shown to play a determinant role at very low metallicity, bringing heavy mass loss where almost none was expected. Is this still true when the metallicity strictly equals zero? The aim of our study is to get an answer to this question, and to determine how rotation changes the evolution and the chemical signature of the primordial stars.} {We have calculated seven differentially-rotating stellar models at zero metallicity, with masses between 9 and 200 $M_{\\sun}$. For each mass, we have also calculated a corresponding model without rotation. The evolution has been followed up to the pre-supernova stage.} {We find that $Z=0$ models rotate with an internal profile $\\Omega(r)$ close to local angular momentum conservation, because of a very weak core-envelope coupling. Rotational mixing drives a H-shell boost due to a sudden onset of CNO cycle in the shell. This boost leads to a high $^{14}$N production, which can be as much as $10^6$ times higher than the production of the non-rotating models. Generally, the rotating models produce much more metals than their non-rotating counterparts. The mass loss is very low, even for the models that reach the critical velocity during the main sequence. It may however have an impact on the chemical enrichment of the Universe, because some of the stars are supposed to collapse directly into black holes. They would contribute to the enrichment only through their winds. While in that case non-rotating stars would not contribute at all, rotating stars may leave an imprint in their surrounding. Due to the low mass loss and the weak coupling, the core retains a high angular momentum at the end of the evolution. The high rotation rate at death probably leads to a much stronger explosion than previously expected, changing the fate of the models. The inclusion of our yields in a chemical evolution model of the Galactic halo predicts log values of N/O, C/O and $^{12}$C/$^{13}$C ratios of -2.2, -0.95 and 50 respectively at log O/H $+12 = 4.2$.} {} ",
        "introduction": "Population III (PopIII) stars occupy a key position in the evolution of the Universe. They are the first sources of light that have re-ionised the Universe. Also, they are the first producers of the heavy elements in the early Universe\\footnote{According to the standard Big Bang Nucleosynthesis, only elements with nucleon number $A\\leq 7$ are formed at a significant level during the first minutes of the Universe \\citep[\\emph{e.g.}][]{iocco07}.}.  Because of the peculiar conditions of the early Universe, especially the lack of coolants provided by the metals, the very first metal-free stars that formed are supposed to be very massive, with a typical mass $M\\approx 10^2 - 10^3\\ M_{\\sun}$ \\citep{abn02,bcl02}. \\citet{nakum01} propose a bimodal IMF, with one peak at $M\\approx 10^2\\ M_{\\sun}$ and the other at $M\\approx 1\\ M_{\\sun}$. The bimodality is due to a threshold in the initial density of the cloud. Since this first suggestion that PopIII may not all be very massive objects, several studies show that the radiation field of a pre-existing star have a strong impact on subsequent star formation \\citep{omyosh03, oshabwn05, greifbr06}. This would allow lower-mass metal-free stars to appear. Moreover \\citet{silk06} show that even if the early Universe was free of magnetic fields, magnetic seeds can be generated in the disc surrounding an accreting stellar progenitor. These seeds can be amplified through a dynamo process and become strong enough to play an important role in the accretion rate. This again would reduce substantially the final mass range of the first stars. Though it has once been suggested that the most metal-poor stars known to this day may be low-mass PopIII that have undergone surface pollution \\citep{shitsuyo03}, it is now widely accepted that metal-free stars were too massive to still be present nowadays ($M > 0.8 M_{\\sun}$). The observation facilities are not yet able to reach far enough for a direct detection \\citep[except maybe for the detection of pair-instability supernovae at intermediate redshifts, see][]{scan05}. The only way to constrain their physical properties is indirect observations of their chemical or radiative signature. Though it is not firmly established whether the most metal-deficient stars of the halo of our Galaxy are truly second generation stars \\citep{salva07}, most of the authors consider they are born from a cloud enriched by only a few primordial progenitors \\citep{christ02,umnom03}. Recently, accurate measurements of their surface abundances have started to allow some constraints to be put on the theoretical results of PopIII evolution \\citep{cayrel04,barklem05,cohen07,aoki07}.  In view of their importance, PopIII stars have been the subject of many evolutionary calculations. Amongst the latest ones, most are presenting the evolution until AGB phase or central C ignition \\citep{chieff01,mar01,mar03,siess02,gil07,lau08}. Others are focused on final stages \\citep{HW02,UN05}. To our knowledge, there are two studies that include rotating models: the one by \\citet{mar03}, and the one by \\citet{hegerFS1}.  \\citet{mar03} study very massive objects with masses between 120 and 1000 $M_{\\sun}$. In their work, rotating stars are treated as rigid bodies and the effects of rotation intervenes mainly through a correction factor applied on the radiative mass loss (see Sect. \\ref{sphymdot}). In these models, no rotational mixing is accounted for. \\citet{hegerFS1} follow the evolution of differentially-rotating stellar models of $15-250\\ M_{\\sun}$ from the zero-age main sequence (ZAMS) to core collapse, but the detailed results of these models have not been published.  The present paper explores the effects of differential rotation on primordial massive stars. We think it is worth doing such a study for two main reasons. First, for metallicities between 0.004 and 0.040, the inclusion of the effects of differential rotation in stellar evolution codes has allowed to improve the agreement between the models and observations in many aspects: it has allowed to better reproduce the observed surface abundances \\citep{heglang00,mm5}, the ratio of blue-to-red supergiants at low metallicity \\citep{mm7}, the variation with the metallicity of the WR/O ratios and of the numbers of Type Ibc to Type II supernovae \\citep{mm10,mm11}. It is interesting to see what a code that has been successfully confronted to observations can predict in a metallicity domain where no direct observations are available. Second, rotation has been shown to have important effects on low metallicity stars \\citep{mem06}. During the Main Sequence (MS), the radiative winds are very low because of the lack of metals, so the star cannot get rid of its angular momentum. If it starts its evolution with a sufficiently high equatorial velocity, it may reach the critical limit. In that case, the outer layers are no more bound and some mass is lost, probably through a \\emph{decretion} disc \\citep{ow05} that may be quickly erased by the radiation. Further in the evolution, the rotational mixing brings freshly synthesised metals at the surface. The opacity of the outer layers are enhanced and the radiative mass loss is increased. Whereas the low metallicity stars were supposed to lose almost no mass at all, it has been shown that fast rotation can dramatically enhance the mass loss and thus modify drastically the evolution of these stars. How does a total absence of metals alter the above picture? It is interesting to extend the study to strictly $Z=0$ models.  The paper is organized as follow: in Sect. \\ref{sphymod}, we describe the physical ingredients of the modelling; Sect. \\ref{sevol} is dedicated to the presentation of the evolution of the models; in Sect. \\ref{sfinstruc}, we present the pre-SN structure and the yields obtained; we study the effect of their inclusion in a galactic chemical evolution code in Sect.~\\ref{sgalevol}, and then a discussion follows in Sect. \\ref{sdiscu}.  Let us mention that electronic tables for the evolutionary sequences and the final structure files are available at \\url{http://obswww.unige.ch/Recherche/evol/Rotating-primordial-stars}. ",
        "conclusions": "}  We computed a set of metal-free stellar models with masses between 9 and 200 $M_{\\sun}$, following their evolution until the pre-supernova stage. The set was composed of seven differentially rotating models with initial velocity 800 km s$^{-1}$ (except for the 9 $M_{\\sun}$ model which started the main sequence with 500 km s$^{-1}$), and a comparison set of non-rotating models. We showed that in PopIII stars, although the effects of rotation are important, they are less strong than in stars with even a tiny amount of metals ($Z=10^{-8}$). There are two main causes for this: \\begin{enumerate} \\item the meridional circulation is extremely weak at $Z=0$. The dependance of the outer cell in $1/\\rho$ (see Sect.~\\ref{sveq}) makes it a factor 100 lower compared to the value at standard metallicity, and still a factor 4 compared to $Z=10^{-5}$. The evolution of the angular momentum is close to local conservation, so the critical limit is reached late. Moreover, only a very small amount of mass needs to be lost in order to bring the surface back to sub-critical rotation; \\item there is little structural readjustment at the end of the MS, because the core is already hot enough to burn some He during the main sequence, so the nuclear burnings are continuous. The transition between H and He burning is smooth and the models remain on the blue side of the Hertzsprung-Russell diagram for a long part of the core He burning. When they eventually reach the red side, their outer convective zone remains thin, so the dredge-up is not efficient and the surface enrichment stays low. The envelope remains transparent to radiation and the mass loss is negligible. \\end{enumerate} There is actually a mass loss where none was expected (because of the wind dependance on $Z$), but it remains modest even for the most massive of our models.  Could this picture be modified by changing the physical inputs in the models? We have shown that the situation is qualitatively the same when we take a faster initial rotation rate. The main problem is the weakness of the meridional circulation, which does not efficiently feed the outer layers with angular momentum taken from the fast spinning core. So, even if the critical limit is reached a little earlier, the models still do not lose a lot of mass. For this work, we did not take into account the anisotropy of the winds in our models. We do not think it would change the picture presented here, because the loss of angular momentum by radiative mass loss is negligible. There is still an interesting possibility: the implementation of magnetic fields. They have been shown to change many features in stellar evolution \\citep{hws05,mmB3}, and to provide the strong core-envelope coupling that is lacking in the $Z=0$ models. With such a coupling, the mass loss becomes much larger and the evolution may be strongly altered \\citep{sekauai07}.  With the initial velocities considered here, our models do not follow a chemically homogeneous evolution \\citep{Maed87,heglang00}. In the frame of the physics we used, homogeneous evolution for Pop III stars would need faster rotation rates to occur. When the magnetic dynamo mechanism proposed by \\citet{Spruit99,Spruit02} is taken into account, chemical mixing becomes more efficient and homogeneous evolution can be obtained for smaller initial rotation rates \\citep{yoon05,yoon06,woosheg06}. Note however that for extremely low metallicities, the weakening of the meridional circulation becomes a real handicap for the models to evolve chemically homogeneously. \\citet{hir08} find that at $Z=10^{-8}$, a 40 $M_{\\sun}$ model (computed with magnetic fields) starting its evolution with $\\upsilon/\\upsilon_{\\rm crit}=0.55$ does not rotate fast enough to follow such an homogeneous evolution. We expect that at $Z=0$ the problem would be still more acute. The consequences of such extreme effects of rotation is left for a future work.  It is important to underline here that the only prescription for stellar dynamo existing to date, the Tayler-Spruit formalism \\citep{Spruit99,Spruit02}, is still debated:  recent works have shed some doubts on the efficiency of the Tayler-Spruit dynamo mechanism. For instance, \\citet{zbm07} have shown that although the Pitts and Tayler instability indeed develop, it does not succeed in driving a dynamo. Moreover, this dynamo needs a pre-existing magnetic-field seed, and it is not yet clear whether magnetic fields were present in the early Universe or not \\citep[see][for a detailed discussion on this subject]{giov04}.  We presented pre-supernova yields for our stellar models. We included them in a chemical evolution model of the Galactic halo. Our results predict that the N/O and the C/O ratios should drop at extremely low metallicity ($Z\\leq 10^{-6}$). For the time being though, the observational discrimination between the two curves is far beyond our observational possibilities. We also gave the rotational characteristics of the iron core of our models and showed that since they do not lose much angular momentum during their evolution, they keep high rotation rates in the core, so their final explosion may be much more energetic than sometimes assumed. We showed that the most massive of our models may contribute to the emission of the gravitational wave background.  Many interesting questions and issues were not treated in this study. Our mass coverage was too coarse to constrain accurately the limiting masses for undergoing a CNO shell boost, for example, or for the different fates at the end of the evolution. Even without knowing if magnetic fields were indeed present in the early Universe, it would be interesting to study the role they could play in primordial stellar evolution. As mentioned above, they could drastically change the picture presented here. We have also left for a forthcoming paper the question of the re-ionisation. Rotation changes the luminosity, the $T_\\mathrm{eff}$ and the lifetimes of the models, so it should change the ionising power of the first stars."
    },
    "0807/0807.0059_arXiv.txt": {
        "abstract": "We present \\cha-\\hetgs\\ observations of the Herbig Ae star \\hd\\ and the associated young stars comprising lower mass stars, in the 0.15-1.75\\msol\\ mass range, in their pre-main sequence phase. The brightest X-ray source in the association is the central system harboring the Herbig Ae primary, and a K3 companion. Its X-ray variability indicates modulation possibly on time scales of the rotation period of the Herbig Ae star, and this would imply that the primary significantly contributes to the overall emission. The spectrum of the Herbig Ae+K3 system shows a soft component significantly more pronounced than in other K-type young stars. This soft emission is reminiscent of the unusually soft spectra observed for the single Herbig Ae stars HD~163296 and AB~Aur, and therefore we tentatively attribute it to the Herbig Ae of the binary system. The \\hetgs\\ spectrum shows strong emission lines corresponding to a wide range of plasma temperatures. The He-like triplet of \\mgxi\\ and \\neix\\ suggest the presence of plasma at densities of about $10^{12}$~cm$^{-3}$, possibly indicating accretion related X-ray production mechanism.  The analysis of the zero-order spectra of the other sources indicates X-ray emission characteristics typical of pre-main sequence stars of similar spectral type, with the exception of the T~Tauri \\hd-D, whose extremely soft emission is very similar to the emission of the classical T~Tauri star TW~Hya, and suggests X-ray production by shocked accreting plasma. ",
        "introduction": "\\label{s:intro}   Nearby young associations provide ideal targets for the study of star and planet formation, and have sparked growing interest in the past few years. The early evolution of young stars is strongly dependent on stellar mass (e.g., \\citealt{Larson72}). For low-mass stars ($M_{\\star}\\lesssim 2$\\msol) the accretion timescales are shorter than the evolution time to the zero-age main sequence, and these stars can be observed in their pre-main sequence (PMS) T Tauri phase \\citep{Dantona94}. In massive protostars ($M_{\\star}\\gtrsim 10$\\msol), however, core hydrogen burning already starts during the accretion phase \\citep{Appenzeller94,Bernasconi96}.  The behavior at the boundaries between high and low mass ranges is yet unclear \\citep{Palla93}.  Herbig Ae stars (HAe; \\citealt{Herbig60}) are PMS stars of intermediate mass ($\\sim$2--10\\msol), and share several characteristics with T Tauri stars (TTS), such as IR excess emission, irregular photometric variability, and disk properties \\citep{Waters98,Mannings97}. HAe stars show evidence for a disk-like geometry of the circumstellar material, and the disk properties are similar to those of TTS \\citep{Mannings97,Grady99}. There is some evidence of magnetic accretion in HAe stars, analogous to the magnetospheric accretion scenario generally accepted for classical TTS (CTTS, i.e.\\ TTS that are still actively accreting) and are therefore thought to be more massive analogs of CTTS (e.g., \\citealt{Muzerolle04,Grady04,Guimaraes06}). While magnetic fields have been detected in a few HAe stars (e.g., \\citealt{Donati97,Hubrig04}), their origins are still unknown; however, a shear magnetic model does allow generation of magnetic fields in the early evolutionary stages of A-stars \\citep{Tout95}.  Systematic X-ray studies have shown that HAe stars are moderately bright X-ray sources \\citep{Damiani94,Zinnecker94,Hamaguchi05,Stelzer06}. Since main sequence A-stars are not known to be luminous X-ray sources due to lack of strong enough winds or magnetic-dynamo driven coronae, the findings may imply that the physical characteristics of young A-stars stars differ from those of main sequence A-stars.  HAe X-ray emission mechanisms remain unknown. Early studies based on ROSAT and {\\em Einstein} observations proposed the standard wind shock model as for early-type massive stars \\citep{Damiani94,Zinnecker94}. However, typical \\lx/\\lbol\\ ratios found for HAe stars are significantly higher than the empirical luminosity ratio of $10^{-7}$ for early type stars \\citep{Skinner04}, casting some doubt for this paradigm. Also, the high X-ray temperature ($> 10^7$~K) component found for a large fraction of HAe X-ray sources (see e.g., \\citealt{Skinner04,Hamaguchi05,Stelzer06}) cannot be produced through this mechanism. Others works attribute the emission to magnetic activity (either coronal activity as in late-type stars, or magnetic activity due to star-disk interaction), or to unresolved late-type companions \\citep{Hamaguchi05,Skinner96,Skinner04,Stelzer06}.  Another emission mechanism possibly at work in HAe stars is X-ray emission from shocks in accreting plasma, proposed as the production mechanism for soft X-rays in several CTTS (TW~Hya, BP~Tau, V4046~Sgr, MP~Mus; \\citealt{Kastner02,Stelzer04,Schmitt05,Gunther06,Argiroffi07}). For these CTTS, \\cha\\ and \\xmm\\ high resolution spectra have revealed peculiar characteristics with respect to all other observed stellar spectra, in particular showing (i) soft excess---prominent cool component at temperatures of $\\sim 3 \\times 10^6$~K, TW~Hya being an extreme case, where this soft component largely dominates the X-ray spectrum; (ii) unusually high density (\\ne$\\gtrsim 10^{12}$~cm$^{-3}$) for the cool ($T \\sim 2-4 \\times 10^6$~K) plasma, as diagnosed through the analysis of the He-like triplets line ratios; (iii) peculiar abundances---extremely high Ne, low metal abundances \\citep{Kastner02,Stelzer04,Drake05}. Analogously, \\cha\\ imaging spectra of the HAe star HD~163296 revealed a very soft spectrum ($\\sim 0.5$~keV) possibly suggesting accretion dominated X-ray emission  \\citep{Swartz05}, as for the CTTS TW~Hya \\citep{Kastner02}; \\cite{Grady07} also find extremely soft X-ray emission ($\\sim 0.25$~keV) from the Herbig Ae star HD~169142. Recently, \\xmm\\ observations of the HAe star AB~Aur have shown a very similar soft spectrum \\citep{Telleschi07}. In this case, the high resolution spectrum provides additional diagnostics: the high forbidden/intercombination ratios ($f/i$) of the He-like triplets, especially \\ovii, indicate the absence of high density and high UV field, expected if X-rays are produced in the accretion shocks, and rule out accretion-related X-ray emission for AB~Aur. For AB~Aur, which has spectral type A0, the photospheric UV emission is strong enough to alter the O He-like triplet up to several stellar radii from the surface; the large measured \\ovii\\ $f/i$ line ratio, constrains the emission to originate high above the photosphere ($\\gtrsim 3$\\rstar) and therefore it implies that also coronal emission is unlikely for AB~Aur. A plausible alternative for this source is X-ray emission from magnetically confined winds, as in the scenario of \\cite{Babel97}. It is important to keep in mind that some of the different X-ray emission mechanisms might not be mutually exclusive and instead they might coexist, as seems to be the case for some CTTS \\citep{Schmitt05,Gunther06,Argiroffi07} where both magnetically confined coronal plasma and accretion streams likely contributes to the overall X-ray emission.  \\hd\\ is one of the nearest known ($d=116$~pc; \\citealt{Perryman97}) and well studied HAe stars, and it is also a strong X-ray source (\\lx~$ \\sim 2 \\times 10^{30}$~erg~s$^{-1}$). \\cha\\ observations have spatially resolved the X-ray emission from \\hd, showing that the brightest X-ray source is coincident with the HAe source, and identifying 4 previously unknown nearby sources, likely TTS forming a young stellar group associated with \\hd\\ (\\citealt{Feigelson03}, hereafter \\FHa; see \\S\\ref{s:targets}). \\cha\\ (\\FHa) and \\xmm\\ \\citep{Skinner04} imaging spectra of this HAe star have shown the presence of hot plasma ($\\gtrsim 2$~keV), taken as evidence of magnetic activity. \\hd\\ with its proximity, high \\lx, and low line-of-sight extinction is a unique target to study the X-ray emission mechanism in HAe stars, through high resolution spectroscopy. High resolution spectra provide plasma diagnostics which are the only means for detailed study of the physical conditions of the emitting material.  Cool, dense plasma ($\\log T\\mathrm{[K]}\\sim 6.5$, $\\log n_e [\\mathrm{cm^{-3}}] > 12$) is now thought to be a signature of accretion, while variable hot emission is indicative of magnetic activity.  In this paper we present \\cha-\\hetgs\\ observations of \\hd. The relatively long exposure of $\\sim 145$~ks provides detailed information on the X-ray variability of the HAe star, as well as of other stars in the field, and it yielded a well exposed \\hetgs\\ spectrum of the HAe star providing detailed diagnostics for the emitting plasma.  We present the characteristics of the targets in \\S\\ref{s:targets}. The observations and our techniques of line flux measurement and spectral analysis are briefly described in \\S\\ref{s:obs}. The results are presented in \\S\\ref{ss:results}. We discuss the results and draw our conclusions in \\S\\ref{s:discuss}, and \\ref{s:conclusions}. ",
        "conclusions": ""
    },
    "0807/0807.1480_arXiv.txt": {
        "abstract": "{We analyse the line and continuum spectra of the symbiotic system CH Cygni.} {To show that the colliding-wind model is valid  to explain this symbiotic star at different phases.} {Peculiar observed features such as flickering, radio variation, X-ray emission, as well as the distribution of the nebulae and shells throughout the system are investigated by modelling the spectra at different epochs. The models account consistently for shock and photoionization and are constrained by absolute fluxes.} {We find that the reverse shock between the stars leads to the broad lines observed during the active phases, as well as to radio and hard X-ray emission, while the expanding shock is invoked to explain the data during the transition phases.} {} ",
        "introduction": "The symbiotic star (SS) CH Cygni (CH Cyg) has been observed since the late '80s of the XIX century.  It was identified as an anonymous red semiregular variable  by Gaposchkin's 1952 analysis of Harvard patrol plates, and classified as a M6-M7 spectrum star with a 90-100 day period and about one-magnitude amplitude. This classification seemed to be suitable until 1963 when a hot blue continuum appeared, along with a late-type spectrum and  emission lines of H, He, [FeII] and [CaII], evidence of a composite nature. This drove Deutsch (1964) to recognize in CH Cyg the brightest SS at visual wavelengths. Since then, active and quiescent phases have been alternating without a predictable trend.  At a distance of 270$\\pm$66 pc (according to the Hipparcos parallax measurements - Munari et al. 1997) CH Cyg is one of the most studied, yet poorly understood SSs. It has been observed in the last thirty years both photometrically and spectroscopically from radio to X-ray wavelengths (Kenyon 2001).  The basic configuration of the CH Cyg system is still debated, despite the large amount of multifrequency observations carried out since the first recorded symbiotic activity. In particular, it is discussed whether it is a binary or a triple star, and, in the latter case, whether composed of  a red giant/white dwarf pair with a G star companion (Hinkle et al. 1993) or a red giant/white dwarf pair with an additional red giant companion (Skopal et al. 1996). The explanation of the total hot component luminosity in a binary star scenario might be a crucial problem: specifically, the suggested binary orbital parameters (e.g. Mikolajewski et al. 1992) would imply an accretion rate onto the white dwarf not in agreement with that requested by the observed hot component luminosity. On the other hand, adopting the binary star scenario the orbital solution of the triple model (Hinkle et al. 1993) would require a low inclination angle. This is in  contrast  with the jet  appearing  in the plane of the sky, implying an accretion disk and/or  an orbital plane virtually edge-on. Furthermore, the low inclination implied by Hinkle et al. (1993) would result in  an unrealistically low mass (0.2 \\msol) white dwarf. Finally, the similarity  of the 756 day period of CH Cyg with non-radial pulsations in a few semi-regular variables drove Schmidt et al. (2006) to interpret this periodic trend more as a non-radial pulsation of the cool component rather than a real orbital period.  Several theoretical  papers have appeared presenting SSs as colliding-wind (CW) binary systems  (e.g. Kenny \\& Taylor 2005). The observational evidence that actually both the component stars of a SS lose mass through strong stellar winds (Nussbaumer et al. 1995), along with the presence in CH Cyg of radio, optical and X-ray jets (Karovska et al. 2007), have confirmed the importance of shocks for a real understanding of such outstanding system.  CH Cyg is characterized  by the presence of an eruptive white dwarf, showing episodes of accretion fed by the high-density wind  coming from the M star, and of wind-wind interaction leading to shock waves. The nebulae downstream of the shock fronts  show  the spectroscopic marks of a turbulent kinematics, resulting in an inhomogeneous and  highly disrupted envelope.  In this paper we review the recent history of CH Cyg, in the light of a colliding wind scenario (e.g. Girard \\& Willson 1987, Kenny \\& Taylor 2005)  that  has been used  to explain many SSs (e.g. Formiggini, Contini, \\& Leibowitz 1995, Contini 1997, Contini \\& Formiggini 1999) and  recently updated including a detailed analysis of dust features (Angeloni et al. 2007a,b,c).  We interpret the observational data throughout both the active and quiescent phases, from the 1978 great outburst up to date,  with the aim of investigating the physical  and morphological conditions of the system by modelling the line and continuum spectra emitted  from the different nebulae.  We adopt the binary scenario,  considering the $\\sim$ 756 day period as due to the giant pulsation mode, and the $\\sim$ 15 year period as the actual orbital period.  Following  the method used in previous SS investigations, we analyse the continuum spectral energy distribution (SED) in different epochs, instead of the light curve through the years. At each epoch we model the data collected from radio to X-ray constraining the results by the analysis of the line spectra. The continuum accounts for bremsstrahlung and dust reprocessed radiation  from the shocked nebulae, radiation from dust shells, as well as  radiation from the stars. Our method is particularly suitable for disentangling the multiple contributions of gas and dust in different conditions at different phases.  We focus on the $nebulae$  within the shocked wind-wind interaction regions and on the $shells$ ejected by the giant star, constraining the models by the observed $absolute$ fluxes.  A short summary of observational history of CH Cyg is presented in Sect. 2. The colliding wind model and  our application to CH Cyg are presented in Sect. 3. The next sections are  devoted to the analysis of the individual active and quiescent phases. As CH Cyg has not been observed regularly in the different wavelength ranges, except of the optical one, the SED profiles in some epochs present some data gaps:  for instance, there are no UV and X-ray observations for the 1988-1990 quiescent phase, nor mid-IR data for the 1998-2000 active phase. This  led us to focus on specific physical processes in each phase: the basic models which result by disentangling the different contributions to the SEDs and to the line ratios at different days during the active phase 1978-87 are presented in Sect. 4; in Sect.5, different types of IR obscurations occurred during the 1988-90 quiescent phase are indicated; near-UV and IR variability during the 1991-95 active phase are discussed in Sect. 6, while Sect. 7 (1996-97 quiescent phase) deals with the complex nature of the radio emission. The very expanded nebulae and the X-ray emission are  modelled throughout the 1998-2000 latest active phase (Sect.8) up to date (Sect.9). Concluding remarks and a summary of the results  are presented in Sect. 10. \\begin{figure} \\begin{center} \\includegraphics[width=0.45\\textwidth]{Fig1.eps} \\caption{AAVSO light curve (10-day means of observations) of CH Cyg from 1948 to 2000.\\label{AAVSO}} \\end{center} \\end{figure} ",
        "conclusions": "The  complex nature of the SS CH Cygni has been investigated by modelling  the line and continuum  spectra at different epochs, from the  impressive outburst of 1978 up to date. Our results give an approximated picture of the shocked nebulae throughout the system because our approach  implies a plane parallel geometry  for the different shock fronts. This is constrained by the composite models which account for both photoionization from the WD and shocks created by the collision of the winds. However, the plane-parallel geometry can be justified for relative large radii and in particular situations characteristic of SS.   The present analysis has revealed that  similar developing trends may be recognised throughout the active  phases and throughout the quiescent phases. In the light of a colliding-wind scenario, we have specifically focused on the role of the nebulae created by collision of the winds from the hot and the cool stars, in order to build up a comprehensive view of the system as a whole, namely to self-consistently interpret the variable line spectra, the radio emission, the IR photometric properties, the fluxes in the U B V bands and the X-ray observations. The most important result of our method  that can be applied to all the SSs, consists in recognizing  the different sources of radiation (stars, nebulae, dust) throughout the continuum SED.  We have found that bremsstrahlung from the nebula downstream of the  shock front facing the WD between the stars  explains the continuum slopes observed during the active states. These fast shocks seem to be always present during the outbursts,  proving that the wind collision is playing a primary role in the spectral properties of the symbiotic phenomenon (e.g. radio and X-ray wavelength region). The circumstellar dusty shells define the optic-UV SED , also contributing to the line spectra during the active phases.  The WD may be related to flickering during the active phases, either by appearing directly (1991-95) or by photoionising the shells throughout a disrupted medium (1978-87). Towards the end of the active phases, the expanding shock propagating outward the system is visible both in the continuum SED and in the line spectra.  During the quiescent phases  the circumbinary expanding shocks  dominate. We have found that these shocks propagate  throughout a relatively dense medium which may be interpreted as matter ejected in previous outbursts. This leads to a relative strong  synchrotron radiation   from the shock front at some epochs.  The jets observed by Corradi et al (2001)   imply the presence of an accretion disk. The modelling of the low-ionisation lines (e.g. [NII, [OII]), that are strong in their images, confirms that the jets expanding in the ISM, have reached a radius $>$ 0.01 pc from the SS.  We have shown,  by consistent modelling,  that soft and hard X-rays are emitted downstream of shock fronts corresponding to quite different velocities, namely, the expanding shock (\\Vs $\\sim$ 100-150 \\kms) propagating  outwards the binary system would explain the soft X-ray emission, while the hard X-rays  correspond to the high velocity shock between the stars, in agreement with the recent results discussed by Karovska et al. (2007) on the basis of radio, optical and X-ray imaging.  Dust grains within the nebulae are collisionally heated by the  gas in a strong shock velocity regime ($\\geq$ 600 \\kms), while, in the dusty shells, grains are mainly heated by radiation from the hot star. The dust expanding  disrupted shells might reach, in some cases, the external edge of the symbiotic system (r$>$ 10$^{14}$ cm),  becoming circumbinary and therefore contributing to the complex obscuration episodes of CH Cyg.  It should be emphasized that the interpretation of CH Cyg spectral evolution is further complicated by the relatively short orbital period ($\\sim$14.5 years): in fact, the orbital motion combined with the unpredictable recurrent activity leads to  different system configuration as seen from Earth. We have tried to explain them by  combining the radiation fluxes from the several emitting contributors with different, variable weights.  For the sake of clarity, we schematically summarise the results of our modelling: \\subsection*{1978-1987 active phase} \\begin{itemize} \\item Emission lines with FWHM of 400 \\kms - 1200 \\kms correspond to the reverse shock; \\item before 1984 November, low-ionisation lines come from the shells; \\item after 1994, higher level lines are emitted both by the expanding nebula and the dusty shells; \\item at days JD2446212 and JD2446776,  UV and optical lines from the expanding nebulae appear; \\item at the end of the active phase, shock dominated models explain both line and continuum spectra. \\end{itemize} \\subsection*{1988-1990 quiescent phase} \\begin{itemize} \\item The  expanding nebula  propagates throughout  matter ejected in previous eruptions; \\item in the IR the system is obscured independently from wavelength.\\\\ \\end{itemize} \\subsection*{1991-1995 active phase} \\begin{itemize} \\item In 1991, at the onset of the new active state, the lines display FWHM$<$ 100 \\kms, indicating a disrupted expanding nebula; \\item in August 1991, broad \\Ha and \\Hb are seen in absorption, suggesting that the reverse shock is absorbing the WD flux; \\item by 1992 absorption lines disappear and M-giant molecular bands are observed. In May 1993,  UV lines  are emitted from an expanding nebula with a radius of 10$^{16}$ - 10$^{17}$ cm. \\end{itemize} \\subsection*{1996-1997 quiescent phase} \\begin{itemize} \\item The radio flux shows a synchrotron component created by the Fermi mechanism at the shock front. \\end{itemize} \\subsection*{1998-2001 active phase} \\begin{itemize} \\item  . A fast reverse shock of 1200 \\kms may explain the X-ray  Chandra observations in 2001; \\item the HST images by Eyres et al. (2002) reveal strong [OIII], which can be explained by the expanding nebulae   in agreement with  the  fit of the SED; \\item in 1998 Sept the OI, HeI, HeII, FeII, and TiII lines  are emitted from the shells. Moreover, [OIII] 5007 is evident and [OIII] 4363 is  absent both in and out of the eclipse. These lines come from    shock-dominated expanding nebulae  with relatively low velocities ($\\sim$70-100 \\kms) and relatively low densities (\\n0$\\sim$ 5000 \\cm3). In fact  the corresponding [OIII] 4363 lines are relatively low; \\item the 1999 images of Corradi et al. (2001) indicate strong [OII] and [NII] lines: they are characteristic of  a shock dominated nebula with a relative low \\n0=5 10$^3$ \\cm3. \\end{itemize} \\subsection*{2001-up to date} \\begin{itemize} \\item three different slopes are observed in the optical-IR range. They are modelled by  the combination of different expanding models with the bb flux from the cool component star. \\item the decline observed by Taranova \\& Shenavrin in the IR is consistent with the decline in the optical-UV range and is explained by the eclipse of the whole system. \\end{itemize}"
    },
    "0807/0807.4570_arXiv.txt": {
        "abstract": "We present a search for Trojan companions to 25 transiting exoplanets. We use the technique of Ford \\& Gaudi, in which a difference is sought between the observed transit time and the transit time that is calculated by fitting a two-body Keplerian orbit to the radial-velocity data. This technique is sensitive to the imbalance of mass at the L4/L5 points of the planet-star orbit. No companions were detected. The median 2$\\sigma$ upper limit is 60~$M_\\oplus$, and the most constraining limit is 2.5 $M_\\oplus$ for the case of GJ~436. ",
        "introduction": "Trojan companions are bodies in a 1:1 mean-motion resonance with a planet, librating around one of the two Lagrange points (L4 and L5) of the planet's orbit around the star. Several methods have been proposed to detect Trojan companions to exoplanets (Laughlin \\& Chambers 2002, Croll et al.~2007, Ford \\& Holman 2007,  Ford \\& Gaudi 2006). In this paper, we present a search for Trojan companions to 25 known transiting exoplanets for which suitable data are available, using the method of Ford \\& Gaudi (2006, hereafter, ``FG''). An important virtue of this method is that a sensitive search for Trojans can be performed using only the RV and photometric data that are routinely obtained while confirming transit candidates and characterizing the planets. ",
        "conclusions": ""
    },
    "0807/0807.3020_arXiv.txt": {
        "abstract": "The study of protoplanetary disks, where the planets are believed to form, will certainly allow the formation of our Solar System to be understood. To conduct observations of these objects at the milli-arcsecond scale, infrared interferometry provides the right performances for T Tauri, FU Ori or Herbig Ae/Be stars.  However, the only information obtained so far are scarce visibility measurements which are directly tested with models. With the outcome of recent interferometers, one can foresee obtaining images reconstructed independently of the models. In fact, several interferometers including IOTA and AMBER on the VLTI already provide the possibility to recombine three telescopes at once and thus to obtain the data necessary to reconstruct images.  In this paper, we describe the use of MIRA, an image reconstruction algorithm developed for optical interferometry data (squared visibilities and closure phases) by E.~Thi\\'ebaut. We foresee also to use the spectral information given by AMBER data to constrain even better the reconstructed images. We describe the use of MIRA to reconstruct images of young stellar objects out of actual data, in particular the multiple system GW Orionis (IOTA, 2004), and discuss the encountered difficulties. ",
        "introduction": "\\label{sec:intro}  The process of star formation triggered by the collapse and fragmentation of a molecular cloud is believed to lead to a protostar surrounded by an accreting disk undergoing important mass lost (wind, jets).  The study of the circumstellar disks during this pre-main sequence evolution is important because these disks are probably the birth location of future planetary systems. Therefore, the understanding of their physical conditions is necessary before we can understand the planets formation process.  The details of the physical processes which are taking place in the circumstellar disk are not well known because of a lack of data to constrain them.  The spatial distances at which the planetary, gas accretion and wind/jet formation occur, correspond to astronomical unit scales requiring milli-arcesecond spatial resolution to be probed unambiguously. Therefore, we used infrared interferometry which is one of the techniques capable to reach this spatial resolution at those wavelengths.  Interferometry consists of a coherent beam combination of the light received from an astronomical source and collected by two or more independent telescopes. As the light beams overlapped, we observe an interference pattern which is composed of fringes. An ideal interferometer measures the complex visibility $V(u,v)$ which is, according to the Van Cittert-Zernicke theorem, the Fourier transform of the object brightness distribution $I(x,y)$.  In reality, the atmosphere turbulence induces random phase shifts and this phenomena does allow the Fourier phases to be properly calibrated. Therefore, we need to compute interferometric measurements that are insensitive to this perturbation : \\begin{itemize} \\item the squared visibility : $V^{(2)}(u,v) = \\hspace{0.2cm} \\mid~V(u,v)~\\mid^2$ \\item the bispectrum : $B(u_i,u_j,v_i,v_j) = V(u_i,v_i) \\times V(u_j,v_j) \\times V(-u_i-u_j,-v_i-vj)$. The phase of the bispectrum is called the closure phase (CP), independant of the atmosphere delay phase. This method allows us to retrieve a part of the Fourier phase. \\end{itemize}  From these interferometric data, the astronomers up to now are fitting simple parametric or homemade models in order to directly constrain their parameters to these visibilities and closure phases and in order to better understand the structure of the inner part of the disk. Yet, although important observations from these analyzes have lead to a profound revision of the inner disk structure at spatial scale corresponding to the planet formation, the lack of model independent imaging, without any knowledge of the observed structure, still forbids unambiguous diagnostics of the physical processes at work.  With the outcome of interferometers able to combine 3 beams, like IONIC on IOTA and AMBER on the VLTI, and later of full imaging array like the second generation of instruments at VLTI, one foresees to obtain images reconstructed independently of the models. In fact, the AMBER instrument at the VLTI provides the possibility to recombine three telescopes at once and thus to obtain many data necessary to reconstruct images. ",
        "conclusions": "\\label{sec:conclu}  With only a few algorithms in the world, the image reconstruction in optical interferometry is really at its first steps especially for YSOs. The problem is much more complex than the image reconstruction in radio interferometry: in fact, in optical interferometry, we do not measure the complex visibility directly, the presence of the atmosphere leads to a lack of information on the phase and finally we have less measurements at the same time. Therefore, the optical interferometry needs its own method to retrieve images from the interferometric observables (squared visibilities and closure phases).  With one of these algorithms, MIRA, we show in this paper the first images from real long-baselines interferometric data of the environment of young stellar objects. The results have been analyzed with a new approach: we foung good correspondence between the reconstructed image and the parametric images, especially for GW Orionis. This demonstrates that this technique allows us to raise the degenerancy which appears when we analyse the interferometric data with model fitting. However, as the image reconstruction is too heavy to do statistics on the reconstructed image and as there is not a method yet which can compute a significant error in order to estimate the rightness of the reconstructed image, we have difficulties to easily derive scientifical results from the reconstructed image.  The image reconstruction in optical interferometry needs therefore to make progress but it already shows that the technique will give great facilities in the analysis of the interferometric data. With future instrument as the VLTI Spectro Imager VSI, a second generation instrument for the VLTI, we can make image reconstruction of several astronomical objects\\cite{FilhoSPIE}  : with the possibility of observing with 4 to 6 telescopes and with a groovy decrease of the error bars, the $(u,v)$ plane coverage is well much better and the reconstruction made by MIRA is easier (see Fig.~\\ref{fig:VSI} for an example on YSOs)  \\begin{figure}[tp] \\begin{center} \\begin{tabular}{c} \\includegraphics[width=0.4\\textwidth]{yso_planuv.eps} \\end{tabular} \\begin{tabular}{cc} \\includegraphics[width=0.4\\textwidth]{yso_depart_log.eps} & \\includegraphics[width=0.41\\textwidth]{yso_MIRA_yarg.eps} \\end{tabular} \\begin{tabular}{cc} \\includegraphics[width=0.4\\textwidth]{yso_vis2.eps} & \\includegraphics[width=0.4\\textwidth]{yso_cp.eps} \\end{tabular} \\end{center} \\caption{\\label{fig:VSI} Reconstructed image from simulated data for the science case of the VLTI Spectro-Imager (VSI): YSOs. First line: the $(u,v)$ coverage of the data set. Second line: the simulated image (left, logarithmic scale) and the reconstructed image by MIRA (right). Third line: the squared visibilities (left) and the closure phases (right) in red squared for the data and in green triangle for the MIRA results.} \\end{figure}"
    },
    "0807/0807.0215_arXiv.txt": {
        "abstract": "We have found the fluorine lines {\\fiv} $\\lambda\\lambda$3996.92,4059.90 in the extremely metal-poor ([Ar/H] = $-$2.10$\\pm$0.21) halo planetary nebula (PN) BoBn 1 in high-dispersion spectra from the 8.2-m VLT/UVES archive. Chemical abundance analysis shows that the fluorine abundance is [F/H] = +1.06$\\pm$0.08, making BoBn 1 the most fluorine-enhanced and metal-poor PN among fluorine-detected PNe and providing new evidence that fluorine is enhanced by nucleosynthesis in low mass metal-poor stars. A comparison with the abundances of carbon-enhanced metal-poor (CEMP) stars suggests that BoBn 1 shares their origin and evolution with CEMP-{\\it s} stars such as HE1305+0132. BoBn 1 might have evolved from a binary consisting of $\\sim$2 M$_{\\odot}$ primary and $\\sim$0.8 M$_{\\odot}$ secondary stars. ",
        "introduction": "Currently, over 1,000 objects are regarded as planetary nebulae (PNe) in the Galaxy (Acker et al. 1992), while 14 of them have been identified as halo members from their location and kinematics since the discovery of K 648 (Ps 1) in M 15. All halo PNe are metal-poor objects ($\\langle$[S/H]$\\rangle$=$-$1.3; Pereira \\& Miranda 2007) and are suspected to be descendants of stars formed in the early history of the Galaxy. Halo PNe should convey important information for the study of the evolution of metal-poor stars as well as the early physical and chemical conditions of the Galaxy. Their origins are, however, not well understood yet.  BoBn 1 (PN G108.4--76.1) is a halo PN that is located in the direction of the south Galactic pole. The heliocentric distance to this object is estimated to be between 16.5 kpc (Henry et al. 2004) and 29 kpc (Kingsburgh \\& Barlow 1992) and the heliocentric radial velocity is +191.9$\\pm$0.8 {\\kms} (Otsuka 2007). Zijlstra et al. (2006) have associated BoBn 1 with the leading tail of the Sagittarius Dwarf Spheroidal Galaxy, which traces several halo globular clusters. BoBn 1 is an extremely metal-poor ([Ar/H] = $-$2.10$\\pm$0.21; Otsuka 2007) and C- and N-rich halo PN ([C/O] = +1.50$\\pm$0.27, [N/O] = +1.02$\\pm$0.13, [O/H] = $-$0.96$\\pm$0.05, according to Otsuka 2007; [C/O] = +1.32, [N/O] = +1.07, [O/H] = $-$1.07, according to Howard et al. 1997), which composes a class of PN together with K 648 and H 4-1.  These C- and N-rich halo PNe have some unresolved issues on their origin. For example, halo PNe can become N-rich, but not C-rich, if the initial mass of the progenitor is $\\sim$0.8 M$_{\\odot}$, which is the turn-off star mass of M 15, according to the theoretical models of single metal-poor stars (e.g., Fujimoto et al. 2000). To become C-rich PNe, the third dredge-up (TDU) must take place in the late asymptotic giant branch (AGB) phase. But, the TDU takes place in stars with initial masses $\\gtrsim$ 1.2--1.5 M$_{\\odot}$ (e.g., Straniero et al. 2006). Also, current stellar evolutionary models predict that the post-AGB evolution of stars with initial masses $\\sim$0.8 M$_{\\odot}$ proceeds too slowly for a visible PN to be formed. These issues must be resolved before one proceeds to discuss the evolution of metal-poor stars and the early chemical evolution of the Galaxy through the observation of halo PNe.  Chemical abundance analysis of PNe is a key tool to investigate these issues. Comparing the observed abundances with those from theoretical evolution and nucleosynthesis models for metal-poor stars will shed light on the origin and evolution of C- and N-rich halo PNe. In this paper, we present the first estimation of the fluorine (F) abundance of BoBn 1. The abundances of F and heavy elements via the slow neutron capture process (the {\\it s}-process), which are synthesized in the He-rich intershell of AGB stars and are brought to the surface by the TDU, would be an effective tool to examine how BoBn 1 became C- and N-rich, when combined with the C abundance. We discuss the origin and evolution of BoBn 1 by comparing the F and other elemental abundances with the abundances of carbon-enhanced metal-poor (CEMP) stars. ",
        "conclusions": "In low-mass AGB stars (initial mass of 1--4 M$_{\\odot}$; Herwig 2005) suffering the TDU and not the hot bottom burning, the F abundance at the stellar surface is efficiently enhanced. Its only stable isotope, $^{19}$F, is synthesized via the $^{14}$N($\\alpha$,$\\gamma$)$^{18}$F($\\beta^{+}$)$^{18}$O({\\it p},$\\alpha$)$^{15}$N($\\alpha$,$\\gamma$)$^{19}$F reaction chain during He-shell burning inside thermal pulses (TPs). Fluorine is further carried to the surface of the star by the TDU, together with other products of He burning: mostly $^{12}$C, produced via partial He burning, and $^{22}$Ne produced via double alpha-captures on $^{14}$N (Jorissen et al. 1992; Lugaro et al. 2004). The required protons are provided by the $^{14}$N({\\it n},{\\it p})$^{14}$C reaction, where the neutrons are produced by the $^{13}$C($\\alpha$,{\\it n})$^{16}$O reaction, which is activated at temperatures $\\gtrsim$ 8.9$\\times$10$^{7}$ K (Herwig 2004). The other neutron source reaction, $^{22}$Ne($\\alpha$,{\\it n})$^{25}$Mg, is not a plausible candidate to provide neutrons for the F production because it is activated at higher temperatures, $\\gtrsim$ 2.8$\\times$10$^{9}$ K, at which $^{14}$N is already completely destroyed and $^{19}$F is also quickly consumed by $^{19}$F($\\alpha$,{\\it p})$^{22}$Ne reactions.  The required $^{13}$C and $^{14}$N nuclei are ingested in the convective TPs where He burning takes place, together with the ashes of the previous H-burning. Note that the abundances of the $^{13}$C and $^{14}$N nuclei in the H-burning ashes have a primary component at low metallicity due to the fact that efficient TDU carries primary $^{12}$C to the stellar envelope, which is then converted into $^{13}$C and $^{14}$N during H-burning. According to the current stellar models, in low-mass AGB stars partial mixing of the bottom of the H-rich convective envelope into outermost region of the $^{12}$C-rich intershell layer leads to synthesis of extra $^{13}$C and $^{14}$N at the end of the TDU (e.g., Werner \\& Herwig 2006). This is necessary to produce enough heavy elements via the {\\it s}-process in order to match observations of enhancements of, e.g., Zr, Ba, and Pb, in these stars (e.g., Busso et al. 2001). During the interpulse phase the $^{13}$C($\\alpha$,{\\it n})$^{16}$O reaction occurs and neutrons are released. The presence of a partial mixing zone also contributes some extra F production. However, due to the uncertainties associated with the formation of this mixing zone (e.g., Herwig 2005), its contribution to the final F production should be considered as part of the model uncertainties (e.g., Lugaro et al. 2004; Karakas et al. 2007). Jorissen et al. (1992 for MS/S stars) and Zhang \\& Liu (2005) showed a strong correlation between $^{12}$C and $^{19}$F abundances, which is expected if both $^{12}$C and $^{19}$F are carried to the stellar surface by the TDU.  In Figure \\ref{diag}(a), we show the diagram of [C/Ar] vs. [F/Ar] for PNe, MS/S stars, and HE 1305. It clearly shows that [F/Ar] increases with [C/Ar]. Figure \\ref{diag}(b) shows a correlation between [Ne/Ar] and [F/Ar], which again is to be expected if $^{22}$Ne is also carried to the surface by the TDU. Figure \\ref{diag}(c) further shows a strong correlation between [Ar/H] and [F/Ar]. This is predicted by the models because of the primary contribution to the F production due to the effect of the TDU, as discussed above.  The [F/H] abundance of BoBn 1 is comparable to NGC 40, which has a WR-type central star ([WC8]; De Marco \\& Barlow 2001). The strong stellar wind from WR stars may also be important for the F injection. In BoBn 1, strong and narrow C~{\\sc iii}, C~{\\sc iv}, and N~{\\sc iii} lines have been detected (Otsuka 2007), suggesting that the central star is a WELS (weak emission-line star) known as a class of H-deficient stars. The central star might have suffered a very late thermal pulse (VLTP) and become H-deficient. Zijlstra et al. (2006) noted that the Ne enhancement of BoBn 1 could be caused by VLTP. The strong stellar wind and VLTP might contribute to the enhancements of F, C, and Ne in the nebula.    The [F/H] abundance of BoBn 1 is compatible with the final [F/H] predicted by the AGB single star models of Karakas \\& Lattanzio (2007) and Lugaro et al. (2004) for initial masses of $\\sim$2 M$_{\\odot}$ stars with {\\it Z} = 10$^{-4}$ ([Fe/H] $\\simeq$ $-$2.3). The single metal-poor star models of Fujimoto et al. (2000) and Suda et al. (2004) predict that stars with [Fe/H] $\\lesssim$ $-$2.5 and initial masses of 1.2--3.5 M$_{\\odot}$ will become C, N, and {\\it s}-process elements rich by the partial mixing and subsequent dredge-up. Such a star, however, would not have survived in the Galactic halo up to now.  To circumvent this issue, BoBn 1 might have evolved from a binary and have experienced binary mass transfer from a massive companion. The abundances of BoBn 1 are comparable to those of HE 1305 and other CEMP stars as shown in the diagrams of [C/Fe] and [N/Fe] versus [Fe/H] (Figure \\ref{cfe}). As currently defined, CEMP stars fall into two categories, those that exhibit large enhancements of {\\it s}-process elements (CEMP-{\\it s}) and those that do not (CEMP-{\\it no}) (Aoki et al. 2007). Most CEMP-{\\it s} stars show large enhancements of C and N abundances. The evolution models for CEMP stars of Lau et al. (2007) have demonstrated that the C and N overabundances would be reproduced by binary interactions. In fact, about 70$\\%$ of CEMP stars are identified as long-period (0.3--100 yr) binaries (Beers \\& Christlieb 2005). Schuler et al. (2007) concluded that HE 1305 might have experienced mass transfer and predicted that {\\it s}-process elements should be enhanced in this object. Considering the model of Karakas \\& Lattanzio (2007), Lugaro et al. (2008) concluded that HE 1305 consists of $\\sim$2 M$_{\\odot}$ (primary) and $\\sim$0.8 M$_{\\odot}$ (secondary) stars with {\\it Z} = 10$^{-4}$ and the enhanced C and F are explained by binary mass transfer from the primary star via Roche lobe overflow and/or wind accretion. BoBn 1 might also have evolved from a binary consisting of $\\sim$2 M$_{\\odot}$ and $\\sim$0.8 M$_{\\odot}$ stars, since its C and F abundances and metallicity are comparable to HE 1305. The enhanced C and F of BoBn 1 can be explained by the model of Lugaro et al. (2008) including the extra $^{13}$C and $^{14}$N, or the upper limit of the $^{18}$F($\\alpha$,{\\it p})$^{21}$Ne reaction. The chemical similarities between BoBn 1 and CEMP-{\\it s} stars suggest that this PN shares a similar origin and evolutionary history, although we have not detected any {\\it s}-process elements in this object. The contradiction on the evolutionary time scale of this object can be avoided if BoBn 1 has indeed evolved from a CEMP-{\\it s} star such as HE 1305.  Abundances of {\\it s}-process elements and F should be enhanced in K 648, H 4-1, and BoBn 1, if they have evolved from CEMP-{\\it s} stars. For K 648 and H 4-1, there have been no reports on the detection of F and {\\it s}-process elements. High sensitivity spectroscopic observations are necessary in the optical and near-infrared to investigate those elements. We would be able to settle the origin and evolution of halo PNe, if we find enhanced abundances for those elements."
    },
    "0807/0807.2998_arXiv.txt": {
        "abstract": "We observed a sample of evolved stars in the Large and Small Magellanic Clouds (LMC and SMC) with the Infrared Spectrograph on the {\\it Spitzer Space Telescope}.  Comparing samples from the SMC, LMC, and the Galaxy reveals that the dust-production rate depends on metallicity for oxygen-rich stars, but carbon stars with similar pulsation properties produce similar quantities of dust, regardless of their initial metallicity. Other properties of the oxygen-rich stars also depend on metallicity.  As the metallicity decreases, the fraction of naked (i.e.\\ dust-free) stars increases, and among the naked stars, the strength of the 8~\\mum\\ absorption band from SiO decreases.  Our sample includes several massive stars in the LMC with long pulsation periods which produce significant amounts of dust, probably because they are young and relatively metal rich.  Little alumina dust is seen in circumstellar shells in the SMC and LMC, unlike in Galactic samples.  Three oxygen-rich sources also show emission from magnesium-rich crystalline silicates.  Many also show an emission feature at 14~\\mum.  The one S star in our sample shows a newly detected emission feature centered at 13.5~\\mum.  At lower metallicity, carbon stars with similar amounts of amorphous carbon in their shells have stronger absorption from molecular acetylene (C$_2$H$_2$) and weaker emission from SiC and MgS dust, as discovered in previous studies. ",
        "introduction": "%  In the Galaxy, stars on the asymptotic giant branch (AGB) may be responsible for nearly 90\\% of the dust injected into the interstellar medium (ISM), with supergiants accounting for most of the remainder \\citep{geh89}.  How much these relative contributions change in the early Universe is not well understood.  This question has grown even more important, now that the Infrared Spectrograph \\citep[IRS][]{hou04} on the {\\it Spitzer Space Telescope} \\citep{wer04a} has detected dust in galaxies at redshifts of 2.6--2.7 \\citep{tep07,lut07}, corresponding to look-back times of $\\sim$11 Gyr in a Universe 13.7 Gyr old.  Supernovae have been suggested as a possible source of dust in the early Universe, \\citep[e.g.][]{cmc67,ds80}, but the amount of dust they contribute remains controversial \\citep[e.g.][]{sug06,mei07}.  The most massive stars that can evolve to the AGB will begin producing dust only $\\sim$100 Myr after the formation of the first stars in a galaxy. This initial population will be metal poor.  Nearby galaxies within the Local Group give us the opportunity to study how the dust produced by individual supergiants and AGB stars depends on metallicity. The Large and Small Magellanic Clouds (LMC and SMC) lie at distances of 50 and 61 kpc, respectively \\citep[e.g.][]{alv04,kw06}.  Measurements of the metallicity ([Fe/H]) of the LMC vary between $-0.3$ and $-0.4$, while the range in the SMC is between $-0.6$ and $-0.8$ \\citep[][and references therein]{kw06}.  The {\\it Infrared Space Observatory} ({\\it ISO}) obtained the first mid-infrared spectra of evolved stars in the LMC. \\cite{tra99} published spectra of 23 targets using the PHOT-S spectrometer, which covered the 2--5 and 6--12~\\mum\\ wavelength ranges at a resolution $R\\sim$90 ($R=\\lambda/\\Delta\\lambda$), and the CAM-CVF spectrometer, which observed the 7--14~\\mum\\ wavelength range at a resolution of $\\sim$40.  The formation of CO molecules in stellar outflows will exhaust either C or O from the material available to form dust and divides the dust chemistry around these objects into oxygen-rich or carbon-rich.  The \\iso\\ sample includes spectra from both chemistries, and \\cite{tra99} showed that they could separate these two groups photometrically with color-color diagrams.  The spectrometers on \\iso\\ allowed observations of only the brightest targets.  The improved sensitivity of the IRS on {\\it Spitzer} has made possible more detailed studies of mass-loss and dust injection in both the LMC and SMC.  The five major programs have observed a combined total of 176 sources in the LMC and 63 in the SMC, most of them evolved stars.  \\cite{slo08a} reviewed these programs and the different criteria used to select targets.  Program 3426 (P.I.\\ Kastner) focused on the brightest sources in the LMC, using the classes defined from observations of the {\\it Mid-Course Space Experiment} ({\\it MSX}) by \\cite{msx01} \\citep{kas06,buc06}.  Program 3591 (P.I.\\ Kemper) targeted objects covering the stages of evolution of oxygen-rich stars from the AGB to planetary nebulae \\citep{lei08,mk08}.  Program 3277 (P.I.\\ Egan) selected targets in the SMC based on their infrared colors \\citep{kra05,kra06,slo06b}.  Program 3505 (P.I.\\ Wood) spanned a range of mass-loss rates in the LMC and the SMC. It included a young stellar object \\citep[YSO;][]{vl05b}, but nearly all of the rest are carbon stars \\citep{zij06,mat06,gro07,lag07}. The other three programs also observed more carbon stars than initially planned, with sources classified photometrically as OH/IR stars turning out to be carbon stars instead.  This paper focuses on Program 200, (P.I.\\ Sloan; hereafter referred to as the MC\\_DUST program).  The next section describes the sample and the ground-based support observations.  In \\S 3, we describe the IRS observations, the data reduction, and the classification of the spectra. The MC\\_DUST sample contains a rich variety of object types, each of which requires a different method of analysis.  We examine each of these types in turn in the following sections:  \\S 4---stars with no associated dust (naked stars), \\S 5---dusty oxygen-rich stars, \\S 6---carbon stars, \\S 7---red sources with emission from polycyclic aromatic hydrocarbons (PAHs).  Several targets show characteristics which place them in multiple categories, and as a consequence, they will be analyzed in more than one of the following sections.  In each of these sections, we examine individual sources and unusual spectral features, but we save a more general comparison of the various populations for \\S 8. Finally, we draw conclusions in \\S 9. ",
        "conclusions": "%  The similar dust production rates among the samples in Figure 29 have important consequences.  The production of amorphous carbon is limited by the amount of excess carbon compared to oxygen, but carbon is produced in the triple-$\\alpha$ process in AGB stars.  The production of silicates and other oxygen-rich grains, on the other hand, is limited by the abundances of heavy elements such as Si and Al, and these are not produced in significant quantities on the AGB.  The ready explanation for the results presented in Figure 29 is that the carbon produced and dredged up by AGB stars dominates the carbon with which they formed.  It should follow, then, that even Population III stars, if they are in the mass range to become carbon stars, will produce approximately the same amount of carbon-rich dust as their Population I counterparts.  \\cite{bbm78,bbm80} discovered that in the SMC, more AGB stars became carbon rich than in the LMC \\citep[see also][]{ch03}.  \\cite{rv81} explained this phenomenon as a metallicity effect.  Stars with solar abundances will become carbon stars if their initial mass $\\ga$1.7~M$_{\\sun}$, but at 1/5 solar metallicity, the mass limit has dropped to 1.2~M$_{\\sun}$, primarily because of the lack of oxygen in the atmosphere of the metal-poor star to balance the dredged up carbon.  Assuming a Salpeter initial mass function \\citep{sal55} and integrating from these carbon limits to 8~M${_\\sun}$, we find that in the SMC, 70\\% more of the AGB should be carbon-rich compared to a sample with solar metallicity.  Since we see no diminishment of carbon-rich dust production at decreasing metallicity, we conclude that the more metal-poor an initial population of stars, the more carbon-rich dust it will contribute to the ISM, simply because it produces more carbon stars.  Many other properties of carbon stars do depend on metallicity.  Our small sample reinforces results from previous studies of {\\it Spitzer} spectroscopy in the Magellanic Clouds, which show that as the metallicity decreases, emission from SiC and MgS dust decrease, while absorption from acetylene gas increases.  The increased acetylene absorption probably arises from higher C/O ratios in Magellanic carbon stars \\citep{vl99,mat02,mat05}, which in turn arise from a combination of increased dredge-up efficiency \\citep{woo81} and less initial oxygen in the atmosphere \\citep{vw93}. The additional acetylene we observe may add even more to the amount of carbon injected into the ISM by metal-poor populations.  The behavior of the oxygen-rich dust stands in stark contrast to the carbon-rich dust.  While the quantity of carbon-rich dust shows no discernable dependence on metallicity, oxygen-rich dust shows many clear trends.  As metallicity decreases, (1) the fraction of naked stars increases, (2) SiO absorption decreases, (3) the silicate dust emission also decreases, and (4) alumina-rich dust emission disappears almost entirely.  The lower metallicity of the Magellanic Clouds provides a ready explanation for all of these effects. Si is not produced on the AGB, so its abundance scales with overall metallicity.  As a result, we observe less SiO gas, which is the fundamental building block of silicate dust, and less silicate dust as well.  The abundance of Al scales even more strongly with metallicity (\\S 5.4.), so the lack of alumina dust should not be a surprise.  Not only do we see few spectra in the SE1--3 range, but we see little or no emission from the 13~\\mum\\ feature, which arises from an Al--O bond.  Standard condensation sequences propose the formation of alumina-rich grains before silicates \\citep[e.g.][]{tie90,ste90}.  The lack of alumina dust in our sample suggests the formation of alumina is not a prerequisite to the formation of silicates.  In other words, silicates can form without pre-existing alumina grains to serve as seeds.  The oxygen-rich sample contains a wealth of interesting dust features.  Five sources show crystalline silicate features at 10~\\mum\\ or longer wavelengths, and these are all consistent with Mg-rich pyroxene and olivine.  There is no evidence for Fe in the silicates, much like the Galactic sample.  Several sources have an unidentified 14~\\mum\\ emission feature, and one source, the only true S star in our sample, shows a previously unknown 13.5~\\mum\\ emission feature.  These two features are newly discovered by the IRS, and their identification awaits further investigation.  The Magellanic Clouds present us with a scenario of steady or increasing carbon-rich dust production as systems grow more metal-poor, compared to declining rates of oxygen-rich dust production.  This result measures what goes into the ISM, but it contradicts observations of what is currently present in the ISM in metal-poor systems.  Most lines of sight in the SMC have unusual extinction curves without the commonly observed 2175~\\AA\\ bump seen in the Galaxy, in common with starburst galaxies and quasars \\citep[e.g.][and references therein]{gc98,pit00}.  The 2175~\\AA\\ bump is commonly associated with graphitic carbonaceous material \\citep[as first proposed by][]{sd65}.  \\cite{ds98} showed that certain PAHs and PAH clusters can reproduce the 2175~\\AA\\ bump and that these particles can result from interstellar processing of hydrogenated amorphous carbon.  The recent study of unusual PAH-like spectra by \\cite{slo07} demonstrates that hydrocarbons can exist with a range of aromatic-to-aliphatic ratios.  Many of the sources of unusual PAH-like emission are carbon-rich post-AGB objects.  All of these observations support an evolutionary relationship from amorphous carbon produced by carbon stars on the AGB and to interstellar PAHs to the more graphitic material responsible for the 2175~\\AA\\ bump.  The results in this paper demonstrate that carbon-rich dust production does not decrease in more metal-poor samples, and yet, metal-poor systems like the SMC are missing the 2175~\\AA\\ bump that should result.  Either the missing carbon-rich dust did not form, or it has been destroyed.  \\cite{gc98} suggested that harsher radiation fields in metal-poor systems can destroy the carrier of the 2175~\\AA\\ bump.  Stellar modelling shows that decreasing the metallicity of a star results in less UV line blanketing and an emitted spectrum with more UV emission \\citep{giv02}, which readily explains the harsher interstellar radiation observed in metal-poor \\ion{H}{2} regions \\citep{mh02}.  \\cite{gal08} recently reached the opposite conclusion.  They used galactic evolution models to investigate the absence of PAH emission in blue compact dwarf galaxies and other metal-poor galaxies \\citep{eng05,wu06}, and concluded that these systems produce significantly less carbon-rich dust than the Galaxy.  Our sample of Magellanic carbon stars provides no evidence for less carbon-rich dust formation in more primordial systems, which leads to the conclusion that the lack of carbon-bearing dust in these systems results not from reduced production, but from enhanced destruction."
    },
    "0807/0807.0954_arXiv.txt": {
        "abstract": "{ The inflaton field responsible for inflation may not be a canonical fundamental scalar. It is possible that the inflaton is a composite of fermions or it may have a decay width. In these cases the standard procedure for calculating the power spectrum is not applicable and a new formalism needs to be developed to determine the effect of short range interactions of the inflaton on the power spectrum and the CMB anisotropy. We develop a general formalism for computing the power spectrum of curvature perturbations for such non-canonical cases by using the flat space K\\\"{a}ll\\'{e}n-Lehmann spectral function in curved quasi-de Sitter space assuming implicitly that the Bunch-Davis boundary conditions enforces the inflaton mode functions to be plane wave in the short wavelength limit and a complete set of mode functions exists in quasi-de Sitter space. It is observed that the inflaton with a decay width suppresses the power at large scale while a composite inflaton's power spectrum oscillates at large scales. These observations may be vindicated in the WMAP data and confirmed by future observations with PLANCK.} ",
        "introduction": "In the generic inflation model \\cite{guth}, inflation is caused by a slow roll of the inflaton scalar field and the perturbations of the inflaton field give rise to density perturbations \\cite{guth2} and CMB anisotropies observed at cosmological scales.  The two-point correlation function of the inflaton perturbation during inflation or the power spectrum of this two-point correlation in momentum space determines the CMB anisotropy of the universe at last scattering which we observe today. The generic slow roll model of inflation are characterized by the inflaton potential and its derivatives and a large variety of particle physics potential have been studied \\cite{lyth-riotto}. Inflation may be caused by more than one scalar field and these multifield models have interesting consequences in the CMB anisotropy like isocurvature perturbation \\cite{gordon} or large non-gaussianity in curvaton models \\cite{curvaton}. In models of inflation with elementary scalar fields, the perturbations of the inflaton obey the Klein-Gordon equation in the quasi de Sitter space \\cite{riotto}, whose solutions are used for calculation of the two point correlation and the curvature power spectrum.  However, it may be possible that the inflaton field is a composite of fermions and we can ask if the compositeness changes the perturbation spectrum which can be observed in the CMB anisotropy.  Similarly if the inflaton is unstable with a decay width $\\Gamma \\sim H/N$ (such that the inflaton decays after N e-foldings of inflation are over) then again we can ask if the decay of the inflaton is reflected in the power spectrum and CMB anisotropy. For such situations the standard methods of calculating the power spectrum do not work as not all forms of the short range structure of the scalar field are reflected in the inflaton potential. If the length scale of the scalar perturbation is of the same order as the compositeness scale then the effective theory description of the scalar potential breaks down. Similarly if the inflaton is a resonance with a lifetime of the same order as the duration of inflation $\\tau \\sim N/H$ i.e a width $\\Gamma \\sim H/N$ then there are corrections to the two point correlation that are not reflected in the inflaton potential.  In this paper we find a general method for computing the power spectrum of inflaton perturbations if the inflaton has non-trivial interactions like a decay width or if the inflaton is not an elementary scalar but a composite of fermions. \\emph{We show in general that the two-point correlations of the interacting field can be written in terms of the two-point function of the free field (in the de Sitter background) by use of the K\\\"{a}ll\\'{e}n-Lehmann spectral function \\cite{Kallen,lehman} if the short wavelength limit of the mode functions are plane wave states $\\frac{1}{\\sqrt{2k}}e^{-ik\\eta}$, which in the quasi de-Sitter space is enforced by the assumption of the Bunch-Davis boundary conditions and there exists a complete orthonormal set of mode functions of the free theory in curved spacetime which is true in the quasi de-Sitter space relevant for inflation power spectrum calculation.} The two point correlation of an interacting theory can be written as a convolution of the free field correlation function with a spectral function $\\rho(\\sigma^2)$ \\begin{equation} G^{\\left({\\rm int}\\right)}(p)=\\int_0^{\\infty} d\\sigma^2 \\rho(\\sigma^2) G^0(p,\\sigma^2), \\end{equation} where $G^{\\left({\\rm int}\\right)}(p)$ is the two point correlation of the interacting theory and $G^0(p,\\sigma^2)$ is the free-field correlation with mass parameter $\\sigma$. The K\\\"{a}ll\\'{e}n-Lehmann representation holds for all two point correlations like the Feynman propagator $\\Delta(p,\\sigma^2)$ or the equal time Wightman function $W_{\\rm ET}(x-y)$. As mentioned above we show that this result can be generalized to the curved space if we assume that a complete orthogonal basis set of states of the interacting theory exists in curved spacetime.  The power spectrum of the inflaton perturbation is related to the equal-time Wightman function in the de-Sitter space as \\begin{eqnarray} W^{\\rm dS}_{\\rm ET}(x)= \\langle0|(\\delta\\phi({\\mathbf x},t))^2|0\\rangle=\\int\\frac{dk}{k}{\\mathcal P}_{\\delta\\phi}(k). \\label{var1} \\end{eqnarray} The Bunch-Davies boundary condition is that the inflaton perturbations, in the limit where the momentum $k$ is large compared to the inflaton horizon (for a spatially flat de-Sitter space), tend to the free field form $\\delta \\phi(k,\\eta)=\\frac{1}{\\sqrt{2 k}} e^{-i k \\eta}$ (where $\\eta$ is the conformal time). Assuming the Bunch-Davies boundary conditions, if we have short range interactions which dominate at scales smaller than the inflation horizon, we may be justified in using the flat space form of the spectral function in Eq.~(\\ref{var1}) to compute the two point correlation function for interacting theory. Therefore power spectrum of the interacting scalar field can be expressed as \\begin{equation} P^{\\rm \\left(int\\right)}(k)=\\int_0^{\\infty}P^{\\left(0\\right)}(k,\\sigma^2)\\rho\\left(\\sigma^2\\right)d\\sigma^2, \\label{power1} \\end{equation} where $P^{\\left(0\\right)}(k,\\sigma^2)$ is the power spectrum of the free scalar field with a mass parameter $\\sigma$ and $\\rho(\\sigma^2)$ is the KL spectral function which encapsulates all the short distance interactions (like compositeness or resonance) of the scalar field.  In Sec.~(\\ref{power-interacting}) we derive the relation between the power spectrum of interacting theory and the free field theory given in Eq.~(\\ref{power1}). In Sec.~(\\ref{case:decay}) we apply this result to calculate the power spectrum for case of a decaying inflaton field. We find that the TT angular spectrum of CMB anisotropy is suppressed at low $l$. In Sec.~(\\ref{case:composite}) we derive the power spectrum of the composite inflaton field. We find that the power spectrum of the composite field has a resonance which gives rise to oscillatory features in the TT angular spectrum. ",
        "conclusions": "We have derived a general formula for incorporating short-range interactions in the two-point correlation functions and the power spectrum by use of the K\\\"{a}ll\\'{e}n-Lehmann spectral function of flat spacetime. This method is useful if the short wavelength limit of the mode functions are plane wave states $\\frac{1}{\\sqrt{2k}}e^{-ik\\eta}$, follows in the quasi de-Sitter inflation by the assumption of the Bunch-Davis boundary conditions and there exists a complete orthonormal set of mode functions of the free theory in curved spacetime which is true in the quasi de-Sitter space relevant for inflation power spectrum calculation. In interacting inflaton models like the ones studied in this paper we find that there are more interesting variations in the power spectrum due to the modification of the propagators which do not affect the slow roll parameters. We apply our formulation to study inflation with decaying and composite inflatons. We find that the decay of the inflaton results in the suppression of long distance correlations and thereby a loss of the quadrupole anisotropy \\cite{decay-ref}.  This may be related to the observation of low quadrupole power by WMAP \\cite{wmap}.  When the inflaton is taken as a composite of two fermions the power spectrum displays even more interesting features like oscillations. An examination of the WMAP data by wavelet analysis and by the cosmic inversion method reveals that the data may have such features \\cite{oscillations}."
    },
    "0807/0807.2483_arXiv.txt": {
        "abstract": "We probe the geometry of magnetospheric accretion in classical T Tauri stars by modeling red absorption at \\helium\\ via scattering of the stellar and veiling continua.  Under the assumptions that the accretion flow is an azimuthally symmetric dipole and helium is sufficiently optically thick that all incident 1-$\\micron$ radiation is scattered, we illustrate the sensitivity of \\helium\\ red absorption to both the size of the magnetosphere and the filling factor of the hot accretion shock. We compare model profiles to those observed in 21 CTTS with subcontinuum redshifted absorption at \\helium\\ and find that about half of the stars have red absorptions and 1-$\\micron$ veilings that are consistent with dipole flows of moderate width with accretion shock filling factors matching the size of the magnetospheric footpoints. However, the remaining 50\\% of the profiles, with a combination of broad, deep absorption and low 1-$\\micron$ veiling, require very wide flows where magnetic footpoints are distributed over 10-20\\% of the stellar surface but accretion shock filling factors are $<1\\%$. We model these profiles by invoking large magnetospheres dilutely filled with accreting gas, leaving the disk over a range of radii in many narrow ``streamlets'' that fill only a small fraction of the entire infall region. In some cases accreting streamlets need to originate in the disk between several $R_*$ and at least the corotation radius. A few stars have such deep absorption at velocities $>0.5~V_{\\rm esc}$ that flows near the star with less curvature than a dipole trajectory seem to be required. ",
        "introduction": "The classical T Tauri stars (CTTS) are optically revealed, low-mass, pre--main-sequence stars that accrete material from a circumstellar disk and have a well-defined connection between accretion and outflow \\citep[HEG hereafter]{har95}.  Accretion from the disk to the star is thought to be guided by the stellar magnetosphere, where  a sufficiently strong magnetic field truncates the disk at several stellar radii and material follows  field lines that direct it to the stellar surface at high latitudes \\citep{gho78,kon91,col93,shu94}.  Magnetospheric accretion controls the star-disk interaction, and an improved understanding of this process will shed light on outstanding issues in the innermost 10~$R_*$ of CTTS systems, such as the regulation of stellar angular momentum and the launching of the inner wind.  Leading diagnostics of CTTS accretion include the optical/UV continuum excess and the profiles of permitted emission lines \\citep{bou07a}.  Kinematic evidence for infalling gas in CTTS began with the discovery that some CTTS show inverse P Cygni structure in upper Balmer lines extending to velocities of several hundred \\kms\\ \\citep{wal72}.  Later, more sensitive surveys found that redshifted absorption components are relatively common in some lines, especially in the upper Balmer and Paschen series \\citep{edw94,ale00,fol01}.  Although redshifted absorption extending to several hundred \\kms\\ clarifies that material accretes in free-fall from at least several stellar radii, it has been the success of radiative transfer modeling of line formation in magnetospheric accretion flows in a key series of papers culminating with \\citet{muz01} that has provided the strongest underpinning for this phenomenon.  Under the assumption of an aligned, axisymmetric dipole, the models have had reasonably good success in reproducing the general morphology of hydrogen profiles and emission fluxes in some stars.  The complementary assessment of accretion rates follows from interpreting the SED of the optical/UV excess, which has been successfully modeled for wavelengths shortward of 0.5~$\\micron$ as arising in a hot accretion shock, where accreting material impacts the stellar surface after free-fall along funnel flows coupled to the disk \\citep[CG hereafter]{joh00,cal98}. To match the observations, the accretion shock filling factor is less than 1\\% in most cases but can climb to 10\\% in a few of the most active accretors.  The derived accretion rates range from $10^{-10}$ to $10^{-6}$ \\msunyr\\ with a median of $10^{-8}$~\\msunyr\\ \\citep{cal00}.  Additionally, Zeeman broadening of unpolarized CTTS photospheric lines indicates mean surface field strengths in the range 1-3~kG \\citep{joh07}, sufficiently strong to induce disk truncation and drive funnel flows.  However, these strong surface fields are not predominantly dipolar, as photospheric lines show  only weak net circular polarization implying dipole components an order of magnitude smaller \\citep{val04,yan07}. Nevertheless, an extended dipole component is inferred for the accretion flow, since the same authors find significant circular polarization in the narrow component of the \\heopt\\ emission line, thought to be formed in the accretion shock at the base of the funnel flow \\citep{ber01}.  The evidence for magnetospheric accretion is thus compelling; however, the topology of the magnetosphere, the geometry of the accretion flow, and the disk truncation radius remain topics of considerable investigation, since their configuration impacts processes for angular momentum regulation and wind launching.  One form of angular momentum regulation, known as disk locking, invokes a spin-up torque from accreting material just inside the corotation radius balanced by a spin-down torque at larger radii \\citep{col93}. This approach has been questioned by \\citet{mat05,mat08a,mat08b}, who instead suggest accretion-powered stellar winds as a more likely means for stellar spindown, which must occur simultaneously with magnetospheric accretion from the disk (see also \\citealt{sau94}).  Alternatively the X-wind model \\citep{shu94} originally featured a narrow annulus of star-disk coupling close to the corotation radius, where closed field lines develop funnel flows and open field lines drive a centrifugal outflow that carries away angular momentum from accreting material, thus inhibiting stellar spin-up.  The flexibility of this model to maintain its basic properties in the face of complex magnetospheric accretion geometries has recently been demonstrated by \\citet{moh08}. If the stellar and disk fields are parallel, then an intermittent outflow can develop via a Reconnection X-wind, which removes angular momentum from the star as well as the inner disk \\citep{fer06}.  Evidence for non-aligned fields coupled with complex accretion geometries is mounting, coming from a variety of recent studies.  Using time-resolved spectropolarimetry of the mildly accreting CTTS V2129 Oph, \\citet{don07} used Zeeman detections from both photospheric features and emission lines from the accretion shock to construct a Doppler tomographic map of the magnetic topology on the stellar surface.  The dominant field on the star is a misaligned octupole, and accretion is confined largely to a high-latitude spot, covering $\\le5$\\% of the stellar surface. These authors also attempt to reconstruct the 3D field geometry out to the disk interaction region and suggest that the large-scale field funneling accreting material from the disk is more complex than a simple dipole. However, such extrapolation techniques, while tantalizing, must be applied with caution at present, since there are numerous uncertainties in the reconstruction process \\citep{moh08}. Numerical magnetohydrodynamic (MHD) simulations of star-disk interactions demonstrate that non-axisymmetric funnel flows arise if the stellar field is tipped by only a few degrees relative to the rotation axis, breaking into two discrete streams under stable accretion conditions \\citep{rom03}. If the accretion rate is sufficiently high, accretion is predicted to proceed through equatorial ``tongues'' that can push apart field lines \\citep{rom08}. Similarly, accretion spots can appear at a wide range of latitudes, including equatorial belts.  Such behavior has been observed in models featuring accretion along quadrupolar as well as dipolar field lines \\citep{lon07,lon08}, accretion along field lines extrapolated from surface magnetograms \\citep{gre06}, and accretion mediated by a dynamo-generated disk magnetic field \\citep{von06}.  Observational signatures for misaligned dipoles are being explored in radiative transfer models for hydrogen line formation in funnel flows, with subsequent predictions for rotationally modulated profile variations. Initially \\citet{sym05} presented radiative transfer models of hydrogen lines featuring curtains of accretion covering a limited extent in azimuth in geometries consistent with aligned dipoles.  Their model profiles exhibit certain characteristics of the observed line profile variability, such as rotationally modulated line strengths and the appearance of red absorption components at certain phases and inclinations, but the predicted level of variability is higher than observed.  More recently, \\citet{kur08} applied a radiative transfer code for H line formation to the 3-D output from the MHD simulations of \\citet{rom03,rom04}, which prescribe the geometry, density, and velocity of two-armed accretion streams that result from dipoles misaligned with the rotation axis by angles ranging from 10$^\\circ$ to 90$^\\circ$.  Applying temperatures similar to those from the \\citet{muz01} axisymmetric models, they were able to reproduce some of the trends in continuum and profile variability from models. One of the larger discrepancies in comparing the model profiles to observed ones is that the model profiles for Paschen and Brackett lines are a factor of two narrower than the mean value observed by \\citet{fol01}. The problem is likely more complex than simply finding another line broadening mechanism, since Paschen~$\\gamma$ has recently been shown to have line widths that are correlated with the 1-$\\micron$ continuum excess, in the sense that the narrowest lines are found among objects with the lowest disk accretion rates \\citep[EFHK hereafter]{EFHK}. Evidence that some of the hydrogen emission is formed in the accretion shock rather than the funnel flow is now clearly demonstrated from the discovery of circular polarization in the core of Balmer lines \\citep{don08}. The implication is that hydrogren lines are not necessarily a definitive means for probing the properties of funnel flows, so additional probes are desirable.  In this paper, we explore a different means of diagnosing the geometry of the accretion flow, making use of the redshifted subcontinuum absorption in the \\heupper\\ $\\rightarrow$ \\helower\\ transition of neutral helium ($\\lambda10830$), recently demonstrated to be a very sensitive probe of both outflowing gas in the inner wind and infalling gas in the funnel flow due to its frequent display of blue and red absorptions (EFHK).  A subcontinuum absorption feature is a more tell-tale diagnostic of a kinematic flow than an emission profile, since its position (blue or red) indicates the direction of the flow, its width indicates the range of line-of-sight velocities in the flow, and its depth at a particular velocity indicates the fraction of the continuum (stellar plus veiling) occulted by material moving at that velocity.  In the case of a red absorption, the absorption depth signals the fraction of the stellar surface covered by the funnel flow at each velocity, making it an effective probe of the CTTS accretion geometry.  This is the third in a series of papers about 1-$\\micron$ diagnostics of accretion and outflow in CTTS.  The first, EFHK, presented 1-$\\micron$ spectra from 38 CTTS including \\helium\\ profiles, \\pgamma\\ profiles, and measurements of the continuum excess ``veiling''  in the 1-$\\micron$ region.  The second paper, \\citet[KEF hereafter]{kwa07}, modeled blueshifted absorption components at \\helium, which appear in about three quarters of the sample, and found that while some stars have winds best explained as arising from the inner disk, others require an outflow moving radially away from the star in an accretion-powered ``stellar'' wind.  In this paper, we analyze redshifted subcontinuum absorption at \\helium\\ in 21 CTTS that present red absorption in at least one observation.  The following section describes the sample selection, data acquisition, and data reduction.  Section 3 presents the data and discusses variability.  In Section 4, we present model scattering profiles that arise in a dipolar flow geometry, show that they explain only a fraction of the red absorptions, and explore modifications to a dipolar flow that better explain the remaining observations.  Discussion and conclusions follow in Sections 5 and 6. ",
        "conclusions": "We have probed the geometry of magnetospheric accretion in classical T Tauri stars by modeling red absorption at \\helium\\ via scattering of the stellar and veiling continua.  Between 2001 and 2007, we acquired 81 1-$\\micron$ spectra of 38 CTTS spanning the full observed range of mass accretion rates.  Of the 38 stars, 1 of 9 with $r_Y>0.5$ and 20 of 29 with $r_Y\\le0.5$ show red absorption at \\helium\\ that extends below the 1-$\\micron$ continuum in one or more spectra, demonstrating that red absorption from magnetospheric accretion is rare in objects with high veiling but is found in about two thirds of objects with moderate to low veiling. The red absorptions can be strong, deep, and broad, with equivalent widths up to 4.5~\\AA, maximum penetrations into the 1-$\\micron$ continuum up to 61\\%, and widths at one quarter of the absorption minimum up to 320~\\kms; furthermore, they tend to be strongest in stars with the lowest veilings.  We model the red absorption by assuming that an axisymmetric dipolar accretion flow scatters photons from the star and from hot zones in the accretion-heated photosphere that produce the 1-$\\micron$ veiling and have filling factor $f$.  Testing a range of magnetosphere widths and $f$ consistent with shock filling factors from the literature, we find that about half of the absorption profiles can be explained by dipolar flows in which the size of the flow is consistent with the size of the shock filling factor $f$. Weak absorptions in stars with weak veiling and intermediate absorptions in stars with intermediate veiling are explained by such flows, but strong absorptions in stars with little to no veiling are not.  We introduce the concept of dilution as a means of producing a strong red absorption while keeping the filling factor and thus the veiling low.  In a diluted flow, the magnetosphere can extend over a wide range of radii, with a large covering factor on the stellar surface, but this volume is incompletely filled by accreting gas.  Instead of a single thick flow, we posit  multiple nested streamlets with a total filling factor  small enough for a low veiling, but each with an intrinsic thermal or turbulent width sufficient to scatter photons as though the entire volume were filled, thereby yielding a large red absorption.  The multiple streamlets can also explain how helium is ionized through the entire flow, rather than just the skin of a thick flow.  Large, dilutely filled accretion flows are necessary for about half of the objects, some of which require accreting streamlets to connect to the disk over a range from 2~$R_*$ out to or beyond corotation. A few stars show such deep absorption at redward velocities exceeding 50\\% of the stellar escape velocity that flows near the star with less curvature than a dipolar trajectory seem to be required.  The frequency of \\helium\\ red absorption is also informative. Our limited temporal coverage suggests that the frequency of helium absorption differs in stars with high and low veiling.  Red absorption at \\helium\\ is far more common in stars with low veiling.  When it is absent from these stars, it is sometimes because helium emission from another source such as a wind fills it in and sometimes because of  inhomogeneous azimuthal coverage of accreting magnetic columns.  Among stars with high veiling ($r_Y\\ge0.5$), red absorption at \\helium\\ is rarely seen. If these stars had accretion geometries similar to those of the low-veiling stars, they would be expected to have extremely strong red absorptions.  Even if the absorption were filled in by emission from the accretion flow, the stars would still be expected to show red absorption at high velocities. In the high-veiling stars, the paucity of \\helium\\ red absorption, the presence of \\helium\\ emission and blue absorption that suggest formation in accretion-powered stellar winds, and the weakness or absence of narrow-component \\heopt\\ emission from an accretion shock lead us to suggest that the magnetospheric accretion structure may be crunched or otherwise reduced in CTTS with the highest disk accretion rates.  We find the study of \\helium\\ red absorption due to infalling gas projected in front of the star to be complementary to studies of emission lines modeled as arising over the full size of the accretion flow. The proximity of \\helium\\ and P$\\gamma$ offer an excellent pair of lines for deeper investigation of magnetospheric geometries through intensive time-monitoring programs that can track non-aziumuthal structures as stars rotate. Our limited phase coverage of AA Tau demonstrates that this approach will be very effective, particularly when coupled with radiative transfer models that can constrain formation conditions for both lines simultaneously."
    },
    "0807/0807.3429_arXiv.txt": {
        "abstract": "We calculate intensity and angular power spectrum of the cosmological background of synchrotron emission from cold dark matter annihilations into electron positron pairs. We compare this background with intensity and anisotropy of astrophysical and cosmological radio backgrounds, such as from normal galaxies, radio-galaxies, galaxy cluster accretion shocks, the cosmic microwave background and with Galactic foregrounds. Under modest assumptions for the dark matter clustering we find that around 2 GHz average intensity and fluctuations of the radio background at sub-degree scales allows to probe dark matter masses $\\ga100\\,$GeV and annihilation cross sections not far from the natural values $\\left\\langle\\sigma v\\right\\rangle\\sim3\\times10^{-26}\\,{\\rm cm}^3\\,{\\rm s}^{-1}$ required to reproduce the correct relic density of thermal dark matter. The angular power spectrum of the signal from dark matter annihilation tends to be flatter than that from astrophysical radio backgrounds. Furthermore, radio source counts have comparable constraining power. Such signatures are interesting especially for future radio detectors such as SKA. ",
        "introduction": "Whereas it is known from cosmological observations that cold dark matter represents a fraction $\\Omega_m\\simeq0.233$ of the total present energy density of our Universe~\\cite{Komatsu:2008hk,Dunkley:2008ie}, its nature is still elusive. Dark matter can not only be detected directly in dedicated experiments searching for nuclear recoils from the scattering of dark matter particles, or produced in particle accelerators such as the LHC, but can also reveal its existence indirectly~\\cite{Bertone:2004pz}: Although, apart from dilution from cosmic expansion, the density of dark matter does not change significantly after self-annihilations freeze out in the early Universe, residual self-annihilation can give rise to significant fluxes of $\\gamma-$rays, electrons, positrons, neutrinos, and even some antimatter such as anti-protons and positrons, especially in regions with large dark matter densities. The energies of the secondary particles can reach up to the dark matter particle mass which can be of order a few hundred GeV. Secondary electrons and positrons can annihilate and give rise to a 511 keV line emission, and they emit synchrotron radiation in the magnetic fields of galaxies which can be detected in the radio band. Therefore, cosmic and $\\gamma-$ray detectors, neutrino telescopes, and even radio telescopes can be used for indirect dark matter detection as well.  One of the most promising dark matter candidates are weakly interacting massive particles (WIMP) such as they are predicted within supersymmetric extensions of the Standard Model. Such particles have masses $m_X\\ga100\\,$GeV. If they are produced thermally, in order to reproduce the correct average dark matter density, their annihilation cross sections have to be $\\left\\langle\\sigma v\\right\\rangle\\sim3\\times10^{-26}\\,{\\rm cm}^3/$s whereas large values are possible in case of non-thermal production.  Traditionally, indirect dark matter detection has focused on high energy emission, specifically signatures in $\\gamma-$rays~\\cite{Bergstrom:2001jj,Ullio:2002pj,Taylor:2002zd,Elsaesser:2004ap,Ando:2005xg,Ando:2006cr,Hooper:2007gi,Cuoco:2007sh,Ando:2006mt,Mack:2008wu}. High energy emission of neutrinos, the particle the most difficult to detect, can be used to establish conservative constraints on the total annihilation cross section, which are of the order $\\sim10^{-23}\\,{\\rm cm}^3/$s~\\cite{Yuksel:2007ac}. The fluxes of antiprotons and positrons from galactic dark matter annihilations has also been extensively used to constrain dark matter properties~\\cite{Bergstrom:2006tk,Bringmann:2006im,Delahaye:2007fr}. Furthermore, the synchrotron radiation emitted by dark matter annihilation products close to the Galactic centre has also been studied. It has been found that if the dark matter profile close to the central black hole is a spike formed by adiabatic accretion, typical dark matter annihilation cross sections within supersymmetric scenarios can lead to intensities comparable to the radio emission observed from the Galactic centre~\\cite{Bertone:2001jv,Bertone:2002je}. It has further been shown that measurements of the radio flux away from the Galactic centre by the WMAP experiment strongly constrain the annihilation cross sections to values $\\la10^{-25}\\,{\\rm cm}^3/$s for $m_X\\simeq100\\,$GeV~\\cite{Hooper:2008zg,Grajek:2008jb}. A multi-wavelength analysis of dark matter annihilations from the Galactic centre has recently been performed in Ref.~\\cite{Regis:2008ij}.  In the present paper we evaluate the diffuse synchrotron emission from the electrons and positrons produced by dark matter annihilation in the cosmological distribution of dark matter halos. We compute both its overall intensity and its angular power spectrum as well as the distribution of visible dark matter annihilation sources as a function of apparent luminosity. We will find that comparing the resulting signals with other backgrounds and foregrounds under conservative assumptions allows to test annihilation cross sections close to the natural scale $\\left\\langle\\sigma v\\right\\rangle\\sim3\\times10^{-26}\\,{\\rm cm}^3/$s.  In Sect.~2 we provide the general setup of our calculations. In Sect.~3 and~4, we apply it to astrophysical backgrounds and the dark matter induced signal, respectively. In Sect.~5 we compare the overall diffuse signal and its anisotropy to other foregrounds and discuss the resulting dark matter constraints, and in Sect.~6 we conclude. Finally, an appendix presents technical details of the calculations. We will use natural units in which $c=1$ throughout. ",
        "conclusions": "Many different indirect detection signatures have been investigated to constrain the parameter space for dark matter~\\cite{Ando:2005xg,Ando:2006mt,Chen:2003gz,Zhang:2006fr,Zhang:2007zzh,Myers:2007fj,Padmanabhan:2005es}. In the present paper we have calculated intensity and angular power spectrum of the cosmological background of synchrotron emission from the electrons and positrons produced in annihilations of cold dark matter. The resulting radio background around $\\simeq2\\,$GHz and its angular power spectrum for multipoles $200\\la l\\la3000$ has comparable or better sensitivity to dark matter annihilation cross sections than other signatures. Furthermore, a comparison of observed radio source counts with predictions for dark matter annihilation results in the constraint $F_{\\rm dm}\\la10$ for the parameter defined in Eq.~(\\ref{eq:F_dm}). Under reasonable assumptions on dark matter clustering and magnetic fields in the halo environment, the range of annihilation cross sections corresponding to the constraint $F_{\\rm dm}\\la10$ is comparable to constraints from synchrotron emission in an NFW profile~\\cite{Hooper:2007gi,Grajek:2008jb}. Galactic $\\gamma-$ray constraints derived under similar assumptions~\\cite{Dodelson:2007gd} are also comparable. The sensitivity of our signal is considerably better than conservative limits based on annihilation into neutrinos~\\cite{Yuksel:2007ac}, and comparable to limits on annihilation into $\\gamma-$rays from diffuse cosmological emission~\\cite{Mack:2008wu}.  Sensitivities to values of order ten for the parameter $F_{dm}$ defined in Eq.~(\\ref{eq:F_dm}) are interesting for non-thermal dark matter whose annihilation cross sections can be larger than our fiducial value $\\left\\langle\\sigma v\\right\\rangle=3\\times10^{-26}\\,{\\rm cm}^3\\,{\\rm s}^{-1}$, the cross section required for thermal dark matter. An example for a non-thermal dark matter candidate with large cross section is the wino LSP occurring in supersymmetric theories with anomaly mediation~\\cite{Randall:1998uk,Giudice:1998xp}. We believe that radio observations in particular with future instruments such as SKA can provide valuable information on dark matter."
    },
    "0807/0807.4209_arXiv.txt": {
        "abstract": "One way of revealing the nature of the coronal heating mechanism is by comparing simple theoretical one dimensional hydrostatic loop models with observations at the temperature and/or density structure along these features. The most well-known method for dealing with comparisons like that is the $\\chi^2$ approach. In this paper we consider the restrictions imposed by this approach and present an alternative way for making model comparisons using Bayesian statistics. In order to quantify our beliefs we use Bayes factors and information criteria such as AIC and BIC. Three simulated datasets are analyzed in order to validate the procedure and assess the effects of varying error bar size. Another two datasets \\cite{ugarte2005, priest2000} are re-analyzed using the method described above. In one of these two datasets \\cite{ugarte2005}, due to the error estimates in the observed temperature values, it is not posible to distinguish between the different heating mechanisms. For this we suggest that both Classical and Bayesian statistics should be applied in order to make safe assumptions about the nature of the coronal heating mechanisms. ",
        "introduction": "\\label{intro} Magnetically confined plasma loops are the fundamental building blocks of the solar atmosphere. Whole loop structures are observed over an extensive spectral range while extending over a large range of length-scales and dynamical time-scales. In particular, the Solar and Heliospheric Observatory (SOHO), the Transition Region and Coronal Explorer (TRACE) and now Hinode record loop-like features from small-scale brightenings lasting for tens of seconds to large-scale (of the order of the solar radius), apparently static loops that last for many hours.  Recent interest in loops has centred on the definitive determination via remote sensing of the basic parameter values within these features. One dimensional hydrostatic simulations of loop plasma ({\\it e.g.} \\opencite{per00}) produce a temperature ($T$) and density ($\\rho$) structure along the loop that results from a balance between thermal conduction along the field-lines, optically thin radiative loss and a predefined coronal heating term. Generally, this profile extends from a cooler temperature at the loop base (or footpoint) area up to a hotter temperature at the loop apex. However, it has been demonstrated that the temperature gradient ($\\frac{dT}{ds}$ where $s$ is the distance along the loop) at each point along this profile is very dependent on where the energy deposition preferentially occurs ({\\it e.g.} \\opencite{priest2000}). If the heat input is predominantly located at the base, the temperature in the ``coronal'' part of the loop will be relatively constant ({\\it i.e.} $\\frac{dT}{ds} \\approx 0$). In comparison, if the energy is released near the loop apex, a significant temperature gradient ($\\frac{dT}{ds} > 0$) travelling away from the apex can result. By observing the local temperature (and/or density) at successive spatial locations along a well-defined loop and then comparing the resulting profile with one generated from a one-dimensional hydrostatic model, then this could provide a means of constraining the possible preferred spatial location of the heating within that loop structure.  Confirming the dominant energy deposition position has been more difficult than first imagined, including for example the dataset introduced by \\inlinecite{priest2000} (PDS from now on) which has been interpreted by separate authors as uniform \\cite{priest2000}, base \\cite{asch01} and apex \\cite{reale02,mackay2000} heating. Other datasets examine the variation in density rather than temperature (\\opencite{ugarte2005}; UUDS from now on). In many of the above investigations, the important step of model comparison between the hydrostatic simulation and the observed plasma parameter values is undertaken by employing a weighted chi-squared analysis. Possibly this is an adequate, ``quick-look'' approach to tackling the model comparison. However two statistical ``obstacles'' present themselves here: \\begin{enumerate} \\item The precision of the temperature observations may not be sufficient to discriminate one heating function from another. As we shall see, substantive changes in the nature of the heating function can result in only subtle changes in the temperature profile of the loop. \\item The current approach (as in PDS and UUDS) for comparing one heating function model with another is to minimize the well-known statistic: \\begin{equation} \\chi^2 = \\sum_{i=1}^n \\frac {\\left(T_i - \\widehat{T}_i\\right)^2} {\\sigma_i^2} \\label{chi} \\end{equation} where $n$ is the number of grid points, $T_i$ is the observed temperature at distance $s_i$ along the strand, $\\widehat{T}_i$ is the theoretical prediction of the model concerning the temperature and $\\sigma_i$ is the standard deviation of the observed temperature. The procedure, applied to the solar coronal loop heating problem, would be to isolate the correct heating functional form ({\\it e.g.} base or apex heating with an exponential profile) based on its capability to furnish the minimum statistic value. \\end{enumerate}  There are a number of deficiencies to this procedure. First of all, for any heating functional form, there is a continuous range of parameter values, with an infinite set of possible parameter combinations, each of which is a candidate model. Therefore, it is not possible to properly compare one family's performance in fitting the data with another ({\\it e.g.} apex heating with basal heating using an exponential form), without resorting to selecting specific values of the parameters. Such an approach was employed by UUDS, where a grid method of equally spaced parameter combinations was used. Of course, the concern is that somewhere we have missed certain parameter values between grid points which could reverse the conclusion reached in the model comparison assessment.  Furthermore, using a procedure such as the one described above, there is no straightforward way of telling if we have significant evidence of one heating function being superior to another. The issue here is that the minimum $\\chi^2$ statistic approach is not well suited to model comparison problems, since it is primarily a goodness-of-fit statistic.  Another point of interest is that the minimum $\\chi^2$ statistic approach is only strictly valid under the normal errors assumption. As with the UUDS observations, this is often clearly not the case --- the error bars may be asymmetric. This means a minimum $\\chi^2$ based assessment may not be reliable.  Finally, merely taking the model with the minimum $\\chi^2$ statistic does not tell us anything about the quality of the model. One can always improve a model fit by adding increasingly more parameters until, at the point of nonidentifiability, the model fit equates to all observation values exactly (``joins the dots''), thus producing a $\\chi^2$ statistic value of zero, which would be then the model of choice based on the minimum $\\chi^2$ statistic criterion. However, it is clear that what we have done in that case is not come closer to the true picture: rather an artefactual model has been constructed which is unlikely to reflect the true picture. This whole problem is one of {\\it overfitting}, whereby model fit is apparently improved by adding increasingly more parameters, and is not taken into account by simply using the $\\chi^2$ statistic.  The practical way round these issues is to resort to a simulation approach. In this paper we describe the use of a Bayesian Markov chain Monte Carlo (MCMC) analysis to solar coronal loop data, which embeds hydrodynamic modelling techniques (\\opencite{wbh95}; see Section \\ref{model}) within a basic Metropolis-Hastings algorithm \\cite{mrrt53,has70}. Then in Section \\ref{results}, we apply our method to three simulated datasets in order to test the efficiency of the method; in particular we examine the artificial inflation and compression of the error bars with subsequent inflation and compression of the, {\\it e.g.}, parameter credible intervals. Section \\ref{realdata} investigates some real coronal loop datasets (PDS and UUDS) and subsequently presents quantitatively based conclusions concerning the nature of the heating of the loops examined. A discussion on how this work can be further progressed can be found in Section \\ref{discussion}. ",
        "conclusions": "\\label{discussion}  In this paper we have presented a new method for comparing observations with theoretical models for solar astrophysical datasets. Bayesian statistics are generally more powerful than a $\\chi^2$ test for the reasons discussed in Section \\ref{intro}. We have examined the robustness of this method with simulated datasets (Section \\ref{results}), and demonstrated how this can be applied in real datasets (Section \\ref{realdata}).  The results from our analysis of the PDS (Section \\ref{priestdataset}) show conclusively that we have basal heating for that loop. The combination of 74 observations and narrower error bars compared with the UUDS mean that all diagnostic assessments strongly and unequivocally indicate this form of heat input. However, as it has been mentioned earlier in this paper, different authors have found different spatial forms of the heating function for this same loop. This is likely to occur because different analysis techniques could give different temperature profiles that do not resemble to each other.  The results of our stochastic analysis of the UUDS (Section \\ref{uudataset}) are inconclusive. On the one hand, we have found that those techniques which use maximisation to locate or estimate parameters (AIC or BIC) suggest a uniform heating mechanism ($\\beta=0$). However, if we use an integral approach, {\\it e.g.} Bayes factors, a negative value of $\\beta$ is found.  The resolution of this contradiction is probably found in the examination of the 95\\% credible interval for $\\beta$ which straddles 0. The marginal distribution of $\\beta$ is negatively skewed, allowing the mode to positive, whilst the majority of the values are negative (the median value is negative). Essentially this is telling us is that we cannot make a firm decision as to the heating mechanism ({\\it i.e.} $\\beta$  negative or positive) from the UUDS. The error bars are simply too large and/or we have insufficient number of observations to allow us to make a heating mechanism determination even with the methodology we have described in this paper.  This leads onto another important point: we need more than one of these diagnostic assessment techniques described in this paper in order to be able to make a rounded reliable judgment on the nature of the heating mechanism. Our recommendation is to use the 95\\% credible intervals, together with a maximisation and integration based method. Use of say, the $\\chi^2$ method alone may lead to a false conclusion, {\\it e.g.} with the UUDS we may have decided that a basal heating mechanism is appropriate for the loop being observed, when in actual fact we can draw no conclusions as to the mechanism.  The fact that information criteria seem to select $\\beta=0$ (uniform heating) as the hypothesis of choice for UUDS may simply reflect a position of the error bars being too wide in comparison with the magnitude of $\\beta$, which in turn determines the strength of basal or apex heating. Therefore, even if a loop is basally or apex heated, the size of $\\beta$ may simply be too small to be ``detected'' by the available data, and, rather like the null hypothesis in classical statistical testing, a conclusion of ``uniform heating'' may be decided upon. Improved data may then come to a very different conclusion on the same loop!        It is clear from the preceding discussion that the Bayesian MCMC methodology described in this paper is a significant development in the assessment of solar coronal loop data. Simply choosing the model paradigm which happens to furnish the minimum $\\chi^2$ value can be problematic. In particular, this Bayesian MCMC approach allows for a quantitative assessment of the parameters simultaneously.  However, it must be kept in mind that this model comparison is only as good as the analysed data to which it is applied. In determining any spatial variation in the thermal/density structure along a loop and relating this to a model form, the main drivers are the number of observed data points along the structure under observation and the size of the error bar associated with each data point. Of course, it is the case that you would want to maximise one (the number of observations obtained) and minimise the other (to produce the smallest error bar).  Given that the spatial resolution of new (future) instrumentation are (will be) an improvement on that considered in this paper, it is likely that the number of data points along a structure will not be a vital issue, assuming one is dealing with a loop of reasonable lengh ($>100$Mm say). For example, the spatial resolution of Hinode/EIS is over twice that of SOHO/CDS. Similarly, a decrease in the size of the associated error bars should occur with greater instrument sensitivity --- however, it is possible that with greater spatial resolution, longer exposure times may be required.  With this in mind, future work in this area will include examining the density structure along many loop examples observed by Hinode/EIS. Also, the numerical scheme will be extended to include gravity, relevant to longer loops.       \\begin{acks} SA has been supported by an STFC grant. \\end{acks}"
    },
    "0807/0807.4179_arXiv.txt": {
        "abstract": "We describe a maximum likelihood regularized beam deconvolution map-making algorithm for data from high resolution, polarization sensitive instruments, such as the \\textsc{Planck} data set. The resulting algorithm, which we call PReBeaM, is pixel-free and solves for the map directly in spherical harmonic space, avoiding pixelization artifacts. While Fourier methods like ours are expected to work best when applied to smooth, large-scale asymmetric beam systematics (such as far-side lobe effects) we show that our $m$-truncated spherical harmonic representation of the beam results in negligible reconstruction error -- even for $m$ as small as 4 for a polarized elliptically asymmetric beam. We describe a hybrid OpenMP/MPI parallelization scheme which allows us to store and manipulate the time-ordered data from instruments with arbitrary scanning strategy. Finally, we apply our technique to noisy data and show that it succeeds in removing visible power spectrum artifacts without generating excess noise on small scales. ",
        "introduction": "One of the most exciting prospects for the upcoming \\textsc{Planck} satellite is its capability to measure the polarization anisotropies of the CMB over the entire sky in nine frequency channels.  The potential rewards from these measurements are many and include tighter constraints on cosmological parameters, determination of the reionization history of the universe, and detection of signatures left by primordial gravitational waves generated during inflation \\citep{bluebook}.  Measurement of the CMB polarization signal presents a great experimental challenge as it is an order of magnitude smaller than the temperature signal and is especially susceptible to distortions due to optical systematics and foreground contaminants. Indeed, if left untreated, leakage from the much stronger temperature signal will contaminate the polarization maps.  Maps and spectra will also suffer from leakage from E-mode polarization to B-mode polarization, jeopardizing the potential detection of inflationary B-modes.  At the resolution and sensitivity of the next generation of experiments, including the Planck mission, studies of primordial non-Gaussianity may also be sensitive to beam-induced systematics. In this paper, we present a novel technique for both assessing and removing systematic effects due to beams in temperature and polarization maps.  The \\textsc{Planck} satellite is designed to extract essentially all of the information in the primordial temperature anisotropies and to measure the polarization anisotropies to high accuracy for $2\\lesssim\\ell\\lesssim 2500$. This will be achieved by measuring the full-sky signal to an angular resolution of 5', to a sensitivity of $\\Delta$T/T $\\sim 2$x$10^{-6}$, and over a frequency range of 30-857 GHz \\citep{bluebook}. The scientific performance of \\textsc{Planck} depends, in part, on the behavior of systematic effects which may distort the signal.  A primary objective of \\textsc{Planck} is to produce all-sky CMB maps at each frequency.  The process by which the satellite's time-ordered data (TOD) is wrapped back on to the sphere to create an image is known as map-making.  The map-making process becomes difficult due to a number of challenges: distortions in the beam, foreground contamination through far-side lobes, size of the data, and correlated noise effects. It is of critical importance to fully characterize the beam, and use this information during map-making to deconvolve beam effects.  We have previously described a powerful map-making algorithm which implements the beam deconvolution technique for temperature measurements \\citep{AW04}.  In this paper, we will extend that description to include {\\it polarization} measurements.  We refer to this new technique as PReBeaM: Polarized Regularized Beam deconvolution Map-making. While we focus on reconstructing the map with a uniform effective beam and realize corrections to the power spectrum as a consequence, other work by \\citet{Souradeep06} and \\citet{Mitra07} has focused on deriving corrections to the power spectrum due to asymmetric (non-circular) beam effects.  Within the \\textsc{Planck} collaboration, the CTP working group has developed five map-making methods and compared their results using the simulated 30 GHz data in what is known as the Trieste paper \\citep{Trieste}. The Trieste paper assessed the impact of beam asymmetries on the \\textsc{Planck} spectra without attempting to treat the problem of beam asymmetry at the map-making level (an angular power spectrum correction method was developed based on simplifying assumptions). In addition to PReBeaM, another deconvolution map-making technique for \\textsc{Planck} has been established by \\citet{Harrison08}. Both methods allow for arbitrary beam shapes and in both cases the asymmetry of the beam is parametrized by an asymmetry parameter  $m_{max}$ which can vary between 0 and $\\ell_{max}$.  Our method scales computationally as a function of $m_{max}$; this is advantageous when large gains in accuracy can be achieved with small increases in $m_{max}$.  In contrast, the Harrison method incurs a fixed computational expense for arbitrarily large $m_{max}$. The Harrison method takes advantage of the Planck scanning strategy to condense the full TOD into phase-binned rings, thereby achieving a significant reduction in processing time.  A complete characterization of the beam includes both the main beam and the far-side lobes.  Sidelobes are located as far away as $90^{\\circ}$ from the main focal plane beam, and therefore require a large $m_{max}$ parameter for a complete harmonic description.  In \\citet{AW04} we demonstrated the full potential of our method using far-side lobes and maps with foreground signals.  Here, we show the usefulness of PReBeaM for deconvolving main-beam distortions.  In fact, we find that it makes sense to use PReBeaM for main beam effects since only a small $m_{max}$ parameter is needed to capture the azimuthal structure of the main beam.  In this way, we profit from the computational advantage of our method in the case of small $m_{max}$, allowing for the unified treatment of main beam and side lobe effects.  In \\S\\ref{sec:deconv} we describe the deconvolution map-making algorithm for PReBeaM. The simulated data and beams are detailed in \\S\\ref{sec:sims}. We present results in \\S\\ref{sec:results} showing the effectiveness of PReBeaM in removing systematic effects due to beam asymmetry and we discuss computational considerations. We finish with our conclusions from this study in \\S\\ref{sec:conclusion}. ",
        "conclusions": "\\label{sec:conclusion}  We have found that PReBeaM has outperformed the standard binned noiseless map using two measures: spectra and residual maps.  We examined the fractional differences in the spectra and found markedly smaller differences in the PReBeaM spectra versus the binned map spectra across a range of multipole moments.  We find that map-making codes which do not deconvolve beam asymmetries lead to significant systematics in the polarization power spectra measurements.  The temperature-to-polarization cross-coupling due to beam asymmetries is manifested as shifts in the peaks and valleys of the spectra.  These shifts are absent from the PReBeaM spectra.  We translated the errors found in the power spectra to an estimate of the statistical significance of the errors in a parameter estimation resulting from these spectra, which we call $n_{\\sigma}$.  This analysis showed that the worst case parameter bias due to beam-induced power spectrum systematics could be tens of sigma while PReBeaM reduces the risk of parameter bias due to beam systematics to much less than 1 sigma We also found the I, Q, and U component residual maps to be smaller for PReBeaM than for the binned map, implying smaller map-making errors.   We have presented here the first results from PReBeaM for a straightforward test cases of CMB only and CMB plus white noise, and including only the effects of beams in the main focal plane.  However, there is great potential for using PReBeaM to remove or assess systematics due to the combination of foregrounds and beam side lobes.  Systematics introduced by side lobes will appear on the largest scales, potentially impeding the detection of primordial B-modes on the scales where they are most likely to be measured. We have already shown for temperature measurements \\citep{AW04} that our deconvolution technique can be used to remove effects due to side lobes.  Future work will examine the noise properties of PReBeaM maps and will include foregrounds from extragalactic sources and diffuse Galactic emission."
    },
    "0807/0807.2862.txt": {
        "abstract": "Zeeman observations of molecular clouds yield the line-of-sight component $B_{LOS}$ of the magnetic vector {\\bf B}, which makes it possible to test the two major extreme-case theories of what drives star formation -- ambipolar diffusion or turbulence.  However, only one of the three components of {\\bf B} is measurable, so tests have been statistical rather than direct, and they have not been definitive. We report here observations of the Zeeman effect in the 18-cm lines of OH in the envelope regions surrounding four molecular cloud cores toward which detections of $B_{LOS}$ have been achieved in the same lines, and evaluate the {\\em ratio} of mass to magnetic flux, $M/\\Phi$, between the cloud core and envelope. This relative $M/\\Phi$ measurement reduces uncertainties in previous studies, such as the angle between {\\bf B} and the line of sight and the value of [OH/H]. Our result is that for all four clouds, the ratios ${\\cal R}$ of the core to the envelope values of $M/\\Phi$ are less than 1. Stated another way, the ratios ${\\cal R^\\prime}$ of the core to the total cloud  $M/\\Phi$ are less than 1. The extreme case or idealized (no turbulence) ambipolar diffusion theory of core formation requires the ratio of the central to total $M/\\Phi$ to be approximately equal to the inverse of the original subcritical $M/\\Phi$, or ${\\cal R^\\prime} > 1$. The probability that all four of our clouds have ${\\cal R^\\prime} > 1$ is $3 \\times 10^{-7}$; our results are therefore significantly in contradiction with the hypothesis that these four cores were formed by ambipolar diffuson. Highly super-Alfv\\'{e}nic turbulent simulations yield a wide range of relative $M/\\Phi$, but favor a ratio ${\\cal R} < 1$, as we observe. Our experiment is limited to four clouds, and we can only directly test the predictions of the extreme-case ``idealized'' models of ambipolar-diffusion driven star formation that have a regular magnetic field morphology. Nonetheless, our experimental results are not consistent with the ``idealized'' strong field, ambipolar diffusion theory of star formation. Comparisons of our results with more realistic models and simulations that include both ambipolar diffusion and turbulence may help to refine our understanding of the relative importance of magnetic fields and turbulence in the star formation process. ",
        "introduction": "Understanding star formation is a fundamental astrophysical problem. \\citet{MO07} have comprehensively reviewed the field. For thirty years what has sometimes been called the ``standard model'' has been that magnetic fields control the formation and evolution of the molecular clouds from which stars form, including the formation of cores and their gravitational collapse to form protostars. However, in recent years doubts about the validity of this model have been raised by those who argue that turbulence controls the formation of clouds and cores, with cores either dissipating back into the general interstellar medium or collapsing and forming stars if they are self-gravitating when formed. In spite of decades of intense research, there is still not consensus on the role that magnetic fields play in the star formation process.  Detailed theoretical work on the strong magnetic field ``standard model'' has been carried out by a number of groups: \\citet{Setal99} and \\citet{MC99} have reviewed and summarized the state of this theory. The fundamental principle is that clouds are formed with subcritical masses ($M < M_\\Phi = \\Phi/2\\pi\\sqrt{G}$). Here $\\Phi$ is the magnetic flux and G is the gravitational constant, and the expression for $M_\\Phi$ is from \\citet{NN78}; other expressions for $M_\\Phi$ differ slightly from this depending on cloud structure (e.g., \\citealt{MS76}). The magnetic field is frozen only into the ionized gas and dust; neutral gas and dust contract gravitationally through the field and the ions, increasing mass (but not to first order flux) in the cloud cores. This process is known as ambipolar diffusion. When the core mass reaches and exceeds $M_\\Phi$, the core becomes supercritical ($M > M_\\Phi$), collapses, and forms stars. The magnetic flux remains behind in the envelope. Because the ambipolar diffusion time scale for the formation of supercritical cores is fairly long ($\\geq 10^7$ yr), molecular clouds would have long lifetimes. The star formation efficiency (the ratio of mass in stars in a molecular cloud complex to the interstellar mass) is low (as observed) due to the slow rate of star formation and to the fact that much of the mass of the molecular cloud is left behind in the subcritical envelope as the core collapses.  The idea that star formation is primarily regulated by ambipolar diffusion was standard for many years.  However, doubts about the validity of this assumption were raised by the development of a weak-field, super-Alfv\\'{e}nic model of dark clouds \\citep{PN99}. The  new weak-field theory had molecular clouds being intermittent phenomena, with short ($\\sim 10^6$ yr) lifetimes. In this theory clouds form at the intersection of turbulent supersonic flows in the interstellar medium. Generally, clouds do not become gravitationally bound and dissipate; those that are self-gravitating form stars in essentially a free-fall time \\citep{E00}. Star formation occurs only in the small fraction of the molecular gas that is sufficiently dense to be self-gravitating. The star formation efficiency is low due to a small fraction of the mass of clouds becoming gravitationally bound. Magnetic fields are present in this theory, but they are too weak to be energetically dominant. The role of turbulence in the energetics of the interstellar medium has been a very active area.  \\citet{Petal04} presented evidence in favor of the weak field, super Alfv\\'{e}nic model. \\citet{ES04} have written an excellent review of interstellar turbulence, and \\citet{MK04} have extensively reviewed arguments that supersonic turbulence controls star formation.  Early work focused on extreme-case models, either strong magnetic field models that did not (at least directly) include the effects of turbulence, or strong turbulence models that neglected the effects of ambipolar diffusion. More recent theoretical work has introduced supersonic turbulence into numerical star formation models \\citep[e.g.,][]{NL05, KB08, NL08}.  These more recent models apply to magnetically subcritical regions, and they include the effects of ambipolar diffusion.   The role of turbulence proves to be important in these models, significantly shortening, for example, the ambipolar diffusion timescale. In this paper we focus on a direct comparison of observational results with the strong field, ambipolar diffusion model without turbulence. Such models have a smooth morphology of the magnetic field which deforms into an hourglass shape with axis along the mean field direction, and make an analytic prediction that can be tested directly. The more complicated models that come from simulations that include both magnetic fields and turbulence can also be tested against the results of this paper, and we discuss how such tests might be carried out.  Our approach to testing star formation theory has been to measure magnetic field strengths in molecular clouds in order to see whether they are weak or strong.  The crucial parameter is the ratio of the mass to the magnetic flux, $M/\\Phi$, which is of course closely related to $M_\\Phi$. If $M/\\Phi$ is observed to be significantly supercritical, particularly at lower densities, the magnetic support theory is not viable. If it is observed to be subcritical at lower densities, magnetic fields would be too strong for the intermittent, turbulent theory to hold. The $M/\\Phi$ parameter provides in principle a straightforward, direct test to discriminate between the two extreme theories of star formation.  The Zeeman effect provides the only known method to directly measure magnetic field strengths in dense gas. In this paper we first briefly review Zeeman results and argue that they have not yet provided an unambiguous discrimination between the two star formation theories. We then describe a new experiment to attempt to overcome the limitations of existing Zeeman results and to test directly a prediction of the ``standard model''. Finally, we describe the observations made for this experiment and the results, and discuss the implications of those results. ",
        "conclusions": "Previous Zeeman studies of magnetic fields in molecular clouds have not been definitive in testing the two extreme-case models of star formation. The mean mass-to-flux ratios $M/\\Phi$ found from these statistical studies were slightly supercritical -- consistent with either theory. Detailed ambipolar diffusion models for two clouds found excellent agreement with the observations, although both required the field to be nearly in the plane of the sky in order not to produce line-of-sight fields much stronger than observed. Uncertainties in the angle between {\\bf B} and the line of sight and in the total hydrogen column density are inherent in measuring  with Zeeman observations. In order to mitigate these uncertainties, we have measured the ratio of $M/\\Phi$ between the envelopes and cores of four molecular clouds in order to test ambipolar diffusion (strong magnetic fields) versus turbulence (weak magnetic fields) driven star formation theory. The theory of star formation that hypothesizes clouds initially supported by strong magnetic fields, with evolution and core formation being driven by ambipolar diffusion, predicts that the central $M/\\Phi$ must increase as ambipolar diffusion acts. Idealized models predict that the increase in $M/\\Phi$ up to the point when the core becomes supercritical and gravitational collapse proceeds is approximately equal to the inverse of the amount by which the original cloud was subcritical; that is, ${\\cal R^\\prime} > 1$. The probability that all four of our clouds have ${\\cal R^\\prime} > 1$ is $3 \\times 10^{-7}$, a highly significant result. On the other hand, simulations which form clouds and cores by turbulence acting in a weak magnetic field environment preferentially yield a $M/\\Phi$ ratio between core and envelope ${\\cal R} < 1$, in agreement with our results.  Telescope availability limitations allowed only four clouds to be observed; unfortunately the extremely large amount of telescope time required precludes expanding this experiment beyond four clouds for the foreseeable future. The theoretical predictions of ${\\cal R}$ and ${\\cal R^\\prime}$ are based on idealized ambipolar diffusion models and idealized turbulence simulations. Nonetheless, the clear conclusion from our experiment is that at least for these four clouds, the prediction of the idealized ambipolar diffusion models does not agree with our observational results, while the prediction of initially supercritical turbulence-driven simulations does. Still untested is whether simulations that include both significant magnetic fields and turbulence better match the data than either of the extreme cases. We suggest that all theorists who simulate the formation and evolution of molecular clouds and cores test their simulations against the results of this experiment in the manner of \\citet{Letal08}; that is, by calculating Stokes I and V spectra of OH from the simulations and ``observing'' $B_{LOS}$ with our beam patterns (figure 1)."
    },
    "0807/0807.0100_arXiv.txt": {
        "abstract": "{} {The aim of the work presented in this paper is to test and optimise supernova detection methods based on the \\emph{optimal image subtraction} technique.  The main focus is on applying the detection methods to wide field supernova imaging surveys and in particular to the Stockholm VIMOS Supernova Survey (SVISS).} {We have constructed a supernova detection pipeline for imaging surveys. The core of the pipeline is image subtraction using the ISIS 2.2 package. Using real data from the SVISS we simulate supernovae in the images, both inside and outside galaxies. The detection pipeline is then run on the simulated frames and the effects of image quality and subtraction parameters on the detection efficiency and photometric accuracy are studied.} {The pipeline allows efficient detection of faint supernovae in the deep imaging data. It also allows controlling and correcting for possible systematic effects in the SN detection and photometry. We find such a systematic effect in the form of a small systematic flux offset remaining at the positions of galaxies in the subtracted frames. This offset will not only affect the photometric accuracy of the survey, but also the detection efficiencies.} {Our study has shown that ISIS 2.2 works well for the SVISS data. We have found that the detection efficiency and photometric accuracy of the survey are affected by the stamp selection for the image subtraction and by host galaxy brightness. With our tools the subtraction results can be further optimised, any systematic effects can be controlled and photometric errors estimated, which is very important for the SVISS, as well as for future SN searches based on large imaging surveys such as Pan-STARRS and LSST.} ",
        "introduction": "\\label{sec:intro} Imaging surveys to find and characterise variable sources are used to study a number of different astronomical phenomena. Detection and photometry of variable sources can be done by subtracting images taken at different epochs from each other and analysing the subtracted frames, a technique known as \\emph{difference imaging}. In addition to supernovae (SNe), this technique has been used by various authors investigating many different subjects, among which are variable stars in crowded fields \\citep[e.g.][]{2006AJ....132.1014C}, gravitational micro-lensing \\citep[e.g.][]{2001MNRAS.320..341A,2006ApJ...636..240S} and planet detection \\citep[e.g.][]{2007ApJ...655.1103H}.  The most complicated step in the difference imaging process is taking into account the spatial and temporal variability of the point spread functions (PSFs) of the images. The difference imaging technique used in this work is based on the Optimal Image Subtraction (OIS) code first presented in \\citet[ hereafter AL98]{1998ApJ...503..325A}, later improved upon in \\citet[ hereafter A00]{2000A&AS..144..363A} and available in the ISIS 2.2 package. In these papers the code is tested on variable stars (i.e. point sources with varying brightness with little or no background light). When detecting SNe the situation can be quite different.  In this type of data the background light from the host galaxy often dominates the total flux at the location of the supernova.  Difference imaging, and OIS in particular, has been used to detect SNe in many projects. Recently a number of large SN surveys, mainly aiming at finding thermonuclear SNe for determining cosmological parameters, have released their first results. Image subtraction using the OIS algorithm followed by detection of SN candidates in the subtracted frames is a technique that has been used by the ESSENCE survey \\citep[the observations and data analysis are described in][]{2007astro.ph..1043M}, the Supernova Legacy Survey \\citep{2006A&A...447...31A} and \\citet{2007MNRAS.382.1169P}. \\citet{2005A&A...430...83C} and \\citet{2008A&A...479...49B} used the same technique in their SN surveys aiming to determine the core collapse SN rate at intermediate redshift. Difference imaging has also successfully been used in near-infrared SN searches \\citep[e.g.][]{2001MNRAS.324..325M,2002A&A...389...84M} concentrating on individual starburst galaxies. For surveys with the aim of finding supernova rates the detection efficiency is the most important issue, while the SN photometry is of somewhat less importance. However, for surveys targeting Ia SNe the photometry is of utmost importance since the photometric accuracy directly determines how well the cosmological parameters can be constrained. Therefore, in these programmes the photometry is treated very carefully and not necessarily obtained from the subtracted frames. For studies of individual SNe difference imaging has been used to obtain accurate photometry of supernovae with a complicated background.  This is particularly important at late times when the SN has become too faint compared to the background host galaxy in the unsubtracted image for PSF photometry to give reliable results \\citep[e.g.][]{2002A&A...386..944S}. More recently, difference imaging has also successfully been applied to mid-infrared SN observations with the Spitzer Space Telescope \\citep{2006ApJ...649..332M} and to adaptive optics assisted detection of SNe in luminous infrared galaxies \\citep{2007ApJ...659L...9M}.  In this paper we investigate several aspects of difference imaging with ISIS 2.2 when used to detect and study supernovae. Previous papers have already reported some tests done on the image subtraction \\citep[e.g., AL98; A00;][]{2007AN....328...16I} and on the SN detection process \\citep[e.g.][]{2001MNRAS.324..325M}.  The surveys for thermonuclear SNe all have extensive sections where the methods to ascertain the photometric accuracy (both in terms of statistical and systematic errors) is described \\citep[e.g.][]{2006A&A...447...31A}. However, no systematic tests on the subtraction method applied to supernova data have been reported. The main purpose of this paper is to describe our supernova detection pipeline and to study how it performs on wide field imaging surveys, in particular the Stockholm VIMOS Supernova Survey (SVISS). The SVISS is an imaging survey aiming to find core-collapse supernovae at $0.1< z < 1.0$ \\citep[for details see][]{1999A&A...350..349D}. The main scientific goal for this project is to find reliable core collapse supernova rate densities for this redshift range. We thus need to have detailed knowledge of the supernova detection efficiency. In this project the photometric accuracy is important to correctly type the SNe based on colour and light curve information. Another goal of this paper is to investigate how supernova detection and photometry (with ISIS used as the image subtraction tool) are affected by the selection of subtraction and detection parameters and the presence of a host galaxy.  In this paper we discuss the subtraction and detection method, in the following papers we will present the light curves of the detected SNe, the typing methodology and calculations of SN rates at high redshifts.  The first part of the paper (section \\ref{sec:method}) contains a detailed description of the standard image subtraction method used in this work. In section \\ref{sec:data} we describe the data sets, in section \\ref{sec:sim} how the simulations of supernovae have been setup and in section \\ref{sec:results} we present the results from the extensive testing.  Finally, we discuss the results and possible implications in section \\ref{sec:disc} and conclude by giving a short summary in section \\ref{sec:summary}. The Vega magnitude system has been used throughout the paper. ",
        "conclusions": "\\label{sec:disc} \\subsection{Overall pipeline performance} We have studied the detection efficiency and photometric accuracy of our SN detection pipeline using real data from the SVISS. However, these results can likely be applied to any SN ground based imaging survey using ISIS subtracted images for detection and/or photometry. Previous tests of image subtraction software have been done mainly on artificial images \\citep[AL98, A00, ][]{2002A&A...381.1095G,2007AN....328...16I} but also on stellar fields (AL98). To our knowledge the work presented in this paper is the first detailed study of subtraction quality in real supernova survey data. Detection efficiency studies in supernova survey data are regularly carried out to be able to find the supernova rates \\citep[e.g.][]{2005A&A...430...83C,2005AJ....130.2272B,2007astro.ph..1043M}, but this testing is seldom reported in detail.  We have found a systematic flux offset affecting the photometry in subtracted images. In general, with host galaxies chosen at random in the field, the offset is not dependent on the SN brightness, the offset is constant for different simulated SN magnitude bins. The offset is only found for SNe placed inside host galaxies whereas in isolated positions no offset is found. The offsets change between the different epochs, in some cases being positive and in others negative. Pixel statistics for pixels inside apertures centred on the host galaxies also show that there is a shift of the median pixel value in the same direction as the flux difference plots indicate and as expected, a broadening of the distribution. We interpret this as a small flux through/plateau at the locations of galaxies, reaching as far out as to the 80\\% flux radius. Currently we do not have an explanation for this effect, it cannot be solved by changing the convolution kernel sizes or stamp sizes. Further testing is needed to pinpoint if the systematic offset is caused by properties of the dataset or if it is in fact an internal effect of the image subtraction method. It should be noted that this is a small offset, the systematic error on the magnitudes of SNe is smaller for brighter sources (for the brightest SNe considered in this paper, $m_{\\mathrm{I}} = 22.0$, the offset is $\\sim 0.005$) and in general significantly smaller than the photometric scatter for an individual source. However, in a survey of multiple objects, and particularly fainter objects, this systematic effect will likely affect the result unless dealt with.  A simple estimate of the photometric uncertainties using the pixel to pixel noise in the local background underestimates the errors when doing photometry on SNe in subtracted images. To get an accurate noise estimate we have instead computed the photometric scatter for each simulated magnitude bin. This scatter has also been used to construct signal to noise ratios, using this method the $SNR$ is not affected by the flux offset and provides a simple way of rejecting less trustworthy observations.  \\citet{2007astro.ph..1043M} use simulations to check whether image registration and subtraction affect their SN photometry.  They find no systematic offsets to a $SNR$ of 5, but this requires an accurate SN position obtained from multiple points on the light curve. They find the photometric scatter to be on the order of 20 \\% larger than what their noise maps indicate (but the noise maps already include the effects of re-binning and convolution). This seems considerably better than the accuracy we find using our pipeline and simulations on the SVISS data. The method used by these authors is quite different from the one used in this work, therefore it is hard to compare the exact figures. We estimate the $SNR$s differently by using the scatter found from simulations as our noise. It should also be noted that the SN positions in the images used by \\citet{2007astro.ph..1043M} are chosen at random (private correspondence G. Miknaitis), thus their results will not be directly comparable to our inside host galaxies simulations. The main difference is that we find that subtraction residuals from the host galaxies can have a noticeable effect on the photometric accuracy of the survey. This effect does not seem to be present in the work of \\citet{2007astro.ph..1043M}.  The presence of spurious detections is a problem for all supernova surveys, and in particular for surveys with spectroscopic follow-up \\citep[see the discussion in][]{2007arXiv0705.0493B}. For multi-band, multi-epoch imaging surveys (like the SVISS) the spurious detections is less of a problem, but it is still desired to keep the number down. The bright residuals found in less successful image subtractions can easily be mistaken for real varying sources when using automatic source detection. The number of spurious detections in our tests is large, partly because we use a low detection threshold, but in general we find that a single epoch, single filter run results in a spurious detection rate of around 80\\% (i.e. 1 out of 5 are real sources) when the detection threshold is set to find objects brighter than the 5$\\sigma$ limiting magnitude. For the SVISS dataset used in this paper the total number of detected sources for a given epoch is on the order of hundreds, making the number of spurious detections quite large. We find that the use of a controlled and bright stamp list makes the number of spurious detections go down by about 10\\%. It should be noted that residuals at the locations of bright galaxies will be problematic even in a multi-band survey since the residuals often appear at the same positions in all filters.  \\subsection{Host galaxy brightness and supernova position} We find that the typical brightness of the host galaxies does affect the detection efficiencies and photometric accuracy. This is of course as expected and also confirmed by other SN surveys \\citep[e.g.][]{2006AJ....132.1126N}. However, both detection efficiency and photometric accuracy becomes gradually worse when considering host galaxies of increasing brightness. The effect is mainly noticeable for the very brightest host galaxies (i.e. largest contrast between the SN and the host). The fact that the photometry is affected by host brightness can in principle cause the systematic offset found above to vary with the SN brightness, since the SN and host brightness are correlated through the host galaxy distance. In projects where the supernova rate is studied, the effects of host brightness most likely have an impact on the detection efficiencies at low redshifts, where the host galaxies can be very bright, as well as in bright galaxies at high redshifts where the supernovae will be very faint. Whether this effect will be noticeable depends on the characteristics of the galaxy population, and simulations of the type reported in this work can successfully quantify the effect for galaxies of different brightness. This information can then be incorporated into the overall detection efficiencies discussed above. The SNe detections in the SVISS data are likely to be found in faint galaxies and thus in the general case not affected strongly by the host background flux.  The effect of the SN galactocentric distance in the host galaxy is difficult to study in the SVISS data set, since most of the galaxies have small angular sizes, and in many cases are not even resolved. Our results indicate that the flux offset is slightly worse in the inner parts (40\\% in flux units) of the galaxies than in the outer (40-80\\%) regions. Therefore, we cannot draw any conclusions on how the detection efficiencies are affected by the galactocentric distances. However, if the detection efficiency is dependent on radial distance from the host centre, the radial distribution function of SNe needs to be known to find a proper correction factor for the total efficiency.  \\subsection{Image subtraction method} Of the different tested parameters we find that the stamp selection method affects the subtraction quality and detection success the most. The automatic stamp selection in ISIS mostly works fine, but the lack of control in selecting stamps can be problematic when dealing with real SN survey data, due to the presence of saturated objects and the rejection routine sometimes failing for varying objects. Using a set of sub-optimal stamps will result in a subtracted frame with worse subtraction quality.  We note that the optimal selection for stamps is based on SNR, selecting preferentially point-like sources as stamps, i.e. in general fainter sources than in the SNR based selection, degrades the subtraction quality. Varying the widths of the Gaussian basis functions of the kernel within a reasonable range does not significantly alter the detection results, however it should be noted that if all of the functions are too wide, depending on the data in question the image subtraction will fail.We also note that for surveys with small seeing differences, like the SVISS, the default kernel widths used in ISIS might be to wide. Using widths more in line with the theoretically optimal convolution kernel width can marginally improve the subtraction quality.  \\subsection{Implications for SN surveys} The detection efficiency of supernovae in a multi-epoch survey can be studied by Monte Carlo simulations of SN light curves as observed by the telescope used for the survey \\citep[see e.g.][]{2004ApJ...613..189D}.  The light curves, together with the individual detection efficiencies for each epoch can then be used to find an overall detection efficiency for supernovae of a certain type and at a given redshift. The flux offset effect found in this work will affect the detection limits as well as the number of spurious detections. These can be handled in different ways, but their presence will always be a problem and minimising the number of spurious detections is frequently important. A positive flux offset will result in an artificially boosted detection efficiency and an increase in the number of spurious detections. A negative offset will instead cause the detection efficiency to go down, making the survey to lose depth. This effect can possibly be countered by controlling the detection threshold used by SE using a weight map, in principle redefining the threshold for positions considered to be inside galaxies.  It is very unlikely that systematic flux offsets of the kind found in this work affect the results found by the large SN Ia projects. In these projects \\citep[e.g.][]{2006ApJ...637..427B,2006A&A...447...31A, 2007astro.ph..1043M}, the photometric accuracy is studied extensively to make sure that the systematic errors in the cosmological parameter determination is minimised. As shown above the detection efficiency will be influenced by the flux offset. Projects studying the supernova rate and doing detections in images subtracted using methods similar to ours might suffer from this effect and detailed simulations should be run to make sure that no offset is present (or correct for it otherwise). The results from some recent surveys \\citep{2004ApJ...613..189D,2005A&A...430...83C, 2006ApJ...637..427B} are based on supernovae with observed spectroscopy. The limiting magnitudes for these surveys are thus set by the magnitude limit of the spectroscopic instrument used , which is considerably brighter than the corresponding imaging limit, thus the offset problem described here is likely not an issue. However, with present and upcoming surveys based on imaging alone and pushing the detections to the fainter limits this effect needs to be understood and studied thoroughly.  Based on the findings reported in this paper we can make some recommendations for supernova detection and photometry using image subtraction methods.  Testing should be done on at least part of the data to determine optimal parameters for the image subtraction. The stamp selection should be controlled in some way, a good option is to create a master stamp list from the reference image and use that in all subtractions but making sure that none of the stamps contain SNe. If the automatic stamp selection in ISIS is used, care must be taken that the selected stamps are not variable or contain saturated pixels. To make sure that a possible flux offset will not affect the results and to estimate photometric errors, simulations should be done. If photometry is done in the subtracted frames, the measurements for an individual SN should be carefully done. This can be done by simulating SNe of similar magnitude in nearby galaxies of the same brightness and size as the host galaxy and by measuring the possible offset. The photometric errors for each point on the light curve should be estimated by using the magnitude scatter found from simulations. If no offset correction for the detection efficiencies is applied we recommend that the detection efficiencies should be capped by the efficiency for isolated SNe (in the presence of a positive flux offset), the true efficiency can never be better inside than outside galaxies.   We have described the SN detection and photometry pipeline used for the SVISS. We have also presented detailed testing of supernova detections and photometry using image subtraction and automatic detection. Our main results and conclusion are: \\begin{itemize} \\item{Image subtraction using ISIS 2.2 works well for the SVISS, the subtraction quality as measured by noise levels in the location of galaxies is in general good. Simulations show that our pipeline is successful in detecting SNe as well as in doing photometry on them.} \\item{The detection efficiency and photometric accuracy for supernovae in a single epoch is affected by how the stamp selection for ISIS is done.  We have also found that the host galaxy brightness does affect these quantities, likely through the presence of powerful subtraction residuals at the location of bright galaxies.} \\item{Image subtraction using ISIS benefits from using a controlled stamp list. We have found that the best subtraction quality is achieved when the subtraction stamps used have high signal-to-noise ratios, ideally the brightest non-varying sources in the field should be used as stamps. The number of stamps depends on the size and quality of the image, and for our data set $\\sim$60 stamps per ($1k\\times1k$) pixels has proven to be sufficient.} \\item{We have discovered a systematic flux offset in subtracted images. This offset only appears at the locations of galaxies and is quite small. Photometry of simulated SNe in isolated positions show no offset. For individual sources the offset is smaller than the estimated photometric scatter. This effect can cause systematic errors in the photometry if not corrected for. Detection limits of SNe will also be affected by non-zero offset, causing problems with increased numbers of spurious detections for positive offsets and a decrease in survey depth for a negative offset.  It should be noted that, based on the tests run so far, we cannot say for sure whether the offset is something unique to the SVISS data set or if it is due to the image subtraction method.} \\item{Simulations of supernovae both in and outside galaxies are needed to obtain proper estimates of the photometric errors. The simulations can also be used to correct for possible systematic offsets.} \\item{Using automatic source detection (in this case SE) we find that the detection efficiencies can be pushed to the limits of the data. The number of spurious detections in a single epoch subtraction is very large, on the order of 80\\%, but decreases as a function of subtraction quality. To reject spurious detections other methods must be used (e.g. multiple filters, light curve fitting).} \\end{itemize} Understanding possible systematic effects of the subtraction and detection process is very important when interpreting the results from the SVISS, but will also be imperative for the upcoming supernova projects based on large imaging surveys, e.g.  Pan-STARRS\\footnote{http://pan-starrs.ifa.hawaii.edu}, LSST\\footnote{http://www.lsst.org/lsst\\_home.shtml} and SNAP \\citep{2005NewAR..49..346A}."
    },
    "0807/0807.4932_arXiv.txt": {
        "abstract": "Galaxy  clustering and  galaxy-galaxy lensing  probe  the connection between galaxies and their dark matter haloes in complementary ways. Since the clustering of dark matter haloes depends on cosmology, the halo  occupation statistics  inferred from  the  observed clustering properties of  galaxies are degenerate  with the adopted  cosmology. Consequently,  different cosmologies  imply  different mass-to-light ratios for dark matter  haloes.  Galaxy-galaxy lensing, which yields direct constraints on the actual mass-to-light ratios, can therefore be used to break this degeneracy, and thus to constrain cosmological parameters.   In this  paper we  establish the  link  between galaxy luminosity  and   dark  matter  halo  mass   using  the  conditional luminosity  function  (CLF), $\\Phi(L|M){\\rm  d}L$,  which gives  the number of galaxies with luminosities in the range $L \\pm {\\rm d}L/2$ that reside in a halo of  mass $M$.  We constrain the CLF parameters using the  galaxy luminosity function and  the luminosity dependence of the  correlation lengths of  galaxies.  The resulting  CLF models are  used  to  predict  the  galaxy-galaxy lensing  signal.   For  a cosmology  that agrees  with constraints  from the  cosmic microwave background, i.e.   $(\\Omega_{\\rm m},\\sigma_8) =  (0.238,0.734)$, the model accurately  fits the galaxy-galaxy lensing  data obtained from the   SDSS.    For  a   comparison   cosmology  with   $(\\Omega_{\\rm m},\\sigma_8)  =  (0.3,0.9)$,  however,  we can  accurately  fit  the luminosity  function   and  clustering  properties   of  the  galaxy population, but the model predicts mass-to-light ratios that are too high,  resulting in  a  strong overprediction  of the  galaxy-galaxy lensing signal.  We conclude  that the combination of galaxy clustering and galaxy-galaxy lensing is  a powerful probe  of the galaxy-dark  matter connection, with  the  potential  to  yield tight  constraints  on  cosmological parameters.  Since  this method mainly probes  the mass distribution on  relatively small  (non-linear)  scales, it  is complementary  to constraints  obtained from the  galaxy power-spectrum,  which mainly probes the large-scale (linear) matter distribution. ",
        "introduction": "\\label{sec:intro}  With   the  advent of large  galaxy  redshift  surveys, it  has become possible to obtain accurate measurements of the clustering of galaxies as a function of their properties, such  as luminosity, morphology and color (e.g.  Guzzo \\etal 2000; Norberg \\etal  2001, 2002; Zehavi \\etal 2005; Wang \\etal 2007). Since galaxies are believed to form and reside in  dark matter haloes, the clustering  strength of a given population of galaxies can be compared to that of dark matter haloes as predicted by  numerical simulations  or the  extended Press-Schechter formalism. Such a comparison reveals a wealth of  information about the so-called \\textit{galaxy-dark matter connection} (e.g.    Jing, Mo  \\&  B\\\"orner 1998; Peacock \\& Smith 2000; Berlind \\& Weinberg 2002; Yang, Mo \\& van den Bosch 2003; van  den Bosch, Yang \\&  Mo 2003; van den Bosch  \\etal 2007).  Unfortunately, this method of  constraining the link  between galaxies and  dark matter  haloes  using  galaxy clustering   has one important shortcoming: the halo occupation statistics inferred from the observed clustering properties  depend on the cosmological  parameters adopted. More precisely,  models based  on  different cosmologies can   fit the clustering data    equally well by simply   relying  on different halo occupation  statistics    or,  equivalently, different   mass-to-light ratios. In  order to break this  degeneracy between cosmology and halo occupation statistics   independent constraints  on  the mass-to-light ratios are required  (e.g.   van den Bosch,   Mo \\& Yang  2003; Tinker \\etal  2005).   One  method that  can  provide  these  constraints  is galaxy-galaxy lensing (hereafter g-g lensing),  which probes the  mass distributions (and  hence  the halo  masses)   around galaxies.   This implies that  the combination of  clustering and  lensing in principle holds the   potential to put   constraints  on cosmological parameters (Seljak \\etal 2005; Yoo \\etal 2006).  The  first attempt  to  detect g-g  lensing  was  made by  Tyson \\etal (1984), but because of the relatively  poor quality of their data they were unable to detect a  statistically  significant signal.  With  the advent  of wider and deeper  surveys  becoming available, however, g-g lensing  has now  been detected  with very  high significance, and  as function  of   various  properties  of   the  lensing  galaxies  (e.g. Griffiths \\etal  1996; Hudson \\etal  1998; McKay \\etal 2001;  Guzik \\& Seljak 2002;  Hoekstra \\etal 2003, 2004;  Sheldon \\etal 2004, 2007a,b; Mandelbaum \\etal 2006; Heymans \\etal 2006; Johnston \\etal 2007; Parker \\etal  2007; Mandelbaum,  Seljak  \\&  Hirata 2008).   Unfortunately, a proper interpretation  of  these data  in terms  of  the  link between galaxies and dark matter haloes has been hampered by the fact that the lensing signal can typically only be detected when stacking the signal of many  lenses.  Since not all  lenses  reside in haloes of  the same mass, the resulting   signal is a non-trivial  average  of the lensing signal due to  haloes of different masses.   Most studies to date have assumed that the relation between the  luminosity of a lens galaxy and the mass of its halo is given by a simple power-law relation with zero scatter (see Limousin  \\etal 2007 for  a detailed overview).  However, it   has become clear,   recently, that the   scatter in this relation between light and mass can be  very substantial (More \\etal 2008b, and references therein). As shown by Tasitsiomi \\etal (2004), this scatter has a very  significant impact on the actual  lensing signal, and thus has  to be  accounted  for  in  the  analysis.   In addition,  central galaxies (those residing   at the center  of  a dark matter halo)  and satellite galaxies (those orbiting around a central galaxy) contribute very different lensing signals, even when they reside in haloes of the same mass  (e.g.   Natarajan, Kneib \\&  Smail  2002; Yang  \\etal 2006; Limousin \\etal 2007). This has to be  properly accounted for (see e.g. Guzik  \\& Seljak 2002), and  requires  knowledge of both the satellite fractions and of the spatial number  density distribution of satellite galaxies within their dark matter haloes.  Over  the years, numerous techniques have  been developed to interpret galaxy-galaxy lensing measurements  (Schneider \\&  Rix 1997; Natarajan \\&  Kneib  1997; Brainerd  \\&  Wright 2002;  Guzik  \\&   Seljak 2001). Several  authors have also   used numerical simulations to investigate the link between  g-g lensing and  the galaxy-dark matter  connnection (e.g., Tasitsiomi \\etal 2004; Limousin \\etal 2005; Natarajan, De Lucia \\& Springel 2007;  Hayashi \\& White 2007).   It has become  clear from these studies that   g-g lensing  in  principle contains   a wealth of information  regarding  the  mass distributions  around   galaxies; in addition to  simply probing halo  masses, g-g  lensing also holds  the potential  to measure the  shapes,  concentrations  and radii of  dark matter haloes, and the first  observational results along these  lines have already been obtained (Natarajan \\etal 2002; Hoekstra \\etal 2004; Mandelbaum \\etal 2006; Limousin \\etal 2007; Mandelbaum \\etal 2008).  In this paper we use an analytical model, similar to that developed by Seljak (2000) and Guzik \\&  Seljak (2001), to  predict the g-g lensing signal as a function  of the luminosity of the  lenses starting from a model  for the halo occupation statistics   that is constrained to fit the abundances and  clustering  properties of  the  lens  galaxies.  A comparison of these  predictions with the  data thus allows us to test the mass-to-light ratios inferred from  the halo occupation model, and ultimately to  constrain  cosmological parameters.  The  model assumes that   haloes have  NFW  (Navarro,   Frenk  \\&  White   1997)  density distributions, and   that satellite  galaxies follow  a  radial number density distribution that is unbiased with respect to the dark matter. The occupation statistics are described via the conditional luminosity function   (CLF;  see Yang \\etal 2003),    which specifies the average number of galaxies of given luminosity that reside in  a halo of given mass. This CLF is ideally  suited to model g-g  lensing, as it  allows one  to properly  account  for  the  scatter  in the relation  between luminosity and   halo mass,  and  to  split the  galaxy  population in centrals and satellites.  We   demonstrate how these different  galaxy populations contribute  to the lensing  signal in different luminosity bins, and  show    that    uncertainties related to     the   expected concentrations of   dark matter haloes  and the  radial number density distributions  of satellite galaxies  do not have a significant impact on our results.  Assuming a flat   $\\Lambda$CDM cosmology with parameters  supported by the third   year data release  of  the  Wilkinson Microwave Anisotropy Probe (WMAP, see Spergel \\etal 2007), we  obtain a CLF that accurately fits the abundances and clustering properties of SDSS galaxies.  Using our analytical model, we   show that this   same CLF also   accurately matches  the g-g lensing data obtained  from the SDSS  by Seljak \\etal (2005) and Mandelbaum \\etal  (2006) {\\it without any additional tuning of the  model parameters}. However, if we  repeat  the same exercise for a cosmology with a matter density and power-spectrum normalization that are slightly ($\\sim$ 20 percent)  higher, the model that fits the clustering data can no longer simultaneously fit the g-g lensing data. This confirms that  a joint analysis of  clustering and g-g lensing is an extremely promising method to constrain cosmological parameters. In a companion  paper  (Li \\etal  2008), we  use  the  SDSS galaxy  group catalogue  of Yang \\etal (2007)   to  predict the g-g lensing  signal, which  we compare to data  from the SDSS.    Although, Li \\etal obtain their halo occupation statistics from a galaxy group catalogue, rather than  from the galaxy clustering properties,  they obtain very similar results.  The present  paper is  organized as follows.  We review  the necessary formalism of  g-g lensing  in \\S~\\ref{sec:gglensing}, with  a detailed description of the model used to interpret the g-g lensing signal. The CLF, used to describe the  connection between galaxies and dark matter haloes,  is  introduced in  \\S~\\ref{sec:CLF}.  The  properties of  the predicted  g-g lensing signal  are illustrated  in \\S~\\ref{sec:predGG} together with  a comparison  between theoretical predictions  and SDSS data.  A  detailed analysis of  the assumptions entering the  model is presented  in  \\S~\\ref{sec:moddep}.    Conclusions  are  presented  in \\S~\\ref{sec:conclusions}. ",
        "conclusions": "\\label{sec:conclusions}  Galaxy  clustering  and galaxy-galaxy  lensing  probe the  galaxy-dark matter connection in complementary ways.  Since the clustering of dark matter  haloes depends  on cosmology,  the halo  occupation statistics inferred  from  the observed  clustering  properties  of galaxies  are degenerate  with   the  adopted  cosmology.   Consequently,  different cosmologies  imply  different  mass-to-light  ratios for  dark  matter haloes.  Galaxy-galaxy lensing, on the other hand, yields {\\it direct} constraints on the actual mass-to-light  ratios of dark matter haloes. Combined, clustering  and lensing  therefore offer the  opportunity to constrain cosmological parameters.  Although the  advent of  wide and deep  surveys has resulted  in clear detections of galaxy-galaxy lensing,  a proper interpretation of these data in terms of the link  between galaxies and dark matter haloes has been hampered by the fact that the lensing signal can only be detected when stacking the  signal of many lenses. Since  not all lenses reside in  haloes of the  same mass,  the resulting  signal is  a non-trivial average of the  lensing signal due to haloes  of different masses.  In addition, central  galaxies (those  residing at the  center of  a dark matter halo)  and satellite galaxies (those orbiting  around a central galaxy)  contribute very  different  lensing signals,  even when  they reside in haloes of the same mass (e.g., Yang \\etal 2006). This has to be  properly  accounted  for,  and  requires  knowledge  of  both  the satellite fractions and of  the spatial number density distribution of satellite galaxies within their dark matter haloes.  In this  paper,   we  model  galaxy-galaxy   lensing  with   the  CLF, $\\Phi(L|M)$,    which describes the   average   number of galaxies  of luminosity $L$ that reside in a halo of mass  $M$. This CLF is ideally suited   to  model  galaxy-galaxy  lensing.    In   particular, it  is straightforward to  account for the fact  that there is scatter in the relation between  the luminosity of a  central galaxy and  the mass of its dark  matter halo. This  represents an improvement with respect to previous attempts  to model the  g-g lensing signal obtained  from the SDSS, which  typically ignored this  scatter (e.g.  Seljak \\etal 2005; Mandelbaum \\etal 2006).  However,  in agreement with Tasitsiomi  \\etal (2004), we have demonstrated that the scatter  in this relation has an important impact on the  g-g lensing signal  and cannot be ignored. We also improved upon previous studies  by modelling the 2-halo term (the contribution to  the lensing signal   due   to the mass   distribution outside of the halo hosting the lens galaxy), including an approximate treatment for halo exclusion.  Following  Cooray \\& Milosavljevi\\'c (2005),  we split  the CLF in two components: one for  the central galaxies  and one for the satellites. This facilitates a proper treatment  of their respective contributions to the g-g  lensing  signal.  The  functional  forms for the two   CLF components are motivated by results obtained by Yang \\etal (2008) from a large  galaxy group  catalogue.  For  a  given  cosmology, the  free parameters  of the CLF  are constrained using the luminosity function, the correlation lengths as function of luminosity, and some properties extracted from the group  catalogue.  We  have performed our  analysis for two different $\\Lambda$CDM cosmologies: the WMAP1 cosmology, which has  $\\Omega_{\\rm m}=0.3$ and  $\\sigma_8=0.9$  and the WMAP3 cosmology with    $\\Omega_{\\rm m}=0.238$   and    $\\sigma_8=0.744$.    For  both cosmologies we   have   obtained CLFs that    can   accurately fit the abundances and clustering  properties of SDSS galaxies. However, these CLFs   predict mass-to-light ratios    that are very  different.  This reflects  the  degeneracy  between   cosmology   and halo   occupation statistics alluded to above. In order to break this degeneracy, we use these CLFs to predict the g-g  lensing signal (with no additional free parameters), which  is compared to   the SDSS data  obtained by Seljak \\etal (2005) and Mandelbaum \\etal (2006).  While the WMAP3 CLF predictions are in excellent agreement with the data,  the CLF for the WMAP1 cosmology predicts excess surface densities that are much higher than observed.  Although the  cosmological parameters of the WMAP1 and WMAP3 cosmologies only differ at the 20 percent level, the combination of   clustering  and lensing allows   us  to strongly  favor the WMAP3 cosmology over the WMAP1 cosmology.  In a  companion paper by Li \\etal (2008), we use a completely  different technique to model g-g lensing, but nevertheless reach exactly the same conclusion.  In order  to test the  robustness of our  results we have  performed a number of tests. In particular, we have shown that small uncertainties in the expected concentrations of dark matter haloes, or in the radial number density  distributions of satellite galaxies, only  have a very small impact  on the predicted  lensing signal. In  addition, although our  treatment  of  halo   exclusion  is  only  approximate,  we  have demonstrated that it is sufficiently accurate. Finally, as shown by Li \\etal (2008),  making the  oversimplified assumption that  dark matter haloes  are spherical rather  than ellipsoidal  also has  a negligible impact on the  lensing predictions.  We thus conclude  that our method yields accurate and reliable predictions for g-g lensing.  To  summarize,   as  already  discussed  by  Yoo   \\etal  (2006),  the combination  of  clustering and  lensing  can  be  used to  put  tight constraints on  cosmological parameters. In  this pilot study  we have shown  that current  data  from  the SDSS  strongly  favors the  WMAP3 cosmology  over the  WMAP1 cosmology.   In a  follow-up paper  we will present a more detailed  analysis of the cosmological constraints that can be obtained using this technique."
    },
    "0807/0807.1726_arXiv.txt": {
        "abstract": "We survey observational constraints on the parameter space of inflation and axions and map out two allowed windows: the classic window and the inflationary anthropic window. The cosmology of the latter is particularly interesting; inflationary axion cosmology predicts the existence of isocurvature fluctuations in the CMB, with an amplitude that grows with both the energy scale of inflation and the fraction of dark matter in axions. Statistical arguments favor a substantial value for the latter, and so current bounds on isocurvature fluctuations imply tight constraints on inflation.   For example, an axion Peccei-Quinn scale of $10^{16}$\\,GeV  excludes any inflation model with energy scale $> 3.8\\times10^{14}$\\,GeV ($r> 2\\times10^{-9}$) at 95\\% confidence, and so implies negligible gravitational waves from inflation, but suggests appreciable isocurvature fluctuations. ",
        "introduction": "Early universe inflation and the QCD axion provide explanations for otherwise mysterious features of the universe.  Here we argue that assuming both at once leads to very significant constraints on their central parameters, and to highly falsifiable predictions.  \\subsection{Energy Scale of Inflation}\\label{EnInf}  Inflation is the leading paradigm for early universe phenomenology \\cite{Inf1,Inf2,Inf3}. Its mechanism and the values of its central parameters are unknown, however. One central parameter is the energy scale of inflation $E_I$, defined as the fourth root of the inflationary potential energy density, evaluated when the modes that re-enter the horizon today left the horizon during inflation. $E_I$ is subject to both theoretical and observational constraints, as illustrated in Figs.~\\ref{Naive} \\& \\ref{ConsPlot}.  A multitude of inflation models involving a broad range of energy scales have been discussed in the literature, including chaotic inflation \\cite{Inf4}, brane inflation \\cite{Tye,KKLMMT} and others \\cite{LiddleLythBook}. However, very high $E_I$ has been argued to be theoretically problematic, at least for single field slow-roll inflation, because it involves super-Planckian displacements of the inflaton field $\\phi$ \\cite{LythBound}. From Ref.~\\cite{Baumann} the region in which $\\phi$ moves at least two Planck masses is $E\\gtrsim2.4\\times10^{16}$\\,GeV. The intuition that high $E_I$ is problematic seems borne out in many string theory models; an example is D-brane models \\cite{Baumann}.  Very low $E_I$ has been argued to be theoretically problematic also. Naive consideration of families of potential energy functions suggests that $E_I\\lesssim2\\times 10^{16}$\\,GeV ($r\\lesssim 0.01$) is non-generic \\cite{Steinhardt}. One of the most striking successes of high $E_I$ potentials is that they can naturally predict $n_s\\sim 0.96$, and generic low-energy potentials fail to make this prediction. Low $E_I$ potentials have the slow-roll parameter $\\epsilon$ exponentially small, so that the observation $n_s=1-6\\epsilon+2\\eta=0.960\\pm0.013$ \\cite{WMAP5} implies $\\eta=-0.02\\pm0.0065$, leaving us wondering why $\\eta$ is so small when it could just as well have been of order $-1$. This problem is not alleviated by anthropic considerations~\\cite{TegmarkInfP}. By using the observed value of density fluctuations, and setting $\\epsilon<10^{-4}$ as the boundary, $E\\lesssim6.7\\times 10^{15}$\\,GeV defines this problematic low-scale region. These theoretical issues for inflation are indicated by the vertical regions in Fig.~\\ref{Naive}. Although there are inflation models in the literature at energy scales both above and below this naive window, the debate about whether they are generic continues.  With theory in limbo, we turn to observational guidance. High $E_I$ implies a large amplitude for primordial gravitational waves (GWs). $E_I>3.8\\times10^{16}$\\,GeV ($r>0.22$) is ruled out by WMAP5 plus BAO and SN data \\cite{WMAP5}, as indicated by the orange region of Fig.~\\ref{ConsPlot}.  Possible future searches for primordial GWs have rightly been a focus of attention.  In this article we emphasize the additional information that can be learned from isocurvature fluctuations.    \\subsection{Axion Physics}\\label{AxionCos}  The QCD Lagrangian accommodates a gauge invariant, Lorentz invariant, renormalizable term $\\propto \\theta\\,{\\bf E}^a\\!\\cdot\\!{\\bf B}^a$, with $\\theta\\!\\in\\![-\\pi,\\pi]$, that manifestly breaks P and T symmetry.  Precision bounds on the electric dipole moment of the neutron constrain $|\\theta|\\lesssim10^{-10}$. The striking smallness of this parameter, which the standard model leaves unexplained, defines  the strong P and T problem (a.k.a.~``CP problem\"). After introducing a new asymptotic (or alternatively, classical) Peccei-Quinn (PQ) symmetry \\cite{PQ} which is spontaneously broken, the effective $\\theta$ becomes a dynamical variable, and relaxes toward extremely small values.  The consequent approximate Nambu-Goldstone boson is the axion \\cite{WeinbergAxion,WilczekAxion}.  \\begin{figure} \\includegraphics[width=\\columnwidth]{Fig1.pdf} \\caption{{\\em Naive expectations} for the energy scale of inflation $E_I$ and the axion PQ scale $\\faN$. For $E_I\\gtrsim2.4\\times10^{16}$\\,GeV, the inflaton must undergo super-Planckian excursions in field space (in single field models). For $E_I\\lesssim6.7\\times10^{15}$\\,GeV, generic inflation potentials fail to reproduce the observed nearly scale invariant power spectrum. For $f_a\\gtrsim2.4\\times10^{18}$\\,GeV, the PQ breaking scale is super-Planckian. For $f_a\\lesssim10^{15}$\\,GeV (and $f_a\\gg$\\,TeV), the PQ breaking is in the ``desert\" of particle physics and non-trivial to achieve in string theory. This leaves the region labeled ``naive window\".} \\label{Naive}\\end{figure}  \\begin{figure} \\includegraphics[width=\\columnwidth]{Fig2a.pdf}\\vspace{-1.07cm} \\includegraphics[width=\\columnwidth]{Fig2b.pdf} \\caption{{\\em Observational constraints} on the energy scale of inflation $E_I$ and the axion PQ scale $\\faN$ are shown in the top (bottom) panel for inefficient (efficient) thermalization at the end of inflation. The thick red diagonal line is $\\faN=\\mbox{Max}\\{\\TGH=H_I/2\\pi,\\TRH=\\eff\\,E_I\\}$, ($\\eff\\approx0$ in top, $\\eff=10^{-3}$ in bottom, with $\\eff=10^{-1.5}$, 1 indicated). Above this line is the inflationary anthropic scenario and below this line is the classic scenario. The region in which there is too much isocurvature, $\\alpha_a>0.072$, depends on the axion fraction $R_a\\equiv\\xi_a/\\cdm$ of the CDM; the purple region applies for any $R_a$, the blue region is for $R_a>0.25\\%$ (which is expected at 95\\% confidence), and the cyan region is for $R_a=100\\%$. The green region has too much axion CDM: $\\xi_a>2.9$\\,eV. Each constraint is divided into two parts: the darker part is for a conservative value $\\tc=1/20$ and the lighter part is for a moderate value $\\tc=1$. The orange region has excessive GWs amplitude: $Q_t>9.3\\times 10^{-6}$. The yellow region has too much axion interaction in stars (darker is firmly ruled out, lighter is for some analyzes). The brown region is excluded by the laboratory ADMX search. The dashed cyan, orange, and brown lines are future targets for isocurvature, GWs, and ADMX searches, respectively.} \\label{ConsPlot}\\end{figure}   The simplest axion models contain only one phenomenologically significant parameter: $f_a$, the scale at which the PQ symmetry breaks.   The zero temperature Lagrangian for the complex field $\\phi=\\rho\\,e^{i\\theta}/\\sqrt{2}$ is \\begin{equation} \\mathcal{L}=\\frac{1}{2}f_a^2(\\partial_\\mu\\theta)^2+\\frac{1}{2}(\\partial_\\mu\\rho)^2-2\\Lambda^4\\sin^2(\\theta/2)-\\lambda(|\\phi|^2-f_a^2/2)^2\\nonumber \\end{equation} ($\\Lambda\\approx 78$\\,MeV, $\\rho$ is irrelevant at low energies). Accelerator bounds require $f_a$ to be well above the electroweak scale, and stellar astrophysics constraints place considerably higher limits. Given that electroweak values for $f_a$ are ruled out, economy suggests that $f_a$ could be associated with unification or Planck scales, rather than the \\ ``desert\" of particle physics or super-Planckian scales, as indicated by the horizontal regions in Fig.~\\ref{Naive}.  This intuition for $f_a$ seems borne out in string theory, where $f_a$ typically lies at or just above the GUT scale, and much lower values are non-trivial to achieve \\cite{WittenAxion}.\\footnote{For example, in weakly coupled heterotic string theory, the model-independent axion has its PQ scale given by $\\faN=\\alpha_U\\mpl/(2\\pi\\sqrt{2})$. A unified coupling $\\alpha_U=1/25$ then gives $\\faN\\approx 1.1\\times 10^{16}$\\,GeV \\cite{Fox,WittenAxion}.} Such high values of $f_a$ correspond to large contributions from axions to cold dark matter (CDM).  Indeed, it is only after selection effects are taken into account that the ratio of axion density to entropy is small enough to be consistent with observations  \\cite{LindeSelection, WilczekSelection}.  When these effects are included, one finds that the expected density of dark matter in axions is close to the amount of dark matter actually observed \\cite{TegmarkSelection}.    \\subsection{Cosmological Observables, Summary}  Quantum fluctuations in an effective inflaton field give rise to the standard adiabatic fluctuations that have grown into our cosmologically observed large-scale structure. If the PQ symmetry undergoes spontaneous symmetry breaking before the end of inflation, quantum fluctuations in the consequent light axion field give rise to isocurvature fluctuations. The amplitude of the isocurvature fluctuations grows with $E_I$, so upper bounds on the amplitude of isocurvature fluctuations imply upper bounds on $E_I$.  The purpose of this article is to delineate these bounds, extending earlier work on this subject such as \\cite{Fox,Beltran,Burns,Lyth,LythStewart,TW}. The bounds depend sensitively on the fraction $R_a$ of CDM in the form of axions, which in turn depends not only on $f_a$, but also on the local initial misalignment angle $\\theta_i\\!\\in\\![-\\pi,\\pi]$. These constraints are shown in Figs.~\\ref{ConsPlot}, for different choices of the axion CDM fraction. We will estimate this fraction using statistical arguments.   In Section \\ref{Iso} we calculate the production and late time abundance $\\xi_a$ of axions and the amplitude $\\alpha_a$ of isocurvature fluctuations, as well as reviewing the amplitude $Q_t$ of primordial GWs. These three observables depend on two micro-physical parameters:  the PQ scale $f_a$ (or equivalently, the $T=0$ axion mass $m_a$) and the energy scale of inflation $E_I$ (or equivalently, the Hubble scale of inflation $H_I$), and on one ``environmental'' parameter: the misalignment angle $\\theta_i$. We summarize the final formulae here: \\begin{eqnarray} \\xi_a \\amp = \\amp \\M(\\theta_i^2+\\sigma_\\theta^2)f(\\theta_i^2)\\tc F \\label{abundeqn},\\\\ \\alpha_a \\amp = \\amp \\frac{8}{25}\\frac{(\\M/\\xi_m)^2}{{\\langle(\\delta T/T)^2_{\\mbox{\\tiny{tot}}}\\rangle}} \\sigma_\\theta^2 (2\\theta_i^2+\\sigma_\\theta^2) f(\\theta_i^2)^2  \\tc^2 F^2 \\label{isoeqn},\\\\ Q_t  \\amp = \\amp \\frac{H_I}{5\\pi\\,\\mpl}=\\frac{E_I^2}{5\\sqrt{3}\\,\\pi\\,\\mpl^2}, \\label{graveqn}\\end{eqnarray} where \\begin{eqnarray} F  \\amp\\approx\\amp \\left\\{ \\begin{array}{ll} 2.8\\,\\left(\\frac{\\M}{\\Lqcd}\\right)^{2/3} \\left(\\frac{\\faN}{\\mpl}\\right)^{7/6} \\,\\,\\,\\,\\,\\,\\, \\mbox{for}\\,\\,\\,\\faN \\lesssim \\fh,\\\\ 4.4\\,\\left(\\frac{\\faN}{\\mpl}\\right)^{3/2} \\,\\,\\,\\,\\,\\,\\,\\,\\,\\,\\,\\,\\,\\,\\,\\,\\,\\,\\,\\,\\,\\,\\,\\,\\,\\,\\,\\,\\,\\, \\mbox{for}\\,\\,\\,\\faN \\gtrsim \\fh, \\end{array} \\right. \\\\ \\sigma_{\\theta} \\amp = \\amp \\gamma\\frac{H_I}{2\\pi f_a} = \\gamma\\frac{E_I^2}{2\\sqrt{3}\\,\\pi f_a\\mpl} \\label{defs}\\end{eqnarray} (definitions are given below).  The most recent observational bounds from WMAP5 combined with other data are \\cite{WMAP5} \\begin{equation} \\xi_a\\leq 2.9\\,\\mbox{eV},\\,\\,\\,\\,\\alpha_a<0.072,\\,\\,\\,\\,Q_t<9.3\\times 10^{-6}, \\end{equation} thereby constraining the two micro-physical parameters $E_I$ and $f_a$.  These expressions for $\\xi_a$ and $\\alpha_a$ only apply if the PQ symmetry undergoes spontaneous symmetry breaking before inflation and is not restored thereafter. This is true if $f_a$ exceeds the Gibbons-Hawking temperature during inflation and the maximum post-inflationary thermalization temperature, as we discuss in Section \\ref{Fluct}. Inefficient thermalization leads to constraints displayed in  Fig.~\\ref{ConsPlot} (top), while efficient thermalization leads to constraints displayed in Fig.~\\ref{ConsPlot} (bottom).  In Section \\ref{Dist}, we present a statistical estimate for the axion abundance to place additional constraints on out parameter space. In Section \\ref{Disc}, we conclude by discussing the implications for inflationary model building and future prospects. ",
        "conclusions": "\\label{Disc}  We have surveyed observational constraints on the parameter space $\\{E_I, f_a\\}$ of inflation and axions, finding that most of it is excluded, leaving only two allowed regions that we term classic and anthropic windows. Part of the classic window $\\faN\\sim 10^{11}$ -- $10^{12}$\\,GeV will be intensely explored by the ongoing ADMX experiment. The region indicated by the arrow to the horizontal brown lines in Fig.~\\ref{ConsPlot} to $m_a(0)=10\\,\\mu$eV is expected to be explored by the end of ADMX Phase II, and onwards to $m_a(0)=100\\,\\mu$eV some years thereafter \\cite{ADMXfuture}. In this window, comparatively little can be concluded about the scale of inflation. From Fig.~\\ref{ConsPlot} (top), taking $m_a(0)=100\\,\\mu$eV and assuming $\\xi_a>0.25\\%\\,\\cdm$, we rule out $8.4\\times 10^{13}\\,$GeV $\\lesssim E_I\\lesssim 1.3\\times 10^{15}\\,$GeV.\\footnote{The quoted lower end of the ruled out region is the geometric mean of the conservative and moderate $\\xi_a$-scenarios.\\label{Intermediate}} From Fig.~\\ref{ConsPlot} (bottom), the upper end of this ruled out region is reduced due to efficient post-inflation thermalization. Although we can rule out a range of low scale inflation models, these conclusions are not exceedingly strong.     On the other hand, a large $\\faN$ axion has strong implications for inflation. According to both Figs.~\\ref{ConsPlot} (top) \\& (bottom), if $f_a=10^{16}$\\,GeV then $E_I\\gtrsim5.5\\times 10^{14}$\\,GeV ($r\\gtrsim9\\times10^{-9}$) is ruled out at 95\\% confidence for the conservative value $\\tc=1/20$, and $E_I\\gtrsim2.6\\times 10^{14}$\\,GeV ($r\\gtrsim4\\times10^{-10}$) is ruled out for the moderate value $\\tc=1$. The geometric mean is $E_I\\gtrsim3.8\\times 10^{14}$\\,GeV ($r\\gtrsim2\\times10^{-9}$), which is reported in the abstract. This is incompatible with many models of inflation, including ``classic'' models with a single slow-rolling scalar field in a generic potential. For example, monomial potentials $V \\propto \\phi^p$ predict $r=4\\,p/N_e$, where $N_e$ is the number of e-foldings of inflation from when it generated our horizon scale fluctuations to when it ended. For such models, $N_e$ around 50 or 60 is expected, so any reasonable $p$ is ruled out, including $\\phi^2$ chaotic inflation \\cite{Inf4} and the stringy N-flation \\cite{Nflation} and Monodromy \\cite{Silverstein} models. The same is true for exponential potentials $V \\propto \\exp(-\\sqrt{2p}\\,\\phi/\\mpl)$, which predict $r=16p$. %  Evidently there is considerable tension between the theoretically appealing large $f_a$ and high-scale inflation scenario (see Fig.~\\ref{Naive}) and the observational constraints (see Fig.~\\ref{ConsPlot}). Low-scale inflation may be emerging as favored from recent work in string theory. If we consider the small subspace (see \\cite{Delicate,Hertzberg1,Hertzberg2}) of presently constructed string models that both inflate and agree with the observed values of $Q$ and $n_s$, we are left with models that tend to be at rather small energies, typically $r<10^{-8}$ for D-brane models and various other scenarios such as modular inflation \\cite{Conlon}. There are also arguments for very low $r$ in the simple KKLT framework discussed in Ref.~\\cite{KalloshLinde}. This allows $f_a\\sim 10^{16}$\\,GeV to be marginally consistent with present isocurvature bounds.  Although it is highly premature to conclude that very low energy scale is a generic feature of string realizations of inflation, it is intriguing that many string constructions have this feature. (See \\cite{Nflation,Silverstein} for interesting exceptions.)  The Planck satellite, CMBPol, and upcoming suborbital CMB experiments should probe well beyond the current bound on GWs of $r<0.22$, perhaps reaching $r\\sim 0.01$. This is indicated by an arrow toward the vertical dashed red line in Fig.~\\ref{ConsPlot}. If gravity waves are observed in this regime, then the PQ scale $\\faN$ must be in the classic window.  Our considerations emphasize the fundamental importance of improving  bounds on isocurvature fluctuations. For example, an order of magnitude improvement to $\\alpha\\sim 0.007$ would push the isocurvature bounds to the diagonal dashed cyan line in Fig.~\\ref{ConsPlot}. (We have indicated the improvement for the case  where axions comprise all the CDM: $\\xi_a=\\cdm$.) Detection of isocurvature fluctuations in this regime has three important implications: \\begin{enumerate} \\item It could be interpreted as evidence for the existence of the axion field, and assuming this: \\item It would probe low inflation scales $E_I$ far beyond the scope of any foreseeable GW measurements. \\item It would be evidence that we live in a highly atypical Hubble volume, i.e., $\\{E_I,\\faN\\}$ must be in the anthropic window. \\end{enumerate} Isocurvature modes and tensor modes thus provide complementary constraints on fundamental physics, making it fruitful to study dark matter and inflation in a unified way."
    },
    "0807/0807.1456_arXiv.txt": {
        "abstract": "The paradigmatic Luminous Blue Variable R127 in the Large Magellanic Cloud has been found in the intermediate, peculiar early-B state, and substantially fainter in visual light, signaling the final decline from its major outburst that began between 1978 and 1980.  This transformation was detected in 2008 January, but archival data show that it began between early 2005 and early 2007.  In fact, significant changes from the maximum, peculiar A-type spectrum, which was maintained from 1986 through 1998, had already begun the following year, coinciding with a steep drop in visual light.  We show detailed correspondences between the spectrum and light, in which the decline mimics the rise.  Moreover, these trends are not monotonic but are characterized by multiple spikes and dips, which may provide constraints on the unknown outburst mechanism.  Intensive photometric and spectroscopic monitoring of R127 should now resume, to follow the decline presumably back to the quiescent Ofpe/WN9 state, in order to fully document the remainder of this unique observational opportunity. ",
        "introduction": "The Luminous Blue Variable (LBV) phase is believed to represent a critical transition in the late evolution of all stars with initial masses greater than about 50~M$_{\\sun}$, during which they lose sufficiently large amounts of mass in recurrent explosive events, to subsequently become classical WR stars.  Although the physical mechanism producing these outbursts remains to be definitively identified, it is likely related to the Eddington Limit (Humphreys \\& Davidson 1994, Smith \\& Owocki 2006), and there is ample phenomenological evidence to support the opening statement above.  A fundamental observational characteristic of massive stellar populations is that the most luminous blue stars have no red counterparts. The LBVs lie just to the left of the inclined Humphreys-Davidson (HD; 1979, 1984, 1994; Humphreys 1983) Limit in the HR Diagram, beyond which no stars are observed; that is, they appear to define the locus of an instability that prevents further redward evolution.  During their outbursts, LBVs expand and develop denser, optically thick winds that reprocess the photospheric radiation to lower effective temperatures, such that they exhibit horizontal excursions across the HD Limit, with increasingly cooler spectra and brighter visual magnitudes while maintaining constant bolometric luminosities.  Following an outburst, they invert these displacements and return to the hotter, minimum spectrum with fainter visual magnitude.  Most LBVs and numerous related objects now interpreted as quiescent LBVs have nitrogen-rich circumstellar nebulae, produced by prior outbursts thousands of years earlier (Walborn 1982; Stahl \\& Wolf 1986a; Stahl 1987; Nota et~al.\\ 1992, 1995, 1996; Clampin et~al.\\ 1993; Davidson et~al.\\ 1993; Smith, Crowther, \\& Prinja 1994; Smith et~al.\\ 1997, 1998; Pasquali, Nota, \\& Clampin 1999; Lamers et~al.\\ 2001).  $\\eta$~Carinae is the most famous and extreme LBV (Davidson \\& Humphreys 1997), but it is located in a giant H~II region associated with some of the most massive stars known, well above 100~M$_{\\sun}$ (Walborn et~al.\\ 2002). Given the uncertainty about the LBV mechanism, it is likewise unclear whether or not $\\eta$~Car represents a more spectacular version of typical LBVs, which are located in more evolved regions and have initial masses between 50 and 100~M$_{\\sun}$.  Well observed examples of the latter are AG~Carinae (Wolf \\& Stahl 1982, Stahl 1986, Hutsem\\'ekers \\& Kohoutek 1988, Walborn \\& Fitzpatrick 2000, Stahl et~al.\\ 2001); and R71 (Wolf, Appenzeller, \\& Stahl 1981) and S~Doradus (Massey 2000, and references therein) in the Large Magellanic Cloud (LMC).  The known incidence of these objects in other Local Group galaxies has been greatly extended by Massey et~al.\\ (2007).  Radcliffe (R)127 (Feast, Thackeray, \\& Wesselink 1960) = HDE~269858\\footnote{Feast et~al.\\ denoted it as HDE~269858f for ``following'' because of the companion 3$\\arcsec$ (north)west, now measured at $y=12.12$ and classified as B0~Ia (N weak) by Heydari-Malayeri et~al.\\ (2003).  It is important to note that the HDE-based photographic magnitude of 11.4 given by Feast et~al.\\ must refer to the combined light of the pair, whereas their description of it as a ``double'' likely implies comparable magnitudes, consistently with the foregoing values.  Most or all of the photometry discussed in this paper includes both stars; accordingly, all values presented have been corrected for the companion, which is significant at the fainter LBV magnitudes.} = Sanduleak (1970, Sk)~$-69^{\\circ}$~220 is arguably the most important typical LBV. In fact, taking into account this report, it is not an overstatement that it is comparable to $\\eta$~Car and even SN~1987A, with regard to what it may teach us about late massive stellar evolution.  Its highly peculiar spectrum was classified as O~Iafpe by Walborn (1977; observation date 1975 Dec), later reformulated to Ofpe/WN9 by Walborn (1982; echelle observation 1978 Nov).  A number of similar objects were found in the LMC (Bohannan \\& Walborn 1989), which have subsequently been reclassified as WN9-11 types by Crowther \\& Smith (1997).  In addition to the normal Of selective emission lines of He~II $\\lambda$4686 and N~III $\\lambda\\lambda$4634-4640-4642, these spectra display other emission features and P~Cygni profiles including Balmer and He~I lines, together with very narrow N~III and Si~IV absorptions indicative of unusually extended atmospheres.  In 1982, R127 was discovered to have entered a classical LBV outburst (Stahl et~al.\\ 1983, Walborn 1983), thus establishing the Ofpe/WN9 category as a quiescent LBV state.  Several other members of the category were subsequently discovered to possess LBV circumstellar nebulae (references above), supporting the foregoing generalization.  R127 is located in a compact subgroup of a small, evolved cluster (NGC~2055) containing several mid-O through early-B supergiants and devoid of nebulosity (Walborn et~al.\\ 1991; Heydari-Malayeri, Meynadier, \\& Walborn 2003), south of 30~Doradus and not far from SN~1987A.  This paper reports that R127 has likely now begun its final decline from the current outburst; it is thus the largest and most extensive LBV event that will have been covered from start to finish with modern instrumentation. ",
        "conclusions": "The R127 outburst through 1986 was extensively documented by Stahl et~al.\\ (1983), Stahl \\& Wolf (1986b), and Wolf et~al.\\ (1988).  As they discuss, the spectroscopic observation by Walborn (1982) and subsequent photometry limit the onset of the outburst to the interval 1978--1980.  Most of the spectroscopic data in the early papers were photographic; only the digital data are now available in electronic form and are shown here; thus it is essential to refer to the figures in those papers for the complete early record, and also for detailed line identifications in the A-type spectra.  The early outburst photometry twice appeared to show declines from maxima, in 1982 and 1984, but these are now seen to be only inflections during the rise to the true maximum (Fig.~5), when R127 became the brightest star in the LMC.  The first outburst spectroscopic observation, in 1982 January, showed a peculiar early B-type spectrum, similar but not identical to the current state (Stahl et~al.\\ 1983).  Subsequently (Stahl \\& Wolf 1986b), the spectrum changed to $\\sim$B7 (based upon near equality of He~I $\\lambda$4471 and Mg~II $\\lambda$4481; Walborn \\& Fitzpatrick 1990) in 1983 February during the decline from the 1982 lightcurve peak; then to $\\sim$B9 (shown in Fig.~2) coinciding with the peak of 1984 August; and back again to $\\sim$B7 in 1985 January during the decline from the 1984 peak. Then, as the visual light continued to increase, the full A-type spectrum developed in 1986 (Wolf et~al.\\ 1988) and was maintained through 1998 (Figs.~2--3).  These relationships between the spectrum and lightcurve encourage a detailed comparison between the spectral types and apparent magnitudes, which is given in Table~2.  It is seen that the correlation is indeed very well defined, and that it applies equally to both the rise and decline, insofar as they have been covered.  This correlation is, of course, a result of the ``constant'' bolometric luminosity during the outburst, as shown in the table.  The scatter in the derived bolometric magnitudes is not surprising, given the different sources of the corrections and the fact that they are computed for normal supergiants, as well as the uncertainties in the classification of these peculiar spectra.  Even more detailed effects occur, e.g., the apparent doubling of the Mg~II line during 1991 (Fig.~2), when the small decline from the maximum hump (with an upward spike?) to an extended plateau took place (Fig.~5).  It is quite possible that other small differences among the spectra are related to spectroscopically unresolved fine structure in the lightcurve.  A larger spectral change is evidently related to the steep decline in light beginning in 1999 (Figs.~3--4).  There is extensive spectroscopic coverage during 2002, not shown here, which may elucidate the strong spike in light during that year; indeed, the Mg~II line again doubled during September, near the lightcurve peak.  The 2002 observation shown might again be classified as $\\sim$B7, from the He~I/Mg~II absorption ratio, albeit with residual A-type emission lines.  A striking feature of Figure~3 is the apparent monotonic broadening of the A-type lines as they fade, which is caused by the gradual dominance of [Fe~II] further out in the wind over Fe~II at nearly the same wavelengths, as discussed by Stahl et~al.\\ (2001) in AG~Car.  In contrast, the He~I $\\lambda$4471 emission line was very narrow when it first appeared in 2005. The recent spectra also display changing structure in the Balmer P~Cyg absorption components, not shown here, analogous to what was observed during the R127 rise and in AG~Car.  All of these features in R127 will be measured and discussed in subsequent publications.  It is now vital to increase the frequency of observations during the coming years, to document R127's final decline to the minimum, quiescent state more thoroughly.  In particular, it will be of some interest to catch the reappearance of the He~II and N~III Of emission lines--provided that an infrared catastrophe as in $\\eta$~Car does not obscure the star entirely. It is possible that the high-ionization lines may already be very weakly present in the 2008 February and May observations."
    },
    "0807/0807.3386_arXiv.txt": {
        "abstract": "{We report the timing-analysis results obtained for RRAT J1819-1458 from regular timing observations at 1.54 GHz using the Urumqi 25 m radio telescope between 2007 April to 2008 March. RRAT J1819-1458 is a relatively young and highly magnetized neutron star discovered by its sporadic short bursts in the Parkes Multibeam Pulsar Survey data. In 94 hrs of observation data, we detected a total of 162 dispersed bursts of RRAT J1819-1458 with the signal-to-noise ratios (S/Ns) above 5-$\\sigma$ threshold. Among them, 5 bursts clearly show two-component structure. The S/N of the strongest burst is 13.3. The source's DM measured through our data is 196.0 $\\pm$0.4 pc cm$^{-3}$. The timing position, frequency and its first derivative were determined using standard pulsar timing techniques on the arrival times of these individual bursts. The accuracy of the solved rotating parameters are improved comparing with that in previous publication. Our timing position with 2-$\\sigma$ error is consistent with the position of its X-ray counterpart CXOU J181934.1-145804. The effect of timing noise and the phase fluctuation of the individual short bursts on the timing residuals is briefly discussed. The distribution of the timing residuals is bimodal, which cannot be explained readily by timing irregularity.} ",
        "introduction": "Recently, a remarkable new class of radio transient sources, the `` Rotating Radio Transients''(RRATs) was discovered in a search for isolated radio bursts in the Parkes Multibeam Pulsar Survey data \\citep{b15}. A total of 11 RRATs have been detected so far; they are characterized by short radio bursts of a typical from 2 to 30 ms duration, and the average time intervals between bursts range from 3 min to 3 hrs. Although the radio bursts of RRATs are sporadic, the timing analysis of the bursts of 10 RRATs indicates that they are likely to be rotating neutron stars with the periods ranging from 0.4 to 7 sec \\citep{b15}.  RRAT J1819-1458 is the brightest and the most prolific radio-burster of all 11 RRATs. At 1.4 GHz, the source is characterized by dispersed radio bursts of average duration 3 ms with one burst detected about every 3 min; its dispersion measure (DM) is 196$\\pm$3 pc cm$^{-3}$; and the peak flux of the brightest burst detected so far is 3.6 Jy. The distance inferred from its DM and position using the free electron density model of \\citet{b2} is about 3.6 kpc. Although the source is not detectable in standard periodicity searches, the analysis of the spacings between the bursts reveals a spin period, $P$ = 4.263 s, and a spin period derivative, $\\dot P$ = $576\\times 10^{-15}$ s s$^{-1}$. RRAT J1819-1458 has a characteristic age of 117 kyr, and a relatively high inferred magnetic field strength of $5\\times 10^{13} $ gauss \\citep{b15}.  On the $P$-$\\dot P$ diagram, the source is located in the same area occupied by the high magnetic field radio pulsars, which suggests a possible association between RRATs and the magnetars \\citep{b15, b24} or the X-ray dim isolated neutron stars \\citep{b16,b5}. The $Chandra$ observation detected the X-ray counterpart of RRAT J1819-1458 \\citep{b17}. Lately, using the $XMM-Newton$, \\citet{b14} discovered the X-ray pulsations with the period predicted by the timing of radio bursts. The X-ray properties show that RRAT J1819-1458 is a cooling neutron star, and unlikely to be a magnetor \\citep{b17,b16,b14,b19}.   Based on a study of subpulse modulation of pulsars \\citep{b22}, \\citet{b23} note that PSR B0656+14 has very energetic and sporadic radio pulses with the same characteristic as RRATs, and suggest that it could have been identified as a RRAT, were it not so nearby. They point out that these energetic bursts are shown to be unlike giant pulses \\citep{b6,b10}, giant micropulses \\citep{b9,b8} or the pulses of normal pulsars \\citep{b18,b11}. \\citet{b21} note RRATs are not directly related to nulling pulsars which are characterized by turning on and off of the pulsations. Although several models have been suggested to explain RRAT phenomenon \\citep{b25,b12,b4,b13}, the nature of emission of RRATs is as yet unclear.  We observed RRAT J1819-1458 at 1540 MHz by using the Urumqi 25m radio telescope from 10 April 2007 to 29 March 2008. In this paper, we present the timing result of RRAT J1819-1458 based on these observations. The details of the observations are presented in Section 2. The process of burst detection is given in Section 3. The timing process and results are described in Section 4. A brief discussion is given in Section 5. In Section 6 we summarize this study. ",
        "conclusions": "In order to estimate a possible contribution of stochastic timing noise to the timing residuals, we calculated the cumulative phase contributed by timing noise over time interval of our observations ($\\sim$ 355 days). If we express the rotation frequency as a Taylor series, the pulse phase at time $t$ is given by \\begin{equation} \\phi= \\phi_0+\\nu t+\\frac{1}{2} \\dot \\nu t^{2}+\\frac{1}{6}\\ddot \\nu t^3  \\cdots \\label{eq:LebsequeI} \\end{equation} where $\\phi_0$ is the phase in cycles at time $t=0$. The fourth term in this equation can be used to estimate the cumulative phase contributed by stochastic timing noise over time interval $t$. We therefore carried out a timing solution with the second frequency derivative $\\ddot \\nu$ included, and detected no significant decrease of the rms residual. Because $|\\ddot \\nu|<2\\sigma_{\\ddot \\nu}$ in our case, following \\citet{b1}, we quoted an upper limit of the cumulative phase ($\\delta_t$) as \\begin{equation} \\delta_t<\\frac{1}{6\\nu} (2\\sigma_{\\ddot \\nu})t^3 \\label{eq:LebsequeI} \\end{equation} where $\\nu$ = 0.23456 is the rotating frequency, $\\sigma_{\\ddot \\nu}$ = 1.19$\\times $10$^{-24}$ is the formal uncertainty of the second frequency derivative, and $t\\sim$ 30672000 s is the time interval over our observations. We obtained the cumulative phase, $\\delta_t<$ 50 ms, in a one year time span.  Although the cumulative phase reaches 50 ms in a one year duration, the result of a fitting for $\\nu$ and $\\dot \\nu$ to a simple rotating model with the $\\ddot \\nu$ included shows that the rms residual is less than 2 ms. The rms of timing residuals in Table 1 is much larger than that contributed by the stochastic timing noise expected for this source over one year. If we take a typical duty circle value of 3\\% of normal pulsars, the burst-phase span of a 4.26 sec pulsar is expected to be about 130 ms. For these reasons, we may note that the residuals in Fig. 3 are strongly dominated by random fluctuations in the burst phases, and can be good representations of the phases of the individual bursts in the source's radiation window.  The bimodal distribution of the residuals (as shown in Fig. 5) cannot be explained readily by timing irregularity, and may indicate the possible existence of a more sporadic burst activity in the longitudinal region preceding the main radiation region. However, the number of detected bursts is not yet sufficient to make this a firm conclusion. More observations and an accurate phase alignment are needed to obtain a reliable luminosity distribution of the radiation window.  As shown in Table 1, RRAT J1819-1458 is a relatively young pulsar with a long spin period and a high $B_s$. In fact, there are only 16 pulsars in ANTF Pulsar Catalog with the measured $B_s$ values higher than that of this source. It has a modest $\\dot E$ and a low $B_{Lc}$ in the known pulsar population. The unusual radio emission and the detection of X-ray pulsation of RRAT J1819-1458 make it a very interesting object. Its sporadic radio activity may has similar mechanism with one of the two known radio-emission phenomena, i.e. giant-pulse emission in some pulsars \\citep{b6, b10} or the very sporadic strong pulses detected in PSR B0656+14 \\citep{b23}. The known giant-pulse emitters have very high values of $B_{Lc}$ ($> 10^{5}$ gauss) and $\\dot E$ ($\\sim 10^{36 - 38}$) which are suggested to be indicators of giant-pulse emissivity \\citep{b27, b30, b10}. Therefor, the low $B_{Lc}$ and $\\dot E$ of RRAT J1819-1458 is against the suggestion of giant-pulse origin of its radio bursts \\citep{b15, b14}. The characteristic ages of PSR B0656+14 and RRAT J1819-1458 are similar, and the X-ray pulsation characteristics of RRAT J1819-1458 seems to be similar with that of PSR B0656+14 \\citep{b14}. However, we should note that the two sources have very different spin period, $B_s$ and $\\dot E$. The anomalous X-ray pulsars (AXPs) have very high values of $B_s$. Therefor, the transient radio emission detected in AXPs XTE J1810-197 and 1E 1547.0-5408 \\citep{b28, b29} may also has a similar mechanism with the radio emission of this high-$B_s$ RRAT. However, the radio emission characteristics of the two AXPs are quite different from that of RRAT J1819-1458 \\citep{b15}. By studying the X-ray properties, \\citet{b14} note that RRAT J1819-1458 could be a transition object between the pulsar and magnetar. More studies are needed to determine the reason of the source's unusual radio behavior. A study of individual-burst properties of RRAT J1819-1458 mainly based on the bursts presented in this paper will be reported in a subsequent paper.    We have carried out timing observations of RRAT J1819-1458 by using the Urumqi 25 m radio telescope at a center frequency of 1.54 GHz. A total of 47 observtions were made in 25 sessions from 10 April 2007 to 29 March 2008. In total 162 strong bursts, with S/N $\\geq$ 5, were detected through a careful burst-detection processing. 5 bimodal bursts, with the component separations ranging from 6 ms to 16.5 ms, were noted. We obtained a more precise DM = 196.0 $\\pm$0.4 pc cm$^{-3}$ from our data. We have presented the improved timing position, $\\nu$ and $\\dot \\nu$ for RRAT J1819-1458 based on these bursts using standard pulsar timing techniques on the arrival times of individual bursts. The location of the X-ray counterpart (CXOU J181934.1-145804) is now within 2-$\\sigma$ error ellipse of our timing position.  No evidence of significant glitch and timing noise were found in the one year of rotating history of this source, likely due to the phase modulation of the sparse individual pulses of bursts within its radiation windows. Continued monitoring is needed to investigate its rotating irregularity. We found a bimodal distribution of the timing residuals, which may suggest a possible two-component structure of the source's radiation window, with a more sporadic component leading the main one. However, more observations are also necessary to confirm this conclusion."
    },
    "0807/0807.1614.txt": {
        "abstract": "{}{}{}{}{} % 5 {} token are mandatory  \\abstract % context heading (optional) % {} leave it empty if necessary {The study of the thermodynamical and chemical properties of the Intra Cluster Medium (ICM) in high redshift clusters of galaxies is a powerful tool to investigate the formation and evolution of large scale structures.  Here we discuss the X--ray properties of clusters of galaxies optically selected in the Red--Sequence Cluster Survey (RCS) observed with the \\chandra\\ satellite, at redshifts $0.6 < z < 1.2$.} % aims heading (mandatory) {We intended to assess the evolutionary stage of optically selected high--z clusters of galaxies, performing a spectral analysis of the diffuse emission from their ICM.  We also investigated the distribution of AGN in their surroundings.} % methods  heading (mandatory) {The background subtracted spectra were analyzed and fitted with a single temperature model to measure average ICM temperature, X--ray bolometric luminosity and Fe abundance within typical radii between 200 and 350 kpc.  We also analyzed the point source number density and spatial distribution in the RCS clusters fields as a function of the X--ray flux.} % results heading (mandatory) {We detected emission for the majority of the clusters, except for three, for which we have only marginal detection at $\\sim3\\sigma$.  We find that the normalization of the Luminosity--Temperature relation for RCS clusters is a factor of $\\sim2$ lower than the one for X--ray selected clusters. We confirm that the Fe abundance in the detected objects is consistent with that of X--ray selected clusters at the same redshift.  We also found an excess of low--luminosity AGN towards the center of the clusters. } %  conclusions heading (optional), leave it empty if necessary {At $z\\sim 1$, optically selected clusters with a well--defined red--sequence show extended X--ray emission in about 70\\% (8/11) of the cases. Nevertheless, their $\\lxtx$ relation is significantly lower with respect to X--ray selected clusters at the same redshift, possibly indicating an incomplete virialization. The Fe abundance measured in the ICM of RCS clusters is comparable to the values measured for X--ray selected clusters at the same redshift, implying a substantial enrichment by SNe products.  These two evidences add to the previously known studies of high--z, X--ray selected clusters, to confirm that the thermodynamical and chemical properties of the ICM are substantially established already at $z > 1$. Finally, we find the significant excess of medium and low luminosity AGN close to the centroid of the X--ray emission. Their X--ray emission is not dominating the ICM, but their presence may be relevant for studying the interaction between AGN and ICM.} ",
        "introduction": "X--ray selection has been shown to be very effective in finding a very well--defined sample of clusters of galaxies, with little contamination and a well--defined selection function.  These properties make X--ray selection a powerful tool for cosmology \\citep[see][]{rbn02}.  On the other hand, optical selection, based in particular on the presence of a well--defined red--sequence, is very effective (and much cheaper in terms of observing strategies) in finding a large number of candidates.  Recently, this technique has been pushed to high redshift, as in the case of the Red--Sequence Cluster Survey (RCS), which uses the color--magnitude relation of early--type galaxies as cluster finding method \\citep{glad00}.  The RCS is a 100 square degree imaging survey using mosaic CCD cameras on 4m-class telescopes, and is designed specifically to provide a sample of optically selected $0.2<z<1.2$ clusters. The red--sequence method is motivated by the observation that all rich clusters have a population of early--type galaxies which follow a well--defined color--magnitude relation \\citep{yee01}.  The red--sequence represents the nominal reddest galaxy population in any group of galaxies at the same redshift, producing a well--defined signature in the color--magnitude diagram.  So the criterion to find cluster candidates is based on identifying galaxy over--densities in the four--dimension space given by color, magnitude and sky coordinate.  For more details about the implementation of this method and the survey design see \\citet{glad00,glad01,glad05}, \\citet{yee01} and \\citet{yee91}.  On the other hand, the X--ray properties of optically selected clusters are usually different from that of X--ray selected clusters \\citep[see, e.g.][]{D02}.  In particular, lack of X--ray emission in presence of a well--defined red sequence, can be ascribed to uncomplete virialization.  In this respect, a systematic study of the X--ray properties of an optically selected sample of clusters can provide key information in order to assess both the reliability of an optical sample for cosmological studies and the dynamical and thermodynamical status of the ICM in these clusters.  In this Paper we focus on the X--ray properties of clusters of galaxies in the redshift range $0.6 < z < 1.2$ in the Red--Sequence Cluster Survey, observed with the \\chandra\\ satellite.  High redshift clusters are relevant both for cosmology and for studies of large scale structure formation and evolution.  First, we investigate the thermodynamical properties of the ICM, which, given the low number of photons detected for each object, are described with a single--temperature model.  Then we also investigate the distribution of X--ray point sources in the field of RCS clusters. This Paper presents an independent and complementary analysis of X--ray data of the RCS clusters with respect to the Paper of \\citet{hick08}. The Paper is organized as follows.  In \\S 2 we describe data reduction and analysis.  In \\S 3 we present our main results on the thermodynamical and chemical properties of the ICM of RCS clusters.  In \\S 4 we present our results on the AGN distribution in their fields.  Our conclusions are summarized in \\S 5. ",
        "conclusions": "We studied the thermodynamical and chemical X--ray properties of the ICM of a galaxy cluster sample of the Red--Sequence Cluster Survey, in the redshift range \\hbox{$0.6<z<1.2$}. We detected emission for the majority of the clusters, except for three, for which we have only marginal detection at $\\sim3\\sigma$.  In general  we found that the slope of the $\\lxtx$ relation of the RCS clusters is in agreement with that of X--ray selected clusters, while the normalization is a factor of 2 lower at high confidence level, in agreement with the result of \\citet{hick08} in their independent analysis. Only the three marginally detected RCS clusters seem to have lower luminosities with respect to the RCS $\\lxtx$, but only if they have $kT>3$ keV. Unfortunately, for this sample the statistic is too poor to draw any conclusion about possible evolution of the $L_{\\mathrm X}-T_{\\mathrm X}$ relation with redshift. Particularly our data are fitted with $L_{\\mathrm X}/T_{\\mathrm X}^{\\alpha} \\propto (1+z)^{0.2 \\pm 0.2}$, consistent with no evolution.  Concerning the Fe abundance in RCS clusters we found that also for this sample of optically selected clusters, the ICM was already enriched with metals at a level comparable with X--ray selected clusters at high redshift \\citep{bale07}.  Thanks to the high spatial resolution of the \\chandra\\ satellite, we also investigate the point source distribution near the region of diffuse cluster emission for RCS with respect to the field.  The number counts of point sources as a function of the flux show a significant excess at all fluxes with respect to the data of the CDFS.  We quantified this excess between 15\\% and 40\\% with a significance of $2\\sigma$.  This result is in agreement with that found by other authors \\citep{capp01,bran07a}.  The spatial distribution of point sources in the field of RCS clusters shows a factor of $\\sim6$ over--density at $3\\sigma$ confidence level in the inner 150 kpc.  The contribution of these point sources to the X--ray emission is limited to few percent in the soft band, showing that the contamination from AGN in the soft band is not severe at least for cluster with $L_X > 10^{44}$ erg s$^{-1}$ up to $z\\sim 1$."
    },
    "0807/0807.5105_arXiv.txt": {
        "abstract": "The dynamics of planetesimals in the circumbinary debris disc of the quadruple star system HD 98800 are investigated. Evolving a spherical shell of test particles from a million years ago to the present day indicates that both coplanar and retrograde warped discs could exist, as well as a high inclination halo of material. Significant gaps are seen in the discs, as well as unexpected regions of stability due to the retrograde nature of the stellar orbits. Despite a viewing angle almost perpendicular to the direction of the warp of the planetesimal disc it is still intersected by the line of sight for eccentricities of the outer orbit of 0.5 or less. ",
        "introduction": "\\label{sec:intro}  HD 98800 is an interesting and unusual system of four 10 Myr old post T-Tauri K stars -- two spectroscopic binaries A and B in orbit about one another -- located in the TW Hydrae association (\\citealt{T95}, \\citealt{Ka97}). It has a large infrared excess attributed to a circumbinary disc around the B pair \\citep{K00, P01, Fu07}. Substantial extinction towards this pair is evidence that they are observed through some of this material (\\citealt{S98}, \\citealt{T99}). A photometric variability is also seen, but with no definite period \\citep{S98}. Both the apparent absence of CO molecular gas in the disc \\citep{De05} and infrared spectrum modelling indicate that this is a T-Tauri transition disc that is just reaching the debris disc stage, with a collisional cascade having been recently initiated \\citep{Fu07, W07}. The orbits of stars are all highly eccentric and inclined, creating a dynamical environment unlike almost all other known debris discs \\citep{T99,P01}.  The dust disc here is generally agreed to be an annulus around the B binary, but the exact structure varies from model to model. \\citet{K00} estimate a coplanar narrow ring outwards of about $5.0\\pm 2.5$ au from the two stars, themselves separated by approximately 1 au. \\citet{P01} however determine a 1 au high coplanar ring now from 2 and 5 au. \\citet{B05} argue that the line of sight extinction means that the disc is not coplanar with the sub-binary unless it is very flared. \\citet{Fu07} suggest an inner optically thin ring at 2 au and an outer puffed up optically thick wall 0.75 au high at 5.9 au, with a gap between the two. Recently, \\citet{Ak07} showed that a single continuous physically and optically thick coplanar disc between 3 and 10 au can reproduce the observed spectral energy distribution. To explain the extinction, they used dynamical models of test particles to show that the inclined stellar orbits could create a warp in the dust disc, the outer layers of which could then just intercept the line of sight.  The unusually large infrared excess has been argued by \\citet{LBA00} to be caused by one of four reasons. The first is that a dust avalance is currently in progress. Radiation pressure acts to push dust grains outwards. When the disc is very dusty, these can impact other dust grains, creating more particles that are themselves pushed outwards, colliding and creating yet more dust, and so on. The second possibility is there is undetected gas maintaining the dust population against radiation pressure. The third and fourth reasons are applicable when the wide orbit has an eccentricity of near unity, much higher than currently believed. In this case, it is possible that either an encounter with the A pair has heated the outer edges of the disc, resulting in abnormally high infrared radiation, or disrupted a Kupier Belt-like stucture, resulting in collisions and releasing large amounts of dust. As the eccentricity is now known to be much lower, these two possibilities are less likely, although it should be noted that the wide orbit is currently very near to periastron.  The first two possibilities present difficulties in modelling the system. If the system is undergoing a dust avalance or is gas-rich, then radiation pressure, collisions and gas dynamics must be included. For example, \\citet{W07} has calculated the dust collision timescale to very short (0.36 years), so collisional effects are important in modelling the dust evolution.  However, in debris discs, dust is generated from an underlying planetesimal population, usually through a collisional cascade. In this case, the dust population follows the planetesimal population, which is less complex to model. If the disc has large amounts of gas or a dust avalance is occuring, this may, of course, no longer be the case. Models of the planetesimal population remain interesting though, especially in such an unusual stellar environment. Indeed, the high inclination and eccentricities suggest that little stability is likely, yet some must exist. Constraining possible planetesimal locations and then comparing with observations of the dust location may even be able to distinguish between the possible scenarios above. They would also serve as an indication of whether any planetary stability might be possible in this system.  There is one other debris disc known in a similar stellar system to HD 98800. This is GG Tau, which has a circumbinary debris disc of several hundred au radius in a quadruple star system with the same hierarchy \\citep{Gu99}. However, in this case the stars are about ten times more distant, the mass ratios much smaller and their orbits relatively coplanar so more stability would be expected~\\citep{Be05,Be06}. Dynamical modelling has shown that the circumbinary material forms a sharp ring and a more diffuse disc component.  HD 98800 remains an unusual environment, and the high inclination and eccentricity of the stellar orbits will have a significant effect on the dynamics and structure of the planetesimal population of the debris disc.  This effect has yet to be studied in detail, and is investigated here using direct numerical integrations. Section \\ref{sec:method} gives a description of the stellar system and the simulation parameters. The results are then presented in Section \\ref{sec:results} and conclusions given in Section \\ref{sec:conclusion}. ",
        "conclusions": "\\label{sec:conclusion}  Dynamical simulations of a planetesimal population in the debris disc around the Bab stellar pair in HD 98800 have been run. By studying a wide range of inclinations, three distinct stable populations have been identified. These are a prograde disc, a retrograde disc and a high inclination halo. The radial profiles of each component are different and distinct. The discs both have large radial gaps, caused presumably by overlapping resonances from the stellar orbits. The radial extent of the discs are summarised in Table \\ref{tab:rlimits} but are generally from 3 to 15 au for all three orbits, with gaps at around 4.5 au and 8 au.  The line of sight can currently pass through slightly warped material in the prograde disc, and this would account for the observed extinction and variability if the dust distribution followed the planetesimals.  However, this alignment with the line of sight effectively only occurs for the I and II orbits (with eccentricities of $0.3$ and $0.5$ respectively), which would rule out the higher eccentricity case ($e = 0.6$) as a possible orbit for the outer star.  The radial profile is complex, and illustrates that, for discs in multiple systems, dynamical effects purely due to the stars are very important. Indeed, if most triple stars are within the limit for resonance overlap to sculpt the stability zone, then this has large consequences for debris discs and planet formation in such systems.  The simulation results here can be compared to other models of the system and estimates of the dust distribution. Indeed, as previously mentioned, comparisons to observations of the dust could place constraints on the physical processes occuring in the circumbinary disc. The bulk of the dynamically stable planetesimals are between 5 and 7 au. This matches up well with the prediction of \\citet{K00} of a ring outwards from 5 au. The model of \\citet{P01} places the disc from 2 to 5 au, which, apart from the ring around 4.5 au, is unstable here. \\citet{Fu07} suggest an inner ring at 2 au and a thicker puffed up component at 5.9 au, the only model to predict gaps and similar to the prograde discs here. The later of these components matches up to the bulk of the material here, and there is indeed an inner ring seen, just further out at 4 au.  \\citet{P01} estimate the height of the disc as 1 au, and \\citet{Fu07} as 0.75 au. However, here material in the prograde planetesimal disc is not only warped but very flared, reaching heights of 5 to 10 au. This is an important dynamical feature that should be taken into account in models of the observations. Although roughly agreeing in location, these models do not have enough detail as yet to further compare to the planetesimals.  The dynamical model of \\citet{Ak07} finds stability from 3 to 10 au, which does not match up that well to the limits here, and in fact their resulting disc has a very different geometry. Their model uses different orbital parameters for the stellar system, most notably a different mutual inclination -- as their aim is not to reproduce the current configuration but to look at the general morphological structure of debris discs in highly inclined triple star systems. They also find that the line of sight can intercept a warp in the disc, but due to the different geometry it crosses at a different point and at the maximum warp in the disc, whereas here the warp is only marginally orientated towards the line of sight. It is in fact material in the sparser populated outer regions of the prograde disc that intercepts the line of sight here, from 8 to 15 au, which is a region that in a coplanar stellar system is predicted to be unstable.  The possible dynamical reason for the high infrared excess already mentioned in the Introduction is a close pericentre passage of the A binary stiring up planetesimals and resulting in high collision rates. Since the binary's period is about 300 years, many such pericentre passages have occured, and by the end of the simulations the particles appear very stable with regards to this. For example, there appears to be no periodic increase in planetesimal eccentricities as the outer star passes close to the disc. There is, however, another feature that may cause higher than normal collision rates. The extent of the flare and warp in the extended disc decreases as the acute angle between the stellar orbits decreases (and the mutual inclination increases). As this occurs, planetesimals that have been spread out in inclination are now packed into less space. This is likely to result in an increased number of collisions and is particularly relevant as the current stellar configuration is such that mutual inclination is large and the disc is near its narrowest point. However, modelling including collisions is needed to quantify this effect.  The stable high inclination particles in the halo are curious from a dynamical perspective. The Kozai mechanism \\citep{Ko62} might be expected to remove such particles very quickly -- however these particles do not seem subject to this. It is possible that they are in fact an artifact of the numerical method, but checks with a standard Bulirsch-Stoer integrator \\citep{P89} and the similar observation in simulations of GG Tau discount this. The high stellar inclinations involved in the system may be one explanation, but it is worth looking at the stability should one of the binaries be removed. If the outer star is not present, a sharp inner disc edge is seen at around 4 au for all inclinations, with all particles outwards from this remaining stable. If the inner binary is approximated as a single star instead, then a reasonably sharp outer edge is seen at around 7 to 10 au, for the lower inclinations only. The near polar inclinations in this case are now unstable. The high inclination particles are within this region which must be shaped by the combined perturbations from both stellar orbits, so are in a new regime of dynamical behaviour. The mechanism stabilising these particles, as well as that stabilising the extended prograde and retrograde discs, is the subject of a future paper.  Some consideration is needed, however, of how such material would form. T Tauri multiple systems are common and believed to form primordially through fragmentation processes, which could result in non-coplanar discs (\\citealt{Ba00}, \\citealt{BB00}). As mentioned, \\citet{Mo06} find a small number of multiple T Tauri systems with misaligned discs in their polarization survey, but suggest that perhaps these are perturbed disc soon to be realigned. So there is some evidence that non-coplanar particle distributions are plausible. It may be possible as well that particles from the disc can be captured into the polar stability region, and certainly a coplanar particle distribution is quickly perturbed to fairly high inclinations.  There is much further work to be done in modelling this system. A model of the planetesimals including collisions, although expected to be a minor effect on the dynamics, would indicate if the rate was currently high and may explain the unually dusty nature of the system. A model that also included dust and dust collisions, as well as interactions with any gas in the system, would then fully model the system and allow detailed comparisons with the observations. Although not an easy task given the nature of the system it would narrow down the physics and dynamics at work, important to our understanding of this stage of the planetary formation process in multiple star systems."
    },
    "0807/0807.2241.txt": {
        "abstract": "We present infrared and millimeter observations of Barnard 335, the prototypical isolated Bok globule with an embedded protostar.  Using {\\it Spitzer} data we measure the source luminosity accurately; we also constrain the density profile of the innermost globule material near the protostar using the observation of an 8.0~\\micron\\ shadow. {\\it HHT} observations of $^{12}$CO 2 -- 1 confirm the detection of a flattened molecular core with diameter $\\sim\\!10000$~AU and the same orientation as the circumstellar disk ($\\sim\\!100$ to 200~AU in diameter).  This structure is probably the same as that generating the 8.0~\\micron\\ shadow and is expected from theoretical simulations of collapsing embedded protostars.  We estimate the mass of the protostar to be only $\\sim\\!5$\\% of the mass of the parent globule. ",
        "introduction": "\\citet{bok47} and \\citet{bok48} were the first to suggest that ``Bok globules'' could be the sites of isolated star formation. With this in mind, \\cbo\\ (CB199 in the catalog of \\citet{clemens88}) was extensively studied in molecular lines \\citep{martin78,dickman78} and in the far-infrared by \\citet{keene80} and \\citet{keene81}. \\citet{frerking82} discovered a ``pedestal'' source, i.e., a molecular outflow, indicating the possible presence of a protostar within the cloud. The presence of a protostar was quickly confirmed with additional far-infrared observations by \\citet{keene83}.  Assuming a distance of 250 pc, as suggested by \\citet{tomita79} and as adopted by most subsequent observers, \\citet{keene83} determined the luminosity of the source within \\cbo\\ to be about 3 L$_{\\odot}$.  \\cbo\\ has subsequently been observed at far-infrared and submillimeter wavelengths by, e.g., \\citet{gee85}, \\citet{mozu86}, \\citet{chandler90}, \\citet{huard99}, and \\citet{shirley00}; it is the first example of a Class 0 young stellar object to be so thoroughly studied.  It is therefore the template source for many of our concepts about the early stages of the formation of low-mass stars.  Modeling has established that the globule material is still undergoing strong noncircular motions that are most likely infall \\citep[e.g.,][]{zhou90,choi95,velusamy95}.  Although this scenario of infall, and more specifically of inside-out collapse, is plausible, important discrepancies remain \\citep[e.g.,][]{wilner00,harvey01,evans05,choi07}.  Models and observations imply that the core is a star of mass $\\sim\\!0.5$~\\msun \\citep{zhou90} with a circumstellar disk $\\sim\\!100$ to 200~AU in diameter \\citep{harvey03a}.  This source has a strong bipolar outflow \\citep[e.g.,][]{cabrit92,hirano88}, oriented nearly in the plane of the sky, that maintains a cavity in the interstellar medium (ISM) and is associated with the observed shocked material \\citep{galfalk07}.  We describe observations with all three instruments on {\\it Spitzer} and with the Heinrich Hertz Submillimeter Telescope that allow us to 1) measure the density distribution in the center of the globule; 2) constrain the geometry of the outflow and thus the protostellar disk; and 3) measure the source luminosity accurately.  These measurements (with a revised distance of 150~pc) reveal a flattened molecular cloud core, consistent with numerical simulations of cloud collapse; provide tight constraints on the opening angle of the outflow; and demonstrate internal consistency between the globule luminosity (and other characteristics) and those expected for a 0.3 to 0.4~\\msun\\ young star.  \\begin{figure*}[t] \\scalebox{1.15}{\\plotone{f1.eps}} \\caption{Gallery of images of B335 (CB199) at the indicated wavelengths.  Images are shown with a logarithmic intensity scale and are 1$\\farcm$6 on a side.  North is up and east is to the left.} \\label{fig:mips} \\end{figure*} ",
        "conclusions": "We report and analyze new Spitzer 3.6 to 160~\\micron\\ data observations of \\cbo\\, along with new ground based data in $^{12}$CO and $^{13}$CO.  \\noindent 1.)  We revise the estimated distance from 250~pc to 150~pc.  Our preferred distance is more compatible with the extinction models of \\citet{harvey01} and with models of the embedded protostar \\citep{robitaille06,robitaille07}.  \\noindent 2.)  We discover an 8.0~\\micron\\ shadow and derive a density profile for radii ranging from $\\sim\\!3,000$ to 7,500 AU. The shadow opacity at maximum is equivalent to $A_V$ = 100 or more.  \\noindent 3.)  We detect a rotating structure in $^{12}$CO on scales of $\\sim\\!10000$~AU which we term a flattened molecular core.  This structure is coaxial with the central circumstellar disk.  \\noindent 4.)  It is very likely that the 8.0\\micron\\ shadow and the flattened molecular core originate from the same structure.  \\noindent 5.)  We measure the luminosity of the protostar to be 1.2~L$_\\odot$ (at 150~pc),in good agreement with models for young stars of the expected mass ($\\sim\\!0.4$~\\msun) and general characteristics of the object.  As the prototypical isolated, star-forming dark globule, there is a huge literature on Barnard 335.  Nonetheless, our mid- and far-infrared data from {\\it Spitzer} and CO observations from the HHT provide new insights regarding the configuration of the collapse of the globule and the properties of the protostar. We have found that the new observations are roughly consistent with recent models. Further comparisons with the full suite of observations have the potential to improve our understanding of isolated star formation significantly."
    },
    "0807/0807.0028_arXiv.txt": {
        "abstract": "{} {Eclipses and occultations of post-AGB stars provide a powerful method of exploring the near-stellar environment, including close companions and  circumstellar debris disks. Only six eclipsing systems and one dust-occultation system are currently known. New cases are important for our understanding of binary evolution during the AGB mass-loss phase.} {We study the post-AGB central star of the (bipolar) Galactic bulge planetary nebula M\\,2-29. We have obtained additional HST imaging and SAAO spectroscopy of the object.} {The star showed a pronounced, long-lasting occultation with subsequent recovery. The event lasted almost 3 years, with a secondary minimum 9 years later. The photometric behavior of M\\,2-29 resembles the dust-occultation events seen in NGC\\,2346, and is modeled as an occultation by a circumbinary disk, where the binary period is 18 yr. Modulation during the decline shows evidence of another companion with a period of 23 days.} {M\\,2-29 is the first eclipsing disk system among post-AGB stars. Close binaries with periods of around 1 month, as found in M\\,2-29, have been proposed to supply the energy needed to create the tori of bipolar planetary nebulae.} ",
        "introduction": "The role of binary stars in the formation and evolution of planetary nebulae (PNe) is a controversial one \\citep{Zijlstra2007}. In the traditional view, the ejection of the nebula occurs through a radiation-driven superwind, at the tip of the asymptotic giant branch (AGB); however, the outflow may instead be driven via an interacting binary system, providing a source of angular momentum to the outflow \\citep[e.g.][]{Livio1988}. Even if not directly interacting, binary companions can trap some previously ejected material into disks, as commonly seen around cool (A-F)post-AGB stars \\citep{vanWinckel2003} with orbital periods of 100--1000 days.  Information on binary systems and disks within PNe is limited.  Only about 20 compact binary systems are known, and all but one have periods in the range 0.1--5 day \\citep{marco2008}. Six eclipsing systems are currently known \\citep{Miszalski08, marco2008}, and they provide valuable but rare information on close companions. Their periods are between 0.17 and 4.9 days.  A few PNe show evidence of trapped circumstellar material. The VLT has detected a 10\\,AU disk around the central star of the Ant nebula \\citep{Chesneau2007}. Infrequent obscurations by orbiting dust clouds provide evidence of a similar structure around the central star of NGC\\,2346 (V651 Mon) \\citep{mendez}. NGC\\,2346 is a known binary system. With a period of 16 days it is so far the only known wider system, with a period some 5 times longer than the next longest. It approaches the period range seen for the A--F post-AGB stars with disks.  \\citet{Peretto2007} suggest that NGC\\,6302 has an undetected binary companion with a period of about 1 month. Such a binary could have supplied the energy necessary for its expanding, high-mass torus. If this is correct, the implication is that PNe with dense tori have binaries with similar periods and possibly central remnant disks.  \\begin{figure} \\centering \\includegraphics[width=88mm]{0492fig1.ps} \\caption{H$ \\alpha$ HST image of M\\,2-29 obtained in 2003. A logarithmic intensity scale is applied. The nebula is complex, including a ring, a halo, a jet-like structure, and a central star. North is on top, east to the left. The length of the bar is 2 arcsec. } \\label{Halpha_image} \\end{figure}  The planetary nebula M\\,2-29 (PN\\,G\\,004.0$-$03.0) is a 4-arcsec object, previously considered as a halo PN \\citep{pena}, but the revised abundance analysis of \\citet{torres} suggests it to be a member of the bulge. Old HST images are described by \\citet{torres}. New HST Wide Field Planetary Camera 2 images in [O\\,{\\sc iii}] 5007 {\\AA} and $\\rm H \\alpha$ lines and in continuum were obtained on May 20, 2003. The H$\\alpha$ image (Fig. \\ref{Halpha_image}) confirms the features described by \\citet{torres} but at much better sensitivity. (The jet-like feature will be discussed elsewhere: we attribute it to interaction with the interstellar medium  \\citep{Wareing}.) The 2MASS NIR colors of $J-H = 0.41$ and $H-K = 1.17$ differ from those of ionized gas, and are indicative of a hot star plus a 1000\\,K black body \\citep{Whitelock}.  In this paper we present data on a new occultation event in a central star of a planetary nebula, lasting several years. ",
        "conclusions": "\\subsection{The lightcurve}  Between the beginning of June (JD 2449512) and September 1994 (JD 2449610), for $\\sim$100 days, the $I$-band magnitude dropped from $\\rm 14.9$ to $\\rm 16.0$ mag (Fig. \\ref{Ogle}). The drop after JD\\,2449512 was approximately linear on the magnitude scale, showing tiny variations (Fig. \\ref{Decline}) superposed on the linear trend. The final part of the decline was not covered by OGLE. If the drop of the brightness did continue to be linear, the star would have reached the minimum approximately at JD 2449620, with the full decline taking about 110 days.  For the next 210 days, it remained at the minimum level of 16.0 mag. The MACHO data indicate that the minimum remained flat even after the OGLE observations ceased, and the recovery took place somewhere between JD 2450380 and 2450507 and lasted less than 130 days (Fig. \\ref{Macho}). The whole event took $\\approx$ 1000 days.  Out of the minimum, a gradual, linear brightening is seen between 1998 and 2007, by $\\Delta I =0.2\\,$mag. The continuum was brighter before the eclipse event, suggesting a reversal of this trend occurred around the minimum. Superposed on this trend is a secondary minimum, during 2002--2005, with an amplitude of $\\Delta I ={\\rm 0.15}$. This secondary minimum is not sampled well regarding its onset. Its duration was similar to that of the primary minimum, but it lacked the flat minimum. The two minima are separated by  $\\sim 9\\,$yr, giving a plausible orbital period of 18\\,yr. Confirmation is needed.  \\begin{figure} \\centering \\includegraphics[width=88mm]{0492fig5.eps} \\caption{$I$ light curve of M\\,2-29 from OGLE during the decline. Two models are shown for an eclipse by the object by a sharp-edged screen. Dashed line: a total eclipse of the star + constant source of the flux present (or an eclipse by a semi-transparent cloud). Solid line: partial eclipse by a rectangular, opaque body. } \\label{Decline} \\end{figure}  The amplitude of the decline was $\\Delta I \\sim \\rm 1.3$\\,mag, and $\\Delta V \\sim \\rm 1.7 - 1.8$\\,mag. MACHO data shows a decline of $\\sim \\rm 1.0$\\,mag in MACHO blue and $\\sim \\rm 0.9$\\,mag in MACHO red. The star may have shown somewhat redder colors during minimum than during maximum.   \\subsection{Disk occultation}  The notable aspects of the eclipse are the long time scale, the slow ingress and egress, and the relatively shallow depth.  An eclipse by a companion star can be ruled out. The decline took about 100 days, and the whole minimum lasted about 1000 days. This implied a size of the occulting body of $\\sim 10\\times$ larger than the central star, or $\\rm 2 \\times 10^{11} cm$ ($\\sim 3 \\, R_{\\odot}$), comparable to the size of a main-sequence star. However, the relative tangential velocities would be close to $\\rm 2 \\times R_\\star / 100 d = 25 m/s$, corresponding to an orbit of $\\sim 10^6\\,$AU.  The eclipsing body could be on a very eccentric orbit, or not bounded gravitationally with the central star; however, this does not change the extreme unlikelihood of such an eclipse event.  The event is fitted better with an occultation by a large screen. The central star of a bulge PN is expected to have a mass $M_{\\rm f} \\approx 0.6\\, M_\\odot$. Assuming a secondary component with similar mass, the orbital separation for an 18-yr period is $\\sim 7\\,\\rm AU$. The eclipse lasted for about 15\\%\\ of the adopted orbital period: this yields a radius of the occulting body of $R \\sim 1 \\hbox{--}5\\,\\rm AU$, depending on mass ratio and where the occultation occurred on the orbit.  We propose a model where the occultation is caused by a circumbinary (dust) disk. Such disks are now known to be common around post-AGB stars, and especially so around RV Tau stars \\citep{vanWinckel2003, Deroo2007, deRuyter2006}. Their absence around central stars of PNe has become puzzling. In recent years, the first evidence of such disks have begun to appear, e.g. the discovery of an unexpected compact disk in the Ant nebula \\citep{Chesneau2007}, a disk in the helix nebula \\citep{Su2007}, and the possibly related structure in CK Vul \\citep{Hajduk2007}. The NIR colors indicate the presence of 1000\\,K dust in M\\,2-29, as mentioned above, i.e. near the sublimation temperature. The dust sublimation radius is around 10\\,AU, and so the dust is located outside of the orbit of the 18-yr binary.  Among objects with similar events, long-lasting minima are observed in wide binary systems, like symbiotic stars; they are ascribed to dusty clouds. In V\\,Hya such a minimum is observed repeatedly with over a 6000 day period \\citep{Knapp1999}. NGC\\,2346 has shown two long-lasting eclipses by clouds \\citep{Kato2001}. Either $\\epsilon$ Aur or EE\\,Cep are examples of eclipses by a disk placed nearly edge-on to the observer, around an unseen companion \\citep{carroll, mikolajewski}. Such eclipses appear to be flat and almost grey.  Many of the obscuration characteristics of M\\,2-29 are replicated in the `winking' star KH 15D, an eccentric pre-main-sequence binary that is gradually being occulted by an opaque screen \\citep{Winn, Herbst2008}. The binary orbit is inclined with respect to the circumbinary disk, and in the current configuration one of the stars experiences a `sunset' and `sunrise' every orbit behind the disk. The lightcurve of KH 15D is very similar to that of M\\,2-29, although the time scales are very different as KH 15D presents a much more compact system.  The nature of the secondary star in M\\,2-29 is uncertain. We find no evidence that is has been observed. The depth of the secondary minimum corresponds to a flux only a few times less at $I$ than the PN central star, which is much brighter than any plausible main-sequence star. Such a luminous (red) star would have shown up in the spectra. It is possible that the secondary minimum, and in fact the slow brightening prior, are caused by reflected light from the dusty disk. This would explain the lack of spectral features, the equal length of the secondary minimum (implying a similar orbital radius), and the partial nature of the secondary eclipse, which shows that the secondary emission region is larger than the primary star.   \\subsection{Ionized gas}  In the HST images, the star appears much brighter in F656N ($\\rm H \\alpha$) than expected from the F547M continuum image: $V=15.6 \\pm 0.1$ versus $m({\\rm 656N}) \\approx 12.9$, corresponding to an H$\\alpha$ EW of $\\approx$ 250 \\AA. The stellar magnitude was measured by conversion of the data count from the HST image into the flux and subsequently turned into Johnson magnitudes using Bessel zero points. \\cite{torres} also finds that the star is blended with an unresolved, inner emission-line nebula. There may also be an excess in F502N ([O\\,{\\sc iii}]), where we estimate $m \\sim 14.3$, which after correction for extinction of $E(B-V)=0.63$ leaves about 1\\,mag excess. The [\\,O\\,{\\sc iii}] 5007 \\AA\\ excess can only arise from low-density gas and argues against a stellar wind as the origin of the excess. Instead a circumstellar emission region is indicated.  These data suggest that the light source may be an ionized, inner component to the disk. Such a disk can cause the stellar excess [O\\,{\\sc iii}] and H$\\alpha$ emission \\citep{Soker2006}. We also tentatively assign the He\\,{\\sc ii} absorption seen by \\citet{pena2} to this disk/component. As the ionized region would be much larger than the stellar radius, this might also help explain the long duration before reaching the minimum.    \\subsection{A close binary}  We find evidence of semi-regular photometric variations on a timescale of $\\sim$23 days, which appeared during the decline and the minimum and were present at a level of $\\Delta {\\rm m}_I \\approx 0.05-0.1$\\,mag. The amplitude of these variations is a few times greater than the  accuracy of the individual data points, which spans $0.004-0.01$ mag. This is shown in Fig. \\ref{zacmienie}: the initially regular variations lost phase stability before the star was fully hidden. A similar modulation during the decline was seen for NGC\\,2346 \\citep{mendez} where the 16-day photometric variations ceased immediately after the minimum. They were attributed to the geometric motion of the primary star due to an undetected, faint companion. As the star moves around the barycenter, the degree of obscuration varies. This gives a photometric modulation on the orbital time scale, where the phase depends on the structure of the occulting material.  Following the case of NGC\\,2346, we interpret the photometric variability as evidence of a companion to the central star with an orbital period of 23 days. The eclipse may be considered as the result of the orbital motion of the central star behind the obscuring screen. The 18-yr orbital motion is responsible for the overall shape of the decline, which lasts for over 2\\% of the orbital period. The amplitude and period of the photometric modulations of the eclipse can be modeled by an orbital motion of the central star on a $\\sim$23-day orbit (Fig. \\ref{zacmienie}). During each 23-day cycle, there is a phase of nearly constant luminosity, when the velocity vectors cancel each other and the star remains stationary with respect to the obscuring body. Assuming that the outer companion has a mass of $0.5-1 M_\\odot$ (as expected for bulge star) and that the screen is stationary with respect of the barycenter, this indicates a mass of the close companion of $0.1-0.2 M_\\odot$.  In M\\,2-29, the photometric variability re-appeared during the minimum. The phase of the photometric variations is conserved only for 2-3 consecutive periods. This indicates that the optical depth through the eclipsing body is not uniform. Some of the received light may be scattered.  \\begin{figure} \\centering \\includegraphics[width=88mm]{0492fig6.eps} \\caption{$I$ light curve of M\\,2-29 from OGLE during the decline and the minimum. The timescale for points from minimum is shifted to position the data below the decline phase. A sample of variability of the central star out of the minimum is shown at the bottom. The dashed line shows a total eclipse of the central star by a semi-transparent, sharp-edged body with the sharp edge. The solid line includes the $\\sim$ 23-day and 18-yr orbital motions of the central star. } \\label{zacmienie} \\end{figure}  The detection of such a companion to the primary central star is significant. As shown by \\citet{Peretto2007}, the infall of a distant companion to such a period can provide the energy needed to create an expanding torus. The morphology of M\\,2-29 clearly shows the existence of such a torus. In contrast, the much shorter-period binaries that have come through a common-envelope evolution show rings, jets, and irregular structures, but no tori or butterfly morphologies \\citep{Zijlstra2007}. This would therefore greatly strengthen the case that close binaries have a strong effect on the shaping of planetary nebulae."
    },
    "0807/0807.0444_arXiv.txt": {
        "abstract": "Clusters of galaxies are the most impressive gravitationally-bound systems in the Universe and its abundance (the cluster mass function) is one important statistics to probe the matter density parameter ($\\Omega_m$) and the amplitude of density fluctuations ($\\sigma_8$). The cluster mass function  is usually described in terms of the Press-Schecther (PS) formalism where the primordial density fluctuations are assumed to be a Gaussian random field. In previous works we have proposed a non-Gaussian analytical extension of the PS approach with basis  on the $q$-power law distribution (PL) of the non-extensive kinetic theory. In this paper, by applying the PL distribution to fit the observational mass function data from X-ray highest flux-limited sample (HIFLUGCS) we find a strong degeneracy among the cosmic parameters, $\\sigma_8$, $\\Omega_m$, and the $q$ parameter from the PL distribution. A joint analysis involving recent observations from baryon acoustic oscillation (BAO) peak and Cosmic Microwave Background (CMB) shift parameter is carried out in order to break these degeneracy and better constrain the physically  relevant parameters. The present results suggest that the next generation of cluster surveys will be able to probe the quantities of cosmological interest ($\\sigma_8, \\Omega_m$) and the underlying cluster physics quantified by the $q$-parameter. ",
        "introduction": "The recent astronomical observations are strongly suggesting that the expansion of the Universe is speeding up and not slowing thereby revealing the presence of some form of repulsive gravity.  The basic set of experiments includes: observations from SNe Ia\\cite{Riess}, cosmic microwave background (CMB) temperature anisotropies\\cite{Sper07}, large scale structure\\cite{LSS,Eisenstein05}, X-ray data from galaxy clusters\\cite{Allen}, age estimates of globular clusters and old high redshift galaxies\\cite{Ages}. In the present cosmic concordance $\\Lambda$CDM model the Universe is formed of $\\sim$ 26\\% matter (baryonic + dark matter) and $\\sim$ 74\\% of a smooth vacuum energy component. The thermal CMB component contributes only about 0.01\\%, however, its angular power spectrum of temperature anisotropies encode important information about the structure formation process and other cosmic observables.  On the other hand, the number density of collapsed objects for a given mass at a certain time, named the \\emph{mass or multiplicity function}, is a key quantity in the analysis of cosmic structures such as clusters of galaxies. The simplest successful approach to analytically describe this quantity was developed more than three decades ago by Press and Schechter\\cite{PS74} (hereafter PS). This Gaussian PS formalism is extensively  adopted to derive the mass function, $F_{(M)}$, of bounded objects in the observed Universe\\cite{katz94,longair,naga01,reiprich02,reip06,rines08}.  The PS formalism was adopted by Reiprich and Boringer\\cite{reiprich02} in their construction of the cluster mass function based on the X-ray flux-limited sample of galaxy clusters (HIFLUGCS) selected from  the ROSAT All-Sky Survey. As a result, the best fit parameters, $\\Omega_{m}=0.12$ and $\\sigma_{8}=0.96$, were obtained in their analysis. These values are, respectively, very low and very high when compared with the nowadays independent CMB results\\cite{Sper07} thereby leading to some skepticism about the usefulness of clusters as sensitive cosmological probes. However, several authors have recently claimed that there is no tension between cosmological constraints from CMB and clusters\\cite{reip06,rines08}. In particular, Rines and collaborators\\cite{rines08} argued that the dynamical determination of cluster masses were overestimated and even modest values of velocity segregation between galaxies and dark matter are sufficient to match the mass function with the WMAP results.  In this article we discuss a different possibility. We advocate here that such a discrepancy comes out because the matter density fluctuations should be described by an intrinsically  non-Gaussian random field (due to the action of the long-range gravitational interaction) and, therefore, the PS approach should be somewhat modified.  In previous papers\\cite{LM04,ML07}, inspired by the Tsallis $q$-nonextensive statistics\\cite{Tsal88} and kinetic theory\\cite{Plas,Lim2,RS00,ZER}, we have proposed a simple extension of the PS analytical formalism. Instead the Gaussian function of the original PS approach, it was adopted the power law (PL) Tsallis distribution for describing the fluctuations of the density field. A basic attractive feature of the PL distribution is that the resulting model is analytically tractable and the standard Press-Schechter formalism is recovered as particular case. In this way, a detailed comparison between the two approaches is immediate.  As we shall see, by applying the PL distribution to fit the observational mass function data from X-ray highest flux-limited sample (HIFLUGCS) we find a strong degeneracy among the pair of cosmic parameters ($\\sigma_8$, $\\Omega_m$), and the $q$ parameter from PL distribution. Trough a joint analysis involving recent observations from baryon acoustic oscillation (BAO) peak\\cite{Eisenstein05} and Cosmic Microwave Background (CMB) shift parameter\\cite{Davis,Elgaroy,sze1} the degeneracy is broken and the tension between independent determinations is alliviated. ",
        "conclusions": "\\label{Sec3}  In the precision era of cosmology it is very important to show the compatibility of the cosmological parameters as determined from independent observations.  The current standard cosmological model, i.e., a flat, accelerating Universe composed of $\\simeq 1/3$ of matter (baryonic + dark) and $\\simeq 2/3$ of a dark energy component in the form of the vacuum energy density ($\\Lambda$) seems to be fully consistent with a variety of observational data. However, some tension has recently been detected between the determinations of the pair ($\\sigma_8,\\Omega_m$) from galaxy clusters and CMB. In principle, many temporary solutions are possible, but, hopefully, only one will survive to future analyses based on the next generation of cluster surveys.  In this article we have discussed an alternative route which seems interesting to be investigated from a theoretical and observational viewpoint. Initially, inspired by the non-extensive kinetic theory, we have extended the original Gaussian distribution of the primordial density field\\cite{LM04,ML07}. By using this PL distribution to fit the observational mass function from the HIFLUGCS data we have identified a degeneracy involving the cosmic parameters $\\sigma_8$ and $\\Omega_m$ with the $q$ parameter from the PL distribution.  Finally,  a joint analysis involving the baryon acoustic oscillation signature and the  CMB  shift parameter have been applied. As a result, we have shown that the non-extensive PL distribution may alliviate the tension underlying the independent determinations of the cosmic parameters (see right panel of Fig. \\ref{Fig4})."
    },
    "0807/0807.1674_arXiv.txt": {
        "abstract": "We present a theoretical model for supernova (SN) 2008D associated with the luminous X-ray transient 080109. The bolometric light curve and optical spectra of the SN are modelled based on the progenitor models and the explosion models obtained from hydrodynamic/nucleosynthetic calculations. We find that SN 2008D is a more energetic explosion than normal core-collapse supernovae, with an ejecta mass of $\\Mej = 5.3 \\pm 1.0\\ \\Msun$ and a kinetic energy of $\\KE = 6.0 \\pm 2.5 \\times 10^{51}$ erg. The progenitor star of the SN has a $6-8 \\Msun$ He core with essentially no H envelope ($< 5 \\times 10^{-4}\\ \\Msun$) prior to the explosion. The main-sequence mass of the progenitor is estimated to be $\\Mms =20-25\\ \\Msun$, with additional systematic uncertainties due to convection, mass loss, rotation, and binary effects. These properties are intermediate between those of normal SNe and hypernovae associated with gamma-ray bursts. The mass of the central remnant is estimated as $1.6-1.8 \\Msun$, which is near the boundary between neutron star and black hole formation. ",
        "introduction": "\\label{sec:intro}   A luminous X-ray transient was discovered in NGC 2770 in the {\\it Swift} XRT data taken on 2008 January 9 for the observation of SN 2007uy in the same galaxy (Berger \\& Soderberg 2008; Kong \\& Maccarone 2008). The X-ray emission of the transient reached a peak $\\sim 65$ seconds, lasting $\\sim 600$ seconds, after the observation started (Page et al. 2008). The X-ray spectrum is soft, and no $\\gamma$-ray counterpart was detected by the {\\it Swift} BAT (Page et al. 2008).  Given the small total X-ray energy and the soft X-ray emission, Soderberg et al. (2008) and Chevalier \\& Fransson (2008) interpreted the X-ray transient as a supernova (SN) shock breakout. On the other hand, Xu et al. (2008), Li (2008) and Mazzali et al. (2008) considered this transient as the least energetic end of gamma-ray bursts (GRBs) and X-ray flashes (XRFs).  The optical counterpart was discovered at the position of the X-ray transient (Deng \\& Zhu 2008; Valenti et al. 2008a), confirming the presence of a SN, named SN 2008D (Li \\& Filippenko 2008). To study this event in detail, intensive follow-up observations were carried out over a wide wavelength range including X-rays (Soderberg et al. 2008; Modjaz et al. 2008b), optical/NIR (Soderberg et al. 2008; Malesani et al. 2009; Modjaz et al. 2008b; Mazzali et al. 2008), and radio (Soderberg et al. 2008).  \\begin{deluxetable*}{lrrrrrrrrrr} \\tablewidth{0pt} \\tablecaption{Explosion Models} \\tablehead{ Model & $\\Mms$ \\tablenotemark{a} & $M_{\\alpha}$ \\tablenotemark{b} & $R_*$ \\tablenotemark{c} & $M_{\\rm cut}$ \\tablenotemark{d} & $\\Mej$ \\tablenotemark{e} & $\\KE$ \\tablenotemark{f} & $M($\\Nifs$)$ \\tablenotemark{g} & $v_{\\rm He}$ \\tablenotemark{h} & $v_{\\rm Ni}$ \\tablenotemark{i} & $M_{0.1c}$ \\tablenotemark{j} } \\startdata HE4          & $\\approx$ 15 & 4  & 3.5  & 1.3 & 2.7 & 1.1 & 0.07 & $<$3500 & 7900 &  $< 3.0 \\times 10^{-5}$  \\\\ HE6          & $\\approx$ 20 & 6  & 2.2  & 1.6 & 4.4 & 3.7 & 0.065 & 6700    & 7000 &  0.007   \\\\ HE8          & $\\approx$ 25 & 8  & 1.3  & 1.8 & 6.2 & 8.4 & 0.07 & 10500   & 9000 &  0.04   \\\\ HE10         & $\\approx$ 30 & 10 & 1.2  & 2.3 & 7.7 & 13.0 & 0.07 & 12500   & 10600&  0.09   \\\\ HE16         & $\\approx$ 40 & 16 & 0.74  & 3.6 & 12.4& 26.5 & 0.07 & 17500   & 14000&  0.24   \\\\ \\hline \\multicolumn{2}{l}{Soderberg et al. (2008)} & & $\\sim 1$ \\tablenotemark{k} &  & 3-5     &  2-4    & 0.05  &   & &  \\\\ Mazzali et al. (2008) & $\\sim 30$  & & &  & $\\sim$7 & $\\sim$6 & 0.09  &   & & 0.03 \\enddata \\tablenotetext{a}{Main-sequence mass ($\\Msun$) estimated from the approximate formula obtained by Sugimoto \\& Nomoto (1980)} \\tablenotetext{b}{The mass of the He star ($\\Msun$)} \\tablenotetext{c}{Progenitor radius prior to the explosion ($\\Rsun$)} \\tablenotetext{d}{Mass cut ($\\Msun$)} \\tablenotetext{e}{The mass of the SN ejecta ($\\Msun$)} \\tablenotetext{f}{The kinetic energy of the SN ejecta ($10^{51}$ erg)} \\tablenotetext{g}{The mass of ejected \\Nifs\\ ($\\Msun$)} \\tablenotetext{h}{Velocity at the bottom of the He layer (\\kms)} \\tablenotetext{i}{Velocity at the outer boundary of \\Nifs\\ distribution (\\kms)} \\tablenotetext{j}{The ejecta mass at $v > 0.1c$ ($\\Msun$)} \\tablenotetext{k}{Estimated from the photospheric radius and temperature of the early part of the LC ($t \\lsim 4$ days) using the formulae by Waxman et al. (2007). This is consistent with the estimate by Modjaz et al. (2008b) while Chevalier \\& Fransson (2008) derived $\\sim 9 \\Rsun$ using $\\Mej$ and $\\KE$ estimated by Soderberg et al. (2008) and the formulae by Chevalier (1992).} \\label{tab:models} \\end{deluxetable*}   SN 2008D showed a broad-line optical spectrum at early epochs ($t \\lsim 10$ days, hereafter $t$ denotes time after the transient, 2008 Jan 9.56 UT; Soderberg et al. 2008). However, the spectrum changed to that of normal Type Ib SN, \\ie SN with He absorption lines and without H lines (Modjaz et al. 2008c). To date, the SN associated with GRBs or XRFs are all Type Ic, \\ie SNe without H and He absorption.  Soderberg et al. (2008) and Mazzali et al. (2008) estimated the ejected mass and the kinetic energy of SN 2008D. Using analytic formulae, Soderberg et al. (2008) suggested that this SN has the ejecta mass $\\Mej = 3 - 5 \\Msun$ and the kinetic energy of the ejecta $\\KE \\sim 2 - 4 \\times 10^{51}$ erg. Mazzali et al. (2008) did model calculations and suggested that this event is intermediate between normal SNe and GRB-associated SNe (or hypernovae), with $\\Mej \\sim 7 \\Msun$ and $\\KE \\sim 6 \\times 10^{51}$ erg. In their model calculations, hydrodynamic/nucleosynthetic models are not used, and thus, their estimate of the core mass prior to the explosion and the progenitor main-sequence mass is less direct.   In this paper, we present detailed theoretical study of emission from SN 2008D. The bolometric light curve (LC) and optical spectra are modelled based on the progenitor models and the explosion models obtained from hydrodynamic/nucleosynthetic calculations. In \\S \\ref{sec:models}, we show the progenitor and explosion models. The optical LC and spectra are modelled in \\S \\ref{sec:LC} and \\S \\ref{sec:spec}, respectively. We discuss the nature of SN 2008D in \\S \\ref{sec:discussion} and finally give conclusions in \\S \\ref{sec:conclusions}. Throughout this paper, we adopt 31 Mpc ($\\mu = 32.46$ mag) for the distance to SN 2008D (Modjaz et al. 2008b; Mazzali et al. 2008) and $E(B-V)=0.65$ mag for the total reddening (Mazzali et al. 2008). ",
        "conclusions": "\\label{sec:conclusions}  We presented a theoretical model for SN 2008D associated with the luminous X-ray transient 080109. Based on the progenitor models, hydrodynamics and explosive nucleosynthesis are calculated. Using the explosion models, radiative transfer calculations are performed. These models are tested against the bolometric LC and optical spectra. This is the first detailed model calculation for Type Ib SN that is discovered shortly after the explosion.  We find that SN 2008D is a more energetic explosion than normal core-collapse SNe. We estimate that the ejecta mass is $\\Mej = 5.3 \\pm 1.0 \\Msun$ and the total kinetic energy of $\\KE = 6.0 \\pm 2.5 \\times 10^{51}$ erg. The ejected \\Nifs\\ mass is $\\sim 0.07 \\Msun$. To eject the optimal amount of \\Nifs, the mass of the central remnant is estimated to be $1.6 - 1.8 \\Msun$. The error bars include only the uncertainty of the LC and spectral modelling.  Summing up the above masses, it is concluded that the progenitor star of SN 2008D has a $6 - 8 \\Msun$ He core prior to the explosion. There is essentially no H envelope with the upper limit of $5 \\times 10^{-4} \\Msun$. Thus, the corresponding main-sequence mass of the progenitor is $\\Mms =20-25 \\Msun$ under the $\\Mms$-$M_{\\alpha}$ relation by Sugimoto \\& Nomoto (1980, used in Nomoto \\& Hashimoto 1988). We note that there exist additional systematic uncertainties in this relation due to convection, mass loss, rotation, and binary effects. Our estimates of these masses and energy suggest that SN 2008D is near the border between neutron star-forming and black hole-forming SNe, and has properties intermediate between those of normal SNe and hypernovae associated with gamma-ray bursts."
    },
    "0807/0807.3447_arXiv.txt": {
        "abstract": "We combine a semi-analytical model of galaxy formation with a very large simulation which follows the growth of large scale structure in a $\\Lambda$CDM universe to predict the clustering of \\lya\\ emitters. We find that the clustering strength of \\lya\\ emitters has only a weak dependence on \\lya\\ luminosity but a strong dependence on redshift. With increasing redshift, \\lya\\ emitters trace progressively rarer, higher density regions of the universe.  Due to the large volume of the simulation, over 100 times bigger than any previously used for this application, we can construct mock catalogues of \\lya\\ emitters and study the sample variance of current and forthcoming surveys. We find that the number and clustering of \\lya\\ emitters in our mock catalogues are in agreement with measurements from current surveys, but that there is a considerable scatter in these quantities. We argue that a proposed survey of emitters at $z=8.8$ should be extended significantly in solid angle to allow a robust measurement of \\lya\\ emitter clustering. ",
        "introduction": "The study of galaxies at high redshifts opens an important window on the process of galaxy formation and conditions in the early universe. The detection of populations of galaxies at high redshifts is one of the great challenges in observational cosmology. Currently three main observational techniques are used to discover high redshift, star-forming galaxies: (i) The Lyman-break drop-out technique, in which a galaxy is imaged in a combination of three or more optical or near-IR bands. The longer wavelength filters detect emission in the rest-frame ultraviolet from ongoing star formation, whereas the shorter wavelength filters sample the Lyman-break feature. Hence, a Lyman-break galaxy appears blue in one colour and red in the other \\citep{steidel96, steidel99}. By shifting the whole filter set to longer wavelengths, the Lyman-break feature can be isolated at higher redshifts; (ii) Sub-millimetre emission, due to dust being heated when it absorbs starlight \\citep{smail97, hughes98}. The bulk of the energy absorbed by the dust comes from the rest-frame ultra-violet and so the dust emission is sensitive to the instantaneous star formation rate; (iii) Ly-$\\alpha$ line emission from star forming galaxies, typically identified using either narrowband imaging \\citep{hu98, kudritzki00, gawiser07, ouchi07} or long-slit spectroscopy of gravitationally lensed regions \\citep{ellis01, santos04, stark07}. The Ly-$\\alpha$ emission is driven by the production of Lyman-continuum photons and so is dependent on the current star formation rate.   The Lyman-break drop-out and sub-millimetre detection methods are more established than Ly-$\\alpha$ emission as a means of identifying substantial populations of high redshift galaxies. Nevertheless, in the last few years there have been a number of \\lya\\ surveys which have successfully found high redshift galaxies e.g. \\citep{hu98, kudritzki00}. The observational samples have grown in size such that statistical studies of the properties of \\lya\\ emitters have now become possible: for example, the SXDS Survey \\citep{ouchi05,ouchi07} has allowed estimates of the luminosity function (LFs) and clustering of \\lya\\ emitters in the redshift range $3<z<6$, and the MUSYC survey \\citep{gronwall07,gawiser07} has also produced clustering measurements at $z \\sim 3$. Furthermore, the highest redshift galaxy ($z=6.96$) robustly detected to date was found using the \\lya\\ technique \\citep{iye}.  Taking advantage of the magnification of faint sources by gravitational lensing, \\citet{stark07} reported 6 candidates for \\lya\\ emitters in the redshift range $8.7<z<10.2$, but these have yet to be confirmed. The DAzLE Project \\citep{horton} is designed to find \\lya\\ emitters at $z=7.73$ and $z=8.78$.  However, the small field of view of the instrument ($6.83 \\arcmin \\times 6.83 \\arcmin$) makes it difficult to use to study large scale structure (LSS) at such redshifts.  On the other hand, the ELVIS Survey \\citep{kim07a, kim07b} would appear to offer a more promising route to study the LSS of very high redshift galaxies ($z=8.8$).   Despite these observational breakthroughs, predictions of the properties of star-forming \\lya\\ emitting galaxies are still in the relatively early stages of development. Often these calculations employ crude assumptions about the galaxy formation process to derive a star formation rate and hence a \\lya\\ luminosity, or use hydrodynamical simulations, which, due to the high computational overhead, study relatively small cosmological volumes. \\citet{hai} made predictions for the escape fraction of \\lya\\ emission and the abundance of \\lya\\ emitters using the Press-Schechter formalism and a prescription for the dust distribution in galaxies. Radiative transfer calculations of the escape fraction have been made by \\citet{zheng02}, \\citet{ahn04} and \\citet{verhamme06} for idealized geometries, while \\citet{tasitsiomi06} and \\citet{laursen07} applied these calculations to galaxies taken from cosmological hydrodynamical simulations. \\citet{barton04} and \\citet{furlanetto05} calculated the number density of Ly-$\\alpha$ emitters using hydrodynamical simulations of galaxy formation. Nagamine et~al. (2006, 2008) used hydrodynamical simulations to predict the abundance and clustering of \\lya\\ emitters. The typical computational boxes used in these calculations are very small ($\\sim 10-30 h^{-1}$Mpc), which makes it impossible to evolve the simulation accurately to $z=0$. Hence, it is difficult to test if the galaxy formation model adopted produces a reasonable description of present day galaxies. Furthermore, the small box size means that reliable clustering predictions can only be obtained on scales smaller than the typical correlation length of the galaxy sample. As we will show in this paper, small volumes are subject to significant fluctuations in clustering amplitude.     The semi-analytical approach to modelling galaxy formation allows us to make substantial improvements over previous calculations of the properties of \\lya\\ emitters. The speed of this technique means that large populations of galaxies can be followed. The range of predictions which can be made using semi-analytical models is, in general, broader than that produced from most hydrodynamical simulations, so that the model predictions can be compared more directly with observational results. A key advantage is that the models can be readily evolved to the present-day, giving us more faith in the ingredients used; i.e. we can be reassured that the physics underpinning the predictions presented for a high-redshift population of galaxies would not result in too many bright/massive galaxies at the present day.   The first semi-analytical calculation of the properties of \\lya\\ emitters based on a hierarchical model of galaxy formation was carried out by \\citet{dell1}. This is the model used throughout this work, which has been shown to be successful in predicting the properties of \\lya\\ emitters over a wide range of redshifts. The semi-analytical model allows us to connect \\lya\\ emission to other galaxy properties. \\citet{dell2} showed that this model succesfully predicts the observed \\lya\\ LFs and equivalent widths (EWs), along with some fundamental physical properties, such as star formation rates (SFRs), gas metallicities, and stellar and halo masses. In \\citet{kim07a}, we used the model to make further predictions for the LF of very high redshift \\lya\\ emitters and to study the feasibility of current and forthcoming surveys which aim to detect such high redshift galaxies. \\citet{kobayashi07} developed an independent semi-analytical model to derive the luminosity functions of \\lya\\ emitters.  The focus of this paper is to use the model introduced by \\citet{dell1} to study the clustering of high-redshift \\lya\\ emitting galaxies and to extend the comparison of model predictions with current observational data. \\citet{dell2} already gave an indirect prediction of the clustering of \\lya\\ emitters by studying galaxy bias as a function of \\lya\\ luminosity. However, these results depend on an analytical model for the halo bias \\citep{sheth01}, and furthermore the linear bias assumption breaks down on small scales. Here we will present an explicit calculation of the clustering of galaxies by implementing the semi-analytical model on top of a large N-body simulation of the hierarchical clustering of the dark matter distribution. This allows us to predict the spatial distribution of \\lya\\ emitting galaxies, and to create realistic maps of \\lya\\ emitters at different redshifts. These maps can be analysed with simple statistical tools to quantify the spatial distribution and clustering of galaxies at high redshifts. The N-body simulation used in this work is the \\textit{Millennium Simulation}, carried out by the Virgo Consortium \\citep{mill}.  The simulation of the spatial distribution of \\lya\\ emitters is tested by creating mock catalogues for different surveys of \\lya\\ emitting galaxies in the range $3<z<9$. The clustering of \\lya\\ emitters in our model is analysed with correlation functions and halo occupation distributions. Taking advantage of the large volume of the Millennium simulation, we also compute the errors expected on correlation function measurements from various surveys due to cosmic variance.   The outline of this paper is as follows: Section 2 gives a brief description of the semi-analytical galaxy formation model and describes how it is combined with the N-body simulation. In Section 3 we establish the range of validity of our simulated galaxy samples by studying the completeness fractions in the model \\lya\\ luminosity functions.  Section 4 gives our predictions for the clustering of \\lya\\ emitters in the range $0<z<9$.  In Section 5 we compare our simulation with recent observational data and we also make predictions for future measurements (clustering and number counts) expected from the ELVIS Survey. Finally, Section 6 gives our conclusions. ",
        "conclusions": "We have combined a semi-analytical model of galaxy formation with a high resolution, large volume N-body simulation to make predictions for the spatial distribution of \\lya\\ emitters in a $\\Lambda$CDM universe.  Our model for \\lya\\ emitters is appealingly simple. Using the star formation history predicted for each galaxy from the semi-analytical model to compute the production of Lyman continuum photons, we find that on adopting a constant escape fraction of \\lya\\ photons the observed number of \\lya\\ emitters can be reproduced amazingly well over a range of redshifts \\citep{dell2}. Our modelling of \\lya\\ emission may appear overly simplistic on first comparison to other calculations in the literature. For example, \\citet{nagamine08} predicted the clustering of \\lya\\ emitters in a gas-dynamic simulation, modelling the \\lya\\ emission through a \\lya\\ escape fraction or a duty cycle scenario.  However, the fraction of active emitters in the duty cycle scenario needs to be tuned at each redshift, for which there is no physical justification. Since our predictions for \\lya\\ emission are derived from a full model of galaxy formation, it is straightforward to extract other properties of the emitters, such as their stellar mass or the mass of their host dark matter halo (Le Delliou et~al. 2006). In this paper we have extended this work to include explicit predictions for the spatial and angular clustering of \\lya\\ emitters.  We have studied how the clustering strength of \\lya\\ emitters depends upon their luminosity as a function of redshift. Generally, we find that \\lya\\ emitters show a weak dependence of clustering strength on luminosity, until the brightest luminosities we consider are reached. At the present day, \\lya\\ emitters display weaker clustering than the dark matter. This changes dramatically at higher redshifts ($z > 3$), with currently observable \\lya\\ emitters predicted to be much more strongly clustered than the dark matter, with the size of the bias increasing with redshift. We compared the simulation results with analytical estimates of the bias.  Whilst the analytical results show the same trends as the simulation results, they do not match well in detail, and this supports the use of an N-body simulation to study the clustering.  A key advantage of using semi-analytical modelling is that the evolution of the galaxy population can be readily traced to the present day. This gives us some confidence in the star formation histories predicted by the model. The semi-analytical model passes tests on the predicted distribution of star formation rates at high redshift (the number counts and redshifts of galaxies detected by their sub-millimetre emission and the luminosity function of Lyman-break galaxies), whilst also giving a reasonable match to the present day galaxy luminosity function (Baugh et~al. 2005), and also matching the observed evolution of galaxies in the infrared \\citep{lacey08}. Gas dynamic simulations as a whole struggle to reproduce the present-day galaxy population, due to a combination of a limited simulation volume (set by the need to attain a particular mass resolution) and a tendency to overproduce massive galaxies. The small box size typically employed in gas dynamic simulations means that fluctuations on the scale of the box become nonlinear at low redshifts, and their evolution can no longer be accurately modelled. A further consequence of the small box size is that predictions for clustering are limited to small pair separations (e.g. \\citet{nagamine08} use a box of side 33 \\mpc, limiting their clustering predictions to scales $r \\la 3 \\mpc$). By using a simulation with a much larger volume than that of any existing \\lya\\ survey, we can subdivide the simulation box to make many mock catalogues to assess the impact of sampling fluctuations (including cosmic variance) on current and future measurements of the clustering of \\lya\\ emitters.  We made mock catalogues of \\lya\\ emitters to compare with the MUSYC ($z=3$) and SXDS ($z=3-6$) surveys, and to make predictions for the forthcoming ELVIS survey at $z \\sim 9$. In the case of MUSYC and SXDS, we found that the observed number of galaxies lies within the 10-90 percentile interval of the number of \\lya\\ emitters found in the mocks. We find that high-redshift clustering surveys underestimate their uncertainties significantly if they fail to account for cosmic variance in their error budget. Overall, the measured angular correlation functions are consistent with the model predictions. The clustering results in our mock catalogues span a wide range of amplitudes due to the small volumes sampled by the surveys, which results in a large cosmic variance. ELVIS will survey \\lya\\ emitters at very high redshift ($z = 8.8$). Our predictions show that a single pointing will be strongly affected by sample variance, due to the small volume surveyed and the strong intrinsic clustering of the \\lya\\ emitters which will be detected at this redshift. Many ELVIS pointings will be required to get a robust clustering measurement.  We have shown that surveys of \\lya\\ emitters can open up a new window on the high redshift universe, tracing sites of active star formation. With increasing redshift, the environments where \\lya\\ emitters are found in current and planned surveys become increasingly unusual, sampling the galaxy formation process in regions that are likely to be proto-clusters and the progenitors of the largest dark matter structures today. Our calculations show that with such strong clustering, surveys of \\lya\\ emitters covering much larger cosmological volumes are needed in order to minimize cosmic variance effects."
    },
    "0807/0807.1218_arXiv.txt": {
        "abstract": "Based on the assumptions that a fraction of cluster dark matter is composed of degenerate neutrinos and they are in hydrostatic equilibrium with other matter, we predict a relation between the density profile and temperature of the cluster hot gas. The predicted relation agrees with observational data of 103 clusters. ",
        "introduction": "Observational data on rotation curves of galaxies and mass profiles of clusters indicate that dark matter exists. For example, the rotation curves of dwarf galaxies indicate that their total masses are much greater than the visible mass \\citep{Swaters}, \\citep{Salucci}. Also, the integrated total mass of a cluster is several times greater than the visible mass including hot gas and galaxies \\citep{Reiprich}. On the other hand, recent neutrino oscillation experiments indicate that neutrinos have finite but small rest mass. Therefore, at least some fractions of dark matter should compose of neutrinos, which is known as hot dark matter (HDM). It is commonly believed that neutrinos exist in clusters and affect their structures. \\citet{Cowsik} provided a simple model to understand the virial mass discrepancy in the Coma cluster if neutrinos have rest mass. After the non-zero rest mass of neutrinos was confirmed by experiments \\citep{Fukuda,Bilenky}, neutrinos being a candidate cluster dark matter has become a hot topic again. \\citet{Treumann} presented a model to calculate the mass distribution in two clusters, Coma and A119, including cold dark matter, $\\sim$ 2 eV neutrino dark matter and hot gas. Recently, \\citet{Nakajima} presented a model using 1-2 eV degenerate neutrinos in hydrostatic equilibrium to fit the observed flat core in A1689, which contradicts with the results obtained by numerical simulations of cold dark matter particles \\citep{Navarro,Moore}. All the above results indicate that neutrino dark matter can be an important component in the mass distribution of clusters. Neutrinos alone cannot form structures as their free streaming scale is too large. However, with the help of cold dark matter, neutrinos can be gravitationally bound in the clusters and their effects may be observable \\citep{formation}. In this article, we study a possible observable consequence of neutrinos in clusters. We assume that the degenerate neutrinos and hot gas particles are in hydrostatic equilibrium under the gravity of cold dark matter, galaxies and intergalactic hot gas. We derive an approximate relation among the parameters specifying the density profiles of the cluster hot gas and their temperatures in 103 clusters. We also make predictions about the density profiles of neutrinos in clusters. ",
        "conclusions": "Neutrinos exist in clusters and they may form structures with help of cold dark matter \\citep{formation}. By assuming the hydrostatic equilibrium of neutrinos and hot gas particles with total mass in clusters, we obtain the density profile of neutrinos in terms of $\\beta$, $T$ and $r_c$, and we can thereby obtain an approximate relation among these parameters with $m_{\\nu} \\le 2$ eV. If $\\rho_c \\propto \\rho_0^{\\gamma}$, then a linear relationship between $\\ln r_c$ and $\\ln \\beta T$ is obtained which agrees with the observed data with $\\gamma \\approx -3/2$. Our result is also compatible with \\citet{Sanders} that the core profiles in clusters can be explained by neutrinos as dark matter."
    },
    "0807/0807.4117_arXiv.txt": {
        "abstract": "We investigate the interaction of nonthermal ions (protons and nuclei) accelerated in an ultrarelativistic blastwave with the pre-existing magnetic field of the medium into which the blastwave propagates.  While particle acceleration processes such as diffusive shock acceleration can accelerate ions and electrons, the accelerated electrons suffer larger radiative losses.  Under certain conditions, the ions can attain higher energies and reach farther ahead of the shock than the electrons, and so the nonthermal particles can be partially charge-separated.  To compensate for the charge separation, the upstream plasma develops a return current, which, as it flows across the magnetic field, drives transverse acceleration of the upstream plasma and a growth of density contrast in the shock upstream. If the density contrast is strong by the time the fluid is shocked, vorticity is generated at the shock transition. The resulting turbulence can amplify the post-shock magnetic field to the levels inferred from gamma-ray burst afterglow spectra and light curves. Therefore, since the upstream inhomogeneities are induced by the ions accelerated in the shock, they are generic even if the blastwave propagates into a medium of uniform density.  We speculate about the global structure of the shock precursor, and delineate several distinct physical regimes that are classified by an increasing distance from the shock and, correspondingly, a decreasing density of nonthermal particles that reach that distance. ",
        "introduction": "\\label{sec:intro}  The standard model for gamma-ray burst (GRB) afterglow emission requires nonthermal acceleration of electrons to ultrarelativistic energies.  This can take place in the ``external'' relativistic shock wave driven into the circumburst medium by the burst ejecta \\citep[e.g.,][and references therein]{Piran:05a,Meszaros:06}.  Candidate mechanisms for shock acceleration, such as the first order Fermi mechanism---also known as diffusive shock acceleration (DSA; e.g., \\citealt{Bell:78,Blandford:78,Blandford:87})---can accelerate all charged particles, including the positively charged ions, for whose acceleration in GRB afterglows we still do not have direct observational evidence.  While theoretical models for particle acceleration in shocks have largely been based on analytic or semianalytic prescriptions for particle injection into the acceleration cycle and its scattering in the random magnetic field of the shock \\citep[e.g.,][and references therein]{Achterberg:01,Ellison:02,Niemiec:06}, recent two-dimensional first-principles particle-in-cell simulations of unmagnetized pair shocks have delivered the first encouraging direct evidence for the formation of the signature nonthermal power-law tail \\citep{Keshet:08,Spitkovsky:08}.  The power law tail seems to arise in tandem with a self-sustained electromagnetic structure in the shock precursor, where the nonthermal particles that move ahead of the shock excite plasma instabilities.  In their nonlinear development, the instabilities generate a magnetic field that scatters the nonthermal particles back toward the shock and thus closes the acceleration cycle.  Our aim here is to explore the consequences of the formation of a power-law tail of accelerated particles for shock hydrodynamics when a weak, large scale magnetic field is present in the shock upstream.  Inverse Compton cooling of the nonthermal electrons imposes an upper limit on the energy of electrons that are accelerated by DSA \\citep[e.g.,][]{Li:06}.  Since higher energy ions reach farther from the shock into the shock upstream, there may exist a region in the shock upstream devoid of nonthermal electrons that still contains nonthermal ions; we restrict our analysis to this region.  In the reference frame in which the shock upstream is at rest, the nonthermal ions are beamed forward within the narrow angle $\\sim \\Gamma^{-1}$, where $\\Gamma\\gg 1$ is the shock Lorentz factor.  The upstream plasma responds to the forward-beamed nonthermal ions by developing a return current \\citep[see, e.g.,][and references therein]{Thode:75,Spicer:84} which flows toward the shock.  The return current cancels the charge separation between the nonthermal ions that reach farther from the shock and the nonthermal electrons that are confined closer to the shock by the inverse Compton losses.  If the upstream plasma is weakly magnetized, the flow of return current across a weak magnetic field accelerates the upstream plasma transversally, i.e., perpendicular to the direction of shock propagation, which creates inhomogeneities in the plasma \\citep{Reville:06,Pelletier:08}; this is a relativistic version of the mechanism discussed by \\citet{Bell:04,Bell:05}.  We show that even when the upstream magnetic field is  weak, these density inhomogeneities may be sufficient to generate a stronger magnetic field in the shock downstream via vorticity production at the shock transition \\citep{Goodman:08,Sironi:07}. This process converts the bulk kinetic energy of the shock downstream into vortical and magnetic energy. An advantage of this process is that it can take place even when the circumburst medium is uniform and therefore it is relevant for, as required by observations, both long and short GRBs \\citep{Nakar:07}. Alternative sources of vorticity that were previously discussed in the literature are pre-existing density inhomogeneities in situations where an ultrarelativistic blastwave propagates into a clumpy stellar wind \\citep{Sironi:07}, or an initially anisotropic blastwave \\citep{Milosavljevic:08}.  Nonthermal particle streaming may generate or amplify a magnetic field in the shock upstream (\\citealt{Bell:01,Milosavljevic:06}; A. Spitkovsky, priv. comm., but also see \\citealt{Pelletier:08}).  The relative efficiencies of the direct, nonthermal particle-driven field amplification in the shock upstream, and the amplification in the shock transition via vorticity generation in the presence of streaming-generated density inhomogeneity will depend on how the shock energy is partitioned between the thermal fluid and the nonthermal accelerated particles and on the efficiency with which magnetic field is generated by cosmic ray-driven turbulence in the shock upstream.  Throughout this work we assume that in ultraretivistic collisionless shocks, DSA accelerates ions to high nonthermal energies.  While there is direct observational evidence of electron acceleration in such shocks, the assumption that protons are accelerated also may not necessarily be justified and deserves further investigation. Specifically, on theoretical grounds, particle acceleration should be inefficient when a compressed oblique large scale upstream magnetic field dominates particle dynamics in the shock downstream. In this case shock compression leads to a downstream magnetic field that is almost perpendicular to the shock normal. Then, a particle with a gyroradius smaller than the coherence length of the field cannot complete two full DSA cycles across the shock before becoming irreversibly bound to and escaping into the shock downstream \\citep[e.g.,][]{Lemoine:06b}. Therefore, a field that contains large-amplitude magnetic field perturbations on small scales seems to be required for acceleration \\citep[e.g.,][]{Ostrowski:02,Niemiec:04,Niemiec:06,Niemiec:06b}. Even in the conditions in which a shock propagates in a weakly magnetized medium and generates magnetic turbulence that serves as a scattering agent for energetic particles, it cannot be taken for granted that the resulting particle spectrum will extend up to the high Lorentz factors that are required for the mechanism that we present. These potential limitations should be born in mind in what follows.  This work is organized as follows.  In \\S~\\ref{sec:charge_separation} we analyze charge separation in the shock precursor and relate the nonthermal ion density ahead of the shock to the maximum energy to which nonthermal ions are accelerated in the shock.  In \\S~\\ref{sec:eom} we study the motion of the upstream medium in response to the Lorentz force associated with the nonthermal ion streaming across a weak magnetic field.  In \\S~\\ref{sec:dencon} we determine the conditions under which nonlinear density inhomogeneities can result and estimate the energy density of the magnetic field produced by vortical motions that are excited as the shock transition sweeps past the inhomogeneities.  In \\S~\\ref{sec:bell} we examine the production of a small scale reversing magnetic field structure that was assumed in \\S~\\ref{sec:eom}, in the weak large-scale magnetic field of the circumburst medium by the nonresonant streaming instability introduced by \\citet{Bell:04,Bell:05}. In \\S~\\ref{sec:discussion} we speculate about the global structure of the shock precursor, and delineate several distinct physical regimes that are classified by an increasing distance from the shock.  In \\S~\\ref{sec:conclusions}, we summarize our main conclusions. ",
        "conclusions": "  In GRB external shocks that accelerate ions to Lorentz factors $\\gamma_i\\gg 10^3$, inverse Compton cooling in the shock upstream hinders nonthermal electron acceleration to high energies.  Because of this, nonthermal electrons do not penetrate as far into the shock upstream as the highest-energy nonthermal ions do.  Thus, a region exists in the shock upstream far from the shock transition at $R\\sim ct[1-(8\\Gamma^2)^{-1}]$, but still behind the photon shell at $R=ct$, where return current flows to cancel the charge separation between the nonthermal electrons and ions.  As the return current flows across the weak pre-existing magnetic field of the shock upstream, the Amp\\`ere force accelerates the upstream plasma perpendicular to the direction of shock propagation.  If the upstream field reverses on scales $\\lambda\\ll R/8\\Gamma^2$, the transverse acceleration produces a nonlinear density contrast on the same scales.  If the pre-existing field is quasi-uniform (non-reversing), Bell's streaming instability will efficiently convert the quasi-uniform field into a reversing field during the light-crossing time of the shock precursor.  As the shock transition sweeps the resulting density clumps, the bulk kinetic energy of the blastwave is converted into vortical energy.  The vortical eddies  in the shock downstream  turn over many times and  amplify the magnetic field via a dynamo mechanism.  For clumps of peak density contrast of about one, the vortical energy \\citep{Sironi:07} and with it the magnetic energy density can be amplified to $\\epsilon_B \\sim (3\\Gamma)^{-1}$, where in GRB external shocks, $\\Gamma \\lesssim  \\textrm{ few }\\times100$. This is consistent with the observed GRB light curves and spectra. The generated magnetic field will be coherent on scales much larger than the plasma skin depth and will, as such, not be susceptible to rapid collisionless decay. Taken together, these results indicate that the microgauss magnetic field ($\\epsilon_B\\sim10^{-9}$) typical of the interstellar medium is sufficient to drive Bell's mechanism in ultrarelativistic GRB external shocks and deliver nonlinear density inhomogeneities to the shock transition.  \\emph{Therefore, magnetic field amplification through shock vorticity generation is expected even when the circumburst medium is not initially clumpy, and we expect it to operate in long and short GRBs.}  We speculatively propose a structure for the shock precursor consisting of four concentric zones characterized by different forms of plasma self-organization.  First, the outermost region of the precursor is the region populated by streaming, nonthermal protons but devoid of nonthermal electrons.  There, the plasma develops non-linear density inhomogeneities.  Second, a region that still contains a return current but where the nonthermal particle precursor streaming through the shock upstream is susceptible to small-scale kinetic plasma instabilities.  Third, a region containing the relatively low-energy nonthermal electrons that can be accelerated in the shock in spite of the inverse Compton cooling losses; in this region, the return current is absent.  Finally, the fourth region is the shock transition itself, where kinetic instabilities and plasma waves facilitate plasma isotropization and establish the hydrodynamic jump. The relative ordering of the second and the third region may vary."
    },
    "0807/0807.0322_arXiv.txt": {
        "abstract": "We consider a class of inflationary models with small oscillations imprinted on an otherwise smooth inflaton potential.  These oscillations are manifest as oscillations in the power spectrum of primordial perturbations, which then give rise to oscillating departures from the standard cosmic microwave background power spectrum.  We show that current data from the Wilkinson Microwave Anisotropy Probe constrain the amplitude of a sinusoidal variation in the inflaton potential to have an amplitude less than $3\\times 10^{-5}$.  We anticipate that the smallest detectable such oscillations in Planck will be roughly an order of magnitude smaller, with slight improvements possible with a post-Planck cosmic-variance limited experiment. ",
        "introduction": "Cosmic microwave background (CMB) experiments continue to be consistent with the simplest predictions of inflationary models, even as the data become increasingly precise \\cite{CMBexperiments}. Constraints on the amplitude and spectral index of primordial density perturbations, and on the amplitude of the inflationary gravitational-wave background, can now be used to constrain the parameter space of the inflaton potential.  In most of the current analyses the inflaton potential is parameterized in terms of its amplitude $V(\\phi)$ and first and second derivatives, $V'(\\phi)$ and $V''(\\phi)$.  Given the recent advances in data quality, as well as the improvements anticipated with forthcoming experiments (e.g., the Planck satellite \\cite{planck}, to be launched within the next year), it is worth asking whether the data can be used to study more complicated forms of the potential.  For example, Adams et al. \\cite{Adams:1997de} showed that supergravity-inspired models may give rise to inflaton potentials with a large number of steps.  Each step corresponds to a symmetry-breaking phase transition in a field coupled to the inflaton. The inflaton mass then changes suddenly when each transition occurs. These steps are responsible for unusual inflaton dynamics, often represented as a hybrid inflation model \\cite{Adams:2001vc}, and they will create oscillating features in the primordial power spectrum. Oscillations can also be directly imprinted on the inflaton potential itself, due to some trans-Planckian physics \\cite{Brandenberger:2000wr,Easther:2002xe,Martin:2003sg}.  Finally, some other mechanisms may create features in the primordial power spectrum \\cite{Starobinsky:1992}. There may also be empirical motivations to consider more complicated potentials, as several CMB analyses suggest that the CMB power spectrum may be better fit by primordial power spectra with features than by smooth power spectra \\cite{Shafieloo:2003gf,Kogo:2003yb,Sealfon:2005em,Verde:2008ix}.  In this paper, we consider a class of inflationary models that feature periodic oscillations imprinted on a smooth inflaton potential, which then give rise to oscillations in the primordial power spectrum.  We look for these oscillations in the Wilkinson Microwave Anisotropy Probe (WMAP) data, and then determine the smallest oscillation amplitude that will be probed with forthcoming experiments.  Our work is somewhat similar to that in Ref.~\\cite{Covi:2006ci}, which considers oscillations in the primordial power spectrum that arise from a step in the inflaton potential, but differs in that we consider wiggles in the inflaton potential itself. There is also related work in Ref.~\\cite{Hunt:2004vt}, which considers oscillations in the CMB power spectrum from a rapid phase transition during inflation. However, the work most closely related to ours is that in Ref.~\\cite{Wang:2002hf}, which considers oscillations in the inflaton potential in natural-inflation models.  While they focused on constraints from existing data, we forecast also the detectability with future measurements.  The plan of this paper is as follows. Section \\ref{S:themodel} introduces the model and discusses the calculation of the power spectrum.  Section \\ref{S:powerspectra} presents numerical results for the primordial, matter, and CMB power spectra. Section \\ref{S:wmapsearch} presents results of our search for oscillations in the WMAP data.  Section \\ref{S:detectability} discusses the forecasts for the detectability of oscillations in future experiments; we consider here both Planck and a post-Planck cosmic-variance limited experiment.  We summarize and provide some concluding remarks in Section \\ref{S:conclusions}. ",
        "conclusions": "\\label{S:conclusions}  The purpose of our paper has been to study whether features in the inflaton potential, beyond the standard slow-roll parameterization, can be detected with current and forthcoming CMB experiments.  To do so, we have considered the possibility that the inflaton potential varies sinusoidally with the inflaton.  A sinusoidal variation of $V(\\phi)$ with $\\phi$ induces a variation in the matter and CMB power spectra that is sinusoidal in $\\log k$ and $\\log \\ell$, respectively.  Our analysis of the WMAP data indicates that the upper bound on $\\alpha$, the oscillation amplitude, is roughly $3\\times 10^{-5}$.  To derive this bound, we argued that the degeneracy between $\\alpha$ and other cosmological parameters is weak, and then assumed in the data analysis that the other parameters were fixed.  We anticipate that a more complete analysis, including marginalization over other parameters, will weaken this bound, but only slightly.  We then showed that the sensitivity of Planck to the oscillation amplitude will be greater, relative to WMAP, by roughly an order of magnitude.  For both Planck and WMAP, the constraint comes primarily from the temperature. Planck might be improved upon slightly by a cosmic-variance limited experiment. The advantage of such an experiment is that the additional information in the polarization may then be fully capitalized upon.  On the theory side, we used the horizon-crossing approximation to calculate the power spectrum.  The work of Ref.~\\cite{Wang:2002hf} suggests that numerical integration of the perturbation equations will result in a slight suppression of the oscillation amplitude, relative to that we have obtained, but certainly by no more than a factor of two.  If so, then our bounds may be accordingly weakened.  We leave investigation of these data-analysis and theory questions for future research."
    },
    "0807/0807.0863_arXiv.txt": {
        "abstract": "{The mixing length theory (MLT) used to compute the temperature gradient in superadiabatic layers of stellar (interior and atmosphere) models contains in its standard form 4 free parameters. Three parameters are fixed a priori (and define what we denote as the MLT 'flavour') whereas one (the so-called mixing length) is calibrated by reproducing observational constraints. The \\lq{classical}\\rq  B\\\"ohm-Vitense flavour is used in all modern MLT-based stellar model computations and, despite its crude approximations, the resulting $T_{eff}$ scale appears -- perhaps surprisingly -- remarkably realistic, once the mixing length parameter is calibrated with a solar model.} {Model atmosphere computations employ parameter choices different from what is used in stellar interior modelling, raising the question of whether a single MLT flavour and mixing length value can be used to compute interiors and atmospheres of stars of all types. As a first step towards addressing this issue, we study whether the MLT flavour (the so-called ML2) and mixing length choice that have been proven adequate to model white dwarf atmospheres, is able to provide, when used in stellar models, results at least comparable to the use of the \\lq{classical}\\rq  B\\\"ohm-Vitense flavour.} {We have computed solar models and evolutionary tracks for both low- and intermediate-mass Population I and II stars, adopting both solar calibrated B\\\"ohm-Vitense and ML2 flavours of the MLT in our stellar evolution code, and state-of-the-art input physics. } {The two sets of models provide consistent results, with only minor differences. Both calibrations reproduce also the $T_{eff}$ of red giants in a sample of Galactic globular clusters. The ML2 solar model provides a mixing length about half the value of the local pressure scale height, thus alleviating -- but not eliminating -- one of the well known inconsistencies of the MLT employed in stellar models. This mixing length is also consistent with the value used in white dwarf model atmosphere computations.}{} ",
        "introduction": "The equations still largely employed to determine the value of the temperature gradient in the superadiabatic regions of stellar convective envelopes and atmospheres are based on the so called ``mixing length theory'' (MLT). The MLT in its original form is a simple, local, time independent model, firstly applied to stellar modelling by \\citet{bie32}. Its description of convective transport considers the motion of 'average' convective cells, all with the same physical properties at a given radial location $r$ within the star. These convective cells have a mean size $l$ ($l=\\alpha_{MLT} H_P$ is the so-called 'mixing length', where $\\alpha_{MLT}$ is a free parameter and $H_P$ is the local pressure scale height $H_P\\equiv P/g \\rho$) also equal to their mean free path, and an average convective speed ${\\rm v_c}$ at a given value of $r$. There is in principle no reason why $\\alpha_{MLT}$ should be kept constant when considering stars of different masses and/or chemical composition and/or at different evolutionary stages; even within the same star $\\alpha_{MLT}$ might in principle vary from layer to layer, although stellar models usually keep $\\alpha_{MLT}$ constant throughout the convective zone \\citep[but see][for some low-mass stellar models computed with varying $\\alpha_{MLT}$ within the convective envelope]{ped90}. Keeping $\\alpha_{MLT}$ constant still means a change of the cells' mean size and free path due to the variation of $H_p$.  The value of $\\alpha_{MLT}$ is usually calibrated by reproducing some empirical constraint. In case of stellar models, it is typically the calibration of a theoretical solar model that determines the value of $\\alpha_{MLT}$. In all major stellar model databases \\citep[e.g.][]{g00, vdb00, yi01, pcsc04, d07} this solar calibrated $\\alpha_{MLT}$ is kept fixed in all evolutionary calculations of stars of different masses and chemical compositions. There are 3 additional free parameters entering the MLT that are generally fixed a priori, before the calibration of $\\alpha_{MLT}$. They will be denoted here as $a, b, c$ following the formalism of \\citet{tas90}. These three parameters plus $\\alpha_{MLT}$ enter the equations that determine ${\\rm v_c}$, the average convective flux $F_c$ and the convective efficiency $\\Gamma$ (defined as the ratio between the excess heat content of a raising convective bubble just before its dissolution, and the energy radiated during its lifetime) at each point within the convective envelope. Once $a, b, c$ are fixed, the temperature gradient at a given location in the convective envelope of a stellar model and the resulting $T_{eff}$, depend on the value of $\\alpha_{MLT}$. On the other hand, when $\\alpha_{MLT}$ is kept fixed and one or more of the other three parameters are varied, the temperature gradients and $T_{eff}$ are also affected \\citep[e.g.][]{heny65}.  The 'classical' formulation of the MLT by \\citet{bv58} that is nowadays almost universally used in theoretical computations of stellar evolution models (denoted here as the BV58 'flavour' of the MLT) employs a specific set of choices for the values for $a, b$ and $c$. Solar calibrations with the present generation of stellar input physics provide a value of $\\alpha_{MLT}$ typically around $\\sim$2.0 \\citep[e.g.][]{vdb00, pcsc04, bsb06, d07, ml08} the precise value (for a given set of input physics) being dependent on the choice for the boundary conditions at the model surface.  The situation for model atmosphere computations is different. The model grids perhaps most used to produce synthetic spectra and bolometric corrections for evolutionary phases before the White Dwarf (WD) stage are the ATLAS9, MARCS and PHOENIX ones. The recent ATLAS9 grid described in \\citet{pcsc04} employs $\\alpha_{MLT}$=1.25, while the latest MARCS models \\citep{gu08} adopt a different value $\\alpha_{MLT}$=1.5. There are some differences between the MLT flavours used in thse two sets of models, and also compared to the BV58 formalism used in stellar evolution computations \\citep[see][]{cas97,gu08}. ATLAS9 models with a much smaller $\\alpha_{MLT}$=0.5 have also been computed \\citep{hei02}. The recent \\citet{bh05} PHOENIX grid employs $\\alpha_{MLT}$=2.0 and (as communicated to us by our referee) the BV58 flavour of the MLT.  Yet another flavour of the MLT is employed in calculations of model atmospheres, spectra and bolometric corrections for WD stars \\citep[e.g.][]{bwb95, rohr01}. It is the so-called ML2 \\citep{tas90}, with its own specific choice of $a, b$ and $c$, different from the case of BV58. \\citet{bwl95} have shown how for the case of ZZ~Ceti stars, the ML2 with $\\alpha_{MLT}$=0.6 provides overall consistency between temperature estimates from both UV and optical spectrum, observed photometry, gravitational redshift mass estimates and trigonometric parallax. The exact form of the MLT employed in WD interior models seem to be irrelevant \\citep[see, e.g.][]{fbb01} so that effectively one can employ the ML2 in computations of both convective atmospheres and envelopes of WDs.  This variety of choices regarding the 4 free parameters of the MLT should probably not be surprising. As already stated before, given the very approximate nature of the MLT, one cannot a priori expect the same choice of $\\alpha_{MLT}$ -- and possibly also of $a, b$ and $c$ -- to be appropriate for all evolutionary phases in both convective envelope and atmosphere of stars. It is however interesting, if only for heuristic purposes, to investigate whether a single flavour of the MLT with just one choice (or a small range of values) of $\\alpha_{MLT}$ is adequate to model convective envelopes and atmosphere covering all major evolutionary phases. Here we present a first step in this direction, where we have selected as reference choice the ML2. Our aim is to study whether a solar calibration with the ML2 provides a value of $\\alpha_{MLT}$ close to what is employed in WD studies, how this 'new' solar model compares with a BV58 solar model, and how evolutionary tracks computed with the solar calibrated ML2 compare with the BV58 calibration and with empirical determinations of $T_{eff}$ for Red Giant Branch (RGB) stars in Galactic globular clusters. Based on the results by \\citet{gou76} and \\citet{ped90} we should expect that stellar models computed with different choices for $a$, $b$ and $c$ are largely equivalent once $\\alpha_{MLT}$ is appropriately recalibrated (although this has not yet been verified along the upper RGB) but we definitely do not know whether the same value of $\\alpha_{MLT}$ employed in WD model atmosphere studies is adequate also for stellar modelling with the ML2.  We wish to stress that the aim of this note is not to argue for the superiority of the MLT over alternative and less crude prescriptions for stellar convection \\citep[see, e.g.][]{canu91}, rather to investigate whether, within the framework of the MLT -- that in spite of all its intrinsic limitations is still widely used in stellar evolution and model atmosphere calculations -- and with state-of-the-art physics inputs, it is possible to find a single combination of parameters suitable for computing interior and atmosphere models of all types of stars with superadiabatic convective layers. With 'suitable' we mean that its use produces $T_{eff}$, spectral energy distributions and line profiles consistent with observations. The precise numerical value of $\\alpha_{MLT}$ we derive for a given MLT flavour, is tied to the specific set of input physics adopted in the model computation. We employ what we believe are among the best possible physically motivated choices for these inputs. Inaccuracies in the current generation of inputs for stellar models (and model atmosphere) may of course impact the calibration of $\\alpha_{MLT}$.  Section ~2 summarizes the MLT formalisms and the choices of $a, b$ and $c$ in its ML2 and BV58 flavours. A comparison of solar stellar models, low mass stars in the metal poor regime and intermediate mass models computed with these two prescriptions for convection is discussed in Sect.~3, together with the comparison with $T_{eff}$ estimates of RGB stars in globular clusters. Conclusions follow in Sect.~4. ",
        "conclusions": "In the previous section we have shown that solar calibrated BV58 and ML2 stellar models computed with the input physics and boundary conditions described in Sect.~\\ref{model}, are very similar. There are some differences in the physical structures of the superadiabatic layers -- as expected -- together with minor variations in $T_{eff}$, but the extension of the convective regions are the same. Effective temperature differences appear essentially at the RGB and slowly increase towards the RGB tip, reaching at most 40-50~K at the lowest temperatures. This latter result does not conform exactly to \\citet{ped90} conclusion that models computed with different choices for the MLT parameters $a$, $b$ and $c$ should be equivalent once $\\alpha_{MLT}$ is recalibrated, but the differences in the RGB $T_{eff}$ are not large.  Both ML2 and BV58 solar calibrated models provide a good match to $T_{eff}$ estimates for bright RGB stars in Galactic globular clusters. ML2 models appear to perform marginally better when observation-theory differences for each cluster are compared, but the effect of changing the MLT flavour is very small compared to both the rms scatter of the individual values, and realistic estimates of systematic errors in the RGB $T_{eff}$ scale.  In solar models, we have compared in some detail the sound speed, T, $\\rho$ stratifications in the convective layers. If the ML2 improves the agreement between predicted and observed helioseismological properties \\citep[e.g. through the comparison of the solar $p$ mode excitation rates, see][]{sk06} is a matter to be established. On the other hand, due to its crudeness, we do not expect the MLT to be the adequate tool to model local details of superadiabatic convection to the degree of accuracy needed to satisfy helioseismological constraints on the structure of the solar superadiabatic layers.  A crucial point is that $\\alpha_{MLT}$ in the ML2 solar calibration obtained with our code and the adopted state-of-the-art stellar input physics is equal to 0.63, a value consistent with the calibration preferred in WD model atmosphere calculations, that deal with a very different regime of gas properties. The exact value of the solar $\\alpha_{MLT}$ may surely be subject to some variations, when different sources for the input physics are employed, but we expect that with nowadays best (i.e. most accurate and updated) possible choices, values around $\\alpha_{MLT}\\sim 0.6$ should still be found.  To verify whether this unification of the MLT treatment can be pushed even further, it will be necessary -- as a next step --  to compute and test model atmospheres (plus spectra and bolometric corrections) for the pre-WD phases with surface convection, using the ML2 formalism with $\\alpha_{MLT}\\sim$0.6. These ML2-based model atmospheres, when computed and as a further test of this unified MLT treatment, should then be also employed to determine the boundary conditions of interior models, for studying the consistency with results obtained with a (solar) $T(\\tau)$ integration."
    },
    "0807/0807.2057_arXiv.txt": {
        "abstract": "Accurate and homogeneous atmospheric parameters ($T_{\\rm eff}$, $\\log g$, $V_{\\rm t}$, [Fe/H]) are derived for 74 FGK non-variable supergiants from high-resolution, high signal-to-noise ratio, echelle spectra. Extremely high precision for the inferred effective temperatures (10--40 K) is achieved by using the line-depth ratio method. The new data are combined with atmospheric values for 164 classical Cepheids, observed at 675 different pulsation phases, taken from our previously published studies. The derived values are correlated with unreddened {\\it B--V} colours compiled from the literature for the investigated stars in order to obtain an empirical relationship of the form: $(B-V)_{0} = 57.984 - 10.3587(\\log T_{\\rm eff})^{2} + 1.67572(\\log T_{\\rm eff})^{3} - 3.356\\log g + 0.0321V_{\\rm t} + 0.2615[{\\rm Fe/H}] + 0.8833(\\log g)(\\log T_{\\rm eff})$. The expression is used to estimate colour excesses E({\\it B--V}) for individual supergiants and classical Cepheids, with a precision of $\\pm0.05$ mag. for supergiants and Cepheids with $n=1-2$ spectra, reaching $\\pm0.025$ mag. for Cepheids with $n>2$ spectra, matching uncertainties for the most sophisticated photometric techniques. The reddening scale is also a close match to the system of space reddenings for Cepheids. The application range is for spectral types F0--K0 and luminosity classes I and II. ",
        "introduction": "The Cepheid period-luminosity (PL) relation continues to play a key role in the determination of distances within the Local Group and to nearby galaxies. The absolute calibration of the PL relation requires reliable distance measurements for calibrating Cepheids as well as accurate corrections for the effects of interstellar reddening and extinction. Galactic Cepheids are often heavily obscured, the average colour excess E({\\it B--V}) being of order 0.5 mag. Here we propose a new method of accurate colour excess determination that relies on spectroscopically determined stellar parameters.  The traditional methods of establishing the interstellar reddening of Cepheids have involved field reddenings derived from photometric and spectroscopic studies of stars surrounding the Cepheid, the use of reddening-independent spectroscopic indices, and period-colour relations calibrated from spectroscopic analyses of a number of well-studied Cepheids or supergiants. Field reddenings are fairly reliable provided that the surrounding stars are unaffected by circumstellar reddening, although the number of suitable objects has tended to be restricted to cluster Cepheids and visual binaries \\citep[see][]{tu95,tu01,lc07}. Reddening-independent indices include the $\\Gamma$ index \\citep{kr60,sj89}, which measures the depression in stellar spectra caused by the {\\it G} band of CH ($\\lambda 4305$), and the {\\it KHG} index \\citep{mp69,mc70,fe72,tu87}, which combines narrow band photometry of Ca II K ($\\lambda 3933$), Balmer H$\\delta$ ($\\lambda 4101$), and the {\\it G} band. Narrow band photometry of Balmer H$\\beta$ ($\\lambda 4861$) has also been used as a temperature indicator for FGK stars, but has been susceptible to problems arising from colour dependences introduced by mismatches in the effective wavelengths of the narrow and wide H$\\beta$ filters employed \\citep{st79}.  Most Cepheid reddenings have been derived from period-colour relations that are tied to specific Cepheid or supergiant calibrators, in which specific continuum or spectral line features (line ratios, for example) are linked to stellar atmosphere models to infer effective temperatures for the stars \\citep[e.g.,][]{fe87,sl90}. The direct link between effective temperature and broad band ({\\it B--V})$_0$ colour has been demonstrated very effectively by \\citet{gr92} for older, published, stellar atmosphere models, and has been used successfully to test the open cluster calibration of the PL relation \\citep{tb02}. As argued by \\citet{fc99}, {\\it B--V} colours for Cepheids appear to be more closely indicative of stellar temperatures than indices that include a near-infrared magnitude. But recent changes to stellar atmosphere models \\citep{ku92} may affect the effective temperature scale for FGK supergiants significantly enough that a recalibration is necessary. That is the purpose of the present study.  Cepheids also undergo significant changes in effective temperature during their pulsation cycles, which makes it essential to track such changes accurately through close monitoring of individual variables. Photometric monitoring of Cepheids has resulted in a large database of published results for many different colour systems \\citep{be07}, but the availability of high resolution spectra for such stars throughout their cycles has been extremely limited, until now. A goal of the present study is therefore to exploit a recently developed collection of high resolution spectra for Galactic Cepheids in order to present accurate intrinsic parameters for all of the studied objects. ",
        "conclusions": "The present study presents newly-derived parameters, namely effective temperature ($T_{\\rm eff}$), surface gravity ($\\log g$), microturbulent velocity ($V_{\\rm t}$, in km s$^{-1}$), and iron abundance ([Fe/H]), for 74 non-variable FGK supergiants, established from model atmosphere analyses of high resolution spectra for the stars. Colour excesses have been computed for the stars on the basis of a new formulation of the relationship linking such parameters to intrinsic colour, ({\\it B--V})$_0$. The formulation has been extended to the 74 FGK supergiants and to a sample of 164 classical Cepheid variables to derive new estimates for their colour excesses E({\\it B--V}), presented in Tables 1 and 2. The reddening estimates are demonstrated to be of extremely high internal precision, and to agree well with the most accurate estimates from the literature. The Cepheid reddenings, in particular, appear to be closely tied to the system of space reddenings that is presently available. Given the large distances of supergiants, the method opens the possibility for large-scale extinction mapping of the Galaxy, with a sensitivity of 0.08--0.2 magnitude.  \\subsection*{ACKNOWLEDGEMENTS} This work is based on spectra collected with the 1.93-m telescope of the OHP (France) and the ESO Telescopes at the Paranal Observatory under programme ID266.D-5655. We thank the referee for useful suggestions that helped to improve the presentation."
    },
    "0807/0807.2782_arXiv.txt": {
        "abstract": "Absolute flux distributions for eight stars are well measured from 0.8--2.5~$\\mu$m with NICMOS grism spectrophotometry at a resolution of R~$\\sim100$ and an accuracy of 1--2\\%. These SEDs are fit with Castelli \\& Kurucz model atmospheres; and the results are compared with the Cohen-Walker-Witteborn (CWW) template models for the same stars. In some cases, the $T_\\mathrm{eff}$, $\\log g$, and $\\log z$ parameters of the best fitting model differ by up to 1000~K from the earlier CWW model. However, differences in the continua of the modeled IR flux distributions from 0.4--40~$\\mu$m are always less than the quoted CWW uncertainty of 5\\% because of compensating changes in the measured extinction. At wavelengths longward of the 2.5~$\\mu$m NICMOS limit, uncertainties still approach 5\\%, because A-star models are not yet perfect. All of these A~stars lie in the \\emph{JWST} continuous viewing zone and will be important absolute flux standards for the 0.8--30~$\\mu$m \\emph{JWST} wavelength range. ",
        "introduction": "The \\emph{James Webb Space Telescope} (\\emph{JWST}) is a NASA flagship mission and requires flux standards in its continuous viewing zones that are located near the ecliptic poles. The \\emph{JWST} instrumentation covers the 0.8--30~$\\mu$m region; and for reference flux distributions at the longer wavelengths, stellar model atmospheres have traditionally been used to extrapolate from the visible into the mid-IR. The nine A-star standards chosen for \\emph{JWST} are established by Cohen et~al.\\ (2003a, CMHMS) as an extension of the original CWW standard star network (Cohen, Walker, \\& Witteborn 1992; Cohen 2007). These nine stars are fainter than the \\emph{JWST}/NIRSpec saturation limit and have already been observed by the Infrared Array Camera (IRAC) on the \\emph{Spitzer Space Telescope}; two of these nine (HD~165459 and 1812095) are primary standards used for the IRAC flux calibration (Reach et~al.\\ 2005).  In order to compare and cross calibrate \\emph{HST} and the IRAC \\emph{Spitzer} absolute flux scales, the nine \\emph{JWST} standards have been observed by the NICMOS objective grism spectrometer in the 2006--2008 time frame. NICMOS spectra at a resolution R$\\,\\sim$100 are obtained in three modes G096, G141, and G206, which together cover the 0.8--2.5~$\\mu$m range. The absolute flux calibration of these spectra is established on the \\emph{Hubble Space Telescope} (\\emph{HST}) white dwarf (WD) flux scale by observations of the primary pure hydrogen WD stars GD71, GD153, and G191B2B (Bohlin 2000, Bohlin, Dickenson, \\& Calzetti 2001, Bohlin 2003, Bohlin \\& Koester 2008) in these three NICMOS modes. The NICMOS count rate spectrophotometry is extracted from the 256$\\times$256 images per the prescription of Bohlin, Lindler, \\& Riess (2005). The ratios of the count rate spectra for the A--stars to the count rates for the primary WD flux standards determine the absolute fluxes of the A--stars after correcting for the non-linearity of the NICMOS HgCdTe array (Bohlin, Riess, \\& de~Jong 2006; Bohlin 2007). The fluxes for the three NICMOS modes are combined into one measured SED covering the full 0.8--2.5~$\\mu$m range. The modeled flux distributions of the  primary WDs are defined by the stellar temperatures and gravities, which are derived from the Balmer line profiles. As a typical example, the uncertainty of 3000~K in the 61193~K effective temperature of G191B2B means that its relative flux should be correct to better than 1\\% from 0.35 to 2.5~$\\mu$m. ",
        "conclusions": "Absolute flux standards are required for \\emph{JWST} calibration in the \\emph{JWST} continuous viewing zone near the north ecliptic pole. NICMOS grism observations measure the flux distributions of eight A~stars, seven of which fit CK04 model atmospheres from 0.82--2.4~$\\mu$m within an rms scatter of $<$0.8\\%. These models are normalized to the NICMOS SEDs and are used to extend the measured NICMOS fluxes to shorter and longer wavelengths. These seven composite NICMOS plus model SEDs are archived in the CALSPEC database of \\emph{HST} flux standards\\footnote{http://www.stsci.edu/hst/observatory/cdbs/calspec.html/} as \\emph{*\\_nic\\_002.fits}. The eighth well measured NICMOS flux distribution for 1812524 is also included in the same CALSPEC web page as \\emph{1812524\\_nic\\_002.fits} but does not have its wavelength range extended beyond the NICMOS range."
    },
    "0807/0807.1865_arXiv.txt": {
        "abstract": "We present an analysis of the foreground emission present in the \\WMAP\\ 3-year data as determined by the method of Independent Component Analysis. The \\WMAP\\ data averaged at each frequency are used together with the standard foreground emission templates as inputs to the \\fastica\\ algorithm. The returned coefficients can be interpreted as coupling coefficients between the \\WMAP\\ data and foreground templates.  These results are then used to infer the spectral behaviour for three foreground components -- synchrotron, anomalous dust-correlated emission and free-free. For the first two components, we find values consistent with previous results although slightly steeper. We confirm the inconsistency in the scaling between the H$\\alpha$ template and free-free emission at K- and Ka-bands where an electron temperature of $\\sim 4000$~K is indicated. We also see evidence of significantly flatter spectral behaviour to higher frequencies than expected theoretically and previously noted by \\citet{dobler_etal_2008a}, but only when analysing the $Kp2$ sky coverage.  We further apply \\fastica\\ \\lq iteratively', using data pre-cleaned using foreground templates scaled to the \\WMAP\\ frequencies by coupling coefficients determined by a prior \\fastica\\ analysis.  This multi-frequency analysis allows us to determine the presence of residual foreground emission not traced by the templates. We confirm the existence of a component spatially distributed along the Galactic plane and particularly enhanced near the center (the \\lq \\WMAP\\ haze'). This emission is less extended when using the \\WMAP\\ K-Ka data as the synchrotron template confirming that it can be considered a better template for foreground cleaning of the \\WMAP\\ data. However its use complicates the physical interpretation of the nature of the foreground emission and residuals since it contains a mixture of several, physically distinct emission mechanisms.  The good agreement between the extracted CMB component and previous results, as well as the low amplitude of residual foreground emission make \\fastica\\ a viable tool to infer foreground emission properties, via template fitting, and the CMB amplitude. ",
        "introduction": "\\label{intro}  Measurements of the Cosmic Microwave Background (CMB) provide cosmological information by probing the early universe some 380,000 years after the Big Bang. However their precision is limited by the contamination of the cosmological signal due to emission from local astrophysical sources such as our own Galaxy. In particular the Galactic contamination arises from three well-understood processes: synchrotron emission, free-free emission (or \\emph{thermal bremsstrahlung}) and thermal dust emission. In addition, there appears to be an enhanced dust correlated emission at microwave wavelengths that is inconsistent with the thermal mechanism, and the exact nature of which is still open to debate. Theoretically, these components can be distinguished on the basis of their own spectral properties and spatial distribution. The best window for observing the temperature anisotropies lies close to 70 GHz, where the integrated diffuse foreground emission seems to display a minimum. This has permitted accurate mapping of the CMB intensity pattern, without the need to excise too much sky coverage from the analysis due to these foregrounds (except in regions near to the Galactic plane and associated with bright point sources). However, to reach the accuracy required for precision cosmology ($\\sim \\mu$K), component separation techniques and/or foreground cleaning methods are needed.  A component separation technique is one that attempts to decompose multi-frequency input data into separate physical emission components, one of which, for our purposes, must be the CMB. However, the technique may not be successful particularly in regions of the sky where the foreground spectral dependencies are complex, or where the physical emissions have complex morphology and are well-mixed. A foreground cleaning method in the context of this paper simply attempts to generate a map of the CMB as clean of foreground emission as possible, with no concern for the detailed separation of the foreground emission into physical components. It should be clear that a component separation algorithm may have limited success, providing an adequate reconstruction of the CMB but not of the foreground emission, then effectively consituting a foreground cleaning method.   In this paper we will consider the detailed application of one technique -- \\emph{Independent Component Analysis} (ICA) -- to data from the \\WMAP\\footnote{\\emph{Wilkinson Microwave Anisotropy Probe}} satellite, focusing our attention on the properties of diffuse Galactic foregrounds at high and intermediate latitude as traced by specific templates of the emission. Of course, \\fastica\\ is only one of a number of methods that have been developed for the application to data at microwave wavelengths in order to (at least) extract a map of the CMB sky with minimal foreground residuals. An inexhaustive list of other techniques applied to the \\WMAP\\ data includes linear combination methods \\citep{bennett_etal_2003,tegmark_etal_2003,park_etal_2007}, WIFIT \\citep{hansen_etal_2006}, FGFit \\citep{eriksen_etal_2006}, SMICA \\citep{patanchon_etal_2005} and CCA \\citep{bonaldi_etal_2007}. The \\WMAP\\ team has also made use of the Maximum Entropy Method (MEM) \\citep{bennett_etal_2003,hinshaw_etal_2007} to study the detailed properties of the Galactic foregrounds. However, since MEM is not able to produce a CMB map that is easy to use for cosmological analyses, the \\WMAP\\ data was cleaned from foregrounds by applying a template fitting algorithm -- essentially a cross-correlation method also used by \\citet{davies_etal_2006} to study regional variations in the foreground emission.   A general summary of the paper is as follows. We first provide a simple introduction to ICA and describe the use of templates of foreground emission in such an analysis, making a comparison to more standard template fitting techniques. We then describe the data used in this paper for the foreground analysis. Then we calibrate the performance of the method using Monte Carlo simulations and by comparison to the simpler template based method. Sections~\\ref{haslam_analysis} and \\ref{K-Ka_analysis} then attempt to improve our knowledge of the spectral and physical properties of the foreground components based on the scaling properties of the templates. An assessment of the efficiency of the method is then afforded by studying the power spectrum of the cleaned CMB data, whilst a final iterative application of ICA attempts to improve the removal of foreground contamination from the data. Section~\\ref{discussion} summarises our main conclusions. ",
        "conclusions": "\\label{discussion}  In this paper, we have utilised a blind component separation algorithm -- {\\fastica} -- to address the issue of the Galactic foreground emission in the \\WMAP\\ 3-year data. In particular, we have adapted the technique to allow three Galactic foreground templates (that trace synchrotron, free-free and thermal dust emission) to be fitted to the data on a frequency-by-frequency basis.  Specifically, we use the \\citet{finkbeiner_2003} H$\\alpha$-map as a template for the free-free emission, the \\citet{finkbeiner_etal_1999} FDS8 model for thermal dust emission, and for synchrotron emission two alternatives -- the 408~MHz radio continuum all-sky map of \\citet{haslam_etal_1982} as utilised in the first year \\WMAP\\ analysis, or the difference of the \\WMAP\\ K- and Ka-band data preferred in their three year analysis.  Detailed simulations indicate that the algorithm behaves in an unbiassed way only for two ($p$, $g$) of the three non-linear functions typically applied by {\\fastica}. In these cases, the $p$-function also appears to be marginally preferred as a consequence of its smaller error bars, both relative to the $g$-function analysis and the simple $\\chi^2$ method that is usually adopted for template fitting.  When applied to the \\WMAP\\ data, we have quantified the foreground contamination in term of coupling coefficients between the data and the foreground templates. The derived values are sensitive to both the extension of the mask applied to exclude the Galactic plane and on the non-linear function used by the algorithm.  The mask dependence is likely to reflect genuine variations in the spectral behaviour of the foregrounds with sky location, as seen in previous local studies of template fits by \\citet{davies_etal_2006}. The non-linear function dependence is plausibly linked to sensitivity to the actual statistical nature of the foreground components (as represented by the templates), and suggests that it may be useful to study more appropriate approximations to the neg-entropy.  We have considered the spectral behaviour of the derived scaling factors when the \\citet{haslam_etal_1982} data is used as the synchrotron template.  We evaluated the spectral index for the synchrotron emission, the anomalous dust-correlated component, and the free-free emission.  In the first two cases, we found steeper, though statistically consistent, spectral behaviour as compared to previous analysis, eg. \\citet{davies_etal_2006}.  For the free-free emission, we confirmed previous observations about inconsistencies in the fitted scaling to convert the H$\\alpha$ emissivity to a free-free contribution at K- and Ka-bands.  Specifically, the derived factors correspond to thermal electron temperatures of $\\sim 4000\\, K$ compared to the expected value of $8000\\, K$.  In addition, we find independent verification of a flat spectral dependence in the free-free emissivity as reported by \\citet{dobler_etal_2008a}. However, this behaviour is seen only for the $Kp2$ sky coverage. With the $Kp0$ mask, the behaviour is even steeper than the expected $\\nu^{-2.14}$ dependence, suggesting that the anomalous free-free behaviour is associated with bright structures close the Galactic plane.  We have also proposed that the {\\fastica} algorithm can be applied \\lq iteratively', that is, used to analyse a set of multi-frequency maps pre-cleaned using templates for which the coefficients have themselves been derived by a {\\fastica} analysis. This is particularly useful in order to determine the presence of residual foregrounds that arise either due to a mismatch between their spectral dependence and the average high-latitude value determined by the template fits, or because they are new components that are not traced by the adopted templates. In this manner, we have confirmed the existence of a component spatially distributed along the Galactic plane, with pronounced emission near the Galactic center. This is the emission previously noted in the SMICA analysis of \\citet{patanchon_etal_2005}, and is the \\lq \\WMAP\\ haze' of \\citet{finkbeiner_2004}. However, the amplitude and extent of the emission is less when the K-Ka map is adopted as the synchrotron template. This is not surprising if the foreground component is a genuinely new contribution to the sky emission that is untraced by the Haslam data. In this context, the K-Ka map must be considered the better template for cleaning the \\WMAP\\ data from foregrounds. However, the interpretation of the results in terms of Galactic emission components is complicated by the fact that the K-Ka template contains a mixture of several, physically distinct, emission mechanisms.  Finally, the analysis that we have performed with the {\\fastica} algorithm is based on the unrealistic hypothesis that the spectral behaviour of the various foreground emission components is unchanging over the sky. In future work, in order to take into account the spatial variation of the foreground spectra, we will attempt to apply {\\fastica} on smaller patches of the sky where the assumption of uniform spectral behaviour is more realistic. Given the success of our method in demonstrating the presence of residual foregrounds surviving template-based corrections, we can perhaps utilise the information (ie. the residuals maps) from a near-global approach to identify significant regional spectral variations for such a local analysis."
    },
    "0807/0807.4737_arXiv.txt": {
        "abstract": "ACO~1703 is a cluster recently found to have a variety of strongly lensed objects: there is a quintuply-imaged system at $z=0.888$ and several other lensed objects from $z=2.2$ to 3.0 (the cluster itself is at $z=0.28$).  It is not difficult to model the lens, as previous work has already done.  However, lens models are generically non-unique. We generate ensembles of models to explore the non-uniqueness.  When the full range of source redshifts is included, all models are close to $\\rho\\propto r^{-1}$ out to $200\\kpc$.  But if the quint is omitted, both shallower and steeper models (e.g., $\\rho\\propto r^{-2}$) are possible.  The reason is that the redshift contrast between the quint and the other sources gives a good measurement of the enclosed mass at two different radii, thus providing a good estimate of the mass profile in between.  This result supports universal profiles and explains why single-model approaches can give conflicting results.  The mass map itself is elongated in the NW-SE direction, like the galaxy distribution.  An overdensity in both mass and light is also apparent to the SE, which suggests meso-structure. ",
        "introduction": "An important prediction of hierarchical clustering of cold dark matter is a universal profile for halos. Universal profiles were originally indicated by cosmological simulations \\citep{1997ApJ...490..493N,1998ApJ...499L...5M} but more recently phenomenological models have also produced them \\citep[e.g.,][]{2005ApJ...634..756A,2006MNRAS.368.1931L}, as have more specialized simulations \\citep{2006ApJ...653...43M}.  The precise form of a universal profile and the expected variation from halo to halo are still being debated; \\cite{2006AJ....132.2685M} discuss some competing parametrizations.  A further complication arises in that infalling baryons would steepen the dark-matter profile somewhat (adiabatic contraction), but in large clusters this effect is expected to be small exterior to $\\sim 20\\kpc$ \\citep{2004ApJ...616...16G}. There is general agreement that in the inner regions of clusters the dark matter profile would have $\\rho\\sim r^{\\rm -1\\ to\\ -1.4}$ \\citep{2004MNRAS.353..624D}.  Lensing is an obvious way to measure the profiles of clusters, and in this paper we will address specifically multiple-image (i.e., ``strong'') lensing.  Weak lensing by outer regions of clusters is also of interest \\citep[for a recent example, see][]{2008MNRAS.385.1431H}.  Other possibilities are to use generalized virial theorems and the kinematics of cluster galaxies \\citep{2006MNRAS.367.1463L}, or hydrostatic equilibrium of x-ray gas \\citep{2008MNRAS.386.1092L}.  Combinations of these approaches are also possible.  Extracting a mass profile from lensing is not trivial, however.  The reason is that lensing depends on {\\em integrated\\/} mass.  For example, the light deflection near the Sun $4GM_\\odot/(c^2R_\\odot)=1.75''$ says nothing about the mass profile of the Sun.  Similarly, a single Einstein ring measures only the mass enclosed within a cylinder.  Having two or more images at different radii provides some more information, but is still not sufficient to specify the mass profile, because it is possible to redistribute the mass in a way that leaves the image configuration invariant (steepness degeneracy).  Much more information becomes available if the background sources are at different distances behind the lens.  Light from further sources experiences a larger effective lensing deflection and hence a larger Einstein radius. Multiple Einstein radii, resulting from multiple source redshifts or ``redshift contrast'', measure the mass within multiple cylinders, thus providing a robust constraint on the mass profile.  ACO~1703 turns out to be a striking example of this type of constraint.  The non-uniqueness of mass models that fit a given set of observables has been much studied in lensing theory.  The steepness degeneracy was discovered by \\cite{1985ApJ...289L...1F}, and then rediscovered multiple times; \\cite{2000AJ....120.1654S} gives a short history of these (re)discoveries, and a unified picture of degeneracies then known.  Further degeneracies continue to be uncovered from theory \\citep{2008MNRAS.386..307L}.  But in practice, lensing degeneracies are most likely to be noticed in the modeling process, when different kinds of models are tried out. \\cite{1991ApJ...373..354K} was the first to try a range of models and distinguish robust features from model-dependent features.  Subsequent work found parameter degeneracies appearing in models in many different ways \\citep{1994AJ....108.1156W,1999AJ....118...14B,2000ApJ...535..671T}. Recently \\cite{2008ApJ...674..711S} report non-uniqueness in models of clusters.  Faced with the non-uniqueness of models, three strategies are possible. \\begin{enumerate} \\item One could justify the model-type being fit from astrophysical arguments, and then assume it is correct.  This approach goes back to the first models of the first two known lenses \\citep{1981ApJ...244..723Y,1981ApJ...244..736Y}. With the benefit of hindsight, we find that these pioneering models inferred correctly that the ``triple-quasar'' was really a quad, but predicted time delays for the double quasar several times too long. \\item One could try and introduce new information such as stellar kinematics or x-rays, as indeed \\cite{1981ApJ...244..736Y} advocated, to reduce model-dependence in the results. \\cite{2008MNRAS.384..987C} is a recent example. \\item One can explore model degeneracies and quantify uncertainties by generating large ensembles of models, thereby also identifying the best-constrained systems. \\end{enumerate}  With this background in mind, the recent discovery by \\cite{2008arXiv0802.4292L} of 13 different multiple-image systems in ACO~1703 is very exciting.  Though not a well-known cluster, ACO~1703 is a strong x-ray source \\citep{1992MNRAS.259...67A} and hence a natural place to search for lensing. For modeling, the authors adopted strategy~1 above, and their model is a generalized NFW profile, having \\begin{equation} \\rho = \\frac{\\rho_0}{{\\left(\\frac{r}{r_s}\\right)^\\alpha}\\left(1+\\frac{r}{r_s}\\right)^{3-\\alpha}} \\label{eq-nfwlike} \\end{equation} where $r$ is an elliptical radius. The normalization $\\rho_0$ is defined indirectly in terms of the concentration $c$, such that the average density within $r=cr_s$ is 200 times the critical cosmological density.  The models give $\\alpha\\simeq1.1$, $r_s\\simeq170''$ or $\\simeq 700\\kpc$, and $c\\simeq3$.  These results appear to verify the prediction of a universal profile, but they also raise some questions: (i)~Is the model an acceptable fit to the data, according to a $\\chi^2$ or other goodness-of-fit test?  (ii)~Can other, very different, models provide equally good fits?  (iii)~What do the fitted values of $r_s$ and $c$ really mean, when the data go out to only $\\sim50''$ ($\\sim200\\kpc$)?  In this paper we analyze the same multiple-image data following strategy~3.  The technique is basically the same as in our previous work on the inner profiles of J1004+411 and ACO~1689 \\citep{2006ApJ...652L...5S}. But in this work we study more carefully just where the constraints come from. ",
        "conclusions": "This paper has been an analysis of the multiple-image lensing systems reported by \\cite{2008arXiv0802.4292L} in the cluster ACO~1703, and comparison of the inner profile with what hierarchical structure formation predicts.  The main conclusions are as follows. \\begin{enumerate} \\item It is easy to model the lens, and in simple qualitative features all models will agree.  But models are highly non-unique in important details.  Readers of the online version can explore models interactively: see Appendix \\ref{online}. \\item Robust constraints can, however, be derived from ensembles of models.  We find the inner profile is well constrained and supports the prediction of universal profiles.  Similar work earlier on the clusters SDSS~J1004+411 and ACO~1689 led to the same conclusion \\citep{2006ApJ...652L...5S}. \\item It is possible to identify where the inner-slope constraint comes from.  Behind the cluster at $z=0.28$ there is on the one hand a source at $z=0.888$ lensed into a quint, and on the other hand several sources at $z$ from 2.2 to 3.0 that are variously multiply-imaged.  The redshift contrast between the quint and the other sources is responsible for the inner-slope constraint. Without the quint, any constraint on the inner slope would have been much weaker, and even an isothermal-type profile would fit. \\item The projected-density contours (see Fig.~\\ref{fig-mass}) correlate with the cluster galaxies. Since the mass involved is much more than the stellar mass of the galaxies, this suggests meso-structure, i.e., extended dark-matter substructure correlated with galaxies. \\end{enumerate}  These results raise a further question: How can one most effectively compare a reconstructed lensing cluster with simulations and/or phenomenological models for structure?  A non-parametric test would be very welcome.  We suggest that the deflection angle (and specifically its redshift dependence) may form the basis for such a test.  In particular, we note that the shallower the lensing profile, the more sensitive the lensed images are to source-redshift contrast."
    },
    "0807/0807.1112_arXiv.txt": {
        "abstract": "We report the discovery of a strong magnetic field in the unique pulsating carbon-atmosphere white dwarf SDSS J142625.71+575218.3. From spectra gathered at the MMT and Keck telescopes, we infer a surface field of $B_s$ $\\simeq$ 1.2 MG, based on obvious Zeeman components seen in several carbon lines. We also detect the presence of a Zeeman-splitted He I 4471 line, which is an indicator of the presence of a non-negligible amount of helium in the atmosphere of this Hot DQ star. This is important for understanding its pulsations, as nonadabatic theory reveals that some helium must be present in the envelope mixture for pulsation modes to be excited in the range of effective temperature where the target star is found. Out of nearly 200 pulsating white dwarfs known today, this is the first example of a star with a large detectable magnetic field. We suggest that SDSS J142625.71+575218.3 is the white dwarf equivalent of a roAp star. ",
        "introduction": "Recently, \\citet{dufournat} reported the discovery of a new type of white dwarf stars with an atmosphere composed primarily of carbon with little or no traces of hydrogen or helium (the ''Hot DQ'' spectral type). Prior to that discovery, white dwarfs cooler than $\\sim $ 80,000 K were known to come in essentially two flavors: those with an almost pure hydrogen surface composition (forming the DA spectral type), and those with a helium-dominated surface composition (the non-DA stars, which comprise the DO, DB, DC, DZ, and DQ spectral types). Pulsationally unstable stars are found among these two broad families of white dwarfs, in each case occupying a narrow range of effective temperature. Variable white dwarfs situated in these instability strips are, respectively, classified ZZ Ceti stars (hydrogen atmospheres, $\\Te \\sim$ 12,000 K) and V777 Her stars (helium atmospheres, $\\Te \\sim$ 25,000 K). Since each of these instability regions is associated with the presence of a partial ionization zone of the primary atmospheric constituent (H or He), it was naturally expected that some carbon-atmosphere white dwarfs could be unstable as well in a certain regime of effective temperature corresponding to partial ionization of carbon. And indeed, \\citet{fontaine76}  found long ago strong similarities between the partial ionization regions and associated superficial convection zones of white dwarf models with H-, He-, and C-dominated atmospheres/envelopes. Theoretical considerations thus suggested that some of the Hot DQ white dwarfs, because they are located in a narrow range of effective temperature around 20,000 K \\citep{dufournat,dufour08}, could possibly pulsate. Following this, a systematic search for pulsations in carbon-atmosphere white dwarfs carried out by \\citet{montgomery08} successfully discovered the first pulsating carbon-dominated atmosphere white dwarf: SDSS~J142625.71+575218.3 (hereafter SDSS~J1426+5752). They uncovered a single pulsation (and its first harmonic) with a period of 417 s in that star.  In parallel with this observational effort, \\citet{fontaine08} carried out the first detailed stability study, based on the full nonadiabatic approach, to investigate the asteroseismological potential of carbon-atmosphere white dwarfs. They showed that pulsational instabilities in the range of effective temperature where the Hot DQs are found are indeed possible, but only if a fair amount of helium is present in the atmosphere/envelope compositional mixture. White dwarf models with pure carbon envelopes are found to pulsate, but only at much higher temperatures than those characterizing Hot DQ's discovered up to now. However, the SDSS spectra analyzed in \\citet{dufour08}, even though quite noisy, rule out large amounts of helium (from the absence of the He~\\textsc{i} $\\lambda$4471 line) for all objects except SDSS~J1426+5752 which, because of its higher surface gravity and lower effective temperature, could perhaps have an He/C abundance ratio as high as 0.5. So, according to the full nonadiabatic models, SDSS~J1426+5752 should be the only object pulsating, as observed, provided helium is present in a relatively large amount at the surface. However, even with such an abundance (He/C = 0.5), only a tiny depression at the He~\\textsc{i} $\\lambda$4471 line is predicted by synthetic models. Given the noisy SDSS spectrum for this object, no firm conclusion could be reached concerning the presence of helium, although the above mentioned abundance looked quite compatible with the spectroscopic observation (see the SDSS spectrum and fits in Figure 1).  All this motivated us to obtain higher sensitivity observations for this special object in order (1) to confirm, as may be the case, the presence of a relatively large amount of helium (thus increasing our confidence in the nonadiabatic approach), and (2) to be able to eventually carry out a full asteroseismological analysis using better constraints/determinations of the atmospheric parameters ($\\Te$, $\\logg$ and He/C) from spectroscopy. In this Letter, we thus report new high signal-to-noise ratio spectroscopic observations (MMT and Keck) that revealed, to our great surprise, that SDSS~J1426+5752 is the first pulsating white dwarf showing clear evidence for the presence of a strong magnetic field. This is certainly an unexpected result, given that strong magnetism in white dwarfs is generally thought to extinguish, or at least greatly diminish, pulsational activity. We discuss below the implications of this discovery. ",
        "conclusions": "We presented evidence that SDSS~J1426+5752 is the first pulsating white dwarf with the clear presence of a strong (i.e., strong enough for Zeeman splitting to be observed) magnetic field ($\\approx$ 1.2 MG). Such a strong magnetic field probably inhibits the convective motions in this object, but it is unclear yet how the pulsations are affected. We proposed that this strange object could be a white dwarf analog of the rapidly oscillating Ap stars and that a usual kappa-mechanism, even in the absence of convection, might still explain the pulsational instabilities. The confirmed presence of helium in the envelope/atmosphere of SDSS~J1426+5752 is likely to play a key role in this process. Future work testing this hypothesis is underway."
    },
    "0807/0807.0009_arXiv.txt": {
        "abstract": "We use high spatial resolution observations of CO to systematically measure the resolved {\\em size-line width}, {\\em luminosity-line width}, {\\em luminosity-size}, and the {\\em mass-luminosity} relations of Giant Molecular Clouds (GMCs) in a variety of extragalactic systems. Although the data are heterogeneous we analyze them in a consistent manner to remove the biases introduced by limited sensitivity and resolution, thus obtaining reliable sizes, velocity dispersions, and luminosities. We compare the results obtained in dwarf galaxies with those from the Local Group spiral galaxies. We find that extragalactic GMC properties measured across a wide range of environments are very much compatible with those in the Galaxy. The property that shows the largest variability is their resolved brightness temperature, although even that is similar to the average Galactic value in most sources. We use these results to investigate metallicity trends in the cloud average column density and virial CO-to-H$_2$ factor.  We find that these measurements do not accord with simple predictions from photoionization-regulated star formation theory, although this could be due to the fact that we do not sample small enough spatial scales or the full gravitational potential of the molecular cloud. We also find that the virial CO-to-H$_2$ conversion factor in CO-bright GMCs is very similar to Galactic, and that the excursions do not show a measurable metallicity trend. We contrast these results with estimates of molecular mass based on far-infrared measurements obtained for the Small Magellanic Cloud, which systematically yield larger masses, and interpret this discrepancy as arising from large H$_2$ envelopes that surround the CO-bright cores. We conclude that GMCs identified on the basis of their CO emission are a unique class of object that exhibit a remarkably uniform set of properties from galaxy to galaxy. ",
        "introduction": "Giant Molecular Clouds (GMCs) are the major reservoirs of molecular gas and the sites of most star formation in our Galaxy and other galaxies.  Their properties set the initial conditions for protostellar collapse, and may play a role in determining the stellar initial mass function \\citep{MCKEE07}. Moreover, because GMCs provide the bulk of the material for forming new stars their creation may be the limiting mechanism that regulates the rate of star formation in galaxies. Therefore, increasing our understanding of their properties and distribution throughout the different environments of external galaxies is likely to provide further insights into GMC and stellar formation processes. There is a limited amount of information, however, that can be gained from studies that resolve the general distribution of molecular gas but not the individual molecular clouds. Resolving GMCs --- to measure their sizes, velocity dispersions, and luminosities --- is a critical step in understanding the processes that ultimately drive galaxy evolution.  Studies of resolved molecular clouds in the Milky Way find that GMCs are in approximate virial equilibrium and obey scaling relations, commonly known as Larson laws \\citep{LARSON81}, that have their origin in the character of the turbulence in the interstellar medium \\citep{ELMEGREEN04,MCKEE07}.  Large scale surveys of the Milky Way show that GMCs follow uniform scaling relations \\citep{SOLOMON87,ELMEGREEN96}, with few differences present between those located in the inner and the outer disk \\citep*{HEYER01}.  \\citet{LARSON81} established that velocity dispersion, size, and luminosity are correlated in Milky Way GMCs.  Observations indicate that GMC line widths increase as a power of their radius \\citep{SOLOMON87}, such that  \\begin{equation} \\sigma_v\\approx0.72\\, R^{0.5}\\ {\\rm km~s^{-1}}~, \\label{size_linewidth_eq} \\end{equation}  \\noindent where $\\sigma_v$ is the one-dimensional velocity dispersion of the GMC, and $R$ is its radius measured in parsecs.  Equation \\ref{size_linewidth_eq} is known as the {\\em size-line width} relation, and it has been shown to hold even within GMCs down to very small scales \\citep{HEYER04,ROSOLOWSKY08}.  The {\\em size-line width} relation is an expression of the equilibrium turbulence conditions in the molecular ISM.  Detailed modeling of GMC line profiles shows that the emission is macroturbulent (i.e., the scale size of the turbulence is larger than the photon mean free path), corresponding to many optically thick clumps that have a clump-to-clump velocity dispersion similar to the observed line width \\citep{WOLFIRE93}. It is known that the observed line widths in all but a few very compact, quiescent clouds are too large to be thermal; at the typical GMC temperatures of $15-25$~K the thermal CO velocity dispersion would be $\\lesssim0.1$ \\kmpers. They are rather due to supersonic turbulence within the clouds. The dominant sources of turbulence in the molecular ISM remain somewhat controversial. For example, it is not yet known whether turbulence in molecular clouds is primarily internally or externally driven \\citep[e.g.,][]{MCKEE07}. Contributors must include star formation --- in the form of jets and winds for low mass stars, and winds and expanding H{\\small II} regions for massive stars --- and star destruction --- in the form of supernovae shocks and expanding superbubbles. Global processes that couple the large reservoirs of thermal or rotational energy to the turbulent cascade --- such as spiral shocks and several types of instabilities that involve shear, magnetic fields, and self gravity --- are important on the largest spatial scales. The fact that turbulence is present on scales $\\gtrsim100$ pc \\citep[e.g.,][]{BRUNT03,DIB05}, and that the dissipation lifetime of the largest eddies is $\\lesssim10^7$ yr \\citep*{FLECK81,STONE98}, point to constant (or at least frequent) energy injection on large scales.  This in turn suggests that massive stellar death and gobal mechanisms are predominantly responsible for the bulk of turbulent energy injection in the Milky Way, although their respective dominance is a matter of debate \\citep{MACLOW04,PIONTEK04,PIONTEK05}.  The last two Larson relations, the {\\em luminosity-line width} and {\\em luminosity-size} relations, describe correlations between cloud luminosity, $L_{\\rm CO}$, and either velocity dispersion or size (note, as is frequently pointed out, that only two of the three Larson relations are independent). They are \\citep{SOLOMON87}  \\begin {equation} L_{\\rm CO}\\approx130\\, \\sigma_v^{5}\\ \\Kkmperspcsq, \\label{lum_lw_eq} \\end{equation}  \\noindent and  \\begin{equation} L_{\\rm CO}\\approx25\\, R^{2.5}\\ \\Kkmperspcsq, \\label{lum_size_eq} \\end{equation}  \\noindent with $\\sigma_v$ and $R$ expressed in their corresponding units of \\kmpers\\ and pc.  Thus, the observations show that the virial mass, $M_{vir}$, relates to the CO luminosity (which is proportional to the inferred luminous mass $M_{lum}$) in the following manner \\citep{SOLOMON87},  \\begin{equation} M_{vir}\\approx39\\, L_{\\rm CO}^{0.81}\\ {\\rm M}_\\odot \\label{lco_mvir_eq} \\end{equation}  \\noindent where $L_{\\rm CO}$ is expressed in \\Kkmperspcsq. We will call this equation the Galactic {\\em mass-luminosity} relation. Note that the exponent of $L_{\\rm CO}$ is somewhat uncertain, and we take here the value reported by \\citeauthor{SOLOMON87}. Although the relation may not be precisely linear it is almost so, and simple models of molecular clouds can explain the observed linearity \\citep*{DICKMAN86}. This approximate proportionality between $L_{\\rm CO}$ and $M_{vir}$ within the Galaxy allows the definition of an empirical CO-to-H$_2$ conversion factor, \\xco, that relates the intensity of a GMC in the $^{12}$CO \\jone\\ transition to its H$_2$ column density ($N(\\htwo)=\\xco\\Ico$) and ultimately its mass. For GMCs around a median mass of $5\\times10^5$ M$_\\odot$ \\citeauthor{SOLOMON87} obtained a CO-to-H$_2$ factor that, when adjusted to the currently accepted distance of 8.5~kpc to the Galactic Center, yields $\\xco\\approx1.9\\times10^{20}$ \\xcounits\\ \\citep{MCKEE07}.  In the Milky Way, this ``virial mass'' approach yields conversion factors similar to those obtained from non-dynamical measurements, strongly suggesting that clouds are in a dynamical state intermediate between marginal gravitational binding and virial equilibrium \\citep*[e.g., $\\gamma$-ray results and modeling of the Galactic plane dust continuum;][]{BLOEMEN84,STRONG96,DAME01}, with magnetic energy introducing a small but non-negligible correction on the small scales \\citep[$\\lesssim50\\%$;][]{CRUTCHER99}.  Throughout this paper we adopt $\\xco=2\\times10^{20}$ \\xcounits\\ as the Galactic conversion factor. Note, however, that some nearby lines-of-sight at high Galactic latitudes may exhibit significant larger values of this coefficient \\citep*[e.g.,][]{MAGNANI88,GRENIER05}.  Therefore, mass estimates of GMCs show that they are gravitationally bound and in approximate virial equilibrium, which requires $M_{vir}\\sim R\\,\\sigma_v^2\\propto M_{lum}$. As a consequence of Equation \\ref{size_linewidth_eq} and virial equilibrium, all GMCs have approximately the same mean column density: $N_{\\rm H}\\approx 1.5\\times10^{22}$ cm$^{-2}$, or an equivalent mass surface density $\\Sigma\\approx170$ M$_\\odot$~pc$^{-2}$ (after applying a 40\\% correction for the mass contribution of elements heavier than Hydrogen). An equivalent result is that all Milky Way GMCs have approximately the same mean visual extinction from edge to edge, $A_V\\sim7.5$, using the standard dust-to-gas ratio \\citep*{BOHLIN78}.  \\citet{MCKEE99} argues that Equation \\ref{lco_mvir_eq} is a consequence of the interplay between CO chemistry, the Milky Way radiation field, the dust-to-gas ratio, and the local pressure of the ISM. As a result, H$_2$ clouds observable in CO are at least marginally bound. Furthermore, \\citeauthor{MCKEE99} argues that the constancy of $N_{\\rm H}$ stems from the fact that clouds with lower column density would not be bright in CO, while clouds with higher column density would rapidly collapse.  Another consequence of the Larson relations is that Galactic GMCs have low volume-averaged densities, $50<n({\\rm H_2})<500$ cm$^{-3}$. Therefore their clumpiness must be high in order to excite the rotational transitions of CO and of other commonly observed high-density tracer molecules such as HCO$^+$, which require volume densities in excess of 10$^4$ cm$^{-3}$. This clumpiness is probably induced by supersonic turbulence \\citep*{PADOAN98,OSTRIKER99}. Further, the average density of a GMC decreases with increasing mass, and as a result the free fall time in the cloud increases. This led \\citet{KRUMHOLZ05} to predict that larger GMCs will form collapsing cores and stars less rapidly than smaller ones.  The purpose of this study is to determine, {\\em using a consistent methodology}, the relations between size, velocity dispersion, and luminosity for extragalactic GMCs that are encapsulated in the Milky Way by the Larson laws.  We do so by studying GMCs in dwarf galaxies and comparing them with those in spiral galaxies. Dwarf galaxies are both numerous and nearby, and host physical conditions that fall outside those explored by surveys of the Milky Way. The range of parameter space thus probed provides an ideal laboratory to test theories of cloud structure and formation. Particularly, dwarf galaxies have low metallicities and their lack of internal extinction results in intense radiation fields, reminiscent of the conditions in primitive galaxies. They display irregular morphologies and are not dominated by spiral structure. Dwarfs often have slowly rising, nearly solid-body, rotation curves; this translates into small rotational shear and perhaps changes the sources and scales of turbulence in their ISM. The stellar potential wells of dwarf galaxies are weaker than those in large spiral galaxies and the ambient ISM may be correspondingly less dense. These environmental changes may influence the equilibrium conditions of GMCs, and affect their Larson relations.  The integrated properties of dwarf galaxies already show clear evidence that environment affects at least the CO content. Dwarf irregular galaxies are fainter in CO than large spiral galaxies. Galaxies with metallicity $12 + \\log {\\rm O/H} \\lesssim 8.0$ are seldom detected in CO emission \\citep{TAYLOR98}. Dwarf galaxies of higher metallicity detected in CO, however, show similar CO luminosity relative to either their stellar or atomic gas content as larger galaxies \\citep{LEROY05}. Here we investigate whether the resolved properties of GMCs in these galaxies also reflect their changing environments.  \\subsection{Resolving extragalactic molecular clouds} \\label{introresol}  The study of resolved GMC properties in other galaxies remains technically challenging.  A typical size for a Galactic GMC is $\\sim 40$~pc, which at the $\\sim4$ Mpc distance of the M~81 group translates to an angular size of $\\sim2\\arcsec$. Achieving such a resolution requires an antenna 250 meters in diameter operating at $\\lambda=2.6$~mm, the wavelength of the ground transition of $^{12}$CO. This is clearly an impractical proposition, but aperture synthesis techniques used by millimeter-wave interferometers routinely attain this resolution. They do so, however, at the price of reduced surface brightness sensitivity with respect to their filled aperture equivalents, and studies such as the one presented here are ultimately limited by sensitivity. Until the advent of the Atacama Large Millimeter Array (ALMA) these observations will remain challenging.  Because of these difficulties, the resolved properties of GMCs in galaxies other than the Milky Way remain a largely unexplored domain. Because of its proximity, located at a distance where even a modest telescope can attain a resolution of $\\sim30-40$~pc, the first galaxy where observations capable of clearly resolving individual GMCs were possible was the Large Magellanic Cloud \\citep{ISRAEL82}. The first resolved study of a GMC beyond the immediate vicinity of the Milky Way was performed in M~31 \\citep*{VOGEL87}, while the first large scale interferometric survey of GMCs in another galaxy was done on M~33 \\citep{WILSON90}. The first comprehensive analyses in the Magellanic Clouds were performed on the Small Magellanic Cloud \\citep*{RUBIO91,RUBIO93b}.  \\citet{BLITZ07} recently presented a study of GMCs in the Local Group, analyzing observations from CO surveys complete down to a known mass limit over a significant fraction of each galaxy. This study concluded that the Larson relations in the Milky Way disk were approximately those observed throughout the Local Group, although some offsets (possibly methodological) were present. They also concluded that GMCs within a particular galaxy have roughly constant surface density, $\\Sigma_{\\rm GMC}$, and that once corrections for the local CO-to-H$_2$ conversion factor are applied the sample as a whole has a scatter of only a factor of 2 in $\\Sigma_{\\rm GMC}$. Finally, they calculated a typical conversion factor $\\xco\\approx4\\times10^{20}$ \\xcounits, accurate to within 50\\% for most of the sample except the Small Magellanic Cloud, for which they obtain a value $\\sim3$ times larger.  This paper proceeds as follows: in \\S\\ref{data} we describe the data sets and and galaxies studied. In \\S\\ref{method} we summarize the method used to consistently derive GMC properties across the different datasets.  In \\S\\ref{results} we discuss the {\\em size-line width}, {\\em luminosity-line width}, and {\\em luminosity-size} relations in our sample of extragalactic GMCs, and compare them with the Galaxy. In \\S\\ref{discussion} we discuss the implications of these results on the equilibrium of clouds in the Small Magellanic Cloud, the photoionization-regulated star formation theory, the CO-to-H$_2$ conversion factor, and the brightness temperature of extragalactic clouds. In each case we contrast our results in dwarf galaxies with those obtained in the Milky Way and the other Local Group spirals, M~31 and M~33.  In \\S\\ref{summary} we summarize this study and present our conclusions. ",
        "conclusions": "\\label{summary}  We present and discuss a comprehensive set of high spatial resolution observations of CO in low mass galaxies, obtained by a combination of interferometer and single-dish instruments. Although the data are heterogeneous, we analyze them in a consistent manner to obtain GMC sizes, velocity dispersions, and luminosities.  The analysis is performed using the algorithm described by \\citet{ROSOLOWSKY06}, which does a good job at removing the biases due to dissimilar resolution and signal-to-noise.  We compare this uniform dataset of resolved extragalactic molecular cloud properties against those of GMCs in the three disk galaxies in the Local Group.  To do so, we analyze the interferometric maps of M~31 \\citep{ROSOLOWSKY07}, M~33 \\citep{ROSOLOWSKY03}, and the sample of Galactic GMCs discussed by \\citet{SOLOMON87}.  The main result of this study is that, remarkably, {\\em there are only small differences between CO-bright GMCs in the Milky Way and GMCs in galaxies with metallicities as low as $0.2Z_\\odot$ subject to a variety of physical conditions}. Our sample of extragalactic GMCs follows approximately the same {\\em size-line width}, {\\em luminosity-size}, and {\\em luminosity-line width} relations as Galactic clouds. Such uniformity may be in part responsible for the observed invariance of the stellar Initial Mass Function in galaxies. In any case, this result underscores that the Galactic Larson relations provide a remarkably good description of CO-bright Giant Molecular Clouds independent of their environment, at least in the range of environments explored by this study.  Although the Larson relations are approximately Galactic there are some significant departures. GMCs in dwarf galaxies tend to be slightly larger than GMCs in the Milky Way, M~31, and M~33 for a given CO luminosity or velocity dispersion. The largest departures occur in the SMC, the galaxy with the lowest metallicity of the sample. A possible interpretation, viable for most of our objects, is that GMCs in small galaxies have on average a surface density of $\\Sigma_{\\rm GMC}\\sim85$ \\msunperpcsq\\ rather than the canonical $\\Sigma_{\\rm GMC}=170$ M$\\odot$~pc$^{-2}$ observed in the Galaxy. In the case of the SMC, however, the implied surface densities, and consequently the central extinction in the GMCs, would be too low to expect bright CO emission. We explore three possibilities: these clouds are transient structures not supported by turbulence and will collapse in a few Myr, or else they are supported by magnetic fields in a larger proportion than similar Galactic clouds, or maybe the kinematics of the CO cores do not reflect the presence of massive \\htwo\\ envelopes.  Analysis of the properties of our sample of extragalactic GMCs shows that they do not accord with simple predictions from the theory of photoionization-regulated star formation \\citep{MCKEE89}. In the framework of this theory, clouds in lower metallicity environments will have larger surface densities to attain similar extinction in their central regions, which can then decouple from the magnetic field and collapse. Such a trend would be evident in the coefficient of the {\\em size-line width} relation, as expressed by Equation \\ref{photosf}. As Fig. \\ref{met_beta} illustrates, we see no evidence for such a trend in our data. We point out four possible caveats with these results: 1) it is possible that our observations do not probe the spatial scales on which these enhancements may occur.  2) Perhaps other parameters that enter in the theory (such as the cosmic ray ionization rate, for example) change from galaxy to galaxy in a manner that conspires to keep the coefficient of the {\\em size-line width} relation approximately constant. 3) Maybe our identification of $\\delta_{gr}$ with metallicity is incorrect and we sample a smaller range of conditions than we think we do. Or, 4) the CO kinematics do not trace the full potential of the cloud. In any case, we find a $\\sigma_v\\,R^{-0.5}$ product that is approximately constant or even decreasing for decreasing metallicity, contrary to expectations.  Finally, we address the matter of the dependency of the CO-to-\\htwo\\ conversion factor, \\xco, on metallicity. We find that the extragalactic GMCs in our analysis agree very well with the Galactic {\\em luminosity-virial mass} relation. Consequently, in our study of resolved GMCs we find no measurable change of \\xco: over a factor of 5 in metallicity most of our galaxies are compatible with a Galactic $\\xco=2\\times10^{20}$ \\xcounits\\ within a factor of two, and there is no discernible trend with metallicity. We emphasize, however, that this measurement is relevant on the scales of the individual, CO-bright GMCs. Studies in the FIR and at millimeter-waves in low metallicity environments (particularly the SMC) suggest that these clouds are embedded in larger \\htwo\\ envelopes that are not traced by CO emission but contain an appreciable mass of molecular gas, at least in the case of the SMC \\citep{ISRAEL97,RUBIO04,LEROY07}. These FIR results are approximately consistent with virial \\xco\\ estimates obtained on larger scales than the ones discussed in this paper \\citep{RUBIO93b,MIZUNO01,BOLATTO03}. We suggest that the FIR and CO observations can be simultaneously understood in the following terms: in low metallicity gas bright CO emission is relegated to the density peaks. Observations that resolve those density peaks show properties that are similar to those of Milky Way GMCs --- that is the result of this study, and it is supported by studies of the structure of Milky Way GMCs \\citep{HEYER04,ROSOLOWSKY08}. On larger scales, the cloud-to-cloud velocity dispersion as well as the dust opacity suggest that there is a large mass of likely molecular material that is not accounted for by the CO intensity. Such regions would be CO-bright in objects of higher metallicity.  Studying the resolved properties of extragalactic GMCs is a challenging undertaking with the current generation of instruments, particularly for clouds in CO-faint dwarf galaxies. In the near future the deployment of the ALMA interferometer will make it possible to obtain more precise measurements on larger samples at farther distances. Studies such as the one presented here suffer from an important bias: we can only perform them on the brightest CO peaks. The detection of clouds as faint as those in the SMC in CO, for example, is currently impossible beyond the immediate neighborhood of the Local Group. These improvements in instrumentation will soon provide powerful tests of theories of star and molecular cloud formation beyond our own Galaxy."
    },
    "0807/0807.2460_arXiv.txt": {
        "abstract": "This is the first paper of a series where we study the clustering of LRG galaxies in the latest spectroscopic SDSS data release, DR6, which has 75000 LRG galaxies covering over 1 $Gpc^3/h^3$ at $0.15<z<0.47$. Here we focus on modeling redshift space distortions in $\\xips$, the 2-point correlation in separate line-of-sight and perpendicular directions, on large scales. We use large mock simulations to study the validity of models and errors. We show that errors in the data are dominated by a shot-noise term that is $40\\%$ larger than the Poisson error commonly used. We first use the normalized quadrupole for the whole sample (mean z=0.34) to estimate $\\beta=f(\\Omega_m)/b=0.34 \\pm 0.03$, where $f(\\Omega_m)$ is the linear velocity growth factor and $b$ is the linear bias parameter that relates galaxy to matter fluctuations on large scales. We next use the full $\\xips$ plane to find $\\Omega_{0m}= 0.245 \\pm 0.020$ (h=0.72) and the biased amplitude $b \\sigma_8 = 1.56 \\pm 0.09$. For standard gravity, we can combine these measurements to break degeneracies and find $\\sigma_8=0.85 \\pm 0.06$, $b=1.85 \\pm 0.25$ and $f(\\Omega_m)=0.64 \\pm 0.09$. We present constraints for modified theories of gravity and find that standard gravity is consistent with data as long as $0.80<\\sigma_8<0.92$. We also calculate the cross-correlation with WMAP5 and show how both methods to measure the growth history are complementary to constrain non-standard models of gravity. Finally, we show results for different redshift slices, including a prominent BAO peak in the monopole at different redshifts.  The $\\xips$ data on large scales is shown to be in remarkable agreement with predictions and shows a characteristic large region of negative correlation in the line of sight, a BAO ring and a prominent radial BAO peak. The significance of this is presented in paper IV of this series. We include a study of possible systematic effects in our analysis to find that these results are quite robust. ",
        "introduction": "\\label{sec:intro}  Galaxy clustering allows the study of different physical phenomena at each scale. On large scales density fluctuations are small and can be modeled by linear theory to constrain cosmological parameters. We need large surveys with galaxy positions to do this.  The measured redshift distance $s$ of a galaxy differs from the true radial distance $r$ by its peculiar velocity along the line-of-sight $v_r$. These displacements lead to redshift distortions, with two important contributions. The first, on large scale fluctuations, is caused by coherent bulk motion. We see walls denser and voids bigger and emptier, with a squashing effect in the 2-point correlation function along the line-of-sight, known as the Kaiser effect \\cite{kaiser}. At small scales, random velocities inside clusters and groups  of galaxies produce a radial stretching pointed at the observer, known as fingers of God (FOG) \\cite{jackson}.  Although such distortions complicate the interpretation of redshift maps as positional maps, they have the advantage of bearing unique information about the dynamics of galaxies. In particular, the amplitude of the quadrupole redshift distortion on large scales yields a measure of the linear redshift distortion parameter $\\beta$, which is related to the matter density, and gives direct information on the growth of the newtonian gravitational potential $f(\\Omega)$ and bias $b$, which describes how the number density of galaxies traces matter density (see Eq.(\\ref{fb})). We will also show how one can use the anisotropic redshift distortions, broken into line of sight and perpendicular directions, to find  separate measurements for $f(\\Omega)$ and $b$.  In this work we use the most recent luminous red galaxies (LRGs) from the spectroscopic SDSS public data release DR6 \\cite{dr6}, and perform studies of linear bias on large scales to obtain cosmological parameters. As we will see, we can break the degeneracy between $bias$ and $\\sigma_8$ present in the correlation function thanks to redshift distortions anisotropies, and look at the growth history and possible modifications of the gravity (see Linder 2005, Guzzo et al 2008 and the recent Yamamoto et al 2008).  A dark energy dominated universe causes gravitational potentials to evolve, changing the energy of CMB photons passing through them. This is the so called  Integrated Sachs-Wolf (ISW) effect. Here, we have also cross-correlated LRGs with WMAP \\cite{wmappaper} in order to investigate this ISW effect, reproducing a high signal as seen in recent studies  (see Gaztanaga etal 2006, Giannantonio et al 2008), as compared to current $\\Lambda CDM$ model.  In this way, we can also break the $bias-\\sigma_8$ degeneracy  and study the growth history using the cross-correlation between temperature of CMB and galaxies $w_{TG}$, which give independent constraints on bias and $\\sigma_8$ compared with those obtained from redshift space distortions.    The same LRGs (but with reduced area) have been studied from different points of view. Tegmark et al (2006) have done an analysis of the power spectrum at large scales to obtain cosmological parameters. Zehavi et al (2005) study LRGs at intermediate scales (0.3 to 40Mpc/h), where they calculate the projected correlation function, the monopole and real-space correlation function to study mainly the linear high bias, the non-linear bias and the differences between luminosities, remarking that there are differences from a power law for scales smaller than 1Mpc/h.  While we were working on our results, a very interesting paper by Okumura et al (2008) showed the first application of the anisotropy in the 2-point correlation function including the baryonic features, to obtain constraints on cosmological parameters. They use the DR3 spectroscopic sample of LRGs to calculate the 2-point correlation function with a different definition of what parallel and perpendicular directions mean (the line of sight is defined in their paper as the direction toward the galaxy 1, Matsubara 2004). They fit $\\xisp$ for large scales using the linear Kaiser model, from 40Mpc/h to 200Mpc/h, excluding the FOG zone. They point out that a direct measurement of the growth function can be obtained from the distortion to the clustering signal caused by bulk flows (Kaiser anisotropy). Although constraints are weak by now, the fitting of all the anisotropic 2-point correlation function, including the baryonic feature, would enable us to divide the effect of redshift distortions into dynamical and geometrical components. The anisotropy due to geometric distortion contributes to a better estimation of the equation of state of the dark energy. Here we try to obtain new constraints, encouraged by the fact that DR6 has doubled the volume of DR3.  There has already  been significant progress in this direction using LRG. Eisenstein et al (2005) detected the baryon acoustic peak in the 2-point correlation function using LRGs. H{\\\"u}tsi (2006 a,b) use LRGs to constrain cosmological parameters in the power spectrum, including the baryonic peak. Percival et al (2007) have analyzed also the LRGs using both 2dF and SDSS. Padmanabhan et al (2007) used the photometric catalog to work with a bigger set of LRGs with photometric redshifts, obtaining also cosmological constraints, the same as Blake et al (2007), which work with the MegaZ-LRG, a photometric-redshift catalog of luminous red galaxies based on the imaging data of the SDSS DR4.  This is the first and main paper of a series on clustering of LRG. Here we will focus on redshift space distortions in $\\xips$ on larges scales. In Paper II \\cite{paper2} we will look at small scales in $\\xips$ and address the issue of the FOG. In Paper III \\cite{paper3} we look at the 3-point function and in Paper IV \\cite{paper4} we will focus on the significance of the BAO detections.  In this first paper we present the basis for the error analysis and the study of systematics effects that will be used in the rest of the series.   The paper is organized as follows. Section II gives a summary of the theory involved in redshift space distortions, the model used for the 2-point correlation function in separate line of sight and perpendicular directions, and the different multipoles. Section III presents an accurate description of the data selected to work with, LRG galaxies. Section IV is the main part of this paper, where we explain the analysis done and our main results. We use the distortions in redshift space to obtain the distortion parameter $\\beta$, $\\Omega_m$, $\\sigma_8$ and modifications to gravity. We also look at the ISW effects for this LRG data, and explore the differences between redshift slices. As usual, we end with discussion and future work in section V. The Appendix A is devoted to present the mock simulations that we use to validate the errors used, Monte Carlo, jackknife and a new theoretical approach. In this section, we also validate the methods used to obtain constraints (section IV) in real data with our mocks. In the Appendix B we present a study of possible systematic effects. ",
        "conclusions": "\\label{sec:discuss}  In this paper we work with luminous red galaxies (LRGs) from the catalog SDSS DR6 to obtain cosmological parameters from large scales. These are good galaxies to study this range of scales since they trace a big volume of the universe. We use a simple approximation to model the redshift space distortions (see \\S\\ref{sec:model2point}), which work incredibly well with LRGs on large scales, as explained in this paper, but also on small scales, shown in Paper II \\cite{paper2}  of this series.  We have used one of the largest N-body simulation run to date (with box size of L=7680Mpc/h with $2048^3$ particles) to create realistic mock LRG galaxy catalogs to analyze the data (Appendix A). We have checked our results to some extent with smaller simulations. This is a crucial step given the complications in the data and the accuracy needed for its interpretation. We have validated all the methods that we use during the paper with credible mock  simulations. We also test the validity of the models and the errors: Monte Carlo, jackknife and a theoretical approach for $\\xisp$. We find that JK errors does not work properly for large perpendicular distance $\\sigma$.  There are usually two important sources of statistical errors in clustering analysis: Shot-noise error, due to the finite number of galaxy pairs, and sampling errors, due to the finite amount of volume in the survey. We have shown with our mock simulations that the LRG samples in SDSS DR6 are dominated by the shot-noise term in the 2-point $\\xisp$. Usually, this term is assumed to follow the Poisson distribution, with an error given by the inverse of the square root of the number of measurements (ie see Eq.(\\ref{eq:Poisson})). This is in fact the case to a good accuracy in our dark matter mocks. But this is not the case for groups of dark matter particles (halos) which are selected in a way that, by construction, the detection of a group creates a large excluding region around the group position.  This produces what we have called super-Poisson errors, which are about $40\\%$ larger than Poisson errors. Jack-knife errors are shown to be good at detecting this additional shot-noise component and we find that this effect is also present in the real LRG data. This is interpreted here as direct evidence that the dominant fraction of the LRG galaxies populates separate dark matter halos. But this does not exclude a minor fraction of them to be found with other LRGs within the same halo. In fact, in Paper III, \\cite{paper3}, of this series we find evidence for a large quadratic bias which indicates that this is the case \\cite{bhalo}.  We conclude our discussion on errors by emphasizing the importance of using realistic mocks to assess our results. We have demonstrated here that dark matter mocks are not good enough. Diagonal, sampling errors can in principle be scaled from dark matter to real data because the relative error only depends in the sampling volume. But this is not the case for the shot-noise or for the covariance. The difference is important. Errors and covariances are underestimated with dark matter mocks and results will be wrongly interpreted.  We have calculated the distortion parameter $\\beta=[0.310-0.375]$ using the normalized quadrupole Q(s), which also provides us of an effective value for the dispersion of random peculiar velocities $\\sigma_v=[365-415]$ km/s (see Table \\ref{tab:tabbetasigma} for more slices in redshift). The great advantage of the normalized quadrupole introduced by Hamilton (1992) is that it can measure the squashing produced by bulk galaxy motions with independence of the overall amplitude or shape of the galaxy 2-point correlations. We have also checked if this value for $\\beta$ gives the correct ratio $\\xis/\\xir$ on large scales. We have shown that with actual data, we can only recover the real-space correlation function $\\xir$ for scales smaller than 30Mpc/h, and we see that the ratio between the redshift-space monopole $\\xis$ and real-space correlation function $\\xir$, $\\xis/\\xir$, can be described by the approximated Kaiser value at large scales, with the same value of $\\beta$ found in the quadrupole.  We obtain a value for $\\Omega_m$ from the shape  of the redshift-space correlation function in the anisotropic plane $\\pi-\\sigma$ at large scales. The value is $\\Omega_m=[0.225-0.265]$ for the mean slice (see Table \\ref{tab:largescales} for other redshift slices). We also obtain a value for the amplitude $b\\sigma_8$ in the linear regime. We can therefore break the degeneracy between $\\sigma_8$ and $b$ thanks to redshift distortions through the parameter $\\beta$, which results in $\\sigma_8=[0.79-0.91]$ and $b$ around 2.  We next look at the growth history and possible modifications of gravity through the parameter $\\gamma$ which is 0.55 for standard gravity and 0.68 for the modified gravity DGP model. We find consistence with standard gravity for $0.8\\le\\sigma_8\\le0.92$ at 1-$\\sigma$. DGP model is inconsistent with our results if $\\sigma_8 < 0.84$.  We have also cross-correlated LRGs with WMAP in order to investigate the ISW effect, obtaining a high signal, as in other studies, compared to current $\\Lambda CDM$ model. The degeneracy  $b-\\sigma_8$ and the growth history can also be studied using the ISW effect which is independent from the study of redshift distortions. The cross-correlation indicates a preference for higher $\\sigma_8$.  In the Appendix B we look for systematic errors, and arrive to the conclusion that our results are robust in the scales we are doing the analysis."
    },
    "0807/0807.1463_arXiv.txt": {
        "abstract": "{}{To model the chemical evolution of manganese relative to iron in three different stellar systems: the solar neighbourhood, the Galactic bulge and the Sagittarius dwarf spheroidal galaxy, and compare our results with the recent and homogeneous observational data.} {We adopt three chemical evolution models well able to reproduce the main properties of the solar vicinity, the galactic Bulge and the Sagittarius dwarf spheroidal. Then, we compare different  stellar yields in order to identify the best set to match the observational data in these systems.} {We compute the evolution of manganese in the three systems and we find that in order to reproduce simultaneously the [Mn/Fe] versus [Fe/H] in the Galactic bulge, the solar neighbourhood and Sagittarius, the type Ia SN Mn yield must be metallicity-dependent.} {We conclude that the different histories of star formation in the three systems are not enough to reproduce the different behaviour of the [Mn/Fe] ratio, unlike the situation for [$\\alpha$/Fe]; rather, it is necessary to invoke metallicity-dependent type Ia SN Mn yields, as originally suggested by McWilliam, Rich \\& Smecker-Hane in 2003.} ",
        "introduction": "McWilliam, Rich \\& Smecker-Hane (2003) considered the abundance trend of Mn in three contrasting stellar populations: a sample of Galactic bulge K giants, stars in the solar neighbourhood and stars in the dwarf spheroidal galaxy Sagittarius (Sgr dSph). They found that the trend of [Mn/Fe] in the Galactic bulge follows more or less the relation of the solar vicinity, but most stars of Sgr dSph show [Mn/Fe] deficient by $\\sim 0.2$ dex. They concluded that to explain this result one has to assume that the yields of Mn both from type II and type Ia SNe should depend on the initial stellar metallicity. This conclusion was at variance with a suggestion by Gratton (1989) that Mn is overproduced in type Ia SNe, relative to type II SNe. Gratton (1989) observed that [Mn/Fe] versus [Fe/H] in solar vicinity stars show the opposite behaviour of the [$\\alpha$/Fe] ratios: in fact, [Mn/Fe] $\\sim$$-$0.4 dex and constant for stars with [Fe/H]$<-1.0$ dex and increases up to the solar value for stars with [Fe/H]$>$ $-$1.0 dex. As it is well known, first Tinsley (1979) and then Greggio \\& Renzini (1983) and Matteucci \\& Greggio (1986) explained the trend of the $\\alpha$-elements (e.g. O, Mg, Si, Ca, Ti), relative to Fe, as due to the different roles played by type Ia and II SNe in galactic chemical enrichment. In particular,  the higher than solar [$\\alpha$/Fe] ratio in Halo stars is interpreted as due to the pollution of type II SNe, which explode on very short timescales (of the order of Myr to tenths of Myr) and produce roughly constant [$\\alpha$/Fe] ratios. Then the [$\\alpha$/Fe] ratio decreases to the solar value in disk stars because of the bulk of Fe produced by type Ia SNe explode on timescales ranging from 30--40 Myr to 10 Gyr. This interpretation is known as ``time-delay model'' and represents the best interpretation of alpha-element abundance ratios in galaxies so far. Type II SNe originate from the core-collapse of massive stars ($M> 10 M_{\\odot}$), whereas it is believed that type Ia SNe arise from mass-accretion in binary C-O white dwarfs systems, exploding when the white dwarf mass reaches the Chandrasekhar limit. Two main scenarios can lead to this situation: i) the single degenerate, where a white dwarf explodes after accreting material from a red giant or a main sequence companion (Whelan \\& Iben 1973); ii) the double degenerate, where the explosion follows the merger of two C-O white dwarfs (Iben \\& Tutukov 1984), due to loss of angular momentum from gravitational wave emission. There is not yet consensus on which scenario is better or if both are at work in galaxies. In both scenarios type Ia SNe start exploding after 30--40 Myr and continue to do so for a Hubble time. The maximum type Ia SN rate depends on the assumed progenitor model but also on the assumed star formation history, being short in systems with a strong burts of star formation (e.g. 0.5 Gyr in the galactic bulge), 1.0-1.5 Gyr in the solar vicinity and several Gyrs in slow star forming systems, such as irregular galaxies (Matteucci \\& Recchi, 2001). Therefore, the bulk of Fe is always produced with a delay relative to the $\\alpha$-elements produced by type II SNe. Manganese is an iron-peak element produced both in type II and Ia SNe.  Nucleosynthesis calculations for massive stars by Arnett (1971) and Woosley \\& Weaver (1995; hereafter WW95), among others, show increased yields of Mn with increasing metallicity (actually with increasing neutron excess). Metallicity-dependent yields for type Ia SN have been largely unexplored; however, Ohkubo et al. (2006) showed that the yields of both Ni and Mn increase with metallicity in type Ia SNe, although this conclusion is sensitive to the details of the adopted hydrodynamical models.   In this paper, we adopt three detailed chemical evolution models for the Galactic bulge, solar vicinity and Sgr dSph. In all the models we will adopt the same yields but different star formation histories tuned to reproduce the majority of the observational constraints on the three systems. We will show that to explain the [Mn/Fe] simultaneously in the three systems, the time-delay model alone is not enough, but we have to assume that the yields of Mn from type Ia SNe are metallicity-dependent. The paper is organized as follows: in section 2 we briefly describe the observational data, in section 3 we describe the chemical evolution models. In section 4 the nucleosynthesis prescriptions are given. In section 5 we discuss our results and in section 6 we draw our conclusions. ",
        "conclusions": "We show quantitatively, for the first time, that galactic chemical evolution models demand that the type Ia SN Mn yields must be metallicity-dependent, by simultaneously fitting the [Mn/Fe] trend in three independent stellar systems: the Galactic bulge, the solar vicinity and the Sgr dSph. We have adopted detailed chemical evolution models, chosen to reproduce the main chemical features of these three stellar systems, using same element yields but different star formation histories.  We have shown clearly that the ``time-delay model'' alone cannot explain the different behaviour of the [Mn/Fe] in Bulge, solar vicinity and Sgr dSph, as it does for [$\\alpha$/Fe] ratios. In fact, in the case of the [$\\alpha$/Fe] ratios, the same models can well reproduce the observed [$\\alpha$/Fe] versus [Fe/H] trends in the three stellar systems without invoking metallicity-dependent yields. This is because the production of $\\alpha$ and Fe elements is only mildly dependent on the initial stellar metallicity.  We propose a functional form of the metallicity-dependent Mn yields in type Ia SNe ($ \\propto Z^{0.65}$) which should be tested by nucleosynthesis calculations. Our result is supported by Ohkubo et al. (2006), who suggested that the yields of Mn from type Ia SNe have a strong dependence on metallicity, and by a recent finding by Badenes et al. (2008), who showed that the Mn to Cr ratio in the ejecta of type Ia SNe depend on the metallicity of the progenitor star in a very similar way to what we have suggested.   It is worth noting that in this work we have taken the Reddy et al. (2003, 2006) abundances for the solar vicinity stars at face value.  However, we note that McWilliam et al. (2003) found that a zero-point correction of $+$0.10 dex to the Reddy et al. (2003) [Mn/Fe] values was required if the sun is typical of stars in the range $-$0.10 $\\leq$ [Fe/H] $\\leq$< $+$0.10.  We have re-investigated this issue and found that together the Reddy et al. (2003) and Reddy et al. (2006) data indicate a mean [Mn/Fe] of solar [Fe/H] stars at $-$0.10$\\pm$0.01 dex, with an rms scatter about the mean of 0.036 dex.  Since the solar Mn/Fe ratio is 2.77 sigma above the mean for solar [Fe/H] stars (equal to 99.7\\% above the mean in a gaussian distribution), this zero-point shift is most likely due to systematic error in the Reddy et al. (2003, 2006) results, rather than a real Mn enhancement in the Sun.  This is supported by the fact that other abundance studies require quite different zero-point corrections; for example the differential study, relative to the sun, by Feltzing \\& Gustafsson (1998) required a zero-point [Mn/Fe] correction of $-$0.06 dex.  Unfortunately, we are unable to make as good a fit to the measured [Mn/Fe] trends if we employ a 0.10 dex zero-point shift to the Reddy et al. (2003, 2006) data.  To do so would require a more complicated form of the metallicity-dependent Mn yields from type Ia SNe, instead of the simple power-law form adopted here. Alternatively, other complications might also be responsible, such as the type II SNe yields of WW95, or variations in the star formation and/or the mass-loss rates for Sgr dSph.  Rather than explore a myriad of possibilities that might give a perfect fit to all the Mn trends, we prefer to focus on our main quantitative result: that the trend of [Mn/Fe] with [Fe/H] in the Galactic disk, the Bulge, and the Sgr dSph can only be understood if the yield of Mn from type Ia SNe is metallicity-dependent.    Finally, we note that the low [Mn/Fe] values seen in Sgr dSph stars near solar metallicity is due, at least in part, to the lower metallicity of old Sgr dSph stars compared to the solar neighborhood.  It is the low average [Fe/H] of the type Ia SNe in the Sgr dSph that results in lower [Mn/Fe] values, thanks to the metallicity-dependence of the Mn yield.  Thus, the [Mn/Fe] ratios of solar metallicity stars in dwarf galaxies is linked to the mean metallicity of these galaxies. Since the metallicity distribution function of stars in dwarf galaxies is sensitive to the mass outflow and star formation rates, the value of these parameters should be constrained by the observed [Mn/Fe] ratios.  Indeed, we would expect that the more massive dwarf galaxies (e.g. LMC) would show a similar effect on [Mn/Fe], but less extreme, than in lower mass dwarf galaxies; thus, there should be families of low-mass galaxies with different [Mn/Fe] trends."
    },
    "0807/0807.1180_arXiv.txt": {
        "abstract": "We present the first results of a project aimed at following the formation and long-term dynamical evolution of star clusters within the potential of a host galaxy. Here we focus on a model evolved within a simplified potential representing the Large Magellanic Cloud. This demonstrates for the first time the self-consistent formation of a bound star cluster from a giant molecular cloud. The model cluster reproduces the density profiles and structural characteristics of observed star clusters. ",
        "introduction": "\\label{s:intro}  Star clusters are vital components of galaxies. Indeed, extracting information from the massive globular clusters (GCs) of our Galaxy provides fundamental information on the epoch of galaxy formation. Furthermore, systems of extragalactic GCs are used as key determinants for tracing the dynamical, chemical and gaseous evolution of their host galaxies (see Brodie \\& Strader 2006 for a review). The smaller open clusters also have a role to play as well, as understanding their destruction within in the Galactic disk impacts the growing field of Galactic archaeology (Freeman \\& Bland-Hawthorn 2002), for example.  Efforts to model the evolution of individual star clusters have made great advances in the last decade. On the hardware front the advent of special-purpose GRAPE processors for calculating gravitational forces, with the latest incarnation being the GRAPE-6 (Makino 2002), have allowed $N$-body models of up to $N \\sim 100\\,000$ stars to be completed in a reasonable timeframe (Baumgardt \\& Makino 2003). The next generation of GRAPE (available in 2008), the use of massively-parallel supercomputers, and even graphics processing technology, will push this $N$ continually higher towards the realm of direct GC models. Complementing this is the push to make the models as realistic as possible by including algorithms to deal with processes such as stellar and binary evolution in concert with the treatment of the gravitational interactions of the stars. This is the case for both $N$-body (Hurley et al. 2005) and statistical (Fregeau \\& Rasio 2007; Giersz, Heggie \\& Hurley 2008) techniques.  On a grander scale simulations of galaxy formation which identify the formation sites of star clusters have been performed (Bekki \\& Chiba 2005). Furthermore, the formation of stars within turbulent giant molecular clouds (GMCs) based at these formation sites can then be followed with smooth-particle hydrodynamics (SPH) simulations (Bekki \\& Chiba 2007). Our aim is to interface these galactic-scale simulations of star cluster formation with the latest modeling techniques for following the long-term evolution of star clusters. The immediate benefits will be: a) the first self-consistent simulations of star cluster evolution from formation through to destruction, and b) a non-simplistic picture of how GC systems evolve with time and how this impacts the interpretation of observed extragalactic GC systems. This latter advance in particular will be important for understanding the connection between extragalactic GCs and galaxy formation by injecting much needed numerical models in to what has become a data-dominated field (Brodie \\& Strader 2006). We will also be able to explore the origin of observed features of star cluster systems, such as the age and parameter distributions of the Large Magellanic Cloud (LMC) clusters (Mackey \\& Gilmore 2003).  In this letter we provide an overview of the modeling process that will be used to achieve our goals. We illustrate this process by describing the evolution of a prototype model which demonstrates a working interface between the cluster formation and long-term evolution codes. It also entails the first simulation of bound cluster formation from a gas cloud within an external galactic potential. We then discuss future improvements to the model and details of how this fits in to our plans for a full investigation of the long-term evolution of galactic and extragalactic star cluster systems. ",
        "conclusions": "\\label{s:summ}  We have demonstrated the self-consistent formation of a bound star cluster that exhibits all the markings of a regular relaxed star cluster: dynamical equilibrium, core collapse, and an approximately spherical spatial distribution of the stars. This initially unrelaxed model, emerging from its progenitor giant molecular cloud, approaches virial equilibrium on the order of a few crossing times and before a half-mass relaxation time has passed -- a result that is free of the influence of the external tidal field. We have also found that the density profile of the model cluster resembles the distributions derived from observations -- King (1962) and Elson, Fall \\& Freeman (1987) -- from the time that virial equilibrium is reached and onwards. Idealized models that start in virial equilibrium can comfortably adopt these density distributions when setting the initial positions of the stars. Our findings are distinct from previous simulations of the collapse of proto-globular clusters and proto-galaxies -- which also observed the emergence of a core-halo structure at equilibrium -- primarily in that the initial process of fragmentation is explicitly followed using a three-dimensional hydrodynamic scheme and the influence of a galaxy-scale tidal field is modelled throughout the process.  This letter describes a model that demonstrates a working interface between simulations of star cluster formation and long-term star cluster dynamical evolution. In this model we have included several aspects, such as the use of equal-mass stars and a point-mass galaxy, that are commonly utilised in models of star cluster evolution but are not the most realistic approach. Future models will expand our study into a full investigation including: \\begin{itemize} \\item the effect of a spectrum of stellar masses on the formation and evolution of bound clusters -- associated processes such as stellar evolution and mass segregation will affect the binding energy and the evolution timescales; \\item a variety of cluster sizes, in terms of physical size and also a greater number of stars; \\item an exploration of how the SFE assumed in the formation step affects the final outcome, and also a greater understanding of the conditions required for substantial binary formation; and, \\item a variety of GMC locations within the host galaxy and an extension to model the {\\it live} gravitational potential of the host galaxy in concert with the internal star cluster evolution, with dwarf, spiral and elliptical hosts considered. \\end{itemize} This will allow a much more detailed study of the effects of internal and external dynamical processes on the formation and evolution of star clusters in a galactic context."
    },
    "0807/0807.4246_arXiv.txt": {
        "abstract": "Recent findings by $\\gamma$-ray Cherenkov telescopes suggest a higher transparency of the Universe to very-high-energy (VHE) photons than expected from current models of the Extragalactic Background Light. It has been shown that such transparency can be naturally explained by the DARMA scenario, in which the photon mixes with a new, very light, axion-like particle predicted by many extensions of the Standard Model of elementary particles. We discuss the implications of DARMA for the VHE $\\gamma$-ray spectra of blazars, and show that it successfully accounts for the observed correlation between spectral slope and redshift by adopting for far-away sources the same emission spectrum characteristic of nearby ones. DARMA also predicts the observed blazar spectral index to become asymptotically independent of redshift for far-away sources. Our prediction can be tested with the satellite-borne {\\it Fermi}/LAT detector as well as with the ground-based Cherenkov telescopes  H.E.S.S., MAGIC, CANGAROO\\,III, VERITAS and the Extensive Air Shower arrays ARGO-YBJ and MILAGRO. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0807/0807.4136_arXiv.txt": {
        "abstract": "We report the {\\it Chandra} detection of X-ray emission spatially coincident with the kpc-scale radio bubbles in the nearby ($D_{L} \\sim 31$ Mpc) AGN-starburst galaxy NGC\\,6764. The X-ray emission originates in hot gas ($kT \\sim 0.75$ keV), which may either be contained within the radio bubbles, or in a shell of hot gas surrounding them. We consider three models for the origin of the hot gas: (1) a starburst-driven galactic wind, (2) shocked gas associated with the expanding radio bubbles, and (3) gas heated and entrained into the bubbles by jet/ISM interactions in the inner AGN outflow. We rule out a galactic wind based on significant differences from known galactic wind systems. The tight correspondence between the brightest X-ray emission and the radio emission in the inner outflow from the Seyfert nucleus, as well as a correlation between X-ray and radio spectral features suggestive of shocks and particle acceleration, lead us to favour the third model; however, we cannot firmly rule out a model in which the bubbles are driving large-scale shocks into the galaxy ISM. In either AGN-driven heating scenario, the total energy stored in the hot gas is high, $\\sim 10^{56}$ ergs, comparable to the energetic impact of low-power radio galaxies such as Centaurus A, and will have a dramatic impact on the galaxy and its surroundings. ",
        "introduction": "\\label{intro}  Galaxy feedback processes are now thought to be an important ingredient in galaxy formation models (e.g. Croton et al. 2006; Bower et al. 2006), potentially solving the long-standing problem of explaining the deficit of galaxies at the low and high mass ends of the galaxy luminosity function relative to model predictions. Disentangling the feedback contributions from AGN outbursts and from star formation to the energetics of gas in galaxies, galaxy groups and clusters is a key problem in galaxy evolution, as it is clear that both can inject large quantities of energy into their surroundings. While there is growing evidence that AGN outbursts are the dominant feedback process operating at the high-mass end of the galaxy luminosity function, it might be expected that at lower masses, and particularly in late-type galaxies, the energy input from star formation processes will dominate.  Recent work has shown that kpc-scale radio bubbles connected to an active nucleus can be found in many Seyfert galaxies (e.g. Gallimore et al. 2006; Kharb et al. 2006). Although it has been argued that these bubbles may be powered by stellar winds (e.g. Baum et al. 1993), the fact that no convincing examples exist in galaxies without an AGN strongly suggests that they are inflated by the active nucleus (e.g. Hota \\& Saikia 2006, hereafter HS06). Many of the known Seyfert radio bubbles are predicted to be overpressured with respect to their surroundings (e.g. Capetti et al. 1999), and therefore may be shocking the surrounding medium, as has been observed for kpc-scale radio lobes in massive elliptical galaxies (e.g. Centaurus A: Kraft et al. 2003; NGC\\,3801: Croston et al. 2007).  In order to search for evidence for galaxy feedback associated with kpc-scale radio bubbles in spiral galaxies, we carried out a {\\it Chandra} observation of the nearby Seyfert 2/LINER galaxy NGC\\,6764 , which is a barred spiral galaxy (type SB(s)bc) with non-thermal radio bubbles perpendicular to the plane of the galaxy disk that extend for $\\sim 1$ kpc (Fig.~\\ref{ropt}). Here we report on the results of this observation.  Throughout this paper, we adopt a cosmology with $H_{0} = 70$ km s$^{-1}$ Mpc$^{-1}$, $\\Omega_{M} = 0.3$ and $\\Omega_{\\Lambda} = 0.7$. We adopt a luminosity distance for NGC\\,6764 of 31.3 Mpc, obtained by correcting the heliocentric velocity of 2146 km s$^{-1}$ to the CMB frame of reference. This gives an angular scale of $1^{\\prime\\prime} = 0.15$ kpc at the distance of NGC\\,6764. ",
        "conclusions": "We have identified X-ray counterparts to the kpc-scale radio bubbles in the Seyfert galaxy NGC\\,6764, whose morphology traces very closely the radio structures. The energy input from these hot gas features is $\\sim 10^{56}$ ergs, comparable to the energy input from shock heating in the low-power radio galaxy NGC\\,3801. The inferred rate of energy input in NGC\\,6764 is $\\sim 10^{42}$ erg s$^{-1}$, and is expected to have a significant impact on the galaxy and its surroundings. Based on the presence of an inner outflow that appears to emanate from the Seyfert nucleus, seen in both radio and X-ray emission, similarities to X-ray features in previously studied AGN-jet systems, key differences compared to starburst outflows, and the spatial coincidence of a region of flatter radio spectrum and hard X-ray spectrum, suggesting particle acceleration, we conclude that the thermal emission is produced by heating of the ISM due to compression and shocks caused by the AGN outflow. The morphology of the radio outflow in NGC\\,6764 leads us to favour a model in which jet/ISM interactions in the inner part of the galaxy lead to heating of the gas and its entrainment into the bubbles, as well as jet disruption; however, we cannot rule out a model in which the bubbles are driving large-scale shocks into the ISM. The existence of a population of kpc-scale radio lobe sources in Seyfert galaxies (expected to be short-lived) suggests that AGN-jet or AGN-jet+starburst powered galaxy feedback as seen in NGC\\,6764 could be important for the galaxy population as a whole."
    },
    "0807/0807.1239_arXiv.txt": {
        "abstract": "{ The degree of randomness in the cosmic microwave background (CMB) maps has been shown to be measurable and indicative of the hyperbolic effect of voids on the propagation of the photon beams. In terms of an introduced porosity parameter, we now obtain the criterion for a hyperbolicity due to domination of voids in the line-of-sight direction. The criterion seems to be supported by the parameters of 30 Mpc scale voids revealed by the galactic redshift surveys. The importance of the geometrical effect of the voids is also due to their possible role in  the temperature independent ellipticity of the excursion sets, as well as in the nature of dark energy. The role of possible larger scale inhomogeneities, from superclusters to semi-Hubble scale, in this effect remains yet unclear. } ",
        "introduction": "The role of matter inhomogeneities in the Universe is currently being explored with respect to the crucial cosmological empirical phenomena: the dark energy and cosmic microwave background (CMB). The inhomogeneities might affect the observational distance scale indicators (backreaction) responsible for the introduction of dark energy, due to the distortions of the photon trajectories while propagating through the filaments and the voids (Mattsson 2007; Wiltshire 2008; Ellis 2008; Larena et al 2008). The voids are also among the suggested reasons for non-Gaussian anomalies in the CMB; their contribution overlapped with the integrated Sachs-Wolfe effect and primordial non-Gaussianity (e.g. \\cite{anomaly1,IS,HCG}). The alignment of the vectors of the CMB low multipoles and the Cold Spot, the large-scale plane-mirroring in the CMB maps, are examples of non-Gaussianities reported for WMAP data (see \\cite{anomaly3,anomaly2,anomaly4}).  The filamentary structure, i.e. the existence of voids and walls as a basic property of the large scale distribution of galaxies, has been known for some time. The void nature of the Cold Spot is among the discussed possibilities.  The novelty of the current investigation is to assess whether the inhomogeneities can be treated as small scale, or whether semi-Hubble scale inhomogeneities are responsible for the anomalies. This choice of scales is associated with a basic dichotomy, namely, whether our own position can be considered as typical (the Copernican principle) or whether we find ourselves in a particular location e.g. at the center of a void \\cite{CMB_void1,CMB_void2,CMB_void3}.  CMB maps vs. large scale survey cross-correlations are currently more indicative of the integrated Sachs-Wolfe effect (see McEwen et al 2008), rather than for the voids.  It is remarkable that a 20-45\\% decrease in the surface brightness and number counts of NRAO Sky Survey radio sources close to the Cold Spot has been reported \\cite{Ru}, with possible association with a void, while Smith \\& Huterer (2008) found no evidence for a radio flux anomaly.  Observationally, the following parameters are attributed to typical voids: a mean scale of about 30 Mpc, a density contrast of about -1, i.e. almost absence of galaxies in the voids (e.g. \\cite{Hoyle,Cec}). Larger scale filaments, supercluster-void structures, are attributed scales of the order of 100 Mpc \\cite{E1,E2,E3}. In either case, obviously, the issue is in certain distribution of the void scales, rather then in a strict lattice-type structure.  Within the geometrical approach it was shown that voids can induce divergence of null geodesics and hyperbolicity \\cite{GK1}, even in a globally flat or slightly positively curved Universe. The role of the voids acting as divergent lenses was noticed in Das \\& Spergel (2008). Here the randomization in the temperature maps due to the hyperbolicity would be measurable, e.g. via Kolmogorov's stochasticity parameter (Gurzadyan \\& Kocharyan 2008; Gurzadyan et al 2008a); the case of the Cold Spot was also considered.  Hyperbolicity can be related to the temperature independent ellipticity of the excursion sets found in CMB maps (Gurzadyan \\& Torres 1997; Gurzadyan et al 2005, 2007).  If the inhomogeneities are able to influence the dark energy effects and be traceable in the CMB properties, then it is of particular importance, first, to reveal the conditions for such effects in terms of observable descriptors, and second, to examine whether the available observational data from large scale galaxy distribution surveys support those criteria. ",
        "conclusions": "The criterion for hyperbolicity in terms of the porosity parameter $p$ % shows that the large, underdense voids ($L_{void}> L_{wall}$, $\\ \\delta_{void} \\simeq -1$) lead to instability/hyperbolicity in the properties of null geodesics. In other words, the criterion indicates that the domination of voids in the line-of-sight direction will support the hyperbolicity unlike the directions dominated by the walls.  The comparison of the criterion with the available observational data leads to the following conclusions.  First, the surveys of large scale galaxy distribution indicate that the parameters of voids of 30 Mpc scale do support the condition for a hyperbolicity.  Second, the temperature independent ellipticity of excursion sets measured for the CMB maps (Gurzadyan et al 2005, 2007), if due to the void hyperbolicity, would require a porosity parameter $p>0.7$.  Thus, together with the degree of randomness in the CMB maps \\cite{GK2}, the porosity can provide another observational link to the voids' distribution in the Universe.  CMB maps can allow us not only to probe the parameters of the voids themselves but also the effective distance (redshift) scale of such a filamentary Universe via the $\\tau(z,\\Omega_m)$ in Eqs.(20),(21) or e.g. to reveal the role of cosmological black holes in the formation and evolution of the voids \\cite{St,Cap,Ser}.  Other inhomogeneity scales, in principle, also can contribute to this effect. This concerns the reported supercluster/void network  of around 100 Mpc (or 120 Mpc) scale \\cite{E1,E2,E3}, and of larger ones, up to the Hubble scale \\cite{Cec}. The properties and statistics for inhomogeneities of such larger scales in the galaxy distribution surveys, however, are still ambiguous and do not allow any conclusions on the fulfilment of the hyperbolicity condition for them.  {\\it Acknowledgments.} We are thankful to the referee for valuable comments."
    },
    "0807/0807.0904_arXiv.txt": {
        "abstract": "{We present the results of an extensive ground-based photometric and spectroscopic campaign on the $\\gamma$\\,Dor CoRoT target HD\\,49434. This campaign was preparatory to the CoRoT satellite observations, which took place from October 2007 to March 2008. } {Whereas satellite data will be limited to the detection of  low-degree pulsation modes with poor  identification (no filters), ground-based data will provide eventually the identification of additional modes and, through the spectroscopic data, detection of additional high-degree modes as well as an  estimate of the azimuthal number $m$. Our aim was to detect and identify as many pulsation modes as possible from the ground-based dataset of the $\\gamma$\\,Dor star HD\\,49434, to anticipate the CoRoT results. } {We searched for frequencies in the multi-colour variations, the pixel-to-pixel variations across the line profiles and the moments variations of a large dataset consisting of both multi-colour photometric and spectroscopic data from different observatories, using different frequency analysis methods. We performed a tentative mode identification of the spectroscopic frequencies using the Moment Method and the Intensity Period Search Method. We also carried out an abundance analysis. } {The frequency analysis clearly shows the presence of  four frequencies in the 0.2--1.7 \\cd\\,interval, as well as six frequencies in the 5--12\\,\\cd\\, domain. The low frequencies are typical for $\\gamma$\\,Dor variables while the high frequencies are common for $\\delta$\\,Sct pulsators. We propose the frequency 2.666\\,\\cd\\, as a possible rotational frequency. All modes, for which an identification was possible, seem to be high-degree modes ($3 \\leq \\ell \\leq 8$). We did not find evidence for a possible binary nature of HD\\,49434. The element abundances we derived are consistent with the values obtained in previous analyses. } {We classify the $\\gamma$\\,Dor star HD\\,49434 as a hybrid pulsator, which pulsates simultaneously in $p$-- and $g$--modes. This finding makes  HD\\,49434 an extremely interesting target for  asteroseismic modelling. } ",
        "introduction": "\\label{intro}  We currently are in a challenging era of asteroseismology. Observers and theoreticians are preparing for the promising output of the CoRoT satellite mission (Baglin et al. 2006), which has been launched successfully in December 2006. The CoRoT science programme is twofold. A first goal of the mission involves the detection of extrasolar planets by using the transit method. The asteroseismic goal, on the other hand, aims at probing the internal structures of stars by extracting information from detected pulsation frequencies (several tens, even hundreds, of frequencies, depending on the pulsational class), and, in general, to understand better the process of stellar evolution. To this end, a selection of stars of a wide range of masses, with different evolutionary status, of different spectral type and belonging to different pulsation classes are monitored.  The CoRoT satellite mission goes hand in hand with a huge observational effort from the ground. Not only have preparatory observations been a key stone in the selection process of suitable targets for asteroseismology (Poretti et al. 2003, 2005), also now the satellite is operating from space, simultaneous ground-based observations are very important and complementary.  Indeed, as CoRoT provides high-precision time-series in white light, multi-colour photometry provides colour information, which allows identification of the degree $\\ell$ by means of amplitude ratios and phase shifts. High-resolution spectroscopy allows the detection of high-degree modes and the identification of both the degree $\\ell$ and the azimuthal order $m$ by means of, for example, the Moment or Doppler Imaging methods. With the goal of obtaining multi-colour and spectroscopic time-series of a selection of $\\delta$\\,Sct, $\\gamma$\\,Dor, $\\beta$ Cep and Be CoRoT primary and secondary targets, an extended ground-based campaign, involving both high-resolution spectrographs and multi-colour photometric instruments, has been included in the scientific plan of the CoRoT mission (Catala et al. 2006; Uytterhoeven \\& Poretti 2007; Uytterhoeven et al. 2008).  The $\\gamma$\\,Dor stars are a challenging class of pulsators for several reasons. They are a fairly recently discovered pulsational group (Kaye et al. 1999) and, even though the number of class members is rapidly increasing (more than 130 candidates are known), the number of well-studied cases based on extended time-series is small. Obtaining a good phase coverage from the ground is an observationally challenging task as the intrinsic variations show periods of the order of a day (0.4--3 days). Moreover, only a few frequencies have been detected so far in individual stars, and amplitudes are fairly small (below 0.05 mag; 2 \\kms). Nevertheless, the $\\gamma$\\,Dor stars are promising targets for seismic studies as they show $g$--modes, which can be used to probe the deep interior of the star. Over the last few years progress has been made in the understanding of the pulsational mechanism. Probably a flux modulation induced by the upper convective layer is the driving source for the pulsations (Guzik et al. 2000; Dupret et al. 2004; Grigahc\\`ene 2004). However, some  details need clarification, such as the the thickness and the depth of the convection zone and the importance of diffusion.  The $\\gamma$\\,Dor stars (Spectral Type F0V-F2V) occupy an interesting region along the Main-Sequence in the Hertzsprung-Russell (HR-) diagram, being enclosed by the classical instability strip and the instability strip of the solar-like stars.  An interesting investigation concerns the presence of hybrid stars, i.e. stars that exhibit both $p$--mode and $g$--mode oscillations, in the overlap region of the $\\gamma$\\,Dor and $\\delta$\\,Sct instability strips (Dupret et al. 2005).  Observational evidence for self-excited hybrid stars is scarce, and less than a handful genuine hybrid candidates exist, including HD\\,8801 (Henry \\& Fekel 2005), HD114839 (King et al. 2006) and BD+18\\,4914 (Rowe et al. 2006).  The hybrid stars are of particular interest as both the envelope and the deep interior of the star can be probed through the study of $p$--modes and $g$--modes, respectively.  In this paper we focus on the ground-based data of the $\\gamma$ Dor star HD\\,49434 (V=5.75, F1V). As HD\\,49434 is a primary CoRoT target of a 150-days Long Run (LRa1; October 2007 -- March 2008), several teams performed a detailed study of abundances and fundamental parameters of this star in the preparatory framework of the mission. Summarising their fairly consistent results, we find $T_{\\rm eff}=7300 \\pm 200$ K, $\\log g = 4.1 \\pm 0.2$ dex, $[Fe/H] = -0.1 \\pm 0.2$, \\vsini $= 84 \\pm 5$ \\kms (Lastennet et al. 2001; Bruntt et al. 2002, 2004; Mathias et al. 2004; Masana, Jordi \\& Ribas 2006). Gillon \\& Magain (2006) derived the slightly higher values $T_{\\rm eff} = 7632 \\pm 126$ K and $\\log g = 4.43 \\pm 0.20$ dex. In addition, we estimated these parameters from $uvby \\beta$ and Geneva photometry taken from the GCPD catalogue ({\\it General Catalogue of Photometric Data}, Mermilliod et al. 1997).  From the calibration of Napiwotzki et al. (1993) we obtained $T_{\\rm eff,[u-b]} = 7300$~K, $T_{\\rm eff,[b-y]} = 7230$~K and $\\log g = 4.1$\\,dex. From the Geneva indices $(B2-V1)$, $d$ and $m2$ as proposed by K{\\\"u}nzli et al. (1997), we determined $T_{\\rm eff} = 7200 \\pm 60$~K and $\\log g = 4.34 \\pm 0.08$\\,dex.  Taking into account all these values of the atmospheric parameters, we will adopt in our analysis effective temperature and surface gravity equal to $7300$~K and $4.2$~dex, respectively. From 2MASS IR photometry Masana, Jordi \\& Ribas (2006) estimate an angular semi-diameter of $0.186 \\pm 0.002$ mas, and $R = 1.601 \\pm 0.052 R_{\\odot}$. An estimate of the mass of HD\\,49434 $M = 1.55 \\pm 0.14 M_{\\odot}$ is obtained from evolutionary tracks by Bruntt et al. (2004). A first study of the intrinsic variability of HD\\,49434 is presented by Bruntt et al. (2002). In a time series of Str\\\"omgren data they found no clear frequencies, but reported a power excess in the frequency range typical for $\\gamma$\\,Dor pulsators. A line-profile analysis suggested the presence of one or several high-degree modes, in the same low frequency range. Line-profile variations were confirmed by Mathias et al. (2004). In Fig.~\\ref{CMD} we show the position of HD\\,49434 in a Colour-Magnitude diagram, based on the $M_v$ value derived from the HIPPARCOS satellite and $(b-y)_0$, calculated from the photometric indices taken from Hauck \\& Mermilliod (1998) using the TempLogg method (Kupka \\& Bruntt 2001). We used the evolutionary tracks, instability strips and ZAMS as described by Poretti et al. (2003). As can be seen, HD\\,49434 lies close to the blue border of the $\\gamma$\\,Dor instability strip, and inside the instability strip of $\\delta$\\,Sct stars.  \\begin{figure} \\centering \\resizebox{1.0\\linewidth}{!}{\\rotatebox{-90}{\\includegraphics{Uytterhoevenfig1.ps}}} \\caption{Colour-Magnitude ($(b-y)_0$ versus Absolute Magnitude $M_v$) diagram indicating the position of HD\\,49434 (star). The $(b-y)_0$ and $M_v$ values of HD\\,49434  (an error bar is given in the left bottom corner) are taken from Hauck \\& Mermilliod (1998) and from the HIPPARCOS satellite data, respectively. Dashed and solid tracks indicate evolutionary tracks for the overshooting extension distances $d_{over}$=0.1 and 0.2, respectively (see Claret 1995 for details). The ZAMS and the borders of the $\\delta$\\,Sct (longer black lines) and $\\gamma$\\,Dor (shorter gray lines) instability strips are indicated by solid lines.} \\label{CMD} \\end{figure} ",
        "conclusions": "In this paper we presented frequency analysis and preliminary spectroscopic mode-identification results of HD\\,49434, obtained from the most extensive ground-based dataset for a $\\gamma$\\,Dor star to date. Photometric and spectroscopic observations were performed between 2005 and 2008 in the framework of the CoRoT ground-based support campaign.  Our central finding is the discovery of the hybrid nature of the star: HD\\,49434 pulsates simultaneously in $p$- and $g$-modes. The frequency analysis clearly shows the presence of at least four $\\gamma$\\,Dor-type, as well as  six $\\delta$\\,Sct-type of frequencies.  Additional frequencies are expected to contribute to the observed variability, but have amplitudes near or below our current detection limit.   All modes, for which an identification was possible, seem to be high-degree modes ($3 \\leq \\ell \\leq 8$), which make future modelling extremely challenging.   Also the fairly fast rotation of HD\\,49434 ($v_{\\rm eq} \\geq 87$ \\kms) will complicate the modelling since the frequencies in the observers frame are rotationally modified with respect to the ones in the stellar rest frame. Actually, a high-degree, $m=-4$ mode can shift the observed frequencies by 10 \\cd, which means that a $p$--mode may appear with a typical $g$--mode frequency ($m<0$) or vice-versa. We expect that the CoRoT time-series will provide additional information on the rotational frequency through the effect of rotational splitting.  In this context, the frequency 2.666(2)\\,\\cd\\, can possibly be interpreted as the rotational frequency $f_{\\rm rot}$ of the star. Its first harmonic, 5.3311\\,\\cd, is detected in the line-profile variations, and if the $f_{\\rm rot}$ value is confirmed, this could imply the presence of surface inhomogeneities. From the asymmetric behaviour of the amplitude distributions in the IPS diagrams (see Figs.\\,\\ref{fig1} and \\ref{mode-id}) we find indications for local temperature variations on the stellar surface (Schrijvers \\& Telting 1999), but we could not detect a clear periodicity in the corresponding EW variations.  A similar behaviour has been noted in other stars, such as the $\\delta$\\,Sct star FG\\,Vir (Zima et al. 2006) and the $\\beta$ Cep pulsator $\\varepsilon$\\,Cen (Schrijvers et al. 2004).   We encountered several puzzling results. For instance, it is not clear why we failed to detect the dominating frequency from the moment variations, 1.2732 \\cd, in the variations across the line profiles, and if this frequency can be linked with the photometric frequency 1.73480 \\cd.  It is clear that the pulsational behaviour of HD\\,49434 is very complex  and that the fast rotation adds to this. The simultaneous presence of several low-amplitude variations complicates the frequency search, which often results in an ambiguous determination of the frequency values (see Table~\\ref{frequencies}).  Moreover, some technical aspects of the analysis add to the confusion, as the delicate process of the alignment of the data.  The detection of the hybrid nature of HD\\,49434 makes it an interesting asteroseismic target. The asteroseismic importance of hybrid $\\gamma$ Dor/$\\delta$\\,Sct stars is obvious as both the envelope and the deep interior of the star can be probed through the study of $p$- and $g$-modes, respectively, and thus giving a unique opportunity to constrain the macro-physics of the stellar interior.  The amount of hybrid detections is steadily growing, although still only a handful of cases is known. The first discovery of its kind was made in HD209295 (A9/F0V, Handler et al. 2002). However, this primary of an eccentric binary might be not a 'bona fide' hybrid variable as the excitation of the $\\gamma$\\,Dor-type variations is possibly induced by tidal interaction. To date at least four other hybrid $\\gamma$ Dor/$\\delta$\\,Sct stars have been reported, of which two were discovered from MOST satellite data: HD~8801 (A7m, Henry \\& Fekel 2005), HD114839 (Am, King et al. 2006), BD+18~4914 (F5, Rowe et al. 2006), and HD~44195 (F0, Uytterhoeven et al. 2008). The simultaneous presence of $p$- and $g$-modes has also been theoretically predicted (Dupret et al. 2004), and hence might not be as uncommon as it currently seems.  Our results on HD\\,49434 show that the $\\delta$\\,Sct-type modes have high degrees ($3 \\leq \\ell \\leq 8$, see Table~\\ref{IPStable}), and consequently are only detected in spectroscopic variations (Table~\\ref{frequencies}) as high-degree modes tend to cancel out in the integrated photometric light.  This observational selection effect can explain why Handler and Shobbrook (2002) in their systematic survey for $\\delta$\\,Sct variations in $\\gamma$\\,Dor stars, using only photoelectric photometry, failed to detect hybrid variables other than HD209295. We can even speculate that actually a large part of all $\\gamma$\\,Dor variables pulsate both in $p$- and $g$-modes, but the high-degree $p$-modes need spectroscopic detection.  This finding would make the $\\gamma$\\,Dor pulsating class extremely suited for asteroseismic studies.  In a subsequent series of papers we will confront the ground-based data, enlarged with a new spectroscopic time-series obtained in January/February 2008, with the actual CoRoT space data with the goal to perform a deeper investigation and to present a modelling of the star.  The combination of continuous and accurate space photometry, and ground-based spectroscopy promises to be the best way to study the frequency spectrum of HD\\,49434 in detail. Results from the MOST satellite show that not all high-degree spectroscopic modes can be recovered from photometric space data (e.g.\\, $\\zeta$ Oph, Walker et al. 2005), and consequently the spectroscopic observations will play an important part in the unraveling and understanding of the pulsational behaviour of the star.  As a matter of fact, the CoRoT satellite mission will provide an excellent opportunity to investigate the pulsational behaviour of a large amount of $\\gamma$\\,Dor variables, as in the forthcoming years continuous time-series of thousands of stars in the exo-planetary CCD field will become available. Also, with at least two hybrid stars selected in the asteroseismic field of CoRoT, the target of this paper HD\\,49434 and HD~44195, we have good prospects of making progress in modelling and understanding the relation between $p$- and $g$-mode pulsators."
    },
    "0807/0807.2240_arXiv.txt": {
        "abstract": "We explore  the role of the group  environment in the evolution of  AGN at the  redshift interval $0.7<z<1.4$,  by combining deep Chandra observations with extensive optical spectroscopy from the All-wavelength Extended Groth strip International Survey (AEGIS).  The sample consists of 3902 optical  sources and 71 X-ray AGN. Compared to the overall  optically selected galaxy population, X-ray  AGN are more frequently  found in  groups at  the 99\\%  confidence level.   This is partly  because AGN  are hosted  by red  luminous galaxies,  which are known to reside, on average, in dense environments.  Relative to these sources, the excess of X-ray AGN  in groups is significant at the 91\\% level only.  Restricting the  sample to $0.7<z<0.9$ and $M_B<-20$\\,mag in order  to control systematics we  find that X-ray  AGN represent $(4.7\\pm1.6)$ and $(4.5\\pm1.0)$\\% of  the optical galaxy population in groups and  in the field  respectively.  These numbers  are consistent with the AGN fraction in  low redshift clusters, groups and the field. The  results  above, although  affected  by  small number  statistics, suggest that X-ray  AGN are spread over a  range of environments, from groups to the field, once  the properties of their hosts (e.g. colour, luminosity)  are accounted  for.   There is  also tentative  evidence, significant  at the  98\\% level,  that the  field produces  more X-ray luminous  AGN compared  to groups,  extending similar  results  at low redshift to $z\\approx1$.  This trend may be because of either cold gas availability or the nature of the interactions occurring in the denser group environment (i.e.  prolonged tidal encounters). ",
        "introduction": "The environment of Active  Galactic Nuclei (AGN) holds important clues on the physical mechanism(s) responsible for the accretion of material on the  central supermassive black hole  (SBH).  Numerical simulations have  shown that an  efficient way  of channeling  gas to  the nuclear galaxy regions and in the  vicinity of the SBH are galaxy interactions and  mergers  (e.g.  Barnes  \\&  Hernquist  1992;  Mihos \\&  Hernquist 1996). As a result, many models for the formation and the cosmological evolution of AGN assume that the  main growth phase of the central SBH in galaxies occurs in major  mergers (e.g. Kauffmann \\& Haehnelt 2000; Di Matteo, Springel  \\& Hernquist; Hopkins et al.   2005; Bower et al. 2006;  Croton  et  al.   2006).   Recent  observations  of  the  local environment of  luminous QSOs supports  this scenario.  The  excess of optical  neighbours  around $z<0.4$  QSOs  on  small scales  ($\\approx 0.1-0.5$\\,Mpc) compared to $L_{\\star}$  galaxies (Serber et al.  2006) and  the higher  fraction of  QSO pairs  with  separations $<0.1$\\,Mpc compared to  the expectation from large scales  ($>3$\\,Mpc; Hennawi et al.  2006; Myers et al.   2007), both suggest that galaxy interactions may play  an important role in  the evolution of QSOs.   This does not appear to be  the case however, for the more  common but less powerful narrow-line AGN (Seyfert, Low Ionization Nuclear Emission Line Regions -- LINERs) selected in the Sloan Digital Sky Survey (SDSS) at $z<0.3$. The small  scale environment  of these sources  ($<$0.1\\,Mpc) suggests that galaxy interactions alone cannot explain the observed activity in the bulk of  this population (Li et al.   2006).  Modeling work indeed predicts that  low luminosity  AGN, such as  Seyferts and  LINERs, may accrete   cold  gas   stochastically   by  mechanisms   such  as   bar instabilities or  tidal disruptions rather than  major interactions or mergers (Hopkins \\& Hernquist 2006).  On scales  much larger than  0.5\\,Mpc, the clustering of  AGN measures the mass  of their  host dark matter  halo and  how it relates  to the observed  nuclear activity.  At  these scales,  an increasing  body of evidence underlines the potential significance in the evolution of AGN of the group environment.  Early  studies for example, have shown that low redshift  luminous QSOs are associated  with moderate enhancements in the distribution  of galaxies and avoid rich  clusters (Hartwick \\& Schade 1990; Bahcall  \\& Chokshi 1991; Fisher et  al.  1996; McLure \\& Dunlop 2001). Moreover, the evolution of the large scale clustering of broad line QSOs  in the redshift range $0.5-2.5$  is consistent with a parent dark matter halo of  almost constant mass in the range $10^{12} - 10^{13}\\, \\, h^{-1} \\rm M_{\\odot}$  (Porciani et al.  2004; Croom et al.  2005; Hopkins et al.  2007; da Angela et al. 2007; Mountrichas et al.  2008;  but see  Coil et  al.  2007), close  to the  threshold for galaxy groups ($6 \\times 10^{12} \\, \\, h^{-1} \\rm M_{\\odot}$ for DEEP2 groups with velocity  dispersion $\\rm \\sigma_v>200\\, km\\,s^{-1}$; Coil et  al 2006) The  more common  narrow emission-line  AGN in  the local Universe  ($z\\approx0.1$) are  hosted by  massive  early-type galaxies (Kauffmann  et al.  2003),  which are  known to  be more  clustered on Mpc-scales than  the overall galaxy population (Zehavi  et al.  2005). A luminosity  dependence of the  clustering has also been  found, with the most  luminous narrow-line AGN  residing in lower  density regions compared to  less luminous ones (Kauffmann  et al.  2004;  Wake et al. 2004; Constantin \\& Vogeley  2006).  Popesso \\& Biviano (2006) further showed that  not only the luminosity but  also the number of  AGN is a strong  function  of local  density.   In  a  sample of  low  redshift clusters,   they  found   that   the  fraction   of  narrow-line   AGN anti-correlates  with cluster  velocity dispersion,  increasing toward galaxy groups.  At  higher  redshift,  $z\\approx1$,  X-ray surveys  provide  the  most efficient and least biased way  of finding active SBH to explore their demographics.  The large scale structure  of the X-ray selected AGN at $z\\approx1$  is usually determined  using the  angular or  the spatial auto-correlation functions  (e.g. Yang et al.  2004;  Basilakos et al. 2004, 2005; Gilli et al.   2005; Miyaji et al.  2007).  Although there is  substantial scatter  in the  results from  these  studies, largely because  of  observational  uncertainties  (Miyaji et  al.   2007)  or luminosity  dependent clustering  (Plionis et  al. 2007),  the general consensus is that X-ray selected  AGN are hosted by dark matter haloes with  mass of about  $10^{13} -  10^{14} \\,  \\, h^{-1}  \\rm M_{\\odot}$ (Miyaji et  al. 2007), i.e.  in  the range expected  for galaxy groups (e.g.  Zandivarez, Merchan, \\& Padilla 2003).  Why AGN  show a  preference for the  group scale environment  is still unclear.   One  possibility is  that  galaxy  interactions, which  are believed  to be one  of the  AGN triggers,  are suppressed  in massive clusters, while  they proceed more  efficiently in the  lower velocity dispersion groups  of galaxies.  Alternatively,  it might be  that the necessary ingredients  for AGN activity, sufficient  cold gas supplies and a  SBH, are  often found in  moderate density  environments.  More massive galaxies,  which are likely  to host larger black  holes (e.g. Ferrarese \\& Merritt 2000; Gebhardt et al.  2000), are associated with high  density  regions  (e.g.   Hogg  et  al.   2003;  Cooper  et  al. 2006).  Moreover, cold  gas, the  fuel of  AGN, is  likely to  be more abundant in  moderate density environments, while it  is depleted more efficiently in  high density regions,  such as clusters.   Despite the potential significance  of galaxy groups, there are  still few studies that attempt  to directly  constrain the incidence  of active  SBHs in these moderate  density environments.  Shen  et al. (2007)  found that the fraction of AGN (both X-ray and optically selected) in a sample of eight  poor groups  at $z  \\approx 0.06$  is consistent  with  that of clusters or  the field (Martini  et al. 2007).  Although  small number statistics are a  issue in these studies, the  finding above is against scenarios where  the group environment promotes  AGN activity, because of e.g.  the higher efficiency of galaxy-galaxy mergers.  At higher redshift, $z\\approx1$, close  to the peak of the AGN density in the  Universe (e.g.   Barger et al.  2005; Hasinger et  al.  2005), there  is still  no information  on  the association  between AGN  and groups.   This is  mainly  because  it is  hard  to identify  moderate overdensities  at high  redshift, as  it requires  large spectroscopic surveys  of  optically faint  galaxies.   The All-wavelength  Extended Groth-strip  International  Survey (AEGIS;  Davis  et  al.  2007)  has extensive optical spectroscopy that can  be used to identify groups at $z\\approx1$  and deep  {\\it Chandra}  observations, which  provide the most efficient way for finding  AGN, particularly at high redshift. In this paper we combine the  AEGIS group and X-ray catalogues to explore the   fraction   of  X-ray   selected   AGN   associated  with   these overdensities.  We adopt $H_{0} = \\rm 70 \\, km \\, s^{-1} \\, Mpc^{-1}$, $\\Omega_{M} = 0.3$ and $\\Omega_\\Lambda = 0.7$. ",
        "conclusions": "In  conclusion,  it  is  found  that  at  $z\\approx1$,  AGN  are  more frequently found in groups  compared to the overall optically selected galaxy population at the 99  per cent confidence level. This excess is to some level  because X-ray AGN are associated  with red and luminous galaxies which are known to  be more clustered than the overall galaxy population.  After taking into  account the AGN host galaxy properties the significance of the excess of  AGN in groups drops to 91 per cent. This suggests that X-ray AGN live  in a range of environments and that groups do not produce more AGN  than the field, once their host galaxy properties are factored in  the analysis.  There is tentative evidence that AGN in the field are  more powerful, on average, than their group counterparts,   extending   similar  results   at   low  redshift   to $z\\approx1$.  This can be attributed to differences in the gas content of  group/field galaxies and/or  the nature  of interactions  in these environments.  Although  the results  presented  here are  affected  by small  number statistics, it  is the first time  that an attempt is  made to explore the relation between SBH  accretion and optically identified groups at $z\\approx1$.  There are two obvious ways to improve the reliability of our  conclusions.  Either  perform deep  X-ray observations  in fields with extensive optical spectroscopy,  like those surveyed by the DEEP2 project,  or complement  existing  deep X-ray  surveys (e.g.   COSMOS, Chandra Deep Fields) with  optical spectroscopy of similar quality and extent to the DEEP2 spectroscopic survey."
    },
    "0807/0807.0590_arXiv.txt": {
        "abstract": "According to the Hubble law, high redshift objects such as Quasar (QSOs), X-ray Active Galactic Nuclei (AGN) together with the Gamma Ray Burst (GRBs) are the fastest and farthest objects. These characteristics provides strong motivations for they to be used to constrain the cosmological parameters, without the limitations found in the Ia supernovae study and which is restricted to redshift of up to 1.7. However, the variability and behavior in the QSOs and AGNs spectra tell us that they have very complex structures and the standard candle framework can not be applied.  So far the available data of QSOs and AGNs have shown some anomalies observed in their  brightness and metallicities, difficult to understand, under an orthodox point of view. Here, we show that their Hubble diagram flattens for $z\\geq 3 $. The result need further confirmation, because the statistics of high redshift extragalactic objects is still poor. Details and some implications of these results are reported in this work. ",
        "introduction": "The quasars (QSOs) were discovered by radio telescopes in the late 1950s. Many were recorded as radio sources with no corresponding visible object. It is believed that the QSOs are extremely bright and distant active galactic nuclei (AGN) of young galaxies. They were first identified as being high-redshift sources of electromagnetic energy, including radio waves and visible light that were point-like. More than 60,000 quasars are known, all observed spectra have shown considerable redshifts, ranging from 0.06 to the recent maximum of 6.4. Quasars are found to vary in luminosity on a variety of time scales. The Sloan Digital Survey has uncovered the most distant quasars known (up to $z\\sim 6.4$). In addition, there is also the report of high redshift X-ray AGNs with moderate to high quality spectra observed by XMM-Newton and Chandra.  On the other hand, a study made by Gehren et al \\citep{gehren84} has shown that quasars are usually located in groups or clusters of galaxies which are dominated by the luminosity contribution of the quasar host galaxy. It is suggested that most radio-loud quasars are situated in giant elliptical galaxies in which a considerable amount of hot gas and young stars may be present owing to recent collisions or tidal interaction with nearby cluster members. In addition, from a study of pairs of quasars and galaxies, which were very close in the sky, \\citep{stockton78} found that, in a fraction ($\\sim 50\\%$) of the sample of quasar-galaxy, the redshifts of the galaxies agreed with the redshifts of the quasars. In other words, quasars are associated with galaxies that have the same redshift as the quasar, and have just the brightness expected, if the quasars are at their cosmological distances. However, according to the standard cosmology, the study can include only quasars at low or moderate redshifts, because the $(1+z)^{-4}$ dimming in surface brightness makes normal galaxies unobservable at large redshifts \\citep{lehnert92}.  However, if the quasar-galaxy pair has different redshifts ($z_q>z_G$), just as it happens in $\\sim 50\\%$ of the Stockton's sample, the situation becomes controversial. Burbidge and coworkers have shown \\citep{burbidge90} that there is strong evidence that normal galaxies and QSOs tend to clustered, whether or not their redshifts are the same. While if the quasar redshift is a distance indicator, as is advocated by the standard cosmology, the identification of quasars shining through or close to galaxies offer a way to probe gravitational lensing effects. From quasar-galaxy gravitational lensing, is possible to deduce that in some quasar-galaxy pairs with very different redshifts, the association is only apparent. Despite the small angular separation between the quasar and its galaxy companion, the quasar must be beyond the lens galaxy. The method works only for galaxy lenses, with $z\\geq 0.5$ \\citep{dar91}. However, a more accurate analysis on gravitational lensing effect shows that gravitational lens measurements can be a powerful tool, but only if a variety of major difficulties are overcome, such as the multiple lenses and microlensing background and the small variability amplitude typical of most quasar \\citep{shechter04}.  In this paper, we have made an analysis on the Hubble diagram of extragalactic high-z objects (QSOs and AGNs), on the basis of the more recent data as possible. This paper is organized as follows: In Sec.2, extra galactic objects such as QSOs and AGNs at high redshift are described. In Sec.3, the extragalactic objects in a cosmological context is presented, ana Sec.4, contains our conclusions. ",
        "conclusions": "The high redshift objects such as QSOs, AGNs and GRBs open the doors for the study of the universe in another scale, the so called ``high-z universe''. So far, there are many controversial aspects on the nature of these high redshift objects. The difficulties lie in the fact that for instance the QSOs and AGNs searches are still quite inhomogeneous over the sky, and probably only large-scale automated surveys will be able to resolve this, and to answer some questions such as: Is there evidence of quasars occurring in the direction of nearby galaxies? Can such an excess be explained by something like gravitational lensing?  Recently, the concept of standard candle was applied to the GRB studies \\citep{shaefer03,shaefer07,ghirlanda04,ghirlanda06}. In this approach the Hubble diagram and the cosmological parameters were obtained. It's claimed an agreement with the standard cosmology inside  $2\\sigma$ to $3\\sigma$ confidence level. The information that is obtained from GRBs is still dependent on the concept of standard candle. Even so, remarkable effort has been made to calibrate the GRBs and to place them in conditions to constrain the cosmological parameters. In addition, the statistic is still incipient in the high redshift region to reach a more robust conclusion.   So far, QSOs, AGNs and GRBs have been observed with redshifts beyond 6. Unhappily up to now, the QSO and AGN luminosity function is a model dependent parameter. Even so, valuable information can be obtained starting from basically  two observational parameters, the redshift, and the apparent magnitude. In this paper, we have made an independent analysis on the Hubble diagram of extragalactic high-z objects (QSOs and AGNs), on the basis of the more recent data as possible.  It is possible to observer in the SDSS data a ``duck beak'' in the correlation between redshift and the apparent magnitude of quasars at $z\\sim 3$ as a transition for a flat correlation. Similar  upside ``duck beak'' at $z\\sim 3$ is also observed in the   $Veron\\& Veron$ data, in the correlation between redshift and the luminosity function obtained using the $\\Omega$ model. In addition the high quality quasar FBQS sample, has been extended using the data of several X-ray AGNs.  In general, the result for QSOs and X-ray AGNs at high redshift ($z>3$) is  a relative flat correlation between the apparent magnitude and redshift, as well as between the luminosity function vs redshift, and that is reflected in the correlation between the magnitude modulus vs redshift. This result needs further confirmation, because the statistic of high redshift objects is still poor. If confirmed, this result is hard to conciliate with the standard cosmological model."
    },
    "0807/0807.3888_arXiv.txt": {
        "abstract": "The S-Z effect and Faraday rotation from halos are examined over a  wide mass range, an including gas condensation and magnetic field evolution. Contributions to the CMB angular power spectrum  are evaluated for galaxy clusters, galaxy groups and galaxies. Smaller mass halos are found to play a more important role than  massive halos for the $B$-mode polarisation associated with the S-Z CMB anisotropies. The $B$-modes from Faraday rotation dominate the secondary $B$-modes caused by gravitational lensing at $\\ell > 3000$. Measurement of $B$-mode polarisation in combination with the S-Z power spectrum can potentially provide important  constraints on intracluster magnetic field and gas evolution at early epochs. ",
        "introduction": "The baryonic components in galaxy clusters, galaxy groups and galaxies  vary  widely in their evolution. Some fraction of the baryons remains gaseous at the virial temperature and is threaded by microgauss magnetic fields, while a fraction condenses into the disk and forms stars. Therefore, knowing the temperature, density and the condensation rate of  the baryons is an important clue for revealing the history of star formation, galaxy formation and even the magnetic field evolution and the reionisation process of the Universe.  The Sunyaev-Zel'dovich effect (S-Z effect) is an important tool for the investigation of baryon properties in halos \\citep{sunyaev-zeldovich-1972}. The S-Z effect is the scattering of cosmic microwave background (CMB) photons by the hot electron gas in the gravitational potential wells of halos. As a result, the CMB frequency spectrum is distorted from the blackbody spectrum. This distortion depends on the temperature and number density of the electron gas. In combination with X-ray observations, measuring the S-Z effect from galaxy clusters at each redshift gives the evolution of physical properties of the baryons \\citep{morandi-ettori-2007}. Moreover, the temperature and the number density of the electron gas are related to the masses of halos, so that the number counts of detectable S-Z signals for different observational thresholds can provide the mass distribution of galaxy clusters. The distribution is sensitive to the cosmological parameters and to the non-Gaussianity of primordial density fluctuations. Therefore, the number counts of S-Z galaxy clusters can also give a constraint on the cosmological model \\citep{bartlemann-sz-2001,moscardini-bartlemann-2002, schafer-pfmmer-2006, sadeh-rephaeli-2007}. Many future CMB surveys will make important contributions, and in some cases are even  dedicated, to the measurement of S-Z-selected  galaxy clusters (e.g. {\\sc Planck}\\footnote{http://www.rssd.esa.int/index.php?project=Planck}, SPT\\footnote{http://pole.uchicago.edu/}, ACT\\footnote{http://wwwphy.princeton.edu/act}).  The S-Z effect from unresolved S-Z clusters is observed as CMB temperature anisotropies. These anisotropies are one of the major secondary anisotropies fand dominant components for the high $\\ell$ modes ($\\ell >2000$) in the angular temperature power spectrum. The amplitude strongly depends on the amplitude of matter density fluctuations, $\\sigma_8$. Small angular-scale CMB anisotropy experiments, BIMA \\citep{dawson-bima}, CBI \\citep{mason-cbi} and ACBAR \\citep{kuo-acbar}, have measured CMB anisotropies at high $\\ell$ modes and detected the excess signal expected to be due to the S-Z effect. This excess corresponds to the S-Z effect with $\\sigma_8 \\sim 1.0$ \\citep{bond-szsigma, douspis-aghanim-2006}. However, this high value is in conflict with the 5 year WMAP result, $\\sigma_8 \\sim 0.8$, which is obtained from large-scale temperature anisotropies \\citep{wmap-spergel-07}. Note that this conflict is lifted when a possible contribution from unresolved point sources is taken into account \\citep{douspis-aghanim-2006}. In order to check whether the excess indeed is due to  the effect of large $\\sigma_8$, we need detailed observational data on the S-Z number counts by future observations.  In this paper, we re-examine the number counts of S-Z halos and S-Z CMB anisotropies. Our aim is to clarify which halos with different masses contribute to any particular $\\ell$ range. In particular, we focus on the S-Z effect of galaxy groups and galaxies ($M<10^{14} M_\\odot$). In those halos, the condensation of baryonic gas is effective. Clearly then, the S-Z effect is expected to depend on the evolution of gas condensation.  We also study Faraday rotation by magnetic fields in galaxy clusters, galaxy groups and galaxies. Many observations suggest that galaxy clusters and galaxies have magnetic fields whose amplitude is typically of the order of $1$--$10$ $\\mu$Gauss. The coherence length of such magnetic fields can be as large as cluster virial scales, and structured magnetic fields have been found even on larger scales. \\citep{kim-kronberg-1989}. Besides, magnetic fields have been measured with the same amplitudes in high redshift objects at $z>2$ \\citep{athreya-kapahi-1998}. However, the evolution of these fields has hitherto been unclear and is the one of the open questions in modern cosmology. In order to answer this question, we need statistical discussions of magnetic fields. In this paper, we discuss the potential of Faraday rotation as a probe of magnetic fields of galaxy clusters, galaxy groups and galactic halos by studying two observational methods. One is the number counts of Faraday rotation for S-Z halos. We calculate the number counts, and discuss the sensitivity of properties and evolution of magnetic fields in halo objects.  The other is the CMB $B$-mode polarisation observation. Faraday rotation distorts the CMB polarisation fields like gravitational lensing, so that it induces the $B$-mode polarisation from the $E$-mode. The angular power spectrum due to this effect in halos was studied by \\citet{takada-ohno} and \\citet{tashiro-faraday}. In particular, \\citet{tashiro-faraday} pointed out that the angular power spectrum is sensitive to the magnetic field evolution and magnetic fields in galaxies can make a more dominant contribution than those in galaxy clusters. However, their adopted magnetic fields are simple toy models. Therefore, for a more detailed discussion, we study the CMB $B$-mode polarisation due to magnetic fields in halo objects based on observations and numerical simulations.   The paper is organised as follows.  In Sec.~II, we discuss the profiles of electron density and magnetic fields that we use, motivated by observations and numerical simulations. We calculate the number counts of halos for the S-Z effect and Faraday rotation in Sec.~III. We discuss the impact of gas condensation and magnetic field evolution on the number counts. In Sec.~IV, we compute the power spectrum of CMB temperature anisotropies by the S-Z effect and $B$-mode polarisation due to Faraday rotation. Sec.~V is devoted to discussion and summary. Throughout the paper, we use the following cosmological parameters: $h=0.7 \\ (H_0=h \\times 100 {\\rm km/s / Mpc})$, $T_0 = 2.725$K, $\\Omega _{\\rm B} =0.044$ and $ \\Omega_{\\rm M} =0.27$ and $\\sigma_8 =0.8$. ",
        "conclusions": "In this paper, we examine the S-Z effect and Faraday rotation from halos which have a wide mass range ($10^{12}M_\\odot < M <10^{17}M_\\odot$). In particular, we take into account the effects of gas cooling based on recent observations and consider two cases for magnetic field evolution motivated by theoretical considerations and numerical simulations.  First, we calculate the halo number counts for the S-Z effect. We find that the number counts for the S-Z effect are sensitive to gas condensation. The redshift distribution is suppressed by gas condensation over all redshifts. In the mass distribution, gas condensation suppresses the number counts on  small scales. As a result, the peak position of the distributions shifts to larger mass scales.  We also calculate the rotation measurement for S-Z halos. The distribution of number counts strongly depends on the magnetic field evolution. The detectable minimum mass is decided by the threshold of rotation measure. Therefore, by varying the threshold we can obtain number counts of S-Z halos over  the corresponding mass scales and  constrain the magnetic field evolution.  We study the CMB angular power spectrum arising from the S-Z effect and Faraday rotation in galaxy clusters, galaxy groups and galactic halos. For the S-Z effect, the main contribution comes from massive halos ($M>10^{15} M_\\odot$). These set  the peak at around $\\ell \\sim 2000$. Halos with mass $M<10^{14}M_\\odot$ dominate the contribution on small scales. The S-Z effect in halos with mass $10^{14}M_\\odot <M <10^{15}M_\\odot$ and with mass $M<10^{14}M_\\odot$ produces CMB anisotropies over $\\ell \\sim 10000$ and $\\ell >50000$, respectively. Gas condensation is effective on small mass scales so that it suppresses the CMB anisotropies at $\\ell \\sim 10000$.  For the $B$-mode polarisation, small mass halos play more important roles than for the S-Z CMB anisotropies. As a result, the $B$-modes produced in such small halos arise on larger $\\ell$ scales than the S-Z CMB anisotropies. The peak of the $B$-mode power spectrum is produced by halos with $10^{14} M_\\odot< M < 10^{15} M_\\odot$ and  is located at   $1000 < \\ell <5000$. Halos with $M >10^{15} M_\\odot$ are dominant around $\\ell \\sim 1000$ and $B$-modes produced by halos with mass $M<10^{14}M_\\odot$ arise at  $\\ell >20000$, respectively. The evolution of magnetic fields modifies the amplitude produced by halos with $M <10^{14} M_\\odot$. The small mass halos at  high redshift  can collectively generate detectable  $B$-mode polarisation. However, since the halos at high redshift are easily affected  by the evolution of magnetic fields, the $B$-modes by such halos are rapidly damped in the rapid evolution case.  Comparing to other $B$-mode polarisation contributions, the $B$-modes from Faraday rotation dominate the secondary $B$-modes caused by gravitational lensing at $\\ell < 3000$. On these scales, the $B$-modes by Faraday rotation are produced by halos with mass $M<10^{14}M_\\odot$. Therefore, measurement of such $B$-mode polarisation could  put constraints on magnetic field evolution, in combination with the S-Z power spectrum."
    },
    "0807/0807.3070_arXiv.txt": {
        "abstract": "{} {Massive stars in low-metallicity environments may produce exotic explosions such as long-duration gamma-ray bursts and pair-instability supernovae when they die as core-collapse supernovae (CCSNe). Such events are predicted to be relatively common in the early Universe during the first episodes of star-formation. To understand these distant explosions it is vital to study nearby CCSNe arising in low-metallicity environments to determine if the explosions have different characteristics to those studied locally in high-metallicity galaxies. Many of the nearby supernova searches concentrate their efforts on high star-formation rate galaxies, hence biasing the discoveries to metal rich regimes. Here we determine the feasibility of searching for these CCSNe in metal-poor dwarf galaxies using various survey strategies.} {We determine oxygen abundances and star-formation rates for all spectroscopically typed star-forming galaxies in the Sloan Digital Sky Survey, Data Release 5, within $z=0.04$. We then estimate the CCSN rates for sub-samples of galaxies with differing upper-metallicity limits. Using Monte-Carlo simulations we then predict the fraction of these CCSNe that we can expect to detect using different survey strategies. We test survey capabilities using a single 2m telescope, a network of 2m telescopes, and the upcoming all-sky surveys of the Pan-STARRS and LSST systems.} {Using a single 2m telescope (with a standard CCD camera) search we predict a detection rate of $\\sim$1.3 CCSNe yr$^{-1}$ in galaxies with metallicities below $12+\\log({\\rm O/H})<8.2$ which are within a volume that will allow detailed follow-up with 4m and 8m telescopes ($z=0.04$). With a network of seven 2m telescopes we estimate $\\sim$9.3 CCSNe yr$^{-1}$ could be found, although this would require more than 1\\,000\\,hrs of telescope time allocated to the network. Within the same radial distance, a volume-limited search in the future Pan-STARRS PS1 all-sky survey should uncover 12.5 CCSNe yr$^{-1}$  in low-metallicity galaxies. Over a period of a few years this would allow a detailed comparison of their properties. We then extend our calculations to determine the total numbers of CCSNe that can potentially be found in magnitude-limited surveys with PS1 (24\\,000 yr$^{-1}$, within $z\\lesssim0.6$), PS4 (69\\,000 yr$^{-1}$, within $z\\lesssim0.8$ ) and LSST (160\\,000 yr$^{-1}$, within $z\\lesssim0.9$) surveys.} {} ",
        "introduction": "Upon exhausting all of its available fuel, the iron core of a massive star ($\\gtrsim$8$M_\\odot$) collapses, often giving rise to an extremely powerful, bright explosion known as a core-collapse supernova (CCSN). Under very specific circumstances these CCSNe are also thought to be associated with another very powerful and luminous event known as a long-duration gamma-ray burst (LGRB). The circumstances that permit this extraordinary partnership are still not well understood, but it seems likely that progenitors born in lower-metallicity environments (sub-solar) favour production of GRBs \\citep[see][]{2006AcA....56..333S,2006Natur.441..463F,2003A&A...406L..63F,2004MNRAS.352.1073T}.  The most popular theoretical model interprets the LGRB as being caused by the core-collapse of a specific class of massive star. This fundamental theory, coined as the `Collapsar' model, was first conceived by \\citet{1993ApJ...405..273W} \\citep[also see][]{1999A&AS..138..499W}. The criteria for the creation of a LGRB resulting from the collapse of a massive star are that a black-hole must be formed upon core-collapse, either promptly or via fall-back of ejected material, and that an accretion disk is allowed to form around this black-hole which acts as a `central engine' causing the release of collimated relativistic jets along the rotation axis of the black-hole \\citep{2006astro.ph.10276H,2008arXiv0801.4362Y,2008arXiv0803.3807T}. It is these relativistic jets that form the LGRB and the associated afterglow emission. These criteria are found to place very specific constraints on the type of progenitor star that can collapse to produce a LGRB. The core of the progenitor must be massive enough to produce a black-hole, it must be rotating rapidly enough to form an accretion disk, and the star needs to be stripped of its hydrogen envelope, so that the resulting relativistic jets are not inhibited from reaching the stellar surface. All of this points toward rapidly rotating Wolf-Rayet stars as viable progenitors. \\citep{2008arXiv0804.0014D,2006ApJ...637..914W,2006ARA&A..44..507W}  Rapidly rotating O-stars which mix-up their core processed material to produce WR stars through chemically homogeneous evolution have been suggested to satisfy the requirements of the collapsar model \\citep{1987A&A...178..159M,2006ASPC..353...63Y,2006A&A...460..199Y,2006ApJ...637..914W}. This model prevents the depletion of angular momentum due to mass-loss and allows a massive star to become a WR star directly, avoiding the supergiant phase. \\citet{2006A&A...460..199Y} have estimated that for a rapidly rotating star to be the progenitor of an LGRB, prior to collapse it must have a helium core mass greater than 10M$_{\\odot}$, that is an initial mass greater than 25-30M$_{\\odot}$. \\citet{2006ApJ...637..914W} have modelled the WR mass-loss rates of \\citet{2005A&A...442..587V} and predict that the upper metallicity limit for forming LGRB may be as high as 0.3Z$_{\\odot}$. Apart from the single star progenitor model, a rapidly rotating massive star could lose its hydrogen envelope and still retain a large fraction of its initial angular momentum if a close binary companion were to strip it of the envelope rapidly enough for the progenitor to avoid spin down \\citep{2007A&A...465L..29C,2007astro.ph..2652B,2007A&A...471L..29D}.  Whatever the model for the progenitor of the LGRB, prior to collapse we require a very rapidly rotating, massive WR star. These constraints lead to the prediction that LGRBs favour a lower-metallicity environment and that CCSNe found in association with an LGRB will be classified as Type Ib or Type Ic, characterised by an absence of hydrogen in their spectra \\citep{2007ApJ...666.1024B,2003fthp.conf..171F,2003astro.ph..1006H}. To date many LGRBs have revealed `bumps' in their afterglow emission, thought to be evidence of associated SNe \\citep[see][]{1999Natur.401..453B, 2000ApJ...536..185G, 2003A&A...406L..33D,  2005ApJ...624..880L}. A few of these associated SNe have been adequately spectroscopically typed, all have been found to be broad-lined Type Ic SNe, helping to confirm the theory that LGRBs are the result of the collapse of hydrogen-stripped massive stars \\citep{1998Natur.395..670G,2003ApJ...591L..17S,2003Natur.423..847H,2004ApJ...609..952Z}. \\citet{2008AJ....135.1136M} have found that the SNe Type Ic-BL associated with an LGRB are found to inhabit a sample of host galaxies of significantly lower metallicity than the host galaxies of their Type Ic-BL counterparts with no associated observed LGRB. Interestingly, there has not been a SN Type Ic-BL without an associated observed LGRB that has been found to inhabit the same low-metallicity galaxy sample that the LGRB associated SNe inhabit.  \\citet{2006Natur.441..463F} have found that LGRBs are preferentially found in irregular galaxies which are significantly fainter, smaller and hence presumably of lower-metallicity than typical CCSNe hosts. \\citet{2005NewA...11..103S} make a similar point, stating that the majority of LGRBs are found in extremely blue, sub-luminous galaxies, similar to the Blue Compact Dwarf Galaxy population. \\citet{2006astro.ph..9208S} has speculated that at primordial metallicities stars of up to 300M$_{\\odot}$ may form and if they have low mass-loss, they may end their lives still retaining much of their original mass. \\citet{2002ApJ...567..532H} predict that those stars within the mass range 140M$_{\\odot}$-260M$_{\\odot}$ may produce  pair-instability supernovae (PISNe), huge thermonuclear explosions with energies as high as $\\sim$10$^{53}$ ergs. Hence understanding SNe in low-metallicity environments locally could offer insights into high-z GRBs and early Universe explosions. An indication that these exotic PISN events may not be beyond our present day reach is the recent SN 2006gy which has been reported to have a peak luminosity three times greater than any other supernova ever recorded reaching a magnitude of -22 \\citep{2006astro.ph.12617S}. Hypothesised as having had a Luminous Blue Variable as a progenitor with an initial mass of $\\sim$150M$_{\\odot}$ this has been suggested to be the first ever PISNe detected \\citep[also see][]{2007arXiv0708.1970L,2007ApJ...659L..13O,2007Natur.450..390W}. The unusual SN 2006jc defines another class of peculiar low-metallicity event, occurring spatially coincident with an outburst thought to be that of an LBV just two years previously \\citep{2007Natur.447..829P}, which has also been postulated to be related to the PISN mechanism \\citep{2007Natur.450..390W}  To date there have been many surveys specifically designed to search for SNe and depending on the motivation of an individual survey generally one of two different survey strategies has been employed. The first of these strategies is the \\textit{pointed survey strategy}; here we have a fixed catalogue of galaxies which are all individually observed with a cadence of a few days and scanned for fresh SNe. These surveys included the highly successful Lick Observatory Supernova Search (LOSS) which uses the fully robotic 0.75m KAIT telescope to scan a catalogue of 7\\,500-14\\,000 nearby galaxies ($z \\lesssim 0.04$), aiming to observe each individual galaxy with a cadence of 3-5 days \\citep{2001ASPC..246..121F}. With the advent of relatively inexpensive CCD technology, the ability of amateur astronomers to detect nearby SNe using this pointed survey strategy has become more than substantial and a large fraction of nearby SNe are now being discovered by this community \\citep{2006JAVSO..35...42G}. The Southern inTermediate Redshift ESO Supernova Search (STRESS) produced a catalogue of $\\sim$43,000 ($0.05 < z < 0.6$) galaxies within 21 fields observed with the ESO Wide-Field-Imager and searched for SNe in these hosts to determine rates and their dependency on host galaxy colour and redshift evolution \\citep{2008A&A...479...49B}.  The second strategy employed by SN surveys is the \\textit{area survey strategy}; here an area of sky is surveyed repeatedly and image subtraction used to identify transient events including SNe. The SN Legacy Survey uses the CFHT MegaCam imager to image four one-square degree fields to search for SNe Type Ia with the motivation of improving the sampling of these SNe within the redshift range $0.2 < z < 1.0$ \\citep{2005ASPC..339...60P, 2006A&A...447...31A}. The Equation of State SupErNova Cosmology Experiment (ESSENCE) uses the Mosaic Imager on the Blanco 4m telescope to survey equatorial fields, that have preferably been observed with other wide-field surveys, to discover SN Type Ia within the redshift range $0.15 < z < 0.74$ \\citep{2007ApJ...666..674M}.  The Nearby SN Factory uses ultra-wide field CCD mosaic images from the Near-Earth Asteriod Tracking (NEAT) and Palomar-Quest survey programs to perform an area survey in a bid to find SN Type Ia in the redshift range $0.03 < z < 0.08$ \\citep{2002SPIE.4836...61A}. The SDSS-II Supernova Survey takes repeated images of Stripe 82, a 300 square degree southern equatorial strip of the sky, driven by the ambition of identifying and measuring the lightcurves of SNe Type Ia in the intermediate redshift range $0.05 < z < 0.35$ \\citep{2008AJ....135..338F, 2008AJ....135..348S}. The Texas Supernova Search search \\citep{2005AAS...20717102Q, 2007AAS...21110505Y} is somewhat unique in that it has a small aperture telescope (0.45m) but very wide field (1.85 square degrees) and is focussed on finding nearby SNe in wide area searches. This survey discovered SN2006gy along with many other of the brightest known SNe ever found and we shall discuss this further in Sec. \\ref{sect:disc}.   The main driver of SN surveys that employ the pointed survey strategy is basically to find as many SNe as possible regardless of type or characteristic.  To ensure that these surveys are as efficient as possible at finding CCSNe, the galaxies catalogues used by the surveys generally consist only of the most massive galaxies with the greatest star-formation rates. As a result these galaxy catalogues tend to be heavily biased towards galaxies with high metallicity. This over arching high-metallicity bias has more than likely placed us in the situation where the vast majority of CCSNe that have occurred in low-metallicity environments (such as low-luminosity, dwarf, irregular galaxies) have remained undetected.  \\citet{2008ApJ...673..999P} have matched the SAI supernova catalogue to the SDSS-DR4 catalogue of star-forming galaxies with measured metallicities and it is clear that the vast majority of the SNe considered have been detected in areas of relatively high-metallicity; SN Type II occurring in host galaxies with mean metallicity \\mbox{$12+{\\rm \\log(O/H)}=8.94\\pm0.04$} and SN Ib/c at 9.06$\\pm$0.04.  It is inherently interesting then to search for SNe specifically in low-metallicty environments, or at least without biasing the search to look at only high metallicity regimes. In this paper we shall discuss several methods to search for low-metallicity CCSN events. First we consider compiling a catalogue of nearby, low-metallicity galaxies from pre-existing catalogues (taking SDSS DR5 as our primary survey source) and using either a single 2.0m telescope or a network of 2.0m telescopes to perform a pointed survey of low-metallicity galaxies in the hope of detecting a few CCSNe. Secondly we consider using a future all-sky transient survey such as the Panoramic Survey Telescope and Rapid Response System (Pan-STARRS) to perform a volume-limited survey for SNe and estimate how many may be found in low-metallicity galaxies. A third and final method we consider is to use a future all-sky transient survey, limited only by the limiting magnitude of the survey, to search for all CCSNe including those low-metallicity events.While this paper is primarily aimed at determining numbers of low metallicity events that could be found, the latter calculation gives an estimate of the total number of CCSNe that are likely to be harvested from these upcoming surveys. ",
        "conclusions": "Having determined oxygen abundances, star-formation rates and CCSN rates for all spectroscopically typed star-forming galaxies in the Sloan Digital Sky Survey Data Release 5 within $z=0.04$, we have used Monte-Carlo simulations to predict the fraction of these CCSNe that we can expect to detect using different survey strategies. Using a single 2m telescope (with a standard CCD camera) search we predict a detection rate of $\\sim$1.3 CCSNe yr$^{-1}$ in galaxies with metallicities below $12+\\log({\\rm O/H})<8.2$ which are within a volume that will allow detailed follow-up with 4m and 8m telescopes ($z=0.04$). With a network of seven 2m telescopes we estimate $\\sim$9.3 CCSNe yr$^{-1}$ could be found, although this would require more than 1000\\,hrs of telescope time allocated to the network. Within the same radial distance, a volume-limited search with the future Pan-STARRS PS1 all-sky survey should uncover 12.5 CCSNe yr$^{-1}$  in low metallicity galaxies. Over a period of a a few years this would allow a detailed comparison of their properties. We have also extended our calculations to determine the total numbers of CCSNe that can potentially be found in magnitude-limited surveys with PS1 (24\\,000 yr$^{-1}$, within $z \\lesssim 0.6$), PS4 (69\\,000 yr$^{-1}$, within $z \\lesssim 0.8$ ) and LSST (160\\,000 yr$^{-1}$, within $z \\lesssim0.9 $) surveys.  All considered, a final strategy chosen to searching for CCSNe in low-metallicity environments shall realistically involve both a volume-limited and a magnitude-limited all-sky survey in order to include a volume-limited galaxy sample with which to accurately determine relative SN rates and have some prior knowledge of the host galaxy characteristics, but yet not exclude the potential of detecting rare CCSN events that would have otherwise been missed had we only considered the volume-limited survey strategy.  With the huge number of CCSNe predicted to be detected, these all-sky surveys are set to serve as a catalyst concerning our understanding of CCSNe; including their varying characteristics with metallicity, the relative rates of the various types of SNe and of extremely rare events similar to SN 2006jc, SN 2006gy and more than likely events the nature of which have not yet been observed."
    },
    "0807/0807.1769_arXiv.txt": {
        "abstract": "In helioseismology, there is a well-known offset between observed and computed oscillation frequencies.  This offset is known to arise from improper modeling of the near-surface layers of the Sun, and a similar effect must occur for models of other stars.  Such an effect impedes progress in asteroseismology, which involves comparing observed oscillation frequencies with those calculated from theoretical models.  Here, we use data for the Sun to derive an empirical correction for the near-surface offset, which we then apply three other stars (\\acena{}, \\acenb{} and \\bhyi).  The method appears to give good results, in particular providing an accurate estimate of the mean density of each star. ",
        "introduction": "Both helio- and asteroseismology involve comparing observed oscillation frequencies with those calculated from theoretical models.  However, for the Sun there is a long-standing systematic offset between observed and computed frequencies that is known to arise from improper modeling of the near-surface layers \\citep{ChDDL88,DPV88,ChDDA96,ChD+T97}.  This offset is independent of the angular degree of the mode ($l$) and increases with frequency.  A similar offset must occur for models of other stars, and should be taken into account whenever observations of stellar oscillations are compared with theory.  In this Letter we use data for the Sun to derive an empirical correction for these near-surface effects and show how to apply this to other stars.  The method appears to give good results, in particular providing an accurate estimate of the mean density of each star. ",
        "conclusions": "The method outlined here for correcting near-surface effects can be applied to model frequencies before they are compared with observations.  As in the case of the Sun, we expect that the correction is independent of degree at a given frequency, for low-degree acoustic modes, and thus the correction determined from radial modes can be applied to all such modes.  There is, however, an important exception.  In evolved stars, mixed modes may be observed that have the character of gravity modes in the deep interior of the star.  Observational evidence for mixed modes has been found in \\eboo{} \\citep{KBV95,KBB2003,CEB2005} and \\bhyi{} \\citep{BKA2007}. Owing to the larger amplitude of these modes in the stellar interior and hence their higher inertia, their frequencies are less affected by the surface effects.  Techniques need to be developed to take this into account in the application of the surface correction before the frequencies are analyzed.  \\new{The use of a single power law is made plausible by the knowledge that the offsets we are modelling are caused by the properties of the near-surface layers and hence presumably depend only on surface gravity, effective temperature and composition, and not on the details of the internal properties of the star.  This may be particularly true of the exponent which, in the simple analysis by \\citet{ChD+G80}, is determined by an effective average polytropic index (i.e., the relation between pressure and density) in the near-surface layers. }  It is important to note that, owing to its strong frequency dependence, the offset also affects the large frequency separation $\\Delta \\nu$, as is indeed implicit in equation~(\\ref{eq.Deltanu}). \\new{For example, although Model~S is one of the best available models of the Sun, it has a large separation that is 1\\,\\muHz{} greater than the observed value.}  Thus, attempting to fit to stellar models based on $\\Delta \\nu$ will introduce systematic errors, unless the corresponding correction is applied.  We have shown how to identify the model for each star that is closest to the ``best'' model, by requiring that $r$ be as close as possible to unity. This gives us an extremely accurate estimate of the mean stellar density. However, it is important to point out that, while a model with $r$ close to unity gives a good match to the overall structure of the star, it does not necessarily reproduce the structure of the core or give a reliable estimate of the stellar age.  Determination of those properties requires taking into account the frequencies of the non-radial modes (including the small frequency separations).  That is the next stage of model fitting, which can be done after the surface correction has been applied to all modes."
    },
    "0807/0807.1902_arXiv.txt": {
        "abstract": "The grand-design spiral galaxy M51 has long been a crucial target for theories of spiral structure.  Studies of this iconic spiral can address the question of whether strong spiral structure is transient (e.g. interaction-driven) or long-lasting.  As a clue to the origin of the structure in M51, we investigate evidence for radial variation in the spiral pattern speed using the radial Tremaine-Weinberg (TWR) method.  We implement the method on CO observations tracing the ISM-dominant molecular component.  Results from the method's numerical implementation--combined with regularization, which smooths intrinsically noisy solutions--indicate two distinct patterns speeds inside 4 kpc at our derived major axis PA=170$\\degree$, both ending at corotation and both significantly higher than the conventionally adopted global value.  Inspection of the rotation curve suggests that the pattern speed interior to 2 kpc lacks an ILR, consistent with the leading structure seen in HST near-IR observations.  We also find tentative evidence for a lower pattern speed between 4 and 5.3 kpc measured by extending the regularized zone.  As with the original TW method, uncertainty in major axis position angle (PA) is the largest source of error in the calculation; in this study, where $\\delta_{PA}$=$\\pm$5$\\degree$, a $\\sim$20\\% error is introduced to the parameters of the speeds at PA=170$\\degree$.  Accessory to this standard uncertainty, solutions with PA=175$\\degree$ (also admitted by the data) exhibit only one pattern speed inside 4 kpc, and we consider this circumstance under the semblance of a radially varying PA. ",
        "introduction": "\\indent  The large angular size and clear spiral structure of the nearly face-on spiral M51 make it ideal for studies of the nature and origin of grand design spiral structure. Two scenarios dominate the discussion in the literature, each based on opposing theories: strong spiral structure as a quasi-stationary density wave (e.g. \\citealt{linshu}), or as a transient feature due to interaction with nearby companion NGC 5195 (e.g. \\citealt{tully}).  \\\\ \\indent Observations of both the stellar and gaseous components reveal consistencies with the density wave interpretation at some level (see \\citealt{elmegreen1}; \\citealt{vog93}; \\citealt{rand93}).  In accord with the seminal study of \\citet{tully} which attributes the (transient) outer pattern to the interaction with its companion and proposes that the inner arms are likely spiral density waves also driven by the encounter, \\citet{elmegreen1} and \\citet{vog93} find independent evidence for two different pattern speeds.  If the strong spiral's corotation radius overlaps with the ILR of the outer, material pattern, as is suspected (e.g \\citealt{tully}; \\citealt{elmegreen1}), this may indicate the stimulation of an inner spiral wavemode by the outer material spiral via mode-coupling \\citep{tagger}.  But observationally it remains unclear whether, if driven by the outer pattern, the strong spiral structure is transient, or has survived a few rotation times (e.g. as speculated by \\citealt{vog93}).  \\\\ \\indent In the simulations of both \\citet{salo} and \\citet{howardByrd} structure throughout the disk is well reproduced by multiple passages of the (bound) orbiting companion, and the nuclear structure, in particular, seems intimately related to the inward propagation of multiple tidally-induced perturbations \\citep{salob}.  Pursuant to the study of \\citet{toomresq}, such simulations of M51 have proved indispensable for exploring and motivating scenarios in favor of short-lived waves.  In the short-lived wave paradigm, propagating wave packets evolved from kinematic distortions in the outer disk may be swing-amplified \\citep{toomre}, causing a strong response in the inner disk.  As predicted by \\citet{salob}, the wave speed has a complex radial dependence, featuring a constant pattern speed for the dominant $m$=2 structure out to 1.2-1.8 kpc (depending on the disk mass assumed), followed by a superposition of structures described by a pattern speed that decreases with radius, down to $\\sim$10-20 $\\kmskpc$ by $\\sim$4.6 kpc.  \\\\ \\indent Knowledge of the pattern speed of the structure can, in principle, both distinguish between and reconcile the short- and long-lived wave scenarios, and so many studies have focused on measuring and characterizing this parameter (see \\citealt{elmegreen1}, \\citealt{tully}, \\citealt{salob} and \\citealt{garcia}, for example).  The pattern speed of the outer spiral has long been proposed near 10-20 $\\kmskpc$.  In the inner disk, application of the traditional, model-independent method of Tremaine \\& Weinberg (1984; hereafter TW) using CO observations yields a pattern speed $\\Omega_p$=38 $\\kmskpc$ \\citep{zrm04}, in general agreement with the determinations based on resonance locations of \\citet{elmegreen1} and \\citet{tully} (but higher than the pattern speed $\\Omega_p$=27 $\\kmskpc$ found by \\citealt{garcia}).  Although the TW analysis shows evidence for significant departures from the expected relation for a constant pattern speed in both the inner- and outer-most regions of the disk, the method cannot quantitatively account for any suspected radial variation of the pattern speed.\\\\ \\indent The radial TW (hereafter TWR) method (\\citealt{mrm}; \\citealt{meidt}) should prove an invaluable resource in this regard, since with it we can characterize the angular speeds of distinct patterns and their possible radial variation.  For the first time, we are able to observationally address issues related to the complex nature and persistence of spiral patterns and the connection, if any, between multiple pattern speeds in a single disk.  \\\\ \\indent Like its traditional counterpart, the TWR method, summarized in $\\S$ 2, relies on the use of a kinematic tracer found to obey continuity. Here, we consider the ISM-dominant molecular component in the inner disk of M51 as traced by CO observations.  In $\\S\\S$ \\ref{sec:obs} and \\ref{sec:molecdomin} we describe these observations and review the arguments which establish their conformity with the assumptions of the method.  In $\\S$ \\ref{sec:quad} we formulate the TWR quadrature and motivate the models developed for testing.\\\\ \\indent Results of the regularized TWR calculation applied to the inner disk are presented in $\\S$ \\ref{sec:inneronly}.  There we establish a best estimate for the pattern speed(s) of the bright spiral structure by considering the characteristics of solutions over a $\\pm$5$\\degree$ range of disk position angles (PAs) ($\\S$ \\ref{sec:PAchar}); according to the findings of \\citet{debPA}, we can expect uncertainty in the PA to be the dominant source of systematic error in the calculation.  We also compare this estimate to other tested models of the radial dependence in $\\S$ \\ref{sec:state}.  \\\\ \\indent In an effort to authenticate the TWR estimate, in $\\S$ \\ref{sec:corrollary} we relate our measurements to other independent evidence for more than a single pattern speed in the inner disk of M51 and investigate the resonance locations and overlaps that they entail.  We also consider our measurements in light of relevant findings throughout the literature, including those of \\citet{shetty} ($\\S$ \\ref{sec:varyPA}) and \\citet{henry} ($\\S$ \\ref{sec:discussion}).  Final results are summarized in a conclusion section. ",
        "conclusions": "Conclusion} In this paper we present regularized TWR solutions for the pattern speed of the bright spiral structure in the inner disk of M51, derived with velocity and intensity information from the ISM-dominant molecular component traced by high-resolution CO observations.  These solutions are arrived at by isolating the inner disk from errors which evidently originate with both the quality of sampling/detection and the pattern speed-modeling in bins covering the outer, material pattern.  So although our procedure prevents us from constraining the outermost pattern speed, calculating the outer bins without regularization in principle improves the accuracy with which the solution for the inner disk can realize the true pattern speed.  \\\\ \\indent Our primary result with this implementation is the measurement inside 4 kpc of two pattern speeds, both significantly higher, and together fitting the data better, than the constant global measure of \\citet{zrm04} at the nominal PA=170$\\degree$.  A third, lower pattern speed, extending beyond 4 kpc out to (at least) 5.3 kpc is also detected, nearer the speed expected for the material pattern.  Significantly, the transitions between the measured pattern speeds coincide with resonances; the two pattern speeds inside 4 kpc both end at corotation within the uncertainties.  Since it is in no way imposed by the method, this dynamically reasonable scenario tends to give us confidence as to the physical plausibility of the pattern speeds returned by the analysis.  \\\\ \\indent Of course, given that a pattern speed interior to $r$$\\simeq$2 kpc is only weakly detected (if at all) at PA=175$\\degree$, the accuracy of the description provided by two pattern speeds may depend on whether PA=170$\\degree$ or PA=175$\\degree$ is more accurate, an uncertainty raised by \\citet{shetty}, for instance.  If the disk is best described with PA=175$\\degree$, we find evidence that a single constant pattern speed inside 4 kpc best characterizes the bright spiral structure.  Furthermore, as contemplated in $\\S$ \\ref{sec:varyPA}, a radially varying PA which decreases from 180$\\degree$ (and reaches 170$\\degree$ near 3 kpc), perhaps suggested by the results of \\citet{shetty}, also favors a single measured pattern speed interior to 4 kpc.  \\\\ \\indent Again, however, while the analysis presented here cannot resolve the question as to which PA is more appropriate, we find meaningful, independent evidence in favor of the pattern speeds measured at 170$\\degree$, in particular.  For example, consistent with expectations of leading structure at the inner most radii (as in the observations of \\citealt{scoville}), the higher speed inside $r$$\\simeq$2 kpc lacks an ILR.  In addition, attendant to our finding that both speeds interior to 4 kpc terminate at corotation, the transition between the two roughly coincides with an inferred location of resonance overlap wherein the inner's corotation resonance and the outer's inner 4:1 resonance align.  The radial domain of the pattern speed measured at PA=175$\\degree$, in contrast, is not as clearly associated with resonance radii.  Since the bright spiral structure does not appear along the minor axis near $\\sim$2 kpc, the corotation resonance at this location implied by the solution at 170$\\degree$ is unfortunately unconfirmable through inspection of radial streaming velocities under the density wave interpretation.  \\\\ \\indent We also find remarkable agreement between the characteristics of the two speeds inside 4 kpc at PA=170$\\degree$ and other evidence in the inner disk consistent with multiple pattern speeds.  The transition parameterized in our best-fit solution for the inner disk coincides with significant variation in the two-armed spiral pitch angle.  Since a change in the pitch angle is expected to be accompanied by a change in streaming motions, both parameters are presumably attendant to the signatures (streaming or otherwise) of the patterns.  \\\\ \\indent Although the pattern speed interior to 2 kpc in the solution at 170$\\degree$ (or 175$\\degree$) may reflect a unique contribution from the $m$=3 mode observed by \\citet{rix}, as described in $\\S$ \\ref{sec:discussion}, the measurements in Figure \\ref{fig-globalcut9a} (or Figure \\ref{fig-twfit175bestc14}) presumably directly relate to the patterns present in the disk and so (depending on the PA) should provide a fair description of the dynamics therein.  As such, it may be possible that ensuing observations and studies better discriminate between the two seemingly disparate radial dependencies implied for the PAs considered here.  \\\\ \\indent Even at present our TWR solutions yield interpretations with which to observationally address the question of spiral longevity.  That the solutions at both 170$\\degree$ and 175$\\degree$ feature constant pattern speeds would imply that our solutions are indicative of long-lasting spiral structure.  Interestingly, at the innermost radii both qualitatively resemble the model of \\citet{salob} where, characteristic of the isolated evolution of the disk, the dominant $m$=2 component has a constant pattern speed $\\sim$50 $\\kmskpc$ out to $\\sim$1.2-1.8 kpc.  As for the region between $\\sim$1.8 kpc and 4.6 kpc in those models where interaction with the companion introduces a succession of transient structures, our solutions at both 170$\\degree$ and 175$\\degree$ otherwise describe at most two steady patterns in distinct radial zones.  \\\\ \\indent In the immediate future, observations with higher resolution and sensitivity should afford TWR calculations with finer radial bins, thereby allowing for the parameterization of more distinct radial zones, if present.  This will either confirm our solutions for $\\Omega_p(r)$ or perhaps demonstrate that solutions describe a succession of many discrete patterns (similar to the transient structures in the \\citet{salob} models), or simply a winding, material pattern.   Again, though, our multiple-speed and other, quadratic solutions in general more closely follow $\\Omega$ throughout the disk than $\\Omega-\\kappa/2$ characteristic of $m$=2 structure in the models of \\citet{salob}.  \\\\ \\indent Despite the lingering ambiguity in the PA, these TWR solutions present a new picture of the bright spiral structure of M51, one that should prompt tests of long-lived density wave theories in other nearby grand-design spirals.  At the very least, this study marks a successful starting point for continued tests of the relation between multiple spiral pattern speeds in a single disk; investigations into the number and radial domains of pattern speeds and spiral winding in nearby spirals will be the subject of upcoming work.  \\\\    S. E. M. acknowledges support from a NASA-funded New Mexico Space Grant Consortium Graduate Research Fellowship.  We thank Heikki Salo for providing insight into some of the intricacies of the simulations in \\citet{salob}."
    },
    "0807/0807.3314_arXiv.txt": {
        "abstract": "We present the discovery of a triplet of emission-line nuclei in the disturbed disk galaxy NGC 3341, based on archival data from the Sloan Digital Sky Survey and new observations from the Keck Observatory. This galaxy contains two offset nuclei within or projected against its disk, at projected distances of 5.1 and 8.4 kpc from its primary nucleus and at radial velocity separation of less than 200 \\kms\\ from the primary. These appear to be either dwarf ellipticals or the bulges of low-mass spirals whose disks have already been stripped off while merging into the primary galaxy.  The inner offset nucleus has a Seyfert 2 spectrum and a stellar velocity dispersion of $70\\pm7$ \\kms. The outer offset nucleus has very weak emission lines consistent with a LINER classification, and the primary nucleus has an emission-line spectrum close to the boundary between LINER/\\ion{H}{2} composite systems and \\ion{H}{2} nuclei; both may contain accreting massive black holes, but the optical classifications alone are ambiguous.  The detection of an offset active nucleus in NGC 3341 provides a strong suggestion that black hole accretion episodes during minor mergers can be triggered in the nuclei of dwarf secondary galaxies as well as in the primary. ",
        "introduction": "In the standard hierarchical merger paradigm for galaxy formation and evolution, a natural consequence of the merger of two galaxies containing massive black holes is that the two black holes will sink to the center of the host galaxy by dynamical friction, eventually forming a binary and merging together (see Merritt \\& Milosavljevi\\'c 2005 for a review).  As a result, there is a general expectation that some merging and post-merger galaxies should have two massive black holes.  However, definitive observational identification of galaxies containing two (or more) massive black holes is limited to objects showing evidence for accretion-powered activity from spatially resolved active galactic nuclei (AGNs) \\citep{kom03a}.  A galaxy with two widely separated active nuclei is often referred to as a ``dual AGN'' to distinguish it from the case of a ``binary AGN'' in which the two black holes form a tight, gravitationally bound binary.  Several examples of dual AGNs are now known \\citep[e.g.,][]{owe85,kom03b,bal04,gua05,ger07,bia08}, most of which are found in major mergers of massive galaxies.  It is also expected that minor mergers can trigger episodes of nuclear activity \\citep{hm95}, and dual AGNs could also be present in some minor mergers with dwarf secondary galaxies if the secondary galaxies also harbor central black holes.  However, the fraction of dwarf or low-mass galaxies that contain black holes is unknown, and only weakly constrained by AGN surveys \\citep{gh07} and by stellar-dynamical searches for black holes in dwarf galaxies \\citep{val05}.  In this \\emph{Letter}, we present the discovery of a triplet of emission-line nuclei (at several kpc separation) in NGC 3341.  This system appears to contain two dwarf galaxies merging with a massive disk galaxy, and one of the offset nuclei has a Seyfert 2 spectrum. NGC 3341 is therefore a rare example of a disk galaxy with an offset AGN, and an excellent nearby case study for investigations of AGN triggering during a minor merger.  The primary nucleus and the second offset nucleus have optical spectra that indicate possible low-level AGN activity as well, making this a candidate dual or even triple AGN system. ",
        "conclusions": "The starlight subtraction procedure yields radial velocities of $cz = 8325\\pm9$, $8131\\pm6$, and $8145\\pm6$ \\kms\\ for the primary nucleus and objects B and C, respectively.  This confirms the physical association between the three objects and rules out a chance projection of unrelated sources.  Emission-line fluxes were measured by direct integration of the starlight-subtracted spectra; the key diagnostic flux ratios are listed in Table \\ref{table1}.   Figure \\ref{bpt} displays the emission-line flux ratios of [\\ion{O}{3}]/\\hbeta\\ and [\\ion{N}{2}]/\\hal\\ in a \\citet{bpt81} diagnostic diagram, in comparison with galaxies from the \\citet{kau03} sample.  The Keck spectrum of object B confirms its classification as a Seyfert 2, as it falls within the Seyfert branch of the diagnostic diagram.  Its [\\ion{O}{3}] $\\lambda5007$ emission line has FWHM = 116 \\kms, similar to the low [\\ion{O}{3}] linewidths seen in other Seyfert 2 nuclei having low-mass host galaxies \\citep{bgh08}.  In the [\\ion{O}{3}]/\\hbeta\\ vs.\\ [\\ion{N}{2}]/\\hal\\ diagram, the primary nucleus falls in the region of LINER/\\ion{H}{2} ``composite'' or ``transition'' objects, close to the boundary with \\ion{H}{2} nuclei, following the \\citet{kew06} classification scheme.  However, we note that the [\\ion{O}{1}]/\\hal\\ and [\\ion{S}{2}]/\\hal\\ ratios at the nucleus are consistent with classification as an \\ion{H}{2} nucleus \\citep{hfs97, kew06}.  Thus, the classification of the nucleus is somewhat ambiguous.  The narrow-line profiles at the nucleus are double-peaked, with a peak-to-peak velocity separation of 310 km s\\per.  This velocity splitting could arise from either radial motion or rotation; similar double-peaked profiles due to the rotation of a circumnuclear disk are seen, for example, in the transition-type nucleus of NGC 3245 \\citep{hfs95}.  For the weak emission lines in object C, we find [\\ion{N}{2}]/\\hal\\ $=0.93 \\pm 0.27$, indicated as a shaded band in Figure \\ref{bpt}.  The nearly-equal strengths of [\\ion{N}{2}] and \\hal\\ are inconsistent with an \\ion{H}{2} region origin for the emission lines.  Given the non-detection of [\\ion{O}{3}] $\\lambda5007$ emission, the most likely classification for this object is a LINER or LINER/\\ion{H}{2} composite.  The \\hal\\ equivalent width is only $\\sim1$ \\AA.     While the optical data identify object B as an AGN with a high degree of confidence, the source of ionizing photons in the nucleus and object C is less certain.  In the absence of definitive AGN signatures such as broad-line emission, very weak LINER-type emission lines such as those seen in object C could in principle result from photoionization of diffuse gas by sources unrelated to an AGN, such as post-AGB stars \\citep{bin94} or X-ray binaries.  Indeed, \\chandra\\ observations have revealed that some nearby ``Type 2'' LINERs do not have a central AGN-like point source at all, but do have X-ray emission dominated by a circumnuclear population of X-ray binaries \\citep{era02, flo06}.  As for the primary nucleus, the source of ionization in LINER/\\ion{H}{2} transition nuclei has been the subject of much debate \\citep[see][]{ho08}.  A likely interpretation is that they are composite systems containing both a weak AGN and \\ion{H}{2} regions ionized by hot stars \\citep[e.g.,][]{hfs93, vgv97, kew06}. Shock heating could play some role in the narrow-line region excitation as well \\citep[e.g.,][]{ds95, cv01}.  Summarizing results from X-ray surveys of nearby galaxies, \\citet{ho08} concludes that the majority of transition nuclei (perhaps $\\sim75\\%$) do contain AGNs. Thus, the optical spectra of the nucleus and object C can be interpreted as possibly containing low-luminosity accretion-powered AGNs, but the optical data do not provide conclusive evidence for the presence of AGNs in these objects.  A search for X-ray sources in NGC 3341 would potentially provide the best constraints on the AGN content in this galaxy.    \\begin{deluxetable}{lcccc} \\tablewidth{3in} \\tablecaption{Emission-Line Ratios} \\tablehead{ \\colhead{Object} & \\colhead{[\\ion{O}{3}]/\\hbeta} & \\colhead{[\\ion{O}{1}]/\\hal} & \\colhead{[\\ion{N}{2}]/\\hal} & \\colhead{[\\ion{S}{2}]/\\hal}  } \\startdata Nucleus & $0.38\\pm0.04$ & $0.05\\pm0.01$ & $0.54\\pm0.02$ & $0.31\\pm0.01$ \\\\ B       & $6.6\\pm0.2$   & $0.20\\pm0.01$ & $0.85\\pm0.02$ & $0.66\\pm0.02$ \\\\ C       & \\nodata       & $<0.25$       & $0.93\\pm0.27$ & $<0.6$ \\enddata \\tablecomments{Forbidden-line measurements refer to [\\ion{O}{3}] $\\lambda5007$, [\\ion{O}{1}] $\\lambda6300$, [\\ion{N}{2}] $\\lambda6583$, and [\\ion{S}{2}] $\\lambda\\lambda6716,6731$. }  \\label{table1} \\end{deluxetable}   NGC 3341 is not detected in the \\rosat\\ All-Sky Survey, and no source corresponding to its position is listed in the \\rosat\\ Faint Source Catalog \\citep{vog00}.  We estimate an upper limit of 0.02 counts s\\per\\ in the \\rosat\\ image, corresponding to an upper limit of $L$(0.1--2 keV)$<5\\times10^{41}$ ergs s\\per\\ for an unabsorbed source with a photon index of $\\Gamma = 1.8$.  Using typical ratios of X-ray to optical emission-line flux for nearby low-luminosity AGNs \\citep{ho08}, the expected luminosity of each of the three nuclei would be at least an order of magnitude below the \\emph{ROSAT} upper limit.  Thus, much deeper observations would be needed in order to search for X-ray sources in this system.  The primary nucleus is a radio source, detected in the FIRST survey (Becker \\etal\\ 1995) with a flux density of 4.8 mJy and deconvolved major and minor axis sizes of $5\\farcs5\\times5\\farcs0$.     \\begin{figure} \\plotone{f3.eps} \\caption{Diagnostic diagram showing the emission-line ratios of the three nuclei.  Black contours denote emission-line galaxies from the \\citet{kau03} sample.  The dashed and dotted curves represent the classification criteria of \\citet{kew06}, in which \\ion{H}{2} nuclei fall below and to the left of the dashed curve, AGNs lie above and to the right of the dotted curve, and composite objects lie in the intermediate region.  The primary nucleus of NGC 3341 is shown as a blue circle, and object B as a red square.  For object B and for the [\\ion{N}{2}]/\\hal\\ ratio at the primary nucleus, the error bars are smaller than the plot symbols.  The violet shaded band gives the $1\\sigma$ uncertainty range for the [\\ion{N}{2}]/\\hal\\ ratio of object C. \\label{bpt}} \\end{figure}   We measured the stellar velocity dispersions of the three nuclei following the direct-fitting methods described by \\citet{bhs02}, by fitting the spectra over the restricted wavelength range 5020--5500 \\AA\\ containing the strong Mg$b$ and Fe5270 features.  We find $\\sigmastar=182\\pm14$, $70\\pm7$, and $44\\pm6$ \\kms\\ for the nucleus and objects B and C, respectively.  If these objects follow the \\msigma\\ correlation for quiescent galaxies of \\citet{tre02}, the expected black hole masses would be $9\\times10^7$, $2\\times10^6$, and $3\\times10^5$ \\msun\\ in the three nuclei.  \\citet{hec04} show that a typical [\\ion{O}{3}] bolometric correction for Seyfert galaxies is $\\lbol/L$([\\ion{O}{3}])$\\approx 3500$ (with a scatter of approximately 0.38 dex).  With this bolometric correction, object B would have $\\lbol/\\ledd \\approx 0.1$.  From the 2MASS \\citep{skr06} point source and extended source catalogs, the $K_s$-band magnitudes of the primary galaxy and objects B and C are 10.78, 14.26, and 15.07 mag, respectively.  To estimate the stellar mass ratio of the merger, we use $K$-band mass-to-light ratios from \\citet{bell03}, computed using reddening-corrected optical colors from SDSS.  Following the \\citet{bell03} prescriptions, the primary galaxy and objects B and C each have $M/L_K\\approx0.85$.  This yields a stellar mass of $\\sim1\\times10^{11}$ \\msun\\ for the primary galaxy and a stellar mass ratio of 50:2:1 for the merging system (primary:B:C).  We note, however, that this mass ratio does not account for any stellar mass already stripped off of objects B and C during the merger.  For a more detailed examination of the structure of the merging system, higher-resolution imaging would be needed. \\emph{Hubble Space Telescope} data could clarify whether the offset nuclei were originally spheroidal or elliptical galaxies, or the bulges or pseudobulges of low-mass disk galaxies.  It is particularly interesting to find nuclear activity in at least one of the secondary nuclei in this system, since the AGN fraction in low-mass galaxies is known to be very small overall \\citep{kau03, gh07}.  Furthermore, the black hole occupation fraction in dwarf galaxies is expected to be well below unity if the efficiency for black hole seed formation is small in low-mass halos at high redshift \\citep{vol08}, or if gravity-wave kicks during binary black hole coalescence can efficiently expel black holes from dwarf galaxies \\citep{mer04}.  There are examples of AGNs with host galaxy luminosities similar to objects B and C, such as POX 52 \\citep{bar04}, but they are predominantly isolated systems not involved in strong interactions.  Simulations have shown that minor mergers can drive radial inflows of gas to the center of the primary galaxy, potentially triggering AGN accretion episodes \\citep{hm95}, but there is little if any statistical evidence for enhanced AGN activity in the secondary galaxies of minor interaction systems.  In a study of close (separation $<50h^{-1}$ kpc) galaxy pairs in SDSS, \\citet{wg07} found that in minor interactions with $\\Delta m >2$ mag, there was a strong increase in the AGN fraction for the primary galaxies (relative to a matched field galaxy control sample) but there was no detectable enhancement in the AGN fraction of the secondary galaxies in the pairs relative to a matched control sample.  In view of the properties of NGC 3341, we speculate that any enhancement in the AGN fueling rate in the secondary might occur at very small separations ($\\lesssim10$ kpc) that are difficult to probe in statistical surveys.  This unusual AGN host galaxy was discovered because the SDSS fortuitously selected object B as the one spectroscopic target in this field.  If the primary nucleus had been targeted instead, the galaxy would have appeared unremarkable, and follow-up observations of objects B and C would have been unlikely.  What are the prospects for detection of offset AGNs in other nearby galaxies?  \\citet{kd05} argued that some of the most extreme ultraluminous X-ray (ULX) sources are likely to be the nuclei of captured satellite galaxies undergoing AGN accretion episodes, rather than being X-ray binaries with intermediate-mass black holes.  While no ULX sources are currently known to have an origin in captured satellites, the discovery of NGC 3341B in the SDSS suggests that optical selection can be at least as efficient as X-ray selection for discovery of off-nuclear, accreting black holes from satellite galaxies in minor mergers.  However, optical selection depends crucially on the ability to identify the nucleus of the secondary galaxy in imaging data.  As discussed by \\citet{wgb06}, there are severe observational biases that make it very difficult to identify satellite galaxies involved in minor interactions ($\\Delta m > 2$ mag) at very close separations of $\\lesssim10$ kpc from the primary galaxy.  Furthermore, at separations of a few kpc, satellite galaxies may be so severely disrupted that they no longer have identifiable nuclei, and offset, low-luminosity AGNs not embedded in stellar nuclei would likely be inconspicuous in optical surveys.  Nevertheless, further searches for systems like NGC 3341 containing low-mass companions at small separations could provide new constraints on the demographics of dual AGNs and the efficiency of AGN fueling in minor mergers."
    },
    "0807/0807.4228_arXiv.txt": {
        "abstract": "Recent observations with the low threshold Cherenkov telescopes proved that sub-TeV $\\gamma$-rays are able to arrive from active galaxies at relatively large distances in spite of expected severe absorption in the extragalactic background light (EBL). We calculate the $\\gamma$-ray spectra at TeV energies from two example OVV quasars, 3C~273 and 3C~279, assuming that $\\gamma$-rays are injected in the inner parts of the jets launched by the accretion disks. It is assumed that $\\gamma$-rays in the broad energy range (from MeV up to TeV) are produced in these blazars with the power law spectrum with the spectral index as observed from these objects by the EGRET telescope at GeV energies. We take into account the internal absorption of these $\\gamma$-rays by considering a few different models for the radiation field surrounding the jet. The classical picture of the relativistic blob in jet model for the injection of primary $\\gamma$-rays is considered with the injection rate of $\\gamma$-rays as observed by the EGRET telescope in the GeV energy range. The results of calculations are compared with positive detection and the upper limits on the sub-TeV $\\gamma$-ray fluxes from these two sources. It is concluded that even with the Stecker EBL model, the level of the $\\gamma$-ray emission from 3C~279 is close to the recent measurements in the sub-TeV $\\gamma$-ray energies provided that the injected $\\gamma$-ray spectrum extends from the GeV energies over the next two decades with this same spectral index. We also suggest that a few day time scale flare from 3C 273 can be detected by the MAGIC II stereo telescopes. ",
        "introduction": "The TeV $\\gamma$-rays are detected from compact sources (massive binary systems, active galactic nuclei) in which their absorption at the source has to be carefully investigated. In fact, the escape conditions of $\\gamma$-ray photons from the vicinity of the accretion disks surrounding compact objects have been studied in the past in the context of the $\\gamma$-ray production close to the accretion disk radiation starting with the first reports on the TeV $\\gamma$-ray emission from the binary systems in the eighties and the GeV $\\gamma$-ray emission from the active galaxies (see e.g.~early calculations of absorption in the disk radiation by Protheroe \\& Stanev~(1987), Carraminana~(1992), Bednarek~(1993), Backer \\& Kafatos~(1995), and in the quasi-isotropic radiation around the disk-jet by e.g.~ Blandford \\& Levinson~(1995), Dermer \\& Schlickeiser~(1994). The problem of $\\gamma$-ray escape became again very attractive with the discovery of TeV $\\gamma$-ray emission from active galaxies of the BL Lac type (see more recent calculations by e.g.~Protheroe \\& Biermann~(1997), Bednarek~(1997), Donea \\& Protheroe~(2003), Liu \\& Bai~(2006), Reimer~(2007)). Thanks to the low energy threshold the MAGIC telescope, it is at present clear that also quasars of the OVV type at very large distances emit at least sub-TeV $\\gamma$-rays (e.g. quasar 3C~279, Albert et al.~2008).  However, higher energy $\\gamma$-rays (above $\\sim 100$ GeV) should be efficiently absorbed in the soft radiation field filling the intergalactic space. Therefore observations of TeV $\\gamma$-rays from the population of AGNs at different distances might be used to put the upper limits on the Extragalactic Background Light (EBL). The conclusions reached based on this method strongly depend on the assumptions on the production spectrum of the $\\gamma$-rays at the source. Usually a simple power law is applied with the spectral index which in the extreme case can be constrained by the likely radiation mechanisms of particles accelerated at AGNs (for the most recent analysis see e.g.~Aharonian et al.~(2006a), Mazin \\& Raue~(2007)). However, if the internal absorption inside the source is important, the shape of the emerging $\\gamma$-ray spectrum can be quite complex, i.e. the spectrum can become significantly flatter (see e.g. Fig. 9 in~Bednarek~1997 and Aharonian et al.~2008) than considered their limiting values $\\propto E_\\gamma^{-1.5}$ or even $\\propto E_\\gamma^{-2/3}$. Therefore, derived upper limits on EBL are in fact model dependent, strongly depending on the details of the production of $\\gamma$-rays inside the source.  In this paper we calculate the opacities for $\\gamma$-ray photons inside the central regions of AGNs in the most general case, i.e. for arbitrary angles of injection of $\\gamma$-rays in respect to the jet axis and also for arbitrary inclination of the jet in respect to the plane of the accretion disk. The opacities for the $\\gamma$-rays propagating in the spherically symmetric Broad Line Region are also taken into account. Therefore, we can obtain the three-dimensional $\\gamma$-spheres around the central engine which can be specially useful in the case of highly inclined jets towards the observer, e.g. as it might occur in the case of strongly variable TeV $\\gamma$-ray radio galaxy M~87 (Aharonian et al.~2006b, Aciari et al.~2008).  Moreover, we show how the intrinsic spectra of $\\gamma$-rays emitted from AGNs should be modified in the case of two example EGRET blazars, 3C~273 (Hermsen et al.~1993, von Montigny et al.~1993, Lichti et al.~1995) and 3C~279 (Kniffen et al.~1993, Wehrle et al.~1998, Hartman et al.~2001). The second source has been  also recently detected by MAGIC telescope (Albert et al.~2008), being at present the most distant sub-TeV $\\gamma$-ray source. Based on the new observations of 3C~279 in the sub-TeV energies, we estimate the location of the production region inside the jet assuming that: (1) the $\\gamma$-ray flare observed by the MAGIC telescope was comparable to the highest activity state observed by the EGRET telescope at GeV energies; (2) the production spectrum of $\\gamma$-rays at the source extends from GeV energies up to TeV energies without a break. The effects of possible absorption of TeV $\\gamma$-rays during propagation in the EBL are taken into account by applying its extreme estimates derived by~Primack et al.~(2005) and~Stecker et al.~(2006). ",
        "conclusions": "We discuss in detail the effects of internal absorption of $\\gamma$-rays produced in jets of the OVV type blazars concentrating on two famous sources: 3C 279 and 3C 273. The first source has been recently detected in the sub-TeV energies by the MAGIC telescope (Albert et al.~2008), but only the upper limit exists on the TeV emission from the second source (von Montigny et al.~1997).  It is shown that the model for $\\gamma$-ray production in the jet of 3C 279, in which the spectrum is extrapolated from the EGRET energy range (observed up to several GeV) to the $\\sim 1$ TeV energy range, is quite consistent with reported by the MAGIC collaboration $\\gamma$-ray spectrum, even in the case of the upper estimate on the extragalactic background light derived by Stecker et al.~(2006). Note that this is the furthest TeV $\\gamma$-ray source ($z = 0.536$) detected up to now. Therefore, we conclude that the observation of this most distant sub-TeV source does not introduce any new important constraints on the EBL. Such constraints are used to be derived in the case of the population of the closer BL Lac type objects in which case the TeV emission extends to much higher energies. However, if the EBL is closer to the estimate by Primack et al.~(2005) model, then the $\\gamma$-ray spectrum injected from the considered active region in the jet has to show a clear break above several GeV. At present it is difficult to conclude only on the theoretical grounds if the $\\gamma$-ray spectrum in OVV blazars continue to higher energies or have significant break.  We also show that detection of 3C 273 with the present MAGIC I telescope is rather problematic even if the source is captured in the highest up to now observed emission state by the EGRET telescope and the EBL is well described by the Primack et al. model. However, there is a real chance to detect the peak of a few day flare from this source with the MAGIC II stereo system (which should have at least 2 times better sensitivity), provided that the EBL is closer to the Primack et al. model.  We have only shown the results of calculations for the two values of the Doppler factor of the emission region in the jet ($D = 7$ and 11). In the case of emission regions moving with larger Doppler factors, the internal absorption of $\\gamma$-rays in the radiation field around the jet should be reduced due to the larger distance traveled by the emission region and related to this on the average more distant injection of $\\gamma$-rays from the accretion  disk and BLR clouds. We also assumed that the flare in the jet starts to develop already from its base. However, farther away from the accretion disk, the internal absorption is certainly weaker. Therefore, we consider the most pessimistic case for the escape of TeV $\\gamma$-rays from these OVV blazars. Note that location of the injection region of $\\gamma$-rays inside the jet can be in the future constrained by the more sensitive observations which allow detection of much shorter variability time scales of the $\\gamma$-ray emission. Discovery of a very short time scale variability in blazars of the BL Lac type (of the order of a few minutes, see Mrk 501 (Albert et al.~2007), and PKS 2155-304 (Aharonian et al.~2007), and also radio galaxy M87 (Aharonian et al.~2006)), strongly suggest that the $\\gamma$-ray emission region is rather located  very close to the base of the jet, i.e. to the surface of the accretion disk."
    },
    "0807/0807.4702_arXiv.txt": {
        "abstract": "The supermassive black holes (SMBHs) massive enough ($\\gtrsim$ few$\\times 10^{9}\\Msol$) to power the bright redshift $z\\approx 6$ quasars observed in the Sloan Digital Sky Survey (SDSS) are thought to have assembled by mergers and/or gas accretion from less massive ``seed'' BHs.  If the seeds are the $\\sim 10^{2}\\Msol$ remnant BHs from the first generation of stars, they must be in place well before redshift $z=6$, and must avoid being ejected from their parent proto--galaxies by the large (several$\\times 10^{2} \\km\\;\\s^{-1}$) kicks they suffer from gravitational--radiation induced recoil during mergers with other BHs. We simulate the SMBH mass function at redshift $z>6$ using dark matter (DM) halo merger trees, coupled with a prescription for the halo occupation fraction, accretion histories, and radial recoil trajectories of the growing BHs.  Our purpose is (i) to map out plausible scenarios for successful assembly of the $z\\approx 6$ quasar BHs by exploring a wide region of parameter space, and (ii) to predict the rate of low--frequency gravitational wave events detectable by the {\\it Laser Interferometer Space Antenna} ({\\it LISA}) for each such scenario. Our main findings are as follows: (1) $\\sim 100 \\Msol$ seed BHs can grow into the SDSS quasar BHs without super--Eddington accretion, but only if they form in minihalos at $z\\gtrsim30$ and subsequently accrete $\\gtrsim 60\\%$ of the time; (2) the scenarios with optimistic assumptions required to explain the SDSS quasar BHs overproduce the mass density in lower--mass (few$\\times10^{5}\\Msol \\lsim M_{\\rm bh}\\lsim$ few$\\times10^{7}\\Msol$) BHs by a factor of $10^2-10^3$, unless seeds stop forming, or accrete at a severely diminished rates or duty cycles (e.g. due to feedback), at $z\\ltsim 20-30$.  We also present several successful assembly models and their {\\it LISA} detection rates, including a ``maximal'' model that gives the highest rate ($\\sim 30$ yr$^{-1}$ at $z=6$) without overproducing the total SMBH density. ",
        "introduction": "The discovery of bright quasars at redshifts $z\\gsim 6$ in the Sloan Digital Sky Survey (SDSS) implies that black holes (BHs) as massive as several $\\times 10^9{\\rm M_\\odot}$ were already assembled when the age of the universe was less than $\\approx 1$ Gyr (see the recent review by Fan 2006).  These objects are among the oldest detected discrete sources of radiation in the Universe.  The likelihood that all of these quasars are significantly magnified by gravitational lensing, without producing detectable multiple images (Richards et al. 2004), is exceedingly small (Keeton et al. 2005), and if their luminosities are powered by accretion at or below the Eddington rate, the central objects must be $\\sim 10^{9}\\Msol$ supermassive black holes (SMBHs). In particular, the quasar SDSS J1148+5251 (Fan et al.  2003, 2001) is likely to be powered by a SMBH with a mass of $\\approx 10^{9.5}\\Msol$ (Willott et al. 2003).  The mechanism by which such massive BHs formed within 1 Gyr after the Big Bang remains poorly understood.  Generically, these SMBHs are thought to have assembled by mergers with other BHs and/or by gas accretion \\footnote{In this paper, ``accretion'' onto BHs should be assumed to mean gas accretion, unless otherwise noted.  } onto less massive BHs. If the first (``seed'') BHs are the $\\sim 10^{2}\\Msol$ remnant BHs of the first generation of stars (e.g. Heger et al. 2003), they must be in place well before redshift $z=6$.  If accretion onto BHs is limited at the Eddington rate with radiative efficiency $\\epsilon$, defined as the fraction of the rest mass energy of matter falling onto the BH that is released as radiation, then $1-\\epsilon$ of the matter is accreted and the growth of the BH mass $m$ is given by \\begin{equation} \\frac{d \\ln m}{dt}=\\frac{1-\\epsilon}{\\epsilon}\\frac{4\\pi G\\mu m_{\\rm p}} {\\sigma_{e}c}, \\end{equation} where $G$ is the gravitational constant, $c$ is the speed of light, $\\mu\\approx 1.15$ is the mean atomic weight per electron for a primordial gas, and $\\sigma_{e}$ is the Thompson electron cross section.  The e--folding time scale for mass growth is $t_{\\rm Edd}\\approx 4.4\\times 10^7 \\;\\yr$ for $\\epsilon=0.1$.  In the concordance cosmological model (see below) the time elapsed between redshifts $z=30$ (when the first seeds may form) and $z=6.4$ (the redshift of the most distant quasar) is $\\approx 0.77$ Gyr, allowing for a mass growth by a factor of $\\approx 10^{7.7}$.  Therefore, individual $\\sim100\\Msol$ seeds can grow into the SDSS quasar BHs through gas accretion alone, provided the accretion is uninterrupted at close to the Eddington rate and $\\epsilon\\ltsim 0.1$. A higher efficiency and/or a lower time--averaged accretion rate will require many seed BHs to merge together; the number of required mergers increases exponentially for lower time-averaged accretion rates.  The discovery of the bright quasars at $z\\gsim 6$ were followed by the first successful numerical calculations in full general relativity of the coalescence of a BH binary and the corresponding emission of gravitational waves (GWs) (Pretorius 2005; Campanelli et al. 2006; Baker et al. 2006).  These calculations also confirmed a result previously known from post-Newtonian (Kidder 1995) and perturbation-theory treatments (Favata et al. 2004, Schnittman \\& Buonanno 2007): the coalesced product receives a large center-of-mass recoil imparted by the net linear momentum accumulated by the asymmetric gravitational wave emission (see Schnittman et al. 2008 for a recent detailed discussion of the physics of the recoil, and for further references).  Typical velocities for this gravitational recoil (or ``kick'') are in excess of $\\sim 100 \\km~\\s^{-1}$. This is likely more than sufficient to eject the BHs residing in the low--mass protogalaxies in the early Universe, since the escape velocities from the DM halos of these galaxies are only a few $\\km~\\s^{-1}$.  Several recent works have studied the role of gravitational kicks as an impediment to merger-driven modes of SMBH assembly.  Simple semi--analytic models show that if every merger resulted in a kick large enough to remove the seed BHs from halos with velocity dispersions up to $\\approx 50 \\km~\\s^{-1}$, then super-Eddington accretion would be required to build SMBHs of the required mass in the available time (Haiman 2004; Shapiro 2005). Monte--Carlo merger tree models that exclude kicks entirely (Bromley et al. 2004) or which include a distribution of kick velocities extending to low values (e.g. Yoo \\& Miralda-Escud\\'{e} 2004 [hereafter YM04]; Volonteri \\& Rees 2006 [hereafter VR06]) give a slightly more optimistic picture, showing that if seed BHs form in most minihalos in the early Universe, and especially if ejected seeds are rapidly replaced by new seeds (YM04), then SMBH assembly is just possible before $z\\approx 6$ without exceeding the Eddington accretion rate.  These works are encouraging steps toward demonstrating that there are plausible physical models that lead to the timely assembly of SMBHs massive enough to power the $z>6$ SDSS quasars.  At present, we have no direct observational constraints on SMBH assembly at $z>6$, and there is, in principle, a large range of ``physically plausible'' possibilities.  The {\\it Laser Interferometer Space Antenna} ({\\it LISA}) is expected to be able to detect mergers of SMBHs in the mass range $\\sim (10^4$--$10^7)\\,{\\rm M_\\odot}/(1+z)$ out to $z\\sim 30$, and to extract binary spins and BH masses with high precision up to $z\\sim 10$ (Vecchio 2004; Lang \\& Hughes 2006, 2007).  It is also likely that by the time {\\it LISA} is operational, there will be additional independent constraints on the demography of high-redshift SMBHs.  It is therefore a useful exercise to calculate the expected {\\it LISA} event rate from high--redshift SMBH mergers (e.g. Wyithe \\& Loeb 2003; Sesana et al. 2004, 2005, 2007) in a range of plausible models. Note that published estimates (Menou et al. 2001, Heger et al. 2003, Menou 2003, Haehnelt 2003,Wyithe \\& Loeb 2003, Sesana et al. 2004, Islam et al. 2004, Koushiappas \\& Zentner 2006, Micic et al. 2007, Lippai et al. 2008, Arun et al. 2008) for the {\\it LISA} event rate, even at lower redshifts, vary by orders of magnitude, from $\\sim 1$ to as high as $\\sim 10^4$ yr$^{-1}$; there is a large range even among models that explicitly fit the evolution of the quasar luminosity function \\citep{lfh08}.  A related open question is to what degree the {\\it LISA} data stream can help pinpoint the actual SMBH assembly scenario.  One aim of this paper is to understand the model degeneracies that can lead to similar {\\it LISA} data streams. Another is to explore as much as possible the full variety of {\\it LISA} event rates arising from various ``physically plausible'' assembly models.  The physical factors that determine the growth of SMBHs at high redshift fall broadly into four categories: (1) the properties of the initial seed BHs, such as their redshift, mass, and abundance; (2) the time--averaged gas accretion rate of individual seeds; (3) the merger rates of BHs; and (4) effects governing the fate of gravitationally kicked BHs.  The first category determines the number and mass of seed BHs available for assembly, and depends primarily on the behavior of gas in the host DM halos, and the mass and metallicity of the first stars.  The second category measures the subsequent growth through accretion, and depends on the availability of fuel over the Hubble time, and its ability to shed angular momentum and accrete onto the BH. The third category is a combination of the halo merger rate, the rarity of seeds, and the timescale for the formation, orbital decay, and ultimate coalescence of a SMBH binary.  Finally, the recoil velocity of the coalesced binary is determined by the mass ratio and spin vectors of the BHs, and its subsequent orbit -- and whether it is retained or ejected before the next merger -- will be determined by the overall depth of the gravitational potential of the DM host halo, and on the spatial distribution of gas and DM in the central region of the halo, which determine drag forces on the kicked BH.  Our approach to model the above effects closely follows those in earlier works (e.g. YM04; Bromley et al. 2004; Sesana et al. 2004; VR06): we use Monte Carlo merger trees to track the hierarchical growth of DM halos, and a simple semi-analytical model to follow the growth and dynamics of BHs.  We expand over earlier works by adding an explicit calculation of the orbits of kicked BHs, and self--consistently include their corresponding time--dependent accretion rate.  Additionally, we extend the merger tree to redshifts beyond $z>40$, and we examine a large range of different models.  For each set of model parameters, we apply our ``tree-plus-orbits'' algorithm to the entire halo population at $z\\approx 6$ to construct a full population of SMBHs at this redshift, and calculate physical quantities of interest: the mass function, the SMBH-to-halo mass ratio, the fraction of DM halos hosting SMBHs, and the expected detection rate of SMBH mergers by {\\it LISA}.  This algorithm assembles SMBHs through simple prescriptions of the aforementioned four categories of physical contributions to SMBH formation.  We model the seed population by assuming that some fraction $ f_{\\rm seed}$ of DM halos reaching a threshold virial temperature $T_{\\rm seed}$ forms a Pop III remnant BH.  In--between mergers, the BHs are assumed to accrete gas at a rate of $f_{\\rm duty}$ times the Eddington rate $\\dot{m}_{\\rm Edd}\\equiv (1-\\epsilon)/\\epsilon\\times L_{\\rm Edd}/c^{2}$. Here $f_{\\rm duty}$ should be interpreted as the mean gas accretion rate (averaged over time--scales comparable to the Hubble time) in units of the Eddington rate.  Note that this prescription makes no distinction between episodic accretion near the Eddington rate during a fraction $f_{\\rm duty}$ of the time (with no accretion in-between), and constant accretion at all times at a fraction $f_{\\rm duty}$ of the Eddington rate.  We assume that the mergers of BH binaries closely follow the mergers of their host halos (but allow for a delay in the latter due to dynamical friction). Finally, we simulate the orbits of recoiling BHs under different assumptions about the baryon density profile and binary spin orientation.  We discuss the relative importance of assembly model parameters on the final SMBH mass function and the {\\it LISA} data stream, and ask whether {\\it LISA} will be able to uniquely determine the underlying assembly model from data.  We also examine several variants of the above scenario, in which (i) the seed BHs are massive, $\\sim 10^5 \\Msol$, and formed from the super-Eddington accretion of a collapsed gaseous core; (ii) the DM halo is initially devoid of gas when the seed BHs is formed; (iii) seed BHs stop forming below some redshift; and (iv) models which maintain the so--called $M_{\\rm bh}-\\sigma$ relation between BH mass and (halo) velocity dispersion (Ferrarese \\& Merritt 2000; Gebhardt et al. 2000) at all redshifts.  This paper is organized as follows.  In \\S~2, we detail our methodology by describing our assumptions, the assembly algorithm, including the prescriptions of the aforementioned physical effects, and the different assembly scenarios we consider.  We present and discuss our main results in \\S~3.  In \\S~4, we summarize our most important results, and comment on future prospects to understand SMBH assembly at high redshift.  To keep our notation as simple as possible, throughout this paper the capital $M$ will be used to denote halo masses, and $m$ will refer to BH masses.  In this paper, we adopt a $\\Lambda$CDM cosmology, with the parameters inferred by Komatsu et al. (2008) using the five-year data from the {\\it Wilkinson Microwave Anisotropy Probe} ({\\it WMAP5} ): $\\Omega_{\\rm CDM}=0.233$, $\\Omega_{b}=0.046$, $\\Omega_{\\Lambda}=0.721$, $h=0.70$, and $\\sigma_{8}=0.82$.  We use $n_{\\rm s}=1$ for the scalar index. ",
        "conclusions": "In this paper, we have attempted to map out plausible ways to assemble the $\\gtrsim 10^{9}\\Msol$ SMBHs that power the bright redshift $z\\approx 6$ quasars observed in the SDSS, without overproducing the mass density in lower--mass ($\\sim 10^{5-7}\\Msol$) BHs.  We also computed the event rates expected for {\\it LISA} in each of the successful models.  The physical effects governing SMBH assembly depend on the answers to four basic questions: (1) how common are the initial BH seeds; (2) how much mass in gas do they accrete, and therefore how much they contribute individually to the final SMBH's mass; (3) how often do they merge; and (4) what happens to SMBH binaries when they do merge? Currently, we do not have empirical constraints to offer definitive answers to any of these questions.  However, we are capable of predicting the final outcome, starting with a set of assumptions for the underlying physics.  Our trees-plus-orbits algorithm simulates the formation history of SMBHs and the subsequent detection rate expectations for {\\it LISA} by isolating and prescribing answers to the above four questions.  It is a powerful simulation tool, as it can incorporate a detailed modeling of individual physical prescriptions without a significant increase in the computational load, as long as the prescriptions can be described by fitting formulae, tabulated in a lookup table of reasonable size or summarized in a statistical manner.  Using this tool, we have surveyed a wide range of candidate assembly models, and reported on common and distinguishing traits in the resulting SMBH mass functions and the corresponding {\\it LISA} detection rates.  In particular, we have shown that SMBHs can form in a manner consistent with other observational evidence either through the rapid growth of rare, massive seeds, or through ultra--early production of numerous Pop-III remnant seeds, provided these seeds stop forming below a redshift $z_{\\rm cut}\\sim 20-30$.  We reach the pessimistic conclusion that these scenarios do not produce {\\it any} detectable {\\it LISA} events at $z>6$ (at least not in a few year's operation).  An alternative model, in which we assume that the extrapolation of the local $m-M$ relation holds at all redshifts (e.g. due to internal feedback), on the other hand, can produce up to $\\sim 30$ {\\it LISA} events per year, with a characteristic mass spectrum.  Our major findings can be summarized more specifically as follows:  \\begin{itemize}  \\item SMBHs must be continuously surrounded by dense gas that was able to cool at the centers of DM halos.  Feeding holes with the low--density gas in DM halos whose gas was unable to cool does not allow for high enough accretion rates to explain the SDSS quasar BHs.  \\item If embedded in dense gas nearly continuously, $\\sim 100 \\Msol$ seed BHs can grow into the SDSS quasar BHs without super--Eddington accretion, but only if they form in minihalos at $z\\gtrsim30$ and subsequently accrete $\\gtrsim 60\\%$ of the time.  However, these optimistic assumptions, required to explain the SDSS quasar BHs, overproduce the mass density in lower--mass (few$\\times10^{5}\\Msol \\lsim M_{\\rm bh}\\lsim$ few$\\times10^{7}\\Msol$) BHs by a factor of $10^2-10^3$.  We find that two conditions need to be satisfied to alleviate this overprediction: the initial occupation fraction of seed BHs has to be low ($f_{\\rm occ}\\ltsim 10^{-2}$), and new seeds must stop forming, or the seeds must accrete at severely diminished rates or duty cycles, at $z\\ltsim 20-30$.  We argued that models in which BH seeds stop forming at $z\\sim 20$ are attractive because there are physical reasons for the seeds to stop forming below some redshift (such as metal pollution and/or radiative feedback that suppresses pop-III star formation), and because there is independent empirical evidence, from {\\it WMAP} constraints on the reionization history, that star and/or BH formation in high--redshift minihalos was suppressed by a factor of $\\gsim 10$ (Haiman \\& Bryan 2006).  \\item The simplest SMBH assembly scenarios, which have constant accretion rates, but in which BH seed formation stops abruptly at some redshift, and which meet constraints at both the high--mass and low--mass end of the $z=6$ SMBH mass function, predict negligibly low {\\it LISA} event rates.  The reason for the low rates is as follows: in these models, the BHs that grow into the most massive, highest-redshift quasar-SMBHs accrete at the same (exponential) rate as all the other BHs, typically resulting in a vast overproduction of massive ($m\\sim 10^{6}\\Msol$) holes.  In order to offset this overproduction, seeds must be made very rare, and this diminishes the {\\it LISA} rates.  It is difficult to envision a scenario for high ($\\gtrsim 10$ per year per unit redshift) detection rates unless a vast number of SMBHs in the $10^{5-7}\\Msol$ range lurk in the universe at all redshifts, which the current electromagnetic surveys have missed.  \\item A different class of successful models, in which the SMBH masses are self--regulated by internal feedback, can evade this constraint, and produce {\\it LISA} rates as high as $30$ yr$^{-1}$.  The key difference in these models that predict higher {\\it LISA} rates is that the SMBH growth is driven by a large number of seed BHs and far lower gas accretion rates than those required in the constant-accretion models.  The majority of these events occur at $z\\approx 6$ and in the low end ($10^3-10^4$~${\\rm M_\\odot}$) of {\\it LISA}'s mass range for detection. Also, for these models we find the ejected BH mass density can exceed that of the galactic BH population at $z=6$.  Most ejected holes are expected to have masses similar to the seed mass, but an ejected BH can be as massive as $\\sim10^{8}\\Msol$ if large recoil velocities are allowed (e.g. if spins are not always aligned with the orbital angular momentum of the binary).  \\end{itemize}  In addition to the above, our modeling reveals a number of interesting aspects of SMBH assembly.  We find that in the successful models the initial seeds are rare, and the most massive SMBHs grow primarily from the few 'lucky' early seeds that avoided ejection due to kicks.  The precise assumptions regarding the kick velocity distribution (such as the assumed spin orientations or the resulting oscillation of the BH) tend to have only a modest effect on the final results in these models. This is because, at least in our simple prescription, BHs either return quickly to the gas-rich nucleus or are left wandering in the outer regions.  Our results suggest that {\\it LISA} will be capable of narrowing the field of plausible SMBH assembly models from the raw event rate, even without detailed measurements of the binary spins or mass ratios.  The spin and mass ratio measurements will further constrain the evolution of SMBH properties.  While the component prescriptions explored in this paper are admittedly crude, exercises similar to the one performed in our study will be crucial in understanding the limits and possibilities offered by {\\it LISA}, and ultimately to interpret the detected {\\it LISA} events.  The scarcity of empirical constraints on the various pieces of physics that determines the SMBH growth leaves us with a large range of ``plausible'' scenarios and free parameters for SMBH assembly."
    },
    "0807/0807.3778_arXiv.txt": {
        "abstract": "We study the consequences of the $f(R/\\Box)$ gravity models for the Solar system and the large scale structure of the universe. The spherically symmetric solutions can be used to obtain bounds on the constant and the linear parts of the correction terms. The evolution of cosmological matter structures is shown to be governed by an effectively time dependent Newton's constant. We also analyze the propagation of the perturbation modes. Tensor and vector modes are only slightly modified, but two new scalar degrees of freedom are present. Their causality and stability is demonstrated, and their formal ghost conditions are related to a singularity of the cosmological background. In general, the Newtonian limit of these models has no apparent conflicts with observations but can provide useful constraints. ",
        "introduction": "One covariant way of generalizing gravity is to consider the effective Newton's constant as a function of $\\phi=R/\\Box$, the inverse d'Alembertian acting on the scalar curvature. This kind of nonlinear dependence can be parameterized as the following \\be \\label{action} S = \\frac{1}{2\\kappa^2}\\int d^4 x \\left(1+f(\\frac{R}{\\Box})\\right)R. \\ee General relativity is the limit $f=0$. This form of the action can be motivated\\footnote{On the other hand, such models might appear in (stringy) compactification scenarios. This will become clearer when we introduce an equivalent description in terms of kinetically coupled scalar fields (non linear sigma model).} by the fact that quantum effects generically introduce corrections involving different operators and curvature invariants \\cite{Donoghue:1994dn,Shapiro:2008sf}. The exact form of the corrections is of course beyond control, but Eq.(\\ref{action}) provides a simple parameterization, which can be constrained by observations: this is the phenomenological approach we take in the present study.  Previously it has been noted that the most appealing (linear) form, $f(\\phi) \\sim \\phi$, could cure the unboundedness of the Euclidean gravity action \\cite{Wetterich:1997bz}. More recently it was suggested \\cite{Deser:2007jk} that these types of effective gravity actions could incorporate the observed cosmic acceleration with reduced fine tuning. The problem of the magnitude of the cosmological constant might be evaded since $\\phi = R/\\Box$ is a dimensionless combination, and thus no new scales (in addition to the Planck scale) need to be introduced. The coincidence problem could also get a simple resolution, since $R$ and thus $\\phi$ vanish during radiation domination, so the acceleratory effect can be interpreted as a delayed consequence of the onset of the matter dominated era.  Detailed analysis of the background dynamics shows that this could work, even for the simplest forms of $f$ \\cite{Koivisto:2008xf}. Unresolved issues exist though, like an occurrence of a sudden future singularity \\cite{Barrow:2004xh,Bamba:2008ut}. In addition, in the simplest models featuring the late acceleration, one finds large effects taking place also during inflation. On the other hand, this could be used as beneficial effect to end inflation and reheat to universe, or even to drive inflation without the usual scalar fields as attempted in Ref.\\cite{Nojiri:2007uq}. Realizing this possibility would require to find the classes of actions unifying consistently very early and the late universe. The fact that the action (\\ref{action}) is capable of reproducing an arbitrary expansion history \\cite{Koivisto:2008xf}, hints that this is possible. However, in this work omit the background problematics, since our aim is to focus on inhomogeneous evolution of these models, particularly at the important Newtonian limit.  This limit within the $f(R)$ class of models has been intensively studied in connection with the dark matter and dark energy problems, see \\cite{Nojiri:2006ri,Sotiriou:2008rp,Sotiriou:2007yd,Capozziello:2007ec,Lobo:2008sg} for reviews. Even if viable examples would be found in that class of models, there seems to be no reason to prefer those more complicated models to the simple cosmogical constant, since they have not (at least this far) promised any help with the cosmological fine tuning problems. In this regard the $f(R/\\Box)$ scenarios make progress, as they, like mentioned above, could provide both a mechanism to connect the onset of the acceleration with the past matter dominated era, and remove the need for unnatural constants in the theory. Though also a $f(R)$ action can be recasted into a scalar-tensor theory, the structure of the action (\\ref{action}), which we'll soon clarify, is different.  There has also been recently much interest in nonlocal scalar field cosmologies \\cite{Barnaby:2007ve}. These are motivated by the p-adic tachyon and open string field theory actions, which include operators with infinite number of derivatives acting on scalar fields. In particular, the exponent of the d'Alembertian, $e^\\Box$, appears. These kind of theories are equivalent to theories with several (possibly an infinite number of) coupled scalar fields, as soon becomes explicit in a different context. Nonperturbative methods can then be applied to treat localizable systems via higher-dimensional formulation \\cite{Calcagni:2007ru,Calcagni:2007ef,Calcagni:2008nm,Mulryne:2008iq}. Friedmann-Robertson-Walker (FRW) solutions have been studied in string field theory \\cite{Aref'eva:2004vw,Aref'eva:2007uk,Calcagni:2005xc,Barnaby:2008fk}, and in related string-inspired forms of nonperturbative gravity with infinitely many derivatives \\cite{Biswas:2005qr,Biswas:2006bs}.  Finally, let us remark that in the so called degravitation models the effective Newton's constant is promoted to a nonlocal operator, whose filtering effect could explain why gravity is insensitive to a cosmological constant \\cite{ArkaniHamed:2002fu}, leaving the possibility of the present cosmological term appearing as an afterglow of inflation \\cite{Patil:2008sp}. These interesting models are burdened with problems associated to massive gravity \\cite{Dvali:2007kt}, including the loss of covariance, as the change of $G$ is done at the level of field equations and not of the action. Though they are thus fundamentally different types of gravity modifications than we consider in the present paper, one can observe some analogous workings at a phenomenological level. In fact, similar actions could be constructed approximating the degravitation models \\cite{Barvinsky:2003kg}.  The present paper is devoted to study of inhomogeneous solutions of the $f(\\phi)$ case, and in particular in the weak field limit where the inhomogeneities appear as perturbations about more symmetric solutions. We study all types of cosmological perturbations for the first time, and look at the Solar system limit in more detail than previously done. In section (\\ref{sektio}) we present the model in generalized framework. In section (\\ref{sektio2}) we write down the equations for each type of linear perturbations. The post-Newtonian limit, relevant for the Solar system physics and for the evolution of cosmological structure, is determined in section (\\ref{sektio3}). The propagation of the scalar modes is more complicated than that of the vector and tensor modes, and we discuss that in the section (\\ref{sektio4}). The conclusion is made in section (\\ref{sektio5}). ",
        "conclusions": "\\label{sektio5}  We considered a class of cosmological models based on the generalized gravity action (\\ref{action}). Though apparently nonlocal, this action may be recasted into a local biscalar-tensor theory. The theories represent a subclass of more general covariant theories, involving more, possibly infinite number of scalar fields. The simple parameterization of the corrections to Einstein gravity as the function $f(R/\\Box)$ has already several new theoretical features and potential to describe new phenomenology. In particular, one may construct dark energy cosmologies without introducing unnatural parameters and explain the coincidence of the similar matter and dark energy densities just today by the response of $R/\\Box$ to the onset of matter domination.  In this paper we focused on the cosmological perturbations and weak field limit of these models. The scalar, vector and tensor perturbation equations were derived up to linear order in the flat FRW background. To simplify the equations, we fixed a gauge for each type of perturbation, but of course the equations in an arbitrary gauge may be obtained from our results using the standard transformation formulas. The scalar equations we wrote down in the Newtonian (or longitudinal) gauge. We considered the evolution of inhomogeneities in a matter dominated universe. We found that gradient effects are absent, but that the growth rate is modified according to a time evolving effective gravitational constant. Also, as rather generic predictions of modified gravity, the Poisson equation is modified and there is effective anisotropic stress. These effects might provide hints to detect these models from any observed deviation from the general relativistic predictions. In particular, the integrated Sachs-Wolfe effect probes the late time variation of the gravitional potentials \\cite{Jain:2007yk,Bertschinger:2008zb}, and the small scale effects of modified gravity might be probed by weak lensing experiments \\cite{Amendola:2007rr,Schmidt:2008hc}. We hope to quantify these constraints on specific models by a detailed numerical analysis in a future work. In addition, inflation and the generation of the primordial spectrum would be interesting to study with the action (\\ref{action}).  We considered here the Post-Newtonian limit for static spherically symmetric solutions, which is crucial for the local tests of general relativity. We found that there are two PPN type parameters possibly differing from their general relativistic values, with a given dependence of the form of $f$. This dependence allows to constrain tightly the constant and the linear parts of the function $f$. In particular, they are severely restricted by the bounds on the post-Newtonian $\\gamma$ obtained from the tracking of the Cassini spacecraft, resulting in Eq.(\\ref{gammab}). For example, the purely linear model resulting in $\\Omega_m = 0.3$ today is ruled out by this constraint alone (though this model would not accelerate the universe enough anyway \\cite{Koivisto:2008xf}), as well as the simplest linear stabilator of the Euclidean action (though it already had tension with the Big Bang nucleosynthesis \\cite{Wetterich:1997bz}). The constraints on the higher powers in $f$ seem to be some orders of magnitude looser.  The issues of causality and stability were also considered. The rotational perturbations were shown to decay as usually. The gravitational waves evolve with an extra friction term and a modulated response to the matter stress sources due to the coupling to $(1+\\psi)$. The evolution of the two extra scalar degrees of freedom is more delicate. We found that it is possible to separate these degrees of freedom into decoupled modes by considering the theory in the Einstein frame, diagonalizing the kinetic terms and then considering the perturbations $\\delta s$ and $\\delta \\sigma$, which are fluctuations of implicitly defined combinations of the fields appearing in the original Jordan frame Lagrangian. However, since it is possible to write down decoupled, canonical actions for those two perturbations, we believe it is legitimate to regard them as the physically propagating degrees of freedom. This recipe would not probably work in a more general case though. Here our result is that the new scalar modes propagate perturbations with the light speed, and thus they are causal and stable. The formal ghost conditions are directly linked to divergence of classical perturbations (similar to the Gauss-Bonnet cosmology) and here also with the sudden future singularity of the cosmological background. Thus these conditions clearly indicate a pathological point of the solutions. We note that the singularity structure of the $f(R)$ models has also raised recent concerns \\cite{Frolov:2008uf,Kobayashi:2008tq,Dev:2008rx,Bamba:2008ut}.  To conclude, the $f(R/\\Box)$ gravities seem have a viable Newtonian limit, although they may feature significant modifications of the cosmological expansion."
    },
    "0807/0807.3252_arXiv.txt": {
        "abstract": "Recently, magnetic reconnection during collisionless, stressed, X-point collapse was studied using kinetic, 2.5D, fully electromagnetic, relativistic Particle-in-Cell numerical code [D. Tsiklauri and T. Haruki, Phys. Plasmas {\\bf 14}, 112905 (2007)]. Here we finalise the investigation of this topic by addressing key outstanding physical questions: (i) which term in the generalised Ohm's law is responsible for the generation of the reconnection electric field? (ii) how does the time evolution of the reconnected flux vary with the ion-electron mass ratio? (iii) what is the exact energy budget of the reconnection process, i.e. in which proportion initial (mostly magnetic) energy is converted into other forms of energy? (iv) are there any anisotropies in the velocity distribution of the accelerated particles? It has been established here that: (i) reconnection electric field is generated by the electron pressure tensor off-diagonal terms, resembling to the case of tearing unstable Harris current sheet studied by the GEM reconnection challenge; (ii) For $m_i / m_e \\gg 1$ the time evolution of the reconnected flux is independent of ion-electron mass ratio; also, in  the case of $m_i / m_e = 1$ we show that reconnection proceeds slowly as the Hall term is zero; when $m_i / m_e \\gg 1$ (i.e. the Hall term is non-zero) reconnection is fast and we conjecture that this is due to magnetic field being frozen into electron fluid, which moves significantly faster than ion fluid; (iii) within one Alfv\\'en time, somewhat less than half ($\\sim 40$\\%) of the initial total (roughly magnetic) energy is converted into the kinetic energy of electrons, and somewhat more than  half ($\\sim 60$\\%) into kinetic energy of ions (similar to solar flare observations); (iv) in the {\\it strongly} stressed X-point case, in about one Alfv\\'en time, a full isotropy in all three spatial directions of the velocity distribution is seen for super-thermal electrons (also commensurate to solar flare observations). ",
        "introduction": "In many astrophysical or laboratory plasma situations (a) plasma beta is small, indicative of large amounts of energy stored in a form of magnetic field and (b) there is a need to explain or provide plasma heating, as well as  plasma particle acceleration. It is believed that in such situations magnetic reconnection, i.e. change of connectivity of magnetic field lines that penetrate the plasma, can serve as one of the important possible mechanisms. There are different types of magnetic reconnection. One of the key descriptors is plasma collisionality, i.e. if plasma is collisional then magnetic resistivity, $\\eta$ (or more specifically $\\eta {\\vec j}$ term in the generalised Ohm's law) is responsible for breaking the frozen-in condition (enabling field line connectivity change). However, if plasma is collisionless, then other terms in the generalised Ohm's law may be more important. In this context, it is instructive to look at typical spatial scales. Let us consider an example of solar coronal plasma. Typical width of a Sweet-Parker current sheet is given by $\\delta=S^{-1/2} L$ \\cite{biskamp} (p. 54). Fixing coronal temperature at 1 MK, Coulomb logarithm at 18.0948, the Lundquist number (using Spitzer resistivity from \\citet{NRLF}, p. 30) is $5.37933 \\times 10^{12}$, which for $L=10$ Mm yields, $\\delta=4.31157$ m. Another way of looking at $\\delta$ is associating it also with the resistive length scale via Alfv\\'en time scale $\\tau_A = L / V_A$, where $V_A$ is the Alfv\\'en speed ($\\approx 1$ Mm s$^{-1}$): i.e. $\\delta=S^{-1/2} L=\\eta \\tau_A / \\mu_0$. Typical scale associated with the Hall term in the generalised Ohm's law at which deviation from electron-ion coupled dynamics is observed is, $c/\\omega_{pi} =  7.20064$ m (ion inertial length). Here particle density of $n=10^{15}$ m$^{-3}$ is used. The fact that $\\delta / (c/\\omega_{pi}) \\leq 1 $  points to a necessity of going beyond single fluid MHD approximation. The Hall term itself cannot break the frozen-in condition, its inclusion into consideration ensures that the magnetic field is frozen into electron fluid. Below, we shall use this argument to conjecture why reconnection is fast when the Hall term is included. Importance of different terms in the generalised Ohm's law is usually inferred by comparing the spatial scales associated with them to the resistive length scale $\\delta$. E.g. one of the other noteworthy scales is $c/\\omega_{pe} =   0.16804$ m (electron inertial length) on which the electron inertia term operates. When $\\delta / (c/\\omega_{pe}) < 1$ then electron inertia would dominate over resistive diffusion \\cite{biskamp} (p. 200). As can be seen from the above estimates  $\\delta / (c/\\omega_{pe}) = 25.6576 \\gg 1 $ in the solar corona, thus electron inertia effects seem to be negligible. However, electron inertia in contrary to the Hall term (as well as the electron pressure tensor) can break the frozen-in condition and thus change the magnetic field connectivity. Because, of the fact that with increase of $T$ (hot plasmas), $\\delta$ gets progressively smaller thus effects other than the resistivity should be included. Indeed, collisionless (non-resistive) reconnection has recently attracted considerable attention (see \\citet{bp07} for a review).  One of the first studies of magnetic reconnection is stressed, X-point collapse  \\cite{d53} (also see Chap. 7.1 in \\citet{pf00}). The latter was using resistive MHD approach. We recently revisited the problem in the regime of collisionless reconnection \\cite{th07}. In Ref. \\cite{th07} we studied the magnetic reconnection during collisionless, stressed, X-point collapse using kinetic, 2.5D, fully electromagnetic, relativistic Particle-in-Cell numerical code. We considered  two cases of weakly and strongly stressed X-point collapse. Where descriptors weakly and strongly refer to 20 \\% and 124 \\% unidirectional spatial compression of the X-point, respectively. Amongst other interesting outcomes, we established that within about one Alfv\\'en time,  2\\%  and 20\\% of the initial magnetic energy can be  converted into heat and accelerated particle energy in the cases of weak and strong stress, respectively. However, open questions remained: (i) which term in the generalised Ohm's law is responsible for the generation of the reconnection electric field? (ii) how does the time evolution of the reconnected flux vary with the ion-electron mass ratio? (iii) what is the exact energy budget of the reconnection process, i.e. in which proportion initial (mostly magnetic) energy is converted into other forms of energy? (iv) are there any anisotropies in the velocity distribution of the accelerated particles? Here we finalise the study of magnetic, collisionless reconnection of a stressed X-point by providing answers to these questions. ",
        "conclusions": "By and large, the present work closes our initial study of stressed X-point collapse in the collisionless regime started in Ref.\\cite{th07}, by bridging gaps in the understanding of key physical aspects. The main findings can be listed as following:  (i) despite significant differences of the initial setup between tearing unstable Harris current sheet \\cite{pritchett01} and stressed X-point considered here, in both cases source of the reconnection out-of-plane electric field at the magnetic null is provided by off-diagonal terms of the electron pressure tensor.  (ii) we find that when $m_i / m_e \\gg 1$ reconnection rate is independent of the ion-electron mass ratio and it is fast, which is also witnessed by \\citet{hesse99}. However, when electron-ion mass ratio is unity, i.e. the Hall term is switched off, we show that reconnection rate is indeed slow. This broadly agrees with the results of \\citet{birn01} (However see also Refs.\\cite{bb07,dk07,hz07} for alternative view). When the Hall physics is included, we also conjecture that the reconnection is fast because the magntic field (being frozen into electron fluid, which moves significantly faster than ion fluid, as shown in Ref.\\cite{th07}) is transported in and out of the diffusion region much faster than in the case of single fluid resistive MHD. We show that the amount of reconnected flux attained by $ \\omega_{ci} t=25$ in the cases of $m_i / m_e \\gg 1$ and $m_i / m_e = 1$ has the same ratio ($\\approx 4$)  as is the ratio of electron and ion speeds ($\\approx 4-5$).  (iii) we find that within one Alfv\\'en time, roughly $\\sim 40$\\% of the initial total energy (which is mostly stored in the magnetic field) is converted into the kinetic energy electrons, and somewhat more than half ($\\sim 60$\\%) into kinetic energy of ions. In solar flare observations a similar behaviour is seen \\cite{emslie04}.  (iv) When X-point is stressed {\\it strongly}, in about one Alfv\\'en time, a full isotropy in all three spatial directions of the velocity distribution is seen for super-thermal electrons. Again similar behaviour is reported in  solar flare observations \\cite{kontar06}.  Resuming aforesaid, it seems that collisionless, stressed X-point collapse is a viable mechanism for solar flares. Also, its behaviour is remarkably similar to tearing unstable Harris current sheet which is thought to be more relevant for the Earth geomagnetic tail and generally to magnetospheric applications.  We close this paper with some words of caution:  (i) realising that size of the numerically simulated area $40 \\times 40$ electron skin depths ($c / \\omega_{pe}=0.16804$ m in solar corona), i.e. $6.7\\times 6.7$ m is a tiny proportion of the actual solar flare site which can be tens of Mm. Thus, {\\it a priori} it is not clear whether this mechanism is suitable to accelerate enough electrons in a much larger volume.  (ii) whether e.g. the obtained electron-ion kinetic energy partition (as the percentage of total released energy) being roughly 40\\% - 60\\%, would hold for a different mass ratio? We remind the reader that except in subsection III.B the ion-electron mass ratio was kept constant at $m_i/m_e=100$. We can only conjecture that this  40\\% - 60\\% partition holds for a different mass ratio, because we also found that the amount of reconnected flux versus time graph (Fig.~2) showed no mass ratio dependence."
    },
    "0807/0807.3587_arXiv.txt": {
        "abstract": "We study the edge-on galaxy NGC~5775, utilizing a 58.2 ks {\\sl Chandra} ACIS-S observation together with complementary {\\sl HST} ACS, {\\sl Spitzer} IRAC and other multi-wavelength data sets. This edge-on galaxy, with its disk-wide active star formation, is particularly well-suited for studying the disk/halo interaction on sub-galactic scales. We detect 27 discrete X-ray sources within the $D_{25}$ region of the galaxy, including an ultra-luminous source with a 0.3-7 keV luminosity of $\\sim7\\times10^{40}\\rm~ergs~s^{-1}$. The source-removed diffuse X-ray emission shows several prominent extraplanar features, including a $\\sim10\\rm~kpc$ diameter ``shell-like'' feature and a ``blob'' reaching a projected distance of $\\sim25\\rm~kpc$ from the galactic disk. The bulk of the X-ray emission in the halo has a scale height of $\\sim$1.5 kpc and can be characterized by a two-temperature optically thin thermal plasma with temperatures of $\\sim$ 0.2 and 0.6 keV and a total 0.3-2 keV luminosity of $\\sim3.5\\times10^{39}\\rm~ergs~s^{-1}$. The high-resolution, multi-wavelength data reveal the presence of several extraplanar features around the disk, which appear to be associated with the in-disk star formation. We suggest that hot gas produced with different levels of mass loading can have different temperatures, which may explain the characteristic temperatures of hot gas in the halo. We have obtained a sub-galactic scale X-ray-intensity-star formation relation, which is consistent with the integrated version in other star forming galaxies. ",
        "introduction": "It is generally believed that extraplanar diffuse X-ray emission detected around actively star-forming disk galaxies arises primarily from the so-called disk/halo interaction (Cox 2005 and references therein; Wang et al. 1995, 2001, 2003; Strickland et al. 2004a,b). Existing observational work has concentrated on the overall correlation between the extraplanar X-ray emission and the integrated properties of various tracers of star formation in such galaxies, especially those with nuclear starbursts. Systematic studies with {\\sl Chandra} data (Strickland et al.~2004a,b) have revealed the correlation between the presence of extraplanar hot gas and warm ionized gas, which is interpreted as a result of the interaction between star formation-driven outflow of hot gas and the cooler ambient interstellar medium (ISM). The integrated extraplanar X-ray luminosity is found to be quasi-linearly correlated with the star formation rate, implying that the mechanical energy release from supernova explosions (SNe) and massive stellar winds is converted into radiative energy with a certain efficiency. Similar correlations have been demonstrated in an {\\sl XMM-Newton} survey of edge-on galaxies with relatively normal or disk-wide star forming activity (T$\\rm\\ddot{u}$llmann et al.~2006a,b). In addition, there might be an energy input threshold to form a multi-phase halo (Rossa \\& Dettmar 2003; Dahlem et al. 2006). In general, little is yet known about the physical state of the extraplanar hot gas (e.g., the pressure and volume filling factor) and its relationship with other ISM components in the host halo.  The edge-on star-forming galaxy NGC 5775 (see Table \\ref{basicpara} for basic parameters) is well-suited for a detailed study of the disk/halo interaction. Unlike a nuclear starburst galaxy, the star formation in NGC 5775 is spread out across the galactic disk. This allows us to explore the general disk/halo interaction without it being overwhelmed by a powerful galactic nuclear superwind. Indeed, this galaxy is one of the rare cases in which extraplanar emission has been detected from essentially all diffuse interstellar components (e.g., Lee et al. 2001). A number of extraplanar features are resolved and found to exhibit coherent structures in \\ion{H}{1}, radio continuum and H$\\alpha$. NGC 5775 is interacting with its companion NGC 5774, as they are bridged by \\ion{H}{1} gas (Irwin 1994). In addition, the kinematics of the extraplanar diffuse ionized gas has been carefully studied, showing a  vertical gradient in azimuthal velocity, which is steeper than what is predicted by a ballistic model of a disk/halo flow (Heald et al. 2006; Fraternali et al. 2007). Hydrodynamic effects may thus be important (e.g., Barnab$\\grave{e}$ et al. 2006), which may result from the interaction of the flow with pre-existing halo gas. Previous X-ray studies based on {\\sl ROSAT} and {\\sl XMM-Newton} observations of NGC~5775 (Lee et al. 2001; T$\\rm\\ddot{u}$llmann et al.~2006a) have revealed the presence of apparently diffuse X-ray emission which is generally spatially related to the extraplanar features observed at other wavelengths. However, these X-ray data of limited spatial resolution do not allow for a clean detection and removal of discrete sources or for a detailed morphological characterization of the extraplanar hot gas and its connection to other components of this galaxy.  In this paper, we present a {\\sl Chandra} study of NGC~5775 and confront the results with multi-wavelength data obtained from {\\sl Spitzer} and {\\sl HST} as well as various ground-based observations. These data allow us to examine various spatially resolved properties of the disk/halo interaction. We describe the observations and data reduction in \\S~\\ref{sec:datareduction} and present our analysis and results in \\S~\\ref{sec:an_re}. We discuss the implications of our results in \\S~\\ref{sec:discussion} and summarize our findings in \\S~\\ref{sec:summary}. Throughout this work errors are quoted at the 90\\% confidence level. ",
        "conclusions": "\\label{sec:discussion}  \\subsection{Connection of extraplanar hot gas to disk star formation}\\label{sec:XFeature}  In Fig.~\\ref{fig:X_multi} (also see Figs.~\\ref{fig:pointsrcblob} and \\ref{fig:x_hst.ps}), we show the comparative morphology of the soft X-ray emission with images in other wavebands. The in-disk soft X-ray emission generally shows spatial similarities with star formation regions as traced by H$\\alpha$, IR and/or radio emission. It is not surprising to have such similarities, since diffuse hot gas and unresolved young stellar objects, both closely related to star-forming regions, contribute to the soft X-ray emission. The extraplanar X-ray emission arises from truly diffuse hot gas. Thus, in general, we expect a spatial relationship between the extraplanar hot gas and underlying star-forming regions if the extraplanar gas is supplied and regulated by underlying star-forming processes. Differences in spatial resolution for the plots of Fig.~\\ref{fig:X_multi} make specific connections to underlying star forming regions more difficult to trace.  An examination of Fig.~\\ref{fig:x_hst.ps}, however, indeed does show some correlation. For example, as indicated in Sect.~\\ref{subsec:diffuse}, the in-disk NW giant \\ion{H}{2} complex appears connected to the extraplanar X-ray emission and the central X-ray contours are wider around the central star forming regions of the galaxy. The giant \\ion{H}{2} in-disk complex on the far SE end of the disk does not appear to have an extra-planar X-ray counterpart in this figure, but the X-ray results of T{\\\"u}llmann et al. (2006) reveal extraplanar X-ray emission associated with a string of \\ion{H}{2} regions off the plane on this side as well.  Again, this argues for a connection between extraplanar X-ray emission and in-disk star formation.  \\begin{figure*}[!htb] \\centerline{ \\epsfig{figure=fig9.eps,width=1.0\\textwidth,angle=0} } \\caption{{\\sl Chandra} ACIS-S 0.3-1.5 keV intensity contours of the diffuse emission overlaid on the H$\\alpha$ (a; Collins et al.~2000), {\\sl Spitzer} 8 ${\\rm \\mu}$m(b), 20 cm radio continuum (c; Lee et al.~2001) and \\ion{H}{1} column density (d; Irwin 1994) of NGC~5775. The X-ray intensity is the same as in Fig.~\\ref{fig:pointsrcblob}. The contours are at (2.0, 2.5, 3.5, 5, 8, 15, 25) ${\\times}10^{-3}{\\rm~cts~s^{-1}~arcmin^{-2}}$. The ``X-shape'' feature discussed in \\S\\ref{sec:XFeature} is illustrated in (a); the extraplanar regions used to calculated soft X-ray intensities in Fig. \\ref{fig:X_IR} are plotted in (b) with the region numbers marked on the side; the positions of the detected sources are marked with {\\sl crosses} in (c); and the F features of Lee et al. (2001) are marked in (d).} \\label{fig:X_multi} \\end{figure*}  On larger scales, several extraplanar features are resolved with the superb resolution of {\\sl Chandra} and other multi-wavelength observations (Fig. \\ref{fig:X_multi}). Here we discuss the three most prominent plume-like features, originally identified in \\ion{H}{1} observations (Irwin 1994; Lee et al. 2001; labelled F1, F2, and F3 in Fig.~\\ref{fig:X_multi}) and later found to have counterparts or extensions in other wavebands (e.g. Lee et al. 2001; Brar et al. 2003). F1 and F2, being opposite to each other with respect to the disk, apparently protrude from an intense star-forming region, within which the ULX is located. On the other hand, in the disk region underneath F3, all components of the ISM, except for \\ion{H}{1}, show a local minimum, indicating a low current star formation rate (Lee et al. 2001), possibly as a result of a previous star forming episode which has produced an in-disk disruption. F3 seems to be associated with an extraplanar H$\\alpha$ filament, within which several \\ion{H}{2} regions reside (Lee et al.~2001). Together with the disk-halo connections discussed above (cf. Fig.~\\ref{fig:x_hst.ps}), they collectively indicate that local star forming regions away from the nucleus affect halo characteristics above and below the plane.  In addition, however, these features are located near the boundaries of the so-called ``X-shape'' which designates conic structures (marked in Fig. \\ref{fig:X_multi}a) and was claimed to be present from previous X-ray (T$\\rm\\ddot{u}$llmann et al. 2006a) and radio (T$\\rm\\ddot{u}$llmann et al. 2000) observations. The shell-like X-ray feature to the SW of the major axis also seems to be part of the X-shape structure, representing an outflow from the galactic center region to a vertical height of $\\sim$1\\farcm5 ($\\sim$11 kpc) (Fig. \\ref{fig:X_multi}).  Thus, NGC~5775 appears to be similar to the starburst galaxy, NGC~253 (e.g. Vogler \\& Pietsch 1999; Pietsch et al. 2001; Strickland et al. 2000) with its nuclear conical outflow, yet it also has smaller-scale outflows associated with specific in-disk regions.  As a result, it is not always straightfoward to interpret the origin of specific extraplanar features.  An example is the large northeastern blob shown in Fig. \\ref{fig:optical} and Fig. \\ref{fig:pointsrcblob} which seems to have a radio continuum counterpart. It is roughly above the F3 disturbance but could also be a result of the nuclear outflow. It is located about twice as far away from the galactic disk than the shell-like feature on the opposite side.  \\subsection{The sub-galactic scale ``X-SF'' relation}\\label{sec:X-SF}  The relation between the integrated extraplanar soft X-ray intensity and the in-disk star formation intensity (hereafter referred to as an ``X-SF'' relation) has been demonstrated for active star-forming galaxies (Strickland et al.~2004b; T$\\rm\\ddot{u}$llmann et al.~2006b). Fig. \\ref{fig:X_IR} compares this relation with a sub-galactic ($\\sim\\rm kpc$) scale version derived from our current work.  \\begin{figure}[!htb] \\centerline{ \\epsfig{figure=fig10.eps,width=0.7\\textwidth,angle=0} } \\caption{In-disk ($|z|\\leq$~0\\farcm2) star formation intensity calculated from 8 $\\rm \\mu m$ intensity versus extraplanar (-1\\farcm5~$\\leq z\\leq$~-0\\farcm2) soft X-ray intensity. The soft X-ray intensity is obtained in 0.3-1.5 keV band and converted to 0.3-2 keV band for comparison with the Strickland et al. (2004a) sample. The binning slices of the extraplanar regions are shown in Fig. \\ref{fig:X_multi}b, while the region numbers are marked on the side. The star formation intensity is calculated just underneath these slices with $|z|\\leq$~0\\farcm2. See text for details. Filled circles are measurements from this work. Open diamond symbols denote measurements from the Strickland et al. sample galaxies, in which the X-ray luminosity of NGC 4244 and NGC 6503 are estimated with upper limits and NGC 4945 is excluded due to its high foreground absorption. Strickland et al.'s sample has a typical error in the y axis of less than $10^{36}\\rm~ergs~s^{-1}~kpc^{-2}$. Only NGC 1482 (the galaxy with the second highest extraplanar X-ray intensity in their sample) has a larger error of about $3\\times10^{36}\\rm~ergs~s^{-1}~kpc^{-2}$. The two dashed lines show the approximate range of the data, which is $I_{0.3-2.0{\\rm~keV}}(10^{38}{\\rm~ergs~s^{-1}~kpc^{-2}})\\sim(2-8)~I_{SF}~({\\rm M_\\odot~yr^{-1}~kpc^{-2}})$. } \\label{fig:X_IR} \\end{figure}  Since H$\\alpha$ emission may suffer from strong absorption in this edge-on case, we estimate the star formation rate from the 8 $\\rm \\mu m$ intensity using Wu et al. (2005)'s relation. We then convert the X-ray intensity to intrinsic luminosity using the halo spectral model (Table \\ref{specfitpara}). Different in-disk regions should have different line-of-sight depth, so when calculating in-disk area we assume the disk has a circular shape with diameter of $D_{25}$. To correct for projection, the inclination angle is adopted to be $\\sim86^\\circ$ as listed in Table \\ref{basicpara}. Note that, although we use only the X-ray data in the southern halo to minimize absorption effects, Fig. \\ref{fig:X_IR} assumes the two sides have the same intrinsic luminosity.  The consistency of the ``X-SF'' relations found inside NGC~5775 and inferred from integrated galactic properties indicates that the efficiency of converting stellar feedback to the extraplanar hot gas does not change significantly from $\\sim$kpc to galactic scales. The threshold for the presence of the extraplanar X-ray emission on these scales seems to be $\\sim 4\\times10^{-3}~\\rm M_{\\odot~}yr^{-1}~kpc^{-2}$, consistent with previous findings based on galaxy samples (Strickland et al. 2004b; T$\\rm\\ddot{u}$llmann et al. 2006b). Fig.~\\ref{fig:X_IR} also suggests that the conical nuclear outflow is not strongly perturbing the extraplanar X-ray emission that is above and below specific in-disk regions.  \\subsection{Physical and dynamic states of the extraplanar hot gas}\\label{sec:Hotgas_physical}  In \\S\\ref{subsec:diffuse}, we have shown that the halo spectrum consists of two hot gas components. There are now at least three ISM components co-existing in the halo of NGC 5775: the warm ionized gas traced by H$\\alpha$ emission, and the two hot gas components. Inferred from H$\\alpha$ and optical spectroscopic observation, the warm ionized gas has a temperature of $\\sim10^4~\\rm K$ (Collins \\& Rand 2001) and density of $\\sim0.1~\\rm cm^{-3}$ (Collins et al. 2000). The warm gas likely has a low filling factor compared with the hot gas if a thermal pressure balance is established between the warm and hot gas. Strickland et al. (2002) presented several disk/halo interaction models to interpret the multi-phase halo in NGC 253. In their models, H$\\alpha$-emitting warm ionized gas can be produced through interaction between the outflow and the pre-existing gas in the halo, from the cooling of swept up outer shells, or simply by being dragged from the disk. Soft X-ray emission can arise from shock heated halo gas or from the conductive interfaces or turbulent mixing layers between cold halo clouds and the tenuous, low emissivity, and hot superwind.  Some parameters of the two hot gas components are listed in Table \\ref{halopara}. These parameters are inferred from the spectral fitting parameters listed in Table \\ref{specfitpara}, with the assumption that the total filling factor of the hot gas is $\\sim1$ and the two hot gas components reach the pressure balance with each other. We expect that hot gas is produced mainly in the disk and then flows into the halo. The temperature of the gas depends on both the efficiencies of the mass loading as well as the radiative cooling and other energy dissipation processes (e.g., via cosmic rays and kinetic energy). Because the mass efficiency of star formation is only $\\sim 10\\%$ and the remaining molecular gas is dispersed on a time scale of $t \\sim 10^7$ yrs, substantial massloading to hot gas is expected in recent star forming regions. With the presence of the dense gas, radiative and mechanical cooling can also be important. In contrast, the heating of gas by core-collapsed SNe can be more efficient in older star formation regions ($10^7 \\lesssim t \\lesssim 5 \\times 10^7$ yrs; the upper limit corresponds to the lifetime of an 8$M_\\odot$ star). Therefore, gas in an ongoing or recent star forming region can be cooler --- explaining the low-temperature component of the hot gas ---  than in those older star formation regions, which are probably responsible for the hotter component. The lower temperature hot gas can also be produced in the halo at interfaces between the hot gas outflow and the pre-existing cooler gas. Radiative cooling generally has so long time scale for both hot gas components that it cannot be very important in the halo (Table \\ref{halopara}), at least in the vicinity of the galactic disk.  The rim-brightened soft X-ray-emitting bubble likely represents the reverse-shocked outflow in the interior of an outer cool shell (Strickland et al. 2002). This cool gas shell could be formed in or near the disk, and then expand into the halo driven by the high pressure interior. During the expansion of this cool shell, the magnetic field may play an important role in preventing the Rayleigh-Taylor instability which would break the shell, given that the magnetic pressure may be comparable with the thermal pressure of the halo hot gas (see below).  The northeast blob may be compared with a similar feature ``ridge'' detected in M 82 (Lehnert, Heckman \\& Weaver 1999). The blob, elongated along the major axis instead of perpendicular to the major axis, is about twice as high as the M 82 ridge. Given that NGC 5775 does not have as strong a superwind as detected in M 82, and there is no clear sign of a cool gas counterpart for the blob, we speculate that it may represent a bulk of outflowing hot gas --- an evolved outflowing wind bubble --- generated in an early episode of massive star formation in the nuclear region of the galaxy. The relatively high concentration of the hot gas at such a large off-disk distance may be due to the confinement of a magnetic field loop, which would still be anchored in the disk. A magnetic field must be present because of the observed radio continuum counterpart. Data of higher quality is needed to truly understand the nature of the blob.  Although a magneto-hydrodynamic model of the nuclear outflow is beyond the scope of this paper, we can still make some general statements about halo conditions. For example, the magnetic pressure in the halo can be estimated from ${\\rm P}_B\\,\\sim10^{-12}~[B/(5~\\mu G)]^2~f_B^{-4/7}\\rm~dyn~cm^{-2}$, where $B$ is the magnetic field strength in units of $\\mu G$ and $f_B$ is the volume filling factor of the magnetic field.  In this relation, the magnetic field strength is assumed to have been determined from multi-frequency radio continuum maps, adopting an approximate energy equipartition between relativistic particles and the magnetic field. Using the radio continuum and spectral index maps of Duric et al. (1998) and an appropriate geometric model (see Irwin et al. 1999 and Irwin \\& Saikia 2003 for examples)  we find a typical halo magnetic field strength of $5~\\mu G$, implying  ${\\rm P}_B\\,\\sim10^{-12}$ for a unity filling factor, or higher if the filling factor is less than unity. The hot gas pressure in the halo is $\\sim10^{-12}\\rm~dyn~cm^{-2}$, derived from the hot gas parameters listed in Table \\ref{halopara}. Although we expect variation in these quantities and filling factors are unknown, it is reasonable to conclude that the magnetic field will be dynamically important in any outflow model.  Parker instabilities may also be present. Observations of the current field configuration shows roughly vertical fields across the disk-halo interface (see T{\\\"ull}mann et al. 2000) although the field configuration far from the plane has yet to be observationally determined.  The feedback energy of the young and old stellar populations can be estimated from the IR (Heckman, Armus \\& Miley 1990) and K band (Mannucci et al. 2005) luminosity, which is $\\sim3.7\\times10^{42}\\rm~ergs~s^{-1}$ in total. It is not easy to estimate the flow velocity of the hot gas, but we can adopt the sound speed in the higher temperature hot gas as an upper limit, which is $\\sim400\\rm ~km~s^{-1}$. Using this velocity, the flowing time in the soft X-ray emitting halo is $>2.5\\times10^7\\rm~yr$ (adopting an extension of the soft X-ray emitting halo to be 10 kpc). This corresponds to a total energy feedback $>3\\times10^{57}\\rm~ergs$, enough to support the soft X-ray emission and the energy of the hot gas.    We have conducted a detailed {\\sl Chandra} study of NGC~5775 and compared the results to multi-wavelength data obtained from {\\sl Spitzer} and {\\sl HST} as well as high resolution ground-based H$\\alpha$, radio continuum and \\ion{H}{1} data. The main results are as follows:  Ninety-nine discrete X-ray sources are detected in or near NGC 5775, 27 within the $D_{25}$ region, including a ULX with an absorbed (intrinsic) 0.3-7 keV luminosity of about $2.4~(6.5)\\times10^{40}\\rm~ergs~s^{-1}$.  Large-scale diffuse X-ray emission is detected to an extent nearly 10 kpc above the galactic disk in an extended halo region. There are several prominent plume-like features detected in the halo indicating hot gas outflows from the disk, some of which are clearly associated with the in-disk star formation regions. In particular, a huge shell-like feature appears in the southern halo, and has a radius of nearly 5 kpc. The soft X-ray luminosity of this feature is $\\sim6.7\\times10^{38}~\\rm ergs~s^{-1}$ in the 0.3-2 keV band. On the opposite side, a large-scale ``blob'' is detected with a maximum projected distance $\\sim25\\rm~kpc$ from the galactic plane.  Our X-ray spectral analysis reveals that the halo hot gas can be decomposed into two characteristic components, with temperatures of about $0.2\\rm~keV$ and $0.6\\rm~keV$. The total luminosity of the halo hot gas is $\\sim3.5\\times10^{39}\\rm~ergs~s^{-1}$.  The multi-phase halo is energetically consistent with stellar feedback. Current star formation with more mass loading may produce the lower temperature hot gas component, while recent and past star formation with less mass loading may be responsible for the higher temperature component. The soft X-ray shell-like feature likely represents an outflowing wind bubble from the galactic central region. The northeast blob may be an evolved version of such a bubble that is still confined by strong magnetic field. We find that the magnetic pressure is comparable to the hot gas pressure, and will therefore play an important role in the evolution of a multi-phase halo.  We have obtained an ``X-ray-star formation'' relation for sub-galactic scale star formation regions in this disk-wide star forming galaxy, and this relation is consistent with that inferred from the integrated galactic properties."
    },
    "0807/0807.4786.txt": {
        "abstract": "{ X-ray colour-colour (CD) and colour-intensity (HID) diagrams are powerful tools that allow the investigation of spectral variability without the assumption of any spectral model. These diagrams have been used extensively in low-mass X-ray binaries and black-hole candidates but very few applications are found for high-mass X-ray binaries. We have investigated the spectral and timing variability of four accreting  X-ray pulsars with Be-type companions during major X-ray outbursts.  The aim is to define source states based on the properties (noise components) of the aperiodic variability in correlation with the position in the colour-colour diagram. Different spectral states were defined according to the value of the X-ray colours and flux. Transient Be/X-ray binaries exhibit two branches in their colour-colour and colour-intensity diagrams: the horizontal branch corresponds to a low-intensity state and shows the larger fractional $rms$, similar to the the island state in atolls and horizontal branch in Z sources;  the diagonal branch corresponds to a high-intensity state, in which the source spends about 75\\% of the total duration of the outburst. Despite the complexity of the  power spectra due to the peaks of the pulse period and its harmonics, the aperiodic variability of Be/X-ray  binaries can be described with a relatively low number of Lorentzian components. Some of these components can be associated with the same type of noise as that seen in  low-mass X-ray binaries, although the characteristic frequencies are about  one order of magnitude lower.   The analysis of the CD/HID and power spectra results in two different types of Be/X.  While in \\vcas, \\ks\\ and \\exo\\ the hard colour  decreases as the count rate decreases, in \\bq\\ it increases.  The pattern traced by \\bq\\ then results in a Z shaped track, similar to the low-mass $Z$ sources, without the flaring branch. In contrast, the horizontal branch in \\vcas, \\ks\\ and \\exo\\ corresponds to a low/soft state, not seen in other types of X-ray binaries. The noise at very low frequencies follows a power law in \\bq\\ (like in LMXB Z) and it is flat-topped in \\vcas, \\ks\\ and \\exo\\ (like in LMXB atoll). \\bq\\ shows a noise component coupled with the periodic variability that it is not seen in any of the other three sources.  } ",
        "introduction": "The colour-colour (CD) and hardness-intensity (HID) diagrams are very useful tools to investigate the spectral variations of an X-ray source. They are model- and instrument-independent and reflect the intrinsic properties of the system. The first application of this type of plots was the separation of different types of galactic X-ray sources and the classification of X-ray populations.  The spectral hardness was defined by the ratio of counting rates in each of the proportional counters of the {\\em MIT OSO-7} experiment \\citep{mark77,rapp77}. \\citet{whit84} made use of the X-ray colour-colour diagrams, obtained from {\\em HEAO 1} observations, to study the spectral states among the different types of X-ray binaries, including black-hole candidates.  CDs became widely used during the EXOSAT era, which also marks the introduction of HIDs \\citep{pried86,klis87a,hasi87}, where \"hardness\" refers to the ratio of the count rate in two energy bands. Before EXOSAT, CDs contained sources of different origin. The novelty of the EXOSAT CDs and HIDs was that they were applied to individual sources. CD/HIDs became a useful way to investigate the rapid aperiodic variability in correlation with spectral states, as defined by the spectral hardness, of individual low-mass X-ray binaries and black-hole candidates \\citep{schu89,hasi89}.  With the advent of RXTE the use of CDs and HIDs have become part of the standard data analysis. RXTE has provided an unprecedentedly large database of X-ray observations with exceptionally good timing and moderate spectral resolution.  Together with the functional form of the variability components (\"noise\"), CD and HID have become an essential tool in developing a whole new phenomenology on the spectral and timing properties of X-ray binaries by introducing the notion of source states.  A state is defined by the appearance of a spectral (i.e. power-law, blackbody) or variability component (i.e. Lorentzian) associated with a particular and well-defined position of the source in the CD/HID. A transition between states is assumed when the relative strength of the spectral or variability components varies and the source motion in the CD/HID changes direction \\citep{klis06}.  While there are numerous references in the literature on the application of CD/HID analysis on low-mass X-ray binaries \\citep[][and references therein]{klis06} and black-hole systems \\citep[see e.g.][]{bell05},  very little work of this type has been done on high-mass X-ray binaries. \\citet{bell90} performed an aperiodic variability study of 12 high-mass X-ray binaries from EXOSAT archive data. The sample included \\exo\\ and \\bq. However, the few observations analysed of \\bq\\ did not correspond to a major X-ray outburst, while the data of \\exo\\ covered the decay of the 1985 outburst only partially. The other 10 sources were supergiant X-ray binaries in bright states. The vast majority of the rapid aperiodic variability studies in accreting X-ray pulsars have concentrated on the detection of QPOs \\citep{ange89,jern00,qu05}; see also \\citet[][and references therein]{shir02}  Massive X-ray binaries are classified according to the luminosity class of the optical component into Be/X-ray binaries (dwarf or subgiant) and supergiant (luminosity class I-II) X-ray binaries. The former tend to be transient systems while the latter are persistent sources. Be stars are non-supergiant fast-rotating B-type and luminosity class III-V stars which at some point of their lives have shown spectral lines in emission. In the infrared they are brighter than their non-emitting  counterparts of the same spectral type. The line emission and infrared excess originate in extended circumstellar envelopes of ionized gas surrounding the equator of the B star. When the Be star takes part on a binary system where the companion is a neutron star, then the system is referred to as a Be/X-ray binary.  Be/X-ray binaries are variable on time scales from seconds to years. The fastest variability is found in the X-ray band. Virtually all the Be/X-ray binaries with identified optical counterparts are X-ray pulsars. Pulse periods cover the range 1--10$^3$ s. On longer time scales (months to years) the variability is also apparent in the optical and infrared bands and it is attributed to structural changes of the circumstellar disc.  The long-term X-ray variability of the transient Be/X-ray binaries is characterised by two type of outbursting activity  \\begin{itemize}  \\item Type I outbursts. These are regular and (quasi)periodic outbursts, normally peaking at or close to periastron passage of the neutron star. They are short-lived, i.e., tend to cover a relatively small fraction of the orbital period (typically 0.2-0.3 $P_{\\rm orb}$). The X-ray flux increases by about one or two order of magnitude with respect to the pre-outburst state, reaching peak luminosities $L_x \\leq 10^{37}$ erg s$^{-1}$.  \\item Type II outbursts represent major increases of the X-ray flux, $10^{3}-10^{4}$ times that at quiescence. They reach the Eddington luminosity for a neutron star ($\\sim 10^{38}$ erg s$^{-1}$) and become the brightest objects of the X-ray sky. They do not show any preferred orbital phase and last for a large fraction of an orbital period or even for several orbital periods. The formation of an accretion disc during Type II outbursts  \\citep{kris83,motc91,haya04,wils08} may occur. The discovery of quasi-periodic oscillations in some systems \\citep{ange89,take94,fing96} would support this scenario. The presence of an accretion disk also helps explain the large and steady spin-up rates seen during the giant outbursts, which are difficult to account for by means of direct accretion.  \\end{itemize}  In this work we present a systematic study of the aperiodic variability of four Be/X-ray binaries during type II outbursts. First, we give a summary of the relevant information of the sources. Sect.~\\ref{obs} summarises the observations and the characteristics of the instruments used in the analysis. Sect.~\\ref{red} refers to the reduction of the data and describes how the CD/HID and  power spectra were obtained. A detailed description of the results is given in Sect.~\\ref{res}. This section is divided into three subsections describing the outburst profiles, the colour analysis and the aperiodic variability. Our interpretation of the results and the implications of this work are discussed in Sect.~\\ref{disc}.    Final conclusions are drawn in Sect.~\\ref{con}.   %------------------------------------------------------------- \\begin{table*} \\caption{Observations and definition of source states.} \\label{reg} \\centering \\begin{tabular}{lcccccccc} \\hline\\hline Source&Branch$^a$ &Time range\t&Proposal &On-source\t&$<SC>^b$ &$<HC>^b$&Flux$^{c}$\t&Luminosity$^{d}$ \t \\\\ state &\t&MJD\t     \t&ID\t&time (s) \t& \t&\t&\t\t&\t\\\\ \\hline \\hline \\noalign{\\smallskip} \\multicolumn{9}{c}{{\\bf \\vcas}}\\\\ \\hline \\noalign{\\smallskip} Rise1\t&DB\t&53254.0676-53256.4660 &90089  &3872\t&2.08$\\pm$0.05  &0.71$\\pm$0.03  &0.95\t&7.5  \\\\ Rise2\t&DB  \t&53260.1495-53265.0921 &90089  &3536\t&1.91$\\pm$0.05\t&0.81$\\pm$0.04  &1.5\t&11.8   \\\\ Peak\t&DB  \t&53267.1686-53272.9721 &90089  &29744   &1.87$\\pm$0.01  &0.89$\\pm$0.01  &1.4\t&11.0   \\\\ Decay1\t&DB\t&53274.8282-53280.7687 &90089  &9712\t&1.95$\\pm$0.02  &0.85$\\pm$0.02  &1.0\t&7.9  \\\\ Decay2\t&DB\t&53278.6230-53280.7687 &90089  &6128    &2.03$\\pm$0.04  &0.70$\\pm$0.06  &0.8\t&6.0  \\\\ Decay3\t&DB\t&53282.5715-53282.6673 &90014  &5536\t&2.09$\\pm$0.01\t&0.62$\\pm$0.01\t&0.6\t&5.0 \\\\ Decay4\t&HB\t&53284.6063-53304.5730 &90014  &32500\t&1.99$\\pm$0.04\t&0.57$\\pm$0.03\t&0.11\t&0.9  \\\\ \\hline \\noalign{\\smallskip} \\multicolumn{9}{c}{{\\bf \\ks}}\\\\ \\hline \\noalign{\\smallskip} Rise0\t&HB\t&51874.1502-51878.3595\t&50425\t&16720\t&1.57$\\pm$0.07\t&1.24$\\pm$0.02\t&0.01\t&0.12\\\\ Rise1\t&DB\t&51883.8286-51892.2963\t&50425\t&14060\t&1.68$\\pm$0.04\t&1.13$\\pm$0.05\t&0.09\t&1.0 \\\\ Rise2\t&DB\t&51896.1987-51910.9382\t&50425\t&17744\t&1.55$\\pm$0.03\t&1.14$\\pm$0.02\t&0.25\t&2.3 \\\\ Rise3\t&DB\t&51914.1750-51932.3484\t&50425\t&10080\t&1.49$\\pm$0.03\t&1.16$\\pm$0.03\t&0.36\t&4.3 \\\\ Peak\t&DB\t&51941.6795-51968.5845\t&50425\t&23344\t&1.40$\\pm$0.02\t&1.20$\\pm$0.02\t&0.53\t&6.3 \\\\ Decay1\t&DB\t&51970.2073-51990.1352\t&60402\t&15536\t&1.51$\\pm$0.03\t&1.19$\\pm$0.02\t&0.39\t&4.7 \\\\ Decay2\t&DB\t&51992.1719-52010.8263\t&60402\t&9888\t&1.59$\\pm$0.02\t&1.16$\\pm$0.02\t&0.22\t&2.6 \\\\ Decay3\t&DB\t&52014.3856-52050.1873\t&60402\t&9984\t&1.65$\\pm$0.04\t&1.05$\\pm$0.07\t&0.11\t&1.3 \\\\ Decay4\t&HB\t&52052.9510-52078.1324\t&60402\t&14720\t&1.32$\\pm$0.08\t&1.11$\\pm$0.04\t&0.008\t&0.1 \\\\ \\hline \\noalign{\\smallskip} \\multicolumn{9}{c}{{\\bf \\exo}}\\\\ \\hline \\noalign{\\smallskip} Rise1\t&DB\t&53908.6156-53924.7247\t&91089 \\& 92067&43200\t&1.39$\\pm$0.07\t&1.09$\\pm$0.02\t&1.4\t&8.4\t\\\\ Rise2\t&DB\t&53926.6282-53943.2576\t&91089\t&26512\t&1.26$\\pm$0.02\t&1.12$\\pm$0.01\t&2.1\t&12.7\t\\\\ Peak\t&DB\t&53944.0532-53975.4312\t&91089\t&35536\t&1.23$\\pm$0.02\t&1.13$\\pm$0.02\t&2.4\t&14.5\t\\\\ Decay1\t&DB\t&53976.1186-53998.6782\t&91089\t&25168\t&1.32$\\pm$0.05\t&1.09$\\pm$0.03\t&1.6\t&9.7\t\\\\ Decay2\t&DB\t&53999.5902-54015.4534\t&92067\t&8656\t&1.47$\\pm$0.05\t&0.99$\\pm$0.02\t&0.57\t&3.4\t\\\\ Decay3\t&DB\t&54016.5928-54030.9436\t&92067\t&10912\t&1.43$\\pm$0.03\t&1.02$\\pm$0.04\t&0.11\t&0.69\t\\\\ Decay4\t&HB\t&54031.7791-54047.9337\t&92067\t&16144\t&1.41$\\pm$0.03\t&0.96$\\pm$0.03\t&0.06\t&0.37\t\\\\ \\hline \\noalign{\\smallskip} \\multicolumn{9}{c}{{\\bf \\bq}}\\\\ \\hline \\noalign{\\smallskip} Rise1\t&DB  \t&53340.2869-53346.7436 &90089  \t&48752\t&2.18$\\pm$0.02  &0.89$\\pm$0.03  &2.3\t&13.5  \\\\ Rise2\t&DB  \t&53352.7706-53361.1469 &90089  \t&32720\t&2.06$\\pm$0.03  &0.73$\\pm$0.02  &5.0\t&29.3  \\\\ Peak\t&DB\t&53364.9113-53369.5965 &90089 \\& 90427&13040 &1.95$\\pm$0.02  &0.66$\\pm$0.01  &5.5&32.2  \\\\ Decay1\t&DB\t&53374.0267-53381.6610 &90427 \\& 90014&23504 &2.04$\\pm$0.02  &0.74$\\pm$0.01  &4.2&24.6  \\\\ Decay2\t&DB\t&53385.8724-53395.4413 &90014\t&38672\t&2.11$\\pm$0.03  &0.83$\\pm$0.03  &2.6\t&15.2 \\\\ Decay3\t&DB\t&53398.4560-53414.0726 &90014\t&21296\t&2.20$\\pm$0.04  &0.94$\\pm$0.03  &1.5\t&8.8 \\\\ Decay4\t&HB\t&53416.0726-53424.4139 &90014\t&20160\t&2.05$\\pm$0.15  &1.13$\\pm$0.05  &0.29\t&1.7 \\\\ \\hline\\hline  \\noalign{\\smallskip} \\multicolumn{9}{l}{$a$: HB: Low-intensity (Horizontal) Branch; DB: high-intensity (Diagonal) branch}\\\\ \\multicolumn{9}{l}{$b$: Normalised to the Crab. Average Crab values are $<SC>_{\\rm Crab}=0.34$ and $<HC>_{\\rm Crab}=0.47$} \\\\ \\multicolumn{9}{l}{$c$: $\\times 10^{-8}$ erg cm$^{-2}$ s$^{-1}$ in the 3--30 keV band} \\\\ \\multicolumn{9}{l}{$d$: $\\times 10^{37}$ erg s$^{-1}$ in the 3--30 keV band} \\\\ \\end{tabular} \\end{table*} %-------------------------------------------------------------  %---------------------------------outbprof.eps---------------------------- \\begin{figure*} \\centering %   \\includegraphics[width=8cm]{./figures/lcoutb.eps} \\\\ \\includegraphics[width=16cm]{./10021f1.eps} \\caption{The outburst profiles as seen by ASM RXTE. Vertical lines mark the time of the RXTE PCA observations. Different colours distinguish the spectral regions. Bin size is 1 day. 1 Crab corresponds to $\\sim$75 c/s. {\\em See the electronic edition of the Journal for a colour version of this figure}. } \\label{outbprof} \\end{figure*} %------------------------------------------------------------- %----rate_hr_1947.eps-----rate_hr_2030.eps-----rate_hr_0115.eps----rate_hr_0332.eps---------- \\begin{figure*} \\begin{tabular}{cc} \\includegraphics[width=8cm]{./10021f2a.eps} & \\includegraphics[width=8cm]{./10021f2b.eps} \\\\ \\includegraphics[width=8cm]{./10021f2c.eps} & \\includegraphics[width=8cm]{./10021f2d.eps} \\\\ \\end{tabular} \\caption[]{PCA light curves and hardness ratio evolution. Each point represents a 512-s integration. Symbols are as follow: pre-outburst (open black circles, only in \\ks), rise1 (red circles), rise2 (orange squares), rise3 (pink crosses, only in \\ks), peak (black triangle-up), decay1 (violet stars), decay2 (brown diamonds), decay3 (green triangle-down) and decay4 (blue pluses). {\\em See the electronic edition of the Journal for a colour version of this figure}.} \\label{rate_hr} \\end{figure*} %------------------------------------------------------------------------------ ",
        "conclusions": "\\label{con}  We have performed, for the first time, a systematic analysis of the spectral (through the study of hardness ratios) and rapid aperiodic variability of four Be/X-ray pulsars during type II outbursts. Our aim was to characterise the rapid aperiodic variability in correlation with the position in the colour-colour diagram as it has been done for black-hole systems and low-mass X-ray binaries.  The correlated spectral-timing behaviour in Be/X-ray binaries shares a number of similarities with low-mass X-ray binaries, namely, {\\em i)} existence of spectral branches in the CD/HID. At high and intermediate flux the sources move along the diagonal branch; at very low count rate the soft colour decreases, while the hard colour remains fairly constant defining a horizontal branch {\\em ii)} smooth motion, i.e. without jumps, in the CD/HID, {\\em iii)} the low-intensity states are more variable (in terms of fractional rms), {\\em iv)} description of power spectra in terms of a small number of Lorentzian components, {\\em v)} similar BLN namely, broad Lorentzians $Q \\sim 0-0.3$ describing the low-frequency and high-frequency noise ($L_b$, $L_l$, $L_u$) and peaked noise in between ($L_{LF}$) and {\\em vi)} flat-topped noise at the lower frequencies in the horizontal branch that turns into power-law noise in the diagonal branch.  There are also important differences: {\\em i)} different patterns in the CD/HID. Unlike low-mass X-ray binary and black hole systems, the low-intensity state is associated with a soft state (except for \\bq), {\\em ii)} slower motion along the spectral branches, hours to days in LMXBs, weeks to months in HMXBs, {\\em iii)} the characteristic time scales implied by the noise components are about one order of magnitude longer in HMXBs (e.g. mHz QPOs), {\\em iv)} $L_S$ noise in HMXBs and {\\em v)} no apparent correlation between the power spectral parameters (characteristic frequencies, $rms$) and mass accretion rate.  Although we have analysed all Be/X-ray binaries that showed type II outbursts in the period 1996--2007 and had good data coverage and sampling, the number of sources is not large enough to draw definite conclusions on the underlying physical parameters that may explain the differences between the two subgroups. Increasing the number statistics is, however, difficult given the transient nature of their X-ray emission and unpredictability of the type II outbursts.  Even though all four sources analysed in this work are X-ray pulsars, it is not clear whether the high-mass X-ray binaries with supergiant companions show a similar behaviour. In these systems accretion occurs via the strong stellar wind of the massive companion, which contrasts with accretion via a disc, as it is believed to occur in Be/X-ray binaries during outburst. The problem of this type of studies in supergiant X-ray binaries is that they are rather stable in long time scales and do not show a wide range in X-ray flux."
    },
    "0807/0807.4472_arXiv.txt": {
        "abstract": "The $\\gamma$ Cephei system is one of the most closely bound binary planet hosts known to date. The companion ($\\gamma$ Cep B) to the planet-hosting star ($\\gamma$ Cep A) should have truncated any protoplanetary disk around $\\gamma$ Cep A, possibly limiting planet formation in the disk.  We explore this problem by calculating the truncation radii of protoplanetary disk models around $\\gamma$ Cep A to determine whether or not there is sufficient material remaining in the disk to form a planet. We vary the accretion rate and viscosity parameter of the disk models to cover a range of reasonable possibilities for the disks properties and determine that for accretion rates of $\\geq 10^{-7}\\,\\msunperyr$ and low viscosity parameter, sufficient material in gas and solids exist for planet formation via core accretion to be possible.  Disk instability is less favored, as this can only occur in the most massive disk model with an extremely high accretion rate. ",
        "introduction": "Over the last decade, several stellar and substellar companions of exoplanet host stars have been found, mostly in seeing-limited wide field imaging surveys \\citep[e.g.,][]{2007MugrauerSeifahrtNeuhauser,2006Raghavan_etal, 2006Chauvin_etal}. The projected separations are a few tens to several thousand AU, with companion masses $\\sim0.08 - 1.1\\,M_{\\odot}$. \\citet{2007EggenbergerUdry} suggest that exoplanets may be less common in binaries closer than 120\\,AU\\@. This apparent lack of close companions to exoplanet host stars may indicate that planet formation is hampered by the gravitational influence of a close massive companion. However, further investigations are needed to confirm this result.  One of the closest planet host binaries presently known is $\\gamma$\\,Cep. The planet orbits the primary $\\gamma$\\,CepA on a 906 day orbit and has $m\\sin(i)\\sim1.7$\\,$M_{Jup}$ \\citep{1988Campbell_etal,2003Hatzes_etal}.  The stellar companion orbits at a semi-major axis of 20 AU with an eccentricity of 0.4 \\citep{ 2003Hatzes_etal,2007Torres_gcep}.  Given that massive companions can disrupt the protoplanetary disks in which planets form, how feasible is the {\\em in situ} formation of a planet around $\\gamma$ Cep A?  \\citet{2006Haghighipour_gcep} considered the dynamical stability of $\\gamma$ Cep in order to put limits on possibility of the existence of another planet in the system. \\citet{2004Thebault_gcep} considered the problem of planet formation via core accretion in this system using N-body simulations, assuming a density profile consistent with the Minimum Mass Solar Nebula (MMSN) model of \\citet{hayashi}, which is steeper than that produced by viscous accretion disk models. \\citet{2008PaardekooperThebaultMellema} revisited the problem including gas drag and determined that giant planet formation by core accretion is feasible in $\\gamma$ Cep, although they did not address the evolution of the gas in the disk or the truncation radius of the disk in detail. \\citet{2008KleyNelson} assumed a specific disk model and examined the fate of planet cores inserted into the disk. Rather than do a detailed hydrodynamic simulation of a planet embedded in the disk, our objectives are to model the protoplanetary disk in $\\gamma$ Cep as a viscous accretion disk and to explore which disk parameters allow planet formation to occur given that the disk is truncated by the stellar companion.  This analysis is similar to that done in \\citet{HJChd188753} (henceforth Paper 1) for the extremely close triple system HD 188753. Paper 1 concluded that HD 188753 was unlikely to support a disk sufficiently massive to support planet formation.  Indeed, the initial claim of a Jupiter-mass planet in HD 188753 \\citep{HD188753} has since been refuted \\citep{2007Eggenberger_etal}. This does not rule out the possibility that the planet could form around a single star or in a wide binary and then undergo dynamical evolution, such as through close encounters with another star \\citep{2005PortegiesZwart_McMillan,2005Pfahl}, but this is outside the scope of this paper. ",
        "conclusions": "We assessed the feasibility of {\\em in situ} planet formation in $\\gamma$ Cep by examining the properties of a protoplanetary disk around the primary star, given the current orbital parameters of the binary star.  We examined a range of accretion rates and viscosity parameters and determined the truncation radius for each disk model.  We find that $\\gamma$ Cep A can host a truncated disk of sufficient mass to form a giant planet with reasonable accretion rates and viscosity parameters, so {\\em in situ} planet formation is possible. There are sufficient solids in the truncated disk for core accretion to occur for accretion rates higher than $10^{-7}\\,\\msunperyr$ and low values of the viscosity parameter, $\\alpha$.  This is a relatively high accretion rate, which may indicate that giant planet formation must take place very early on, within $10^5 - 10^6$ years, since it appears that accretion rates of T Tauri stars decrease with age \\citep[e.g.][]{2006Sicilia-Aguilar_etal}. Giant planet formation by disk instability is unlikely to have occured in $\\gamma$ Cep.  Disk instability also requires an extremely high accretion rate, $\\sim10^{-4}\\,\\msunperyr$, a rate that is typical of a transient FU Ori phenomenon. On the other hand, this may mean that FU Ori outbursts and giant planet formation by disk instability are correlated.  We have omitted effects such as shock-heating or triggered planet formation by the binary.  Whether these effects inhibit \\citep[e.g.,][]{2000Nelson} or enhance \\citep[e.g.,][]{2006Boss_binary} planet formation depends on whether cooling times are long or short, respectively \\citep[see also][]{2005Mayer_etal}. Low mass disks that are not self-gravitating are not subject to strong shocks, so our results for core accretion are unaffected \\citep{2005Mayer_etal}.  This strengthens our argument that core accretion is the favored mechanism for giant planet formation in close binaries.  The analysis presented here should hold also true for wider binary systems. That is, giant planet formation is possible around any star where the stellar companion has a mass ratio $\\mu\\lesssim0.2$ and has an orbit wider than $\\gamma$ Cep B, with disk instability becoming increasingly feasible in wider systems. The handful of triple systems with exoplanets that have been discovered to date are all hierarchical, where the exoplanet host star and a close stellar pair revolve around a common barycenter. In these cases, the close stellar pair can be treated dynamically as a single object.  The closest planet host triple system with confirmed exoplanets presently known is HD\\,65216\\,A+BC, at $\\sim$250\\,AU \\citep{2007MugrauerSeifahrtNeuhauser}.  Thus, we conclude that there is no barrier to planet formation in the known planet-hosting triple systems.  In \\citet{HJChd188753}, the same analysis was carried out for a hypothesized planet in HD188753, a multiple system which at first glance is only slightly closer than $\\gamma$ Cep: a semi-major axis of 12.3 AU and eccentricity of 0.5. {\\em In situ} planet formation in HD 188753 was ruled out, whereas it is deemed feasible for $\\gamma$ Cep. Somewhere between the parameters of these two systems lies the transition between possiblity and impossibility of planet formation in close binaries.  We will explore this parameter space and put limits on the closeness of planet host binaries in a future paper.  Another interesting case is a system where the massive companion is a white dwarf.  Known white dwarf planet hosts include Gl\\,86\\,B with a projected separation of only 20\\,AU; and HD\\,27442 with a subgiant primary and white dwarf secondary \\citep{2007MugrauerNeuhauserMazeh}. Stellar evolution and mass loss generally widens the orbits of binary stars, so it is likely that these systems were originally much closer \\citep[e.g.][]{2002DebesSigurdsson}. In order to determine whether {\\em in situ} planet formation could have taken place in these sytems, the original orbital configuration must first be determined and the disk truncation radii based on that. We will address this issue in a future paper."
    },
    "0807/0807.4969_arXiv.txt": {
        "abstract": "{One of the most interesting unsolved questions in science today is the question of life on other planets. At the present time it is safe to say that we do not have much of an idea as to whether life is common or exceedingly rare in the universe, and this will probably not be solved for certain unless definitive evidence of extraterrestrial life is found in the future. Our presence on Earth is just as consistent with the hypothesis that life is extremely rare as it is with the hypothesis that it is common, since if there was only one planet with intelligent life, we would find ourselves on it. However, we have more information than this, such as the the surprisingly short length of time it took for life to arise on Earth. Previous authors have analysed this information, concluding that it is evidence that the probability of abiogenesis is moderate ($>$ 13\\% with 95\\% probability) and cannot be extremely small. In this paper I use simple probabilistic model to show that this conclusion was based more on an unintentional assumption than on the data. While the early formation of life on Earth provides some evidence in the direction of life being common, it is far from conclusive, and in particular does not rule out the possibility that abiogenesis has only occurred once in the history of the universe.} ",
        "introduction": "Attempting to make predictions about life elsewhere based on observations about Earth is inherently difficult due to the sample size of 1. It is also fraught with controversial ``anthropic'' considerations \\citep{smolin}. However, there is no reason in principle why it cannot be done. If we use probability theory to model uncertainty \\citep{jaynes}, and data about life on Earth really is uninformative about extraterrestrial life, then probability theory will return wide probability distributions, indicating the large uncertainty.  The surprising fact that life arose on Earth very quickly after its formation \\citep[e.g.][]{early} and at the end of a likely phase of sterilisation due to frequent impacts, has been used to argue that abiogenesis must therefore be easy. \\citet[][hereafter L\\&D]{2002AsBio...2..293L, 2004IAUS..213..259L} have modelled this reasoning with probability theory and concluded with 95\\% confidence (Bayesian posterior probability) that the probability of abiogenesis on an Earth-like planet is greater than 13\\%. This was done by using a model where there was constant hazard (chance of life arising per discrete time interval) $q$. The probability distribution for the time $t_L$ (our $t_L$ corresponds to L\\&D's $\\Delta t_{\\textnormal{biogenesis}})$ at which life arises depends on $q$, and this is also calculated conditional on the fact that $t_L$ must be less than the age of the Earth, to correct for the fact that we couldn't have observed the Earth unless life began. This probability distribution is a likelihood function for $q$ when the actual observed $t_L$ is substituted into it. Combined with a prior distribution for $q$, we can then make inferences about its value.  Whilst it is possible and interesting to calculate such things, the model used by L\\&D contains a flaw that renders the conclusion invalid. Unfortunately, the conclusion quoted above depends on a choice of prior distribution over $q$ that is overconfident and unrealistic as a description of our state of knowledge about abiogenesis. While uniform priors representing ``initial ignorance'' are common in Bayesian Analysis, a uniform prior for an unknown probability such as $q$ is actually quite informative \\citep[][chapter 18]{jaynes}, assigning most of its probability to moderate values of $q$, and ignoring the possibility of extreme values. That the uniform prior is inappropriate can be illustrated using the technique of {\\it elaboration}: are we happy with all of the implied consequences of assuming this prior distribution? For instance, one implication is that $q \\in [0.49,0.51]$ is just as plausible as $q \\in [0, 0.02]$, whereas if we are ignorant, possibilities such as $q \\sim 10^{-6}$ and so on should not be ignored as they almost are by the uniform prior. A more realistic representation of complete prior ignorance would be the improper Haldane prior $\\propto [q(1-q)]^{-1}$ \\citep{jaynes2}, or a modification that removes the singularities at $q=0$ and $q=1$ and makes the prior proper. The Haldane prior corresponds to an improper uniform prior for the ``logit'' $\\log[q/(1-q)]$, representing uncertainty not just about the exact value of $q$ but also about its {\\it order of magnitude}. This paper uses a model that bypasses direct use of $q$ and deals with expected waiting times instead, although its conclusions can be interpreted in the L\\&D framework as well.  The conclusions of \\citet{2002AsBio...2..293L} have been criticised previously on the basis that no observer would ever see a recent abiogenesis due to the large number of intermediate steps required between abiogenesis and the development of intelligent life \\citep{reply}. However, it is still possible to imagine life arising after 5 Gyr on a planet and intelligent observers discovering this at, say, $t=8$ Gyrs. The fact that we are not in this situation could still be considered surprising, and therefore informative \\citep{reply2}. ",
        "conclusions": "The fact that life arose surprisingly early after the formation of the Earth can be used as evidence for the hypothesis that abiogenesis is easy, and hence supports the conclusion that life is common in the universe. However, the evidence is not as conclusive as has been claimed. Specifically, this study has highlighted the fact that knowledge of the early abiogenesis time on Earth is still compatible with the following hypothesis: that life is extraordinarily rare in the universe, perhaps even only on Earth, and we observe early abiogenesis due to chance (we'd have to be moderately lucky, but not obscenely so). This conclusion differs from \\citet{2002AsBio...2..293L} because they unwittingly made overconfident prior assumptions. Hence, unless there is a direct detection, the answer to the perennial question ``are we alone'' remains ``nobody knows''."
    },
    "0807/0807.2853_arXiv.txt": {
        "abstract": "I report on observations of the z=0.01 dwarf galaxy SBS1543+593 which is projected onto the background QSO HS1543+5921. As a star-forming galaxy first noted in emission, this dwarf is playing a pivotal role in our understanding of high-redshift galaxy populations, because it also gives rise to a Damped Lyman Alpha system. This enabled us to analyze, for the first time, the chemical abundance of $\\alpha$ elements in a Damped Lyman Alpha galaxy using both, emission and absorption diagnostics. We find that the abundances agree with one another within the observational uncertainties. I discuss the implications of this result for the interpretation of high-redshift galaxy observations. A catalog of dwarf-galaxy--QSO projections culled from the Sloan Digital Sky Survey is provided to stimulate future work. ",
        "introduction": "We now know a wealth of galaxy populations at high redshifts. Damped Lyman Alpha (DLA) systems are identified from the strong neutral Hydrogen gas absorption they cause in background QSO spectra, while Lyman Break Galaxies, Lyman Alpha Emitters, and Gamma Ray Burst (GRB) host galaxies are discovered through emission from their stars and ionized gas. What we would yet like to understand is how these populations relate to each other, and to the galaxies we see today. While the study of DLA systems has made great contributions to our empirical knowledge of the chemical evolution of galaxies, unfortunately, despite decades of observational searches it has proven extremely difficult to detect the host galaxies directly in emission (cf. Dessauges-Zavadsky, this volume). Our understanding of high redshift galaxies remains biased by how they were selected, via absorption or emission.  Figure~1 shows $\\alpha$ element abundances for star-forming galaxies (SFG) and HII regions derived from emission lines (these are O abundances), as well as for DLA systems determined from absorption lines (O, S, or Si abundances). This figure illustrates several key points which have entered in the debate over the abundances of emission- versus absorption-selected systems. First, emission diagnostics tend to dominate the low-redshift abundance data while absorption diagnostics are still the method of choice for deriving high-redshift abundances. And second, there appears to be very little overlap in the $\\alpha$ element abundances derived using emission versus absorption techniques. What does that mean for the galaxy populations probed? ",
        "conclusions": ""
    },
    "0807/0807.2254_arXiv.txt": {
        "abstract": "The first, self-consistent calculations are presented of the cosmological, ${\\rm H}_2$-dissociating UV background produced during the epoch of reionization by the sources of reionization. Large-scale radiative transfer simulations of reionization trace the impact of all the ionizing starlight on the IGM from all the sources in our simulation volume down to dwarf galaxies of mass $\\sim 10^8 \\,{\\rm M}_\\odot$, identified by very high-resolution N-body simulations, including the self-regulating effect of IGM photoheating on dwarf galaxy formation. The UV continuum emitted below 13.6 eV by each source is then transferred through the same IGM, attenuated by atomic H Lyman series resonance lines, to predict the evolution of the inhomogeneous radiation background in the Lyman-Werner bands of ${\\rm H}_2$ between 11 and 13.6~eV. On average, the intensity of this Lyman-Werner background is found to rise to the threshold level at which dissociation suppresses ${\\rm H_2}$ cooling and star formation inside minihalos, long before reionization is complete. Spatial variations in the Lyman-Werner background are found which result from the clustering of sources associated with large-scale structure formation, such that intensity fluctuations correlate with matter density fluctuations. As a result, the Lyman-Werner background rises to the threshold level for ${\\rm H_2}$ suppression earlier in the vicinity of the reionization sources and their H~II regions. ",
        "introduction": "Simulations suggest that the first stars in the Cold Dark Matter (``CDM'') universe formed inside minihalos of mass $M \\sim 10^{5-6} {\\rm M}_\\odot$ at $z\\gtrsim 20$, when ${\\rm H}_2$ molecules cooled the primordial, metal-free halo gas and gravitational collapse ensued (e.g. \\citealt{2004ARA&A..42...79B} and references therein). This critical role of ${\\rm H}_2$ molecules as the primary coolants responsible for triggering the gravitational collapse that caused stars to form inside minihalos was limited, however, by the fact that ${\\rm H}_2$ can be dissociated by absorbing UV radiation in the ${\\rm H}_2$ Lyman-Werner (``LW'') bands in the energy range $11-13.6\\,$eV. In the presence of a high enough LW-band radiation intensity, $J_{\\rm LW}$, the ${\\rm H}_2$ abundance would have been too low to cool the gas sufficiently to form stars. The threshold level of the intensity, $(J_{{\\rm LW}})_{{\\rm threshold}}$, above which minihalo star formation was suppressed is still uncertain. Early estimates (\\citealt{2000ApJ...534...11H}; henceforth, ``HAR'') found that $(J_{{\\rm LW}})_{{\\rm threshold}}$ depended upon minihalo mass and redshift. In terms of the dimensionless quantity, $J_{{\\rm LW},\\,21}\\equiv J_{\\rm LW}/(10^{-21}\\,{\\rm erg}\\,{\\rm s}^{-1}\\,{\\rm cm}^{-2}\\,{\\rm Hz^{-1}}\\,{\\rm sr}^{-1})$, they found that, to suppress ${\\rm H}_2$ for {\\em all} minihalo masses, $\\left( J_{{\\rm LW},\\,21} \\right)_{\\rm threshold} \\sim 10^{-2}$ to $1$ was required, as redshift varied from $z\\sim 10$ to $50$, respectively. Later estimates, including those based upon 3D, numerical gas dynamical simulation of minihalos evolving in the presence of a LW background from more realistic cosmological initial conditions, found a similar range of threshold values (e.g. \\citealt{2002ApJ...575...49R}; \\citealt{2003ApJ...592..645Y}; \\citealt{2007ApJ...663..687Y}).  Once stars began to form inside minihalos and also inside the rarer, more massive halos with $M\\gtrsim 10^8\\,{\\rm M}_\\odot$ virial temperatures above $10^4 {\\rm K}$, in which radiative cooling by atomic hydrogen was possible even without ${\\rm H_2}$ (to cool the gas down to $10^4 {\\rm K}$, at least, and initiate gravitational collapse), a rising diffuse LW background would have been inevitable, however. Starlight at energies below $13.6\\,{\\rm eV}$, the ionization potential of atomic hydrogen, would have been largely free to escape from these star forming halos into the intergalactic medium (``IGM''), while the ionizing radiation above $13.6\\,{\\rm eV}$ was partially absorbed by the neutral hydrogen within the halos and thereby reduced by an escape fraction, $f_{\\rm esc}$. In that case, given the value of this $f_{\\rm esc}$ and the ratio of the number of ionizing photons to the number of ${\\rm H}_2$-dissociating photons released by the stars, $N_{\\rm i}/N_{\\rm LW}$,  which depends on the mass function and spectra of the stars, the rise of the cosmic LW background can, in principle, be related to the rise of the diffuse ionizing background. The latter is believed to have been responsible for the reionization of the intergalactic medium completed by redshift $z\\gtrsim 6$. Estimates by HAR showed that the sources of this reionization would have caused the mean LW intensity in the IGM to exceed $\\left( J_{{\\rm LW},\\,21} \\right)_{\\rm threshold}$ long before reionization was complete.  This outcome is expected on quite general grounds, in fact. The mean number of LW photons in the background per H atom is related to the LW intensity according to: \\begin{eqnarray} \\left( \\frac{n_{\\rm LW}}{n_{\\rm H}}\\right) &=&\\left( \\frac{4\\pi}{c}\\int_{11.5\\,{\\rm eV}}^{13.6\\,{\\rm eV}} \\frac{J_{\\nu}}{h\\nu}d\\nu \\right)/n_{\\rm H}, \\nonumber \\\\ &\\simeq &\\frac{4\\pi}{ch}\\frac{2.1\\,{\\rm eV}}{12.6\\,{\\rm eV}}\\langle J_{\\nu}\\rangle /n_{\\rm H}, \\nonumber \\\\ &\\simeq & 1.05\\times 10^{-5} J_{{\\rm LW},\\,21}/n_{\\rm H},  \\nonumber \\\\ &\\simeq & 6.3\\times 10^{-3} J_{{\\rm LW},\\,21} \\left( \\frac{1+z}{21} \\right)^{-3}. \\label{eq:nLW_nH} \\end{eqnarray} Hence, if $\\left( J_{{\\rm LW},\\,21} \\right)_{\\rm threshold} <1$, this implies that $(n_{\\rm LW}/n_{\\rm H})_{\\rm threshold} \\lesssim 6\\times 10^{-3} (1+z)_{21}^{-3}\\ll 1$. The mean number of photons in the LW background per H atom at a given epoch is determined by their rate of emission per atom, integrated over time, reduced by a factor which accounts for attenuation and redshifting after emission. As discussed in \\S~\\ref{sec:Methodology}, photons emitted in the LW bands below 13.6 eV are removed from the background when they redshift to the frequency of the nearest H atom Lyman series transition from level $n\\ge 3$ to  $n=1$. Only those sources within the distance from which photons emitted at the Ly$\\gamma$ frequency would be received at the Ly$\\beta$ frequency can contribute at all to the LW background at a given point. On average, roughly a third of the LW photons emitted within this horizon will survive their trip\\footnote{This factor 1/3 comes from the survival rate from the attenuation by Lyman series resonace lines, given approximately by $\\int_{0}^{r_{\\rm LW}} dr_{\\rm os} f_{\\rm mod} (r_{\\rm os}) / \\int_{0}^{r_{\\rm LW}} dr_{\\rm os}$, where the terms used are described in \\S~\\ref{sub:picket-fence}. Compare this expression to equation (\\ref{eq:Jnuhom}) and also see Figure~\\ref{fig:J21comp}.}. The same sources also emit UV photons with energies above 13.6 eV which are destroyed en-route by photoionizing H atoms. Reionization required that at least one ionizing photon was released into the IGM per H atom by the end of reionization. Let $\\xi_{\\rm LW}$ and $\\xi_{\\rm i}$ be the total number of LW and ionizing photons, respectively, released per H atom into the IGM up to some time. The mean number of photons per H atom in the LW background at that time is then roughly given by \\begin{equation} \\frac{n_{\\rm LW}}{n_{\\rm H}}\\simeq \\frac{1}{3}\\left(\\frac{\\xi_{\\rm LW}}{\\xi_{\\rm i}}\\right) \\xi_{\\rm i} . \\label{eq:nn_xixi} \\end{equation} For the stars believed to be responsible for reionization, the intrinsic ratio of ionizing UV photons to LW photons released, $N_{\\rm i}/N_{\\rm LW}$, ranged from $N_{\\rm i}/N_{\\rm LW} \\sim 15$ (Pop III stars, high mass) to $N_{\\rm i}/N_{\\rm LW} \\sim 1$ (Pop II stars, Salpeter IMF), while the escape fraction of ionizing photons was some $f_{\\rm esc} \\ll 1$. In that case, we can write \\begin{equation} \\frac{n_{\\rm LW}}{n_{\\rm H}}\\simeq \\frac{1}{3} f_{\\rm esc}^{-1}\\left(\\frac{N_{\\rm i}}{N_{\\rm LW}}\\right)^{-1} \\xi_{\\rm i}, \\label{eq:nn_NNxi} \\end{equation} and, therefore, if ($J_{\\rm LW,\\,21})_{\\rm threshold}<1$, equations~(\\ref{eq:nLW_nH}) and (\\ref{eq:nn_NNxi}) imply that $\\left( n_{\\rm LW}/n_{\\rm H}\\right)$ reaches the threshold level when $\\xi_{\\rm i}$ is only \\begin{equation} \\left( \\xi_{\\rm i} \\right)_{\\rm LW,\\, threshold} \\simeq 3 \\left( \\frac{n_{\\rm LW}}{n_{\\rm H}} \\right)_{\\rm threshold} \\left(\\frac{N_{\\rm i}}{N_{\\rm LW}}\\right) f_{\\rm esc} \\ll 1 . \\label{eq:xi_th} \\end{equation} Accordingly, since reionization was not complete until the condition $\\xi_{\\rm i}>1$ was reached, the LW threshold for suppressing ${\\rm H}_2$ in the minihalos must have been reached long before the end of reionization. As a result, as HAR suggested, the minihalos were generally ``sterilized'' before they could contribute significantly to reionization.  What level of LW background is required to suppress minihalo star formation when the minihalos are also directly exposed to {\\it ionizing} radiation, as well, is a more complicated question to answer. If a minihalo forms in a region of the IGM which is already photoionized, the pressure of the photoheated IGM there prevents it from collapsing gravitationally into the dark-matter dominated halo. Those minihalos are missing their baryonic component, therefore. This phenomenon, sometimes referred to as ``Jeans-mass filtering'' \\citep{1994ApJ...427...25S,1998MNRAS.296...44G}, ensures that ${\\rm H}_2$ cooling and star formation do not occur inside minihalos in the ionized regions of the IGM unless those minihalos had formed there {\\it prior} to the arrival of the ionizing radiation. The impact of both LW {\\it and} ionizing radiation (including X-rays) on {\\it pre-existing} minihalos inside H~II regions is a subject of ongoing work, beyond the scope of this paper \\citep[e.g.][]{2001ApJ...548..509M,2002ApJ...575...49R, 2003ApJ...592..645Y,2003MNRAS.346..4560,2004MNRAS.348..753S, 2005MNRAS...361..405I,2006ApJ...639..621A,2006ApJ...650....7H, 2006ApJ...645L..93S,2006ApJ...648..835M,2007MNRAS.375..881A, 2007ApJ...671.1559W,2007ApJ...663..687Y,2008ApJ...673...14O}. The presence of dissociating radiation {\\it always} implies the potential for limiting the ${\\rm H_2}$ abundance and, with it, the cooling required to make stars form. This is true even for the atomic-cooling halos above the minihalo mass range, although the level required for suppression might be higher. We shall focus here, however, on the rise of the LW background and its spatial variations contributed by the dominant sources of reionization, but leave the question of how the intensity impacts star formation for future studies.  Previous estimates of the cosmic LW background were limited to the {\\it mean} background and were based upon a {\\it homogeneous} approximation. These calculations assumed that the sources and the IGM were uniformly distributed, with uniform emissivity, given either by analytical approximation (HAR) or by summing over the sources found in small-box simulations, too small to account for the large-scale clustering of sources or to follow global reionization \\citep[e.g.][]{2002ApJ...575...49R, 2003ApJ...592..645Y}. Background of Ly$\\alpha$ pumping radiation was considered semi-analytically by \\citet{2005ApJ...626....1B} and \\citet{2006MNRAS.367.1057P}, including the effects of fluctuations in the background in a linear approximation. The focus of these papers is on the Ly$\\alpha$ intensity originating from photons that are absorbed by hydrogen Lyman resonance lines and converted into Ly$\\alpha$ photons, which then radiatively mixes the 21 cm levels to drive the spin temperature to the gas kinetic temperature (Wouthuysen-Field effect), while we are interested in photons that remain unattenuated by Lyman resonance lines and affect ${\\rm H}_2$ abundance through photo-dissociation. Note that the horizon for sources responsible for the fluctuating Ly$\\alpha$ background is considerably larger than that responsible for the LW background.  Here we present the first self-consistent radiative transfer calculations of the {\\it inhomogeneous} LW background produced by the same sources which reionized the universe in a large-scale radiative transfer simulation of reionization. This problem presents a formidable computational challenge. The horizon for seeing LW photons is $\\sim100$ comoving Mpc, much larger than the size of typical H~II regions ($\\sim10$ Mpc). Since the mean free path for LW photons ($\\sim100$ Mpc) is much larger than that for H-ionizing photons we must account for sources distributed over large volume and lookback time. Finally, the LW band photons are attenuated as they redshift into H-atom Lyman series resonance lines as they travel across the IGM. Hence, it is necessary to perform a multi-frequency radiative transfer calculation from each of the millions of sources in a cosmological volume larger than $\\sim (100\\, \\rm Mpc)^3$, integrated along the light cones from each source to every observation point they intersect, which is computationally prohibitive. As we shall show, a good approximation is possible which reduces the multi-frequency calculation to an equivalent gray opacity calculation. As a result, we are not only able to derive the evolution of the rising globally-averaged mean LW background during the epoch of reionization (EOR), but also to map its pattern of spatial variations over time. In \\S~\\ref{sec:Methodology}, we describe how continuum radiation emitted in the LW range below 13.6~eV is transferred through the IGM along the light cones from sources to observers, and how we solve this problem numerically. In \\S~\\ref{sec:applic}, we apply this method to one of our recent large-scale radiative transfer simulations of self-regulated reionization, described in \\citet{2007MNRAS.376..534I}. We compare the mean LW background evolution thus derived numerically with the homogeneous universe approximation and describe the spatial fluctuations of the LW background in some detail. Our conclusions are summarized in \\S~\\ref{sec:conclusions}. Some of our results on the inhomogeneous LW background during the EOR described here were first summarized in \\citet{2008AIPC..990..374A}. ",
        "conclusions": "\\label{sec:conclusions}  We have, for the first time, calculated the inhomogeneous background of ${\\rm H}_{2}$ dissociating UV radiation caused by the same sources which reionized the Universe in a large-scale radiative transfer simulation of cosmic reionization. The UV continuum emitted below 13.6 eV by each source was transferred through the IGM, attenuated by atomic H Lyman series resonance lines, to predict the evolution of the inhomogeneous radiation background in the energy range of the LW band of ${\\rm H}_{2}$ between 11 and 13.6~eV. This required us to transfer the radiation from the many thousands of source galactic halos in our simulation, which resolved all halos of mass above $10^8M_\\odot$ in a comoving volume of $(50\\,\\rm Mpc)^3$, from each source to each of the millions of grid cells of our reionization simulation. To accomplish this, we developed a novel method to calculate the attenuation of LW band photons from individual sources by H Lyman series resonance lines, an otherwise prohibitively expensive, multi-frequency calculation, in a very fast way, instead, without an explicit multi-frequency operation. This was achieved by a grey opacity approximation, the {}``picket-fence'' modulation factor, a simple function of the comoving distance between a source and an observer, which represents the frequency-averaged attenuation of LW band photons. We also explicitly accounted for the effect of the finite light-crossing time between sources and observers, by constructing the conformal space-time diagram of all reionization sources and observer grid cells and finding the intersections between source world lines and observer past light cones.  Our results here demonstrate that the rise of the cosmic LW background to levels which have previously been identified as the threshold for dissociating $\\rm H_2$ and thereby suppressing star formation inside minihalos occurs well before the epoch of reionization is very advanced. Not only is this the case for the mean LW background intensity, as anticipated by earlier estimates based upon a homogeneous universe approximation, but it is even more the case for the inhomogeneous background. The first regions to form halos are the regions with the highest mean overdensity, and we show that this is where the LW background rises the fastest and is at the highest levels.  This means that our assumption here, for simplicity, that reionization is dominated by the atomic cooling halos, and that minihalos are sterilized by the rising LW background before they can contribute significantly to the ionizing and, accordingly, the LW background, as well, is self-consistent.  On the other hand, since there are also some minihalos that form far from the density peaks around which the halos cluster, we might also expect that there are {\\em some} minihalos that form far from the peaks in the LW background, before their local LW intensity has risen to the threshold level for $\\rm H_2$ suppression.  In the future, it will be interesting to consider the possible role of these minihalos that form in places where the LW background is not high enough to suppress their star formation.  Our result can be used for various applications. For example, fluctuating ${\\rm H}_2$ abundance in IGM and cosmological halos may be calculated inside our simulation box. This would at least require implementing reaction rates of neutral and ionic species of H, He, and ${\\rm H}_2$. One may have to run many small-box simulations similar to those of \\citet{2003ApJ...592..645Y}, in order to calculate ${\\rm H}_2$ abundance and track source formation inside cosmological halos under fluctuating $J_{\\rm LW}$ calculated in this paper. Similarly, once X-ray emitting sources are properly populated in the simulation box, a composite effect of negative (due to photodissociation by UV) and positive (due to partial ionization by X-ray) feedback may be studied as well.  There is the possibility that source formation inside the more massive, atomic-cooling halos is affected by the LW background, too, because molecular cooling takes place inside these halos in the following way. When the gas cools from the ionized state in these halos, it does so out of equilibrium (e.g. \\citealt{1987ApJ...318...32S}; \\citealt{1992ApJ...386..432K}): gas cools faster than it recombines, so even after atomic H cooling has reached its typical saturated phase at $T\\lesssim 10^4 \\,{\\rm K}$, there still remains a significant trace amount of electrons. Gas-phase reactions can then create ${\\rm H}_2$ with the help of these electrons, which can further cool the gas down to $T\\sim 200 \\,{\\rm K}$. Even though ${\\rm H}_2$ at this stage can become self-shielded against the UV dissociating background (\\citealt{2002ApJ...569..558O}), a strong enough background may nevertheless dissociate ${\\rm H}_2$ and suppress star formation to some extent (HAR). The threshold LW intensity for such suppression inside these atomic cooling halos, $(J_{\\rm  LW})_{\\rm threshold, atomic}$, is believed to be much larger than $(J_{\\rm LW})_{\\rm threshold}$ for minihalos (HAR). This might, therefore, somewhat affect the history of cosmic reionization near the end of reionization, when even $(J_{\\rm LW})_{\\rm threshold, atomic}$ has been reached.  Emissivity by the first stars that contribute to the cosmic near-infrared background (NIRB) may also be affected by the fluctuating LW background calculated here. Observations of a strong NIRB excess over the known foreground have been reported (e.g. \\citealt{2005ApJ...626...31M}), and this has been interpreted as a possible signature of the first stars (c.f. \\citealt{2006Natur.440.1002M} and references therein), although the true identity of the contributing sources is currently under debate (e.g. \\citealt{2007ApJ...666L...1K}; \\citealt{2007ApJ...666..658T}). Our work will impact the interpretation of both the homogeneous and inhomogeneous components of the NIRB, because predictions of the first star formation should be strongly affected by our results.  We argued in \\S~\\ref{sec:Methodology} that we are justified in neglecting ${\\rm H}_{2}$ self-shielding by the IGM in calculating the opacity to LW photons. This was justified by the fact that the concentration of $\\rm H_2$ in the pre-reionization IGM was only $\\sim10^{-6}$ and, as such, $\\tau_{H_2}\\lesssim1$. However, we can now use hindsight to see that our neglect of $\\tau_{H_2}$ is even more self-consistent than that estimate would suggest. Hydrogen molecules in the IGM are quickly dissociated by the rising LW background. It is easily seen from Figures~\\ref{fig:nn}, \\ref{fig:J21cuts} and \\ref{fig:J21dev} that $J_{{\\rm LW},\\,{21}} \\gtrsim 10^{-2}$ at $z\\simeq 20$ in the vicinity of atomic cooling halos, and the mean value quickly rises to $\\langle J_{{\\rm LW},\\,{21}}\\rangle \\simeq 10^{-2}$ and above. When the mean IGM prior to reionization (with temperature $T_{\\rm IGM}\\simeq 369(1+z)^{2}/135^{2}$, hydrogen number density $n_{\\rm H}\\simeq 1.7 \\times 10^{-3}\\,(1+z)^{3} / 21^{2}\\,{\\rm cm}^{-3}$, and $y_{{\\rm H}_2}\\simeq 2\\times 10^{-6}$ when there is no LW background) is exposed to the LW intensity $J_{{\\rm LW},\\,{21}} = 10^{-2}$ from $z=20$ onward, for example, its molecular fraction drops to $y_{{\\rm H}_2}\\simeq 2\\times 10^{-7}$ by $z\\simeq 17$. Even if ${\\rm H}_{2}$ self-shielding were marginally important with $y_{{\\rm H}_2}\\simeq 2\\times 10^{-6}$ at high redshift (\\citealt{2002ApJ...575...49R}), therefore, this will quickly become negligible when molecules are dissociated to the level $y_{{\\rm H}_2}\\simeq 2\\times 10^{-7}$, which would occur at $z\\simeq 17$ from the LW background contributed by the atomic-cooling halos alone. Peculiar motion of gas elements will weaken self-shielding even further (e.g. \\citealt{2001ApJ...548..509M}). \\citet{2007ApJ...665...85J}, on the other hand, claim that relic H~II regions created by the first stars can recombine and generate abundant ${\\rm H}_{2}$ before being exposed to other external sources, through the nonequilibrium $\\rm H_2$-formation mechanism described above (\\citealt{1987ApJ...318...32S}; \\citealt{1992ApJ...386..432K}). \\citet{2007ApJ...665...85J} find, however, that this boost of ${\\rm H}_2$ abundance is delayed significantly when $J_{{\\rm LW},\\,{21}} \\gtrsim 10^{-2}$, which is easily satisfied in the mean IGM at $z\\simeq 17$ and in the vicinity of atomic cooling halos at $z\\simeq 20$. This suggests that the proposed nonequilibrium enhancement of the $\\rm H_2$ concentration in the IGM inside relic H~II regions is not likely to provide a significant enough $\\rm H_2$ opacity to shield the IGM through most of the EOR.  \\ack  This study was supported in part by KICOS through the grant K20702020016-07E0200-01610 provided by MOST, NSF grant AST 0708176, NASA grants NNX07AH09G and NNG04G177G, Chandra grant SAO TM8-9009X, Swiss National Science Foundation grant 200021-116696/1, and Swedish Research Council grant 60336701. The authors acknowledge the Texas Advanced Computing Center (TACC) at The University of Texas at Austin for providing HPC resources that have contributed to the research results reported in this paper."
    },
    "0807/0807.0317_arXiv.txt": {
        "abstract": "Natural inflation driven by pseudo-Nambu-Goldstone bosons have a problem that the nearly scale invariant spectrum of density perturbations is attained   only when the symmetry breaking scale is of the order of Planck scale. We show here that if one couples the PNGB to a thermal bath as in warm inflation models, the amplitude and spectral index which agrees with the Wilkinson Microwave Anisotropy Probe (WMAP) data is obtained with the symmetry breaking in the GUT scale. We give a GUT model of PNGB arising out of spontaneously broken lepton number at the GUT scale which gives rise to heavy Majorana masses for the right handed neutrinos which is needed in see-saw models. This model also generates a lepton asymmetry because of the derivative coupling of the PNGB to the lepton current. A characteristic feature of this model is the prediction of large non-gaussianity which may be observed in the forthcoming PLANCK experiment. ",
        "introduction": "Inflation \\cite{Guth} was introduced to solve the horizon and curvature problems of cosmology and in addition it predicted a scale invariant spectrum of density perturbations which was verified by the Cosmic Background Explorer (COBE), WMAP and other CMBR anistropy experiments.The successful model of inflation requires a flat potential and a natural candidate for such a potential is the Pseudo-Nambu-Goldstone potential as first pointed out in \\cite{natural1,natural2,natural3}.  One limitation of natural inflation models is that the symmetry breaking scale $f$ is related to the spectral index $n_s =1-M_p^2/(8 \\pi f^2)$ and observations of microwave anisotropy constrain the symmetry breaking scale to be close to the Planck scale. As discussed in Banks et al \\cite{Banks:2003sx} a symmetry breaking scale larger than $M_P$ makes the theory susceptible to large quantum corrections which can destabilize the flat PNGB potential. There have been several attempts at solving this large $f$ problem in natural inflation. Arkadi-Hamed et al \\cite{ArkaniHamed:2003wu} invoke extra dimensions with the Wilson loop of a gauge field in the extra dimension to explain why $f \\sim M_P$. Similar arguments are also given by Kaplan and Weiner \\cite{Kaplan:2003aj}. Kim et al \\cite{Kim:2004rp} invoke two field natural inflation to bring down the symmetry breaking scale below Planck scale. Kinney and Mahanthappa \\cite{Kinney:1995cc} show that in some special symmetry breaking schemes  the quadratic term in the PNGB field is subdominant compared to the higher order terms  and in these models the symmetry breaking scales can be lower than the Planck scale.  In this paper we show that if the PNGB inflaton of the natural inflation model is coupled to a radiation bath (with a sub-dominant energy density) as in warm inflation models \\cite{warm} the symmetry breaking scale $f$ can be in the GUT scale and be consistent with the observations of the temperature anisotropy spectrum observed by WMAP \\cite{Komatsu:2008hk}. In this model the dissipative coupling of the PNGB inflaton makes it roll slowly even in a steep potential which results when $f$ is lowered from $M_P$ to $M_{GUT} \\sim 10^{16} GeV$.  As a specific model let us consider the SU(5) model where the right handed neutrino $N$ is a singlet. In the see-saw mechanism \\cite{see-saw} one generates a heavy Majorana mass by coupling this right handed neutrino to a SU(5) singlet Higgs, \\be -\\mathcal L_{\\nu}=g H N N^C. \\label{yukawa} \\ee In order to break lepton number spontaneously we have a potential for the Higgs \\be -\\mathcal L_H=\\frac{\\lambda}{8}(H^{\\dagger}H-\\frac{f^2}{2})^2. \\label{vhiggs} \\ee Here $f$ is the spontanous symmetry breaking scale.  At the minima of the potential the Higgs is given by $H=\\frac{1}{\\sqrt 2}\\, f\\, e^{i\\frac{\\phi}{f}}$. Here angular variable $\\phi$ is the Goldstone boson of the spontaneously broken lepton number symmetry. Quantum gravity effects are expected to break global symmetries at the Planck scale. If there is an explicit symmetry breaking due to gravity the Goldstone boson acquires mass. The explicit symmetry breaking term can be of the form \\be -\\mathcal L=\\frac{M^2}{M_P}N N^C+O(\\frac{1}{M_P^2}). \\label{esb} \\ee Because of this explicit symmetry breaking the potential of PNGB is given by \\cite{Frieman} \\be V(\\phi)=\\Lambda^4\\left(1+\\cos\\left(\\frac{\\phi}{f}\\right)\\right). \\label{vpngb} \\ee $\\Lambda$ is related to explicit symmetry breaking scale  $\\mu=\\frac{M^2}{M_P}$. The mass of the PNGB is given by $m_\\phi=\\frac{\\mu^2}{f}=\\frac{\\Lambda^2}{f}$. This implies that $\\Lambda=\\mu=\\frac{M^2}{M_P}$. Now if we take $M\\sim\\, M_{GUT}\\sim\\, 10^{16}-10^{17}$GeV then we have $\\Lambda\\sim\\, 10^{13}-10^{14}$GeV which is the allowed range by WMAP data. ",
        "conclusions": "There has been a long standing problem with utilizing the flat potential of PNGB's for inflation as the nearly scale invariant power spectrum which is consistent with observations generated only when the symmetry breaking scale $f \\sim M_P$ \\cite{natural1,natural2,natural3,Banks:2003sx,ArkaniHamed:2003wu,Kaplan:2003aj,Kim:2004rp}. In this paper we show that by coupling the inflaton to a radiation bath (as in warm inflation models \\cite{warm}) can reduce  $f$ to the GUT scale. The value of the inflaton field $\\phi \\sim f\\sim M_{GUT}$ which makes the inflaton potential stable against Planck scale radiative corrections. We give a model of inflation where the inflaton is the PNGB arising from spontaneous breaking of lepton number which also gives a large Majorana mass for the right handed neutrinos as required in see-saw models \\cite{see-saw}. Since the PNGB's have a derivative coupling to the lepton current this model also generates a lepton asymmetry spontaneously \\cite{cohen} during inflation. We show that with the parameters of the inflation model which give the correct amplitude and spectral index of CMBR also give the required lepton asymmetry of $\\eta_L \\sim 10^{-10}$ which can be converted to a baryon asymmetry of the same order by sphaleron processes in the electro-weak era \\cite{Fukugita:1986hr}."
    },
    "0807/0807.2638_arXiv.txt": {
        "abstract": "Stellar mass black holes (BHs) are expected to segregate and form a steep density cusp around supermassive black holes (SMBHs) in galactic nuclei.  We follow the evolution of a multi-mass system of BHs and stars by numerically integrating the Fokker-Planck energy diffusion equations for a variety of BH mass distributions. We find that the BHs ``self-segregate'', and that the rarest, most massive BHs dominate the scattering rate closest to the SMBH ($\\lesssim 10^{-1}\\,$pc).  BH--BH binaries form out of gravitational wave emission during BH encounters. We find that the expected rate of BH coalescence events detectable by Advanced LIGO is {\\new $\\sim 1-10^2\\,$yr$^{-1}$, depending on the initial mass function of stars in galactic nuclei and the mass of the most massive BHs. We find that the actual merger rate is likely $\\sim 10$ times larger than this due to the intrinsic scatter of stellar densities in many different galaxies.}  The BH binaries that form this way in galactic nuclei have significant eccentricities as they enter the LIGO band ($90\\%$ with $e > 0.9$), and are therefore distinguishable from other binaries, which circularise before becoming detectable. We also show that eccentric mergers can be detected to larger distances and greater BH masses than circular mergers, up to $\\sim 700\\,\\Msun$.  Future ground-based gravitational wave observatories will be able to constrain both the mass function of BHs and stars in galactic nuclei. ",
        "introduction": "\\subsection{Motivation}\\label{sec:motivation} The coalescence of stellar mass black hole-black hole (BH-BH) binaries is one of the most anticipated sources of gravitational waves for ground-based interferometers such as LIGO\\footnote{http://www.ligo.caltech.edu/} or VIRGO\\footnote{http://www.virgo.infn.it/}. Whether formed primordially \\citep{2002ApJ...572..407B, 2004ApJ...611.1068B, 2008ApJ...676.1162S} or in dense star clusters \\citep{2000ApJ...528L..17P, 2004ApJ...616..221G, 2006ApJ...637..937O, 2006ApJ...640..156G, 2007PhRvD..76f1504O, 2008arXiv0804.2783M}, nearly all of these binaries are expected to be circularised by gravitational wave (GW) emission before they are detectable by ground based observatories \\citep[][]{2006ApJ...637..937O, 2006ApJ...640..156G}.  In order to be detected with a high signal-to-noise ratio, matched filtering algorithms sift through the LIGO data stream looking for such circular inspirals \\citep{2008PhRvD..77f2002A}, but might miss many eccentric events {\\new \\citep{PhysRevD.60.124008,2007arXiv0712.3199T,2008ApJ...681.1431M}}. In principle, eccentric inspirals are well suited for detection as well, however they have not been expected to be a significant source for terrestrial detectors \\citep[except see][]{2003ApJ...598..419W}.  In this paper, we propose an important additional source of gravitational wave sources, which preferentially form binaries that are still eccentric as they merge. In galactic nuclei with supermassive black holes (SMBH) of mass $\\msmbh < 10^7\\,\\Msun$, relaxation times are often less than a Hubble time, and can result in the formation of steep density cusp of stars and stellar mass black holes (BHs).  Indeed, as many as $\\sim 20,000$ BHs are expected to have segregated into the inner $\\approx 1\\,$pc of the Milky Way \\citep[][hereafter HA06]{1993ApJ...408..496M,2000ApJ...545..847M, 2006ApJ...649...91F,2006ApJ...645L.133H}.  When two BHs have a close encounter, they can release sufficient amount of energy in GWs to form a tight binary, and merge in less than a few hours.  These mergers are almost always eccentric, in contrast to circular inspirals expected in globular clusters. A similar process can occur in the runaway growth of BHs when no SMBH is present that will also result in eccentric mergers; however, such systems are inherently unstable and not long-lived \\citep{1987ApJ...321..199Q,1989ApJ...343..725Q,1990ApJ...356..483Q,1993ApJ...418..147L}.  The GWs generated during the evolution of these orbits are endowed with a rich structure.  For sufficiently small impact parameters, the GWs can be detected with second generation GW instruments already during first passage \\citep{2006ApJ...648..411K}. Subsequently, the eccentricity is close to unity for several thousand orbits after the formation of the binary, generating a train of short-duration distinct GW bursts with a continuous frequency spectrum.  As the eccentricity gradually decreases, the separation between subsequent bursts decreases, and the waveform becomes continuous in the time-domain and decouples into discrete harmonics in frequency. Interestingly, as we show below, the eccentricity is non-negligible all the way to coalescence.  The full waveform, consisting of a correlated set of distinct GW bursts, which evolves into a continuous eccentric inspiral signal, would not be discovered in the GW data by existing data analysis techniques such as GW burst search algorithms, nor by small eccentricity inspiral templates, since the contribution of individual bursts to the GW power is small and the contribution of higher harmonics is significant throughout the evolution \\citep[see][for a similar discussion for extreme mass ratio GW bursts during encounters with a supermassive black hole]{2008ApJ...675..604Y}.  Eccentric inspirals are more luminous in GWs than circular inspirals and extend to higher frequencies, and so may be detectable to higher redshifts.  Furthermore, these events may be capable of detecting intermediate mass black holes using terrestrial instruments \\citep{2007PhRvL..99t1102B,2008ApJ...681.1431M}.  The evolution of the signal during the initial phase of eccentric inspirals can be described using a small number of parameters, allowing one to construct very efficient and sensitive data analysis algorithms to search for these signals in the LIGO data. Additionally, these signals spend a longer physical time in the LIGO band, possibly reducing the false alarm rate for a fixed number of templates that cover the observational period. The modulation introduced by the Earth's rotation during such an event can be utilised to further improve the measurement accuracy for the source position.  The GW waveforms have been studied for eccentric inspirals through several different approaches. Post-Newtonian (PN) waveforms exist up to 1.5PN order beyond the leading order gravitational radiation effects ($\\propto v^3$) for general mass ratios with spins \\citep{Vasuth:2007gx,2008PhysRevD.77.104005}, {\\new 2PN without spins \\citep[($\\propto v^4$, including the effects of eccentricity and radiation reaction]{2004PhRvD..70f4028D} and 3PN including eccentricity but without spins and also neglecting radiation reaction \\citep[$\\propto v^6$,][and references therein]{2008PhRvD..77f4035A}, 5.5PN for extreme mass ratios without spins \\citep[$\\propto v^{11}$,][]{1996PThPh..96.1087T}. The evolution of a binary is chaotic at (or beyond) the 2PN approximation if the component BHs have spins \\citep{2006PhRvD..74l4027L,2007PhRvD..76l4004W} and also possibly for nonspinning BHs approaching the unstable circular orbit \\citep[i.e. exhibiting zoom-whirl orbits,]{2007CQGra..24...83P,2008arXiv0802.2520W}.} Numerical relativity has mostly focused on the circular inspiral for comparable mass BHs \\citep[and references therein]{\\baker,\\berti,2008arXiv0804.4184B}, or extreme mass ratios \\citep{2004PhRvD..69d4025M,2007CQGra..24...83P,2007PhRvD..75b4005B}. Recently, however, the waveforms have been evaluated for a handful of equal mass eccentric initial conditions \\citep{2007arXiv0710.3823S,2008arXiv0802.2520W,2008arXiv0806.1037H,2008PhRvD..77h1502H}.  In this paper, we perform two separate analyses to two different problems.  First, in \\S~\\ref{sec:massseg}, we determine the multi-mass distribution of BHs in galactic nuclei for a variety of mass models. In \\S~\\ref{sec:GB}, we determine the formation rate of binaries due to gravitational wave capture in a variety of galactic nuclei. In \\S~\\ref{sec:detection} we describe the general features of the generated GW waveform, calculate the expected signal-to-noise ratio for its detection with second generation terrestrial GW instruments and determine the expected detection rate of such mergers.  We also show that the eccentricity distribution of such sources will distinguish binaries formed in the manner outlined in this paper from BH-BH binaries formed dynamically in massive star clusters. ",
        "conclusions": "In this paper, we have analysed two separate problems. First we determined the multi-mass distribution of BHs in galactic nuclei, which we then used to analyse the detection rate and merger of gravitational wave capture binaries.  We integrated the time-dependent Fokker-Planck equations for a variety BH mass distributions until they reached a steady state.  We found, consistent with previous results, that within a relaxation timescale at the radius of influence of the SMBH, the BHs and stars form approximately power-law density cusps ($\\propto r^{-3/2-p_0 m / M_{\\max}}$) within the radius of influence of the SMBH.  Because the BHs are more massive than the stars, they have steeper density profiles ($\\propto r^{-2}$) and dominate the dynamics and relaxation processes in the inner $\\sim 0.1\\,$pc near the SMBH.  In such dense population environments, the probability of close flybys between two BHs is nonnegligible. Using our results for the steady-state distribution of BHs, we calculated the expected rate and distribution of GW capture binaries in galactic nuclei.  We showed that after forming, these binaries rapidly inspiral and merge within hours.  Unlike other sources of merging binaries, the BH binaries in galactic nuclei form with a characteristic GW frequency {\\em inside} the AdLIGO frequency band, and the detectable GW signal duration is much longer compared to circular inspirals.  In addition, for sufficiently small impact parameters, the waveforms for such eccentric inspirals are broadband, and can be detected for much larger masses, up to $\\mtot \\gtrsim 700\\,\\Msun$. This exceeds the BH masses previously considered detectable \\citep{2007PhRvL..99t1102B,2008ApJ...681.1431M}, and as such it opens a new avenue to probe for the existence of a population of intermediate mass black holes in galactic nuclei.  Additionally in \\S~\\ref{sec:clusters}, we also estimated the rate of GW captures in young, massive star clusters where BHs decouple from stars to form a subcluster in the centre. We found that the rate of GW captures in a single cluster may be intrinsically larger than in a galactic nucleus. Typically, most binaries circularise before becoming detectable, however a significant fraction {\\new ($\\sim 10\\%$) }may still merge with residual eccentricity.  Given the vastly different eccentricity distribution of mergers in galactic nuclei and massive star clusters, the dominate source of eccentric inspirals can be determined with only a few detections.  The total rate of such events from massive clusters depends on the relatively unconstrained number of young clusters within the detection limits of AdLIGO.   Finally, we analysed the properties of the GW signal for the eccentric inspiral of comparable mass BHs, starting from the initial highly eccentric phase up to the final eccentric inspiral.  We found that the maximum luminosity distance that such events may be detected using a single AdLIGO-type interferometer is $1.2\\,$Gpc for component masses $m=M\\sim 10\\,\\Msun$, and up to $3.3\\,$Gpc for $m=M\\sim 50\\,\\Msun$.  We then used this to calculate the total detection rate of signals, by using a model merger rate convolved with a realistic population of SMBHs.  We found that the most massive BHs dominate the detection rate of future ground based gravitational wave detectors with $\\sim (10 - 10^3) \\xi_{30}$ events expected per year, where $\\xi_{30}$ is a measure of the expectation value for the square of number densities near the SMBH (see Eq.~\\ref{eq:xi}). The overall detection rate is sensitive to the number fraction of BHs as well as the maximum mass of BHs, and so future observations will be able to constrain both the average star formation properties and upper mass of BHs in galactic nuclei.  Overall, the expected AdLIGO detection rate for GW capture binaries in galactic nuclei is comparable to estimates for other sources of gravitational waves. BH-BH binaries that form dynamically in massive star clusters may be detected by AdLIGO $\\sim 10-10^4\\,$ times per year, depending on the total fraction of star formation that occurs in massive, long lived clusters \\citep{2000ApJ...528L..17P, 2004ApJ...616..221G, 2006ApJ...637..937O, 2006ApJ...640..156G, 2007PhRvD..76f1504O, 2008arXiv0804.2783M}. Estimates for the population of merging binaries that formed in the early evolution of present day globular clusters are $\\sim 10-100\\,$yr$^{-1}$ alone. Although globular clusters contain only a small fraction of the mass of stars in the universe, the merger rate of BHs is sufficiently enhanced by dynamical interactions that the rate is comparable to those expected from standard binary evolution for all other stars (\\citealt{2002ApJ...572..407B, 2004ApJ...611.1068B, 2008ApJ...676.1162S}; however see \\citealt{2007ApJ...662..504B}). Merging neutron stars are also a promising source of gravitational waves, whose progenitors have been directly observed.  Using constraints from the modest population neutron star-neutron star binaries known in our galaxy, \\citet{2004ApJ...601L.179K,2004ApJ...614L.137K} estimate that $\\sim 20 - 600$ events per year may be detected by AdLIGO.  We emphasise that the rates of GW captures calculated in this paper are {\\em independent} of other sources of GWs, and will occur in addition to other sources.  Given the estimates of sources from a variety of environments, gravitational wave detectors promise to open a new and interesting window into the physics, dynamics, and evolution of compact binaries.  So far, we have not looked at the interaction between the BHs and other compact objects such as neutron stars (NSs). The NSs are not expected to segregate significantly in galactic nuclei, however they are intrinsically more common than BHs.  To test their importance we have performed an additional run of Model E-2 to look at the expected detection rate of NS-BH mergers. We found that the total detection rate of such events is about $\\sim 1\\%$ of the total rate ($\\sim 1\\,$yr$^{-1}$), and merits future study. Gravitational waves from BH-NS inspirals can provide interesting constraints on the equation of state for NSs \\citep{2002PhRvL..89w1102F}, given model sources from numerical simulations \\citep{2006PhRvD..73b4012F,2008ApJ...680.1326R}. To date, such analyses have focused on the circular inspiral of BHs with NSs, however recent population synthesis studies suggest the number of BH-NS binaries that form may be as rare as those expected here, $\\sim 1\\,$yr$^{-1}$ \\citep{2007ApJ...662..504B}.  Given that GW capture binaries are predominately eccentric throughout inspiral, we also expect the newly formed BH-NS binaries to be eccentric. However, with BH-NS binaries, the tidal effects of the BH on the NS can provide an additional source of energy dissipation in the binary, which can enrich the GW signal, and cause it to deviate from the point-mass approximations calculated here.  Future numerical simulations with eccentric encounters must be done in order to test the importance of eccentricity on the GW signal for BH-NS binaries.  Another area we neglected is the influence of binarity on the merger rate of BHs in galactic nuclei. However, we can estimate their overall importance by looking at what fraction of binaries are hard enough to survive in galactic nuclei.  Given the distribution of BH binaries in \\citet{2004ApJ...611.1068B}, only a small fraction, $\\lesssim 0.1\\%$ of the BH-BH binaries are sufficiently hard (with an orbital period $\\lesssim 10\\,$days) not to be disrupted by repeated encounters with stars or BHs. Interestingly, this fraction is comparable to the fraction of BHs that merge due GW capture. The properties of the merger of such binaries have not been explored yet, however, given the evidence that such encounters produce few eccentric events in massive star clusters, we expect that they will not have many eccentric mergers in galactic nuclei. Unfortunately, the methods we have used here are not suitable for looking at higher order $N$-body interactions (such as exchanges, or binary-binary scattering); simulations similar to FAK06 are more adequate for this type of study. The rates we calculated might be enhanced if high-dispersion environments of protogalactic cores without a SMBH also exist, although such systems are old and not long lived \\citep{1987ApJ...321..199Q,1989ApJ...343..725Q,1990ApJ...356..483Q,1993ApJ...418..147L}.   {\\new  In order to estimate the detection rate of gravitational wave capture binaries in galactic nuclei we had to make assumptions that introduced uncertainty in our calculations. For the sakeof clarity, we wish to summarise these uncertainties and the order of magnitude effect they may have on the actual detection rate for AdLIGO.  The largest uncertainty in the rate calculation is the distribution of BHs in galactic nuclei, especially the number fraction of BHs ($\\sim 10^2$) and their mass distribution ($\\sim 10$) in the centres of galaxies. Observations of our own galactic centre suggest that the star formation may be top-heavy, and place the rate on the higher end of that reported here.  One effect we neglect in our calculations is resonant relaxation, which may reduce the rate $(\\sim 10)$ by depleting the central cusp of BHs. Our final rate, however, is in practice only logarithmically sensitive to the inner edge of the cusp of BHs, and the densest cusps that dominate our rate will be least affected by resonant relaxation. The final rate is also nearly logarithmically dependent on the population of SMBHs in the local universe (by a factor of at most $\\sim 10$), but is much more sensitive to the distribution of densities at the radius of influence of the BHs (by a factor of $\\sim 10$--$10^2$). Finally, there is some uncertainty in the detection rate given our calculation of the $S/N$ of the mergers ($\\sim 10$) since we do not consider corrections for zoom-whirl orbits and conservatively ignore the contribution of the merger and ringdown waveforms of the binary to the signal.  }  There are many aspects to the merger of GW capture binaries that require future work, which will be aided by advanced numerical simulations of the evolution of these binaries.  In our calculations we used only the leading order (Newtonian) formula to calculate the evolution of the binaries and the GW waveforms, and do not account for the GW signal during the plunge and merger of the binary, nor do we consider the contribution of zoom-whirl orbits. The recent breakthrough of numerical simulations finally allows the full GW signal to be calculated directly. Eccentric mergers promise to have a richer signal than that for circular binaries. Future studies should incorporate the contribution of the final coalescence to the GW signal, which would increase the signal strength generated during the inspiral for the same capture event, and has a potential to further increase the maximum range of detection.  We also found that a fraction ($\\sim 10-20\\%$) of mergers will occur very close to, or within the estimated last stable orbit of BH binaries.  These capture events may be zoom-whirl orbits, which dramatically increases the GW signal compared to a regular inspiral \\citep{2007CQGra..24...83P}.  Such encounters are inherently in the strong GR regime, and again require a full numerical treatment of General Relativity to be studied. A precise assessment of these sources will further increase our estimates of detection rates.  Our estimates for the number of GW capture binaries in galactic nuclei (and perhaps in massive star clusters) suggest that such events may be common enough to be directly detected, and may provide the richest gravitational wave sources in the universe for ground based detectors."
    },
    "0807/0807.1531_arXiv.txt": {
        "abstract": "The Gini coefficient, a non-parametric measure of  galaxy morphology, has recently taken up an important role in the automated identification of galaxy mergers. I present a critical assessment of its stability, based on a comparison of HST/ACS imaging data from the GOODS and UDF surveys. Below a certain signal-to-noise level, the Gini coefficient depends strongly on the signal-to-noise ratio, and thus becomes useless for distinguishing different galaxy morphologies. Moreover, at all signal-to-noise levels the Gini coefficient shows a strong dependence on the choice of aperture within which it is measured. Consequently, quantitative selection criteria involving the Gini coefficient, such as a selection of merger candidates, cannot always be straightforwardly applied to different datasets. I discuss whether these effects could have affected previous studies that were based on the Gini coefficient, and establish signal-to-noise limits above which measured Gini values can be considered reliable. ",
        "introduction": "\\label{sec:intro}  The ever increasing amount of galaxies probed by imaging surveys such as the Sloan Digital Sky Survey \\citep[SDSS,][]{sdssedr}, the Great Observatories Origins Deep Survey \\citep[GOODS,][]{goods}, or the Hubble Space Telescope (HST) Ultra Deep Field \\citep[UDF,][]{udf} triggered the need for automated galaxy classification methods. Ideally, these should not only be faster in classifying thousands of galaxies than any visual examination, but should be more robust and less subjective than the latter.  An important aspect is the evolution of galaxy structure over cosmic time: deep HST surveys such as GOODS or UDF reveal a huge amount of peculiar or irregular galaxies \\citep[see, e.g.,][]{cow95a,elm04b,con08} that do not fit into the standard Hubble scheme, nor into the class of irregular galaxies as defined for the local Universe. This is owed to the fact that more distant galaxies are seen in a significantly younger stage, in which starburst events and mergers occur at a larger rate \\citep[e.g.][]{cow95b,mad98,bell06a}. Moreover, at higher redshifts, bluer rest-frame wavelengths are probed, giving emphasis to the distribution of the young stellar component. Although trust in visual classification remains for both low-redshift \\citep[e.g.][]{sch07b} and high-redshift datasets \\citep[e.g.][]{ferlis05}, the hope was that new structural classification schemes  would also be suitable to describe the various kinds of peculiar galaxies in an appropriate way \\citep{abr94,abr96b}, and to separate them from the more familiar galaxy classes.  ``Non-parametric'' structural measures that have been developed in this context are the concentration index \\citep[$C$,][]{abr94,ber00,tru01,gra01,con03a}, the asymmetry index \\citep[$A$,][]{abr96a,con97,kor98,con00}, and the clumpiness index \\citep[$S$,][]{con03a}. These three are known as the CAS parameters \\citep{con03a}. Two further parameters that are frequently used are the Gini coefficient \\citep{gini,abr03,lot04} and the $M_{20}$ parameter \\citep{lot04}. Gini describes the distribution of flux values among the pixels of an object's image, while $M_{20}$ quantifies the distribution of the brightest 20\\% of pixels. As opposed to the CAS parameters, neither Gini nor $M_{20}$ require the center of a galaxy to be defined, and thus do not need to assume that a well-defined visible center exists at all. All these parameters are commonly dubbed ``non-parametric'', since they do not rely on certain model parameter fits, such as the S\\'ersic index \\citep{ser63} in models of radial surface brightness profiles. A more appropriate term might therefore be ``model-independent''.  One particular application of these model-independent measures was the identification of galaxy mergers. Studying their abundance and properties is crucial to understand the origin of the most massive galaxies in the Universe today \\citep[e.g.][]{bell06b}, of which many must have formed within the last gigayears \\citep{bell04, ferlis05}. \\citet{con03b} selected major galaxy mergers as objects with a rest-frame $B$-band asymmetry index of $A\\geq 0.35$, and used it to derive merger fractions and subsequently galaxy merger rates out to a redshift of $z\\approx 3$. By applying this structural selection to submillimeter-detected galaxies, \\citet{con03c} concluded that many of these are undergoing a major merger. \\citet{lot06} studied the fraction of major mergers and minor merger candidates among high-redshift star-forming galaxies, using a classification based on Gini, $M_{20}$, and concentration. A major merger selection criterion within the  two-dimensional parameter space of Gini ($G$) and $M_{20}$ was established by \\citet{lot08} as $G > -0.14\\cdot M_{20} + 0.33$, enabling them to study the evolution of the galaxy merger rate.  The Gini coefficient was found to correlate strongly with stellar mass \\citep{zam07}. In fact, \\citeauthor{zam07} claim that Gini traces better the overall structure of a galaxy than any other morphological parameter of their study. Apart from describing the overall structure, the Gini coefficient was also used to identify substructure in galaxies: \\citet{celeste} presented a preliminary method to automatically identify large bars, using the radial variation of Gini within a galaxy.  The Gini coefficient therefore plays an important role in present-day attempts to understand the formation and evolution of galaxy structure, and in particular, to unveil the formation process of massive galaxies. Despite simulations and a direct comparison of GOODS and UDF images \\citep{lot04,lot06}, which were done in order to quantify how Gini and other parameters depend on the signal-to-noise ratio and angular size resolution, the stability of the Gini coefficient has not been fully examined yet. In what follows, I present a critical assessment of the stability of Gini with respect to the signal-to-noise ratio and the choice of aperture. ",
        "conclusions": "So far, most authors using the Gini coefficient to describe galaxy morphology applied a reasonable magnitude and/or $S/N$ selection that prevented their investigations from being seriously affected by $S/N$ effects. However, I have shown that the Gini coefficient depends strongly on the aperture within which it is computed, and that it exhibits a strong dependence on $S/N$ below a certain $S/N$ level (which, again, depends on the aperture). Therefore, quantitative selection criteria involving the Gini coefficient -- such as the merger selection of \\citet{lot08} -- cannot be straightforwardly applied to a different dataset than the one for which they have been established \\citep[cf.][]{con08}. Care needs to be taken with the selection of aperture and limiting magnitude, as well as with the comparison of calculated Gini values to those of other studies. From the analysis presented here, the use of the Petrosian aperture is recommended. Larger apertures are strongly disfavored due to their inclusion of many faint pixels.  Typically, the total $S/N$ or the $S/N$ per pixel are quoted as quantitative, but crude estimates of the ``quality'' of a galaxy's image. From these average values, it can of course not be seen whether a given galaxy image actually consists of pixels with very different individual $S/N$ values, or whether the $S/N$ of most pixels is similar --- which might make a huge difference depending on the desired analysis. This pixel flux distribution is exactly what the Gini coefficient, by definition, measures. One could thus conclude somewhat ironically that the Gini coefficient could serve as a sophisticated measure of the $S/N$ distribution within a galaxy's image, if it did not depend so much on the galaxy's morphology."
    },
    "0807/0807.4778_arXiv.txt": {
        "abstract": "A large drift in the rotation rate of Titan observed by Cassini provided the first evidence of a subsurface ocean isolating the massive core from the icy crust.  Seasonal exchange of angular momentum between the surface and atmosphere accounts for the magnitude of the effect, but observations lag the expected signal by a few years.  We argue this time lag is due to the presence of an active methane weather cycle in the atmosphere.  An analytic model of the seasonal cycle of atmospheric angular momentum is developed and compared to time-dependent simulations of Titan's atmosphere with and without methane thermodynamics.  The disappearance of clouds at the summer pole suggests the drift rate has already switched direction, signaling the change in season from solstice to equinox. ",
        "introduction": "Models of the interior of Titan for plausible ammonia concentrations predict the existence of a water ocean under a thin ice I crust of perhaps 100 km thickness \\citep{2003JGRE..108.5130S}.  A subsurface ocean would dynamically isolate the surface from the more massive core, effectively lowering the moment of inertia of the surface.  For the expected magnitude of seasonal angular momentum change in the atmosphere, a thin ice crust experiences measurable changes in length-of-day, or equivalently spin rate, in a matter of a few years \\citep{2005GeoRL..3224203T}.  Cassini RADAR spin rate measurements \\citep{2008AJ....135.1669S}, seasonal winds from an atmospheric model prediction \\citep{2005GeoRL..3224203T} and Huygens wind observations \\citep{Bird:2005ai} were used to infer the moment-of-inertia (MOI) of Titan's surface \\citep{2008Sci...319.1649L}.  The surface MOI is well below that allowed for any solid internal configuration, but is consistent with models of Titan's interior that with a thin surface shell overlaying a subsurface ocean \\citep{2003JGRE..108.5130S}.  Seasonally changing zonal winds in the atmosphere torque the freewheeling surface, introducing a forced libration about synchronous rotation.  The forced libration of the surface, however, unexpectedly lags seasonal forcing by a few years.  Here we suggest this time lag signals an active methane weather cycle in the atmosphere.  Titan's atmosphere is known to superrotate, with angular momentum exceeding that of solid body rotation.  Stratospheric zonal winds exceed 100 m/s in the prograde direction \\citep{1993Icar..103..215H, 2008Icar..194..263A}.  Between 10 and 60 km, the troposphere weakly superrotates with winds limited to 10-20 m/s \\citep{Bird:2005ai}.  The lowest $\\sim$10 km at low latitudes show slight retrograde motion, which is consistent with a seasonal, cross-equatorial circulation, or Hadley cell, driven by convection \\citep{2007P&SS...55.1990T}.  Winds in the lowest kilometer turn prograde again \\citep{2007P&SS...55.1896K}, which is corroborated by the directionality of equatorial dune features \\citep{Lorenz:2006fv}.\\footnote{The presence of prograde winds at the equatorial surface is mysterious given frictional coupling with the surface.  Explanation of this phenomenon is beyond the scope of this work.}  Atmospheric models suggest angular momentum is seasonally redistributed in the atmosphere by a cross-equatorial circulation that changes direction with seasons \\citep{Hourdin.95, Tokano:1999fk, Tokano.et.al.05, Rannou.et.al.06, Mitchell:2006lr, 2007P&SS...55.1990T, Mitchell:2008}.  Though stratospheric zonal winds dwarf those in the troposphere, very little mass is being advected in the stratosphere.  Zonal winds in more than 90\\% of the atmospheric mass are below 20 m/s.  The lowest one-third of the atmospheric mass is subject to the largest seasonal change in angular momentum due to frictional torque from the surface acting on a seasonally reversing Hadley cell.  In the following, we develop a theory for the magnitude of seasonal change in atmospheric angular momentum (AAM) and its effect on the spin rate.  We then develop a full time-dependent model of this process and show by comparison with numerical simulations that a seasonal time lag of a few years is a natural consequence of methane thermodynamics in Titan's atmosphere.  We conclude by offering predictions for current and future Cassini observations. ",
        "conclusions": "We have developed an analytical model for the time-dependent atmospheric angular momentum in an atmosphere with an oscillating Hadley cell which agrees with a simulation of Titan's atmosphere \\citep{2005GeoRL..3224203T}; while both models predict the observed magnitude of spin rate drift, neither can produce the observed time lag of the seasonal drift.  Numerical simulations with methane thermodynamics included show an increase in the thermal inertia of the system, thus introducing a time lag with the seasonal forcing.  The time lag is sensitive to the supply of methane evaporation at the surface.  If methane is allowed to efficiently evaporate from the model surface, deep precipitating convection persists at all seasons and warms the high-heat-capacity atmosphere which increases the thermal inertia of the system.  In this ``moist'' case, the convective updraft of the Hadley circulation is in quadrature with solar forcing, which translates to a $\\sim$ 7 year time lag.  If instead surface evaporation is uniformly restricted, the numerical model has the appropriate amount of time lag, matching the observed lag in the spin rate drift (see Figure \\ref{fig2}).  Our model predicts the spin rate drift has changed directions; current Cassini RADAR observations of Titan may show this trend.  In conclusion, the observed time lag of Titan's drift in rotation rate is consistent with the lag introduced by thermodynamic-dynamic coupling resulting from the methane cycle in the atmosphere of Titan.  We have shown the change in total atmospheric angular momentum is primarily controlled by the latitude of the Hadley cell updraft, and this updraft marks the location of the most persistent clouds features.  The recent disappearance of persistent south polar clouds \\citep{Schaller.et.al.06b} suggests the Hadley cell updraft is moving equatorward in response to the changing season.  The resulting transfer of angular momentum to the atmosphere from the surface should currently be decelerating the spin rate, as shown in Figure \\ref{fig2}.  Other sources of torque might measurably contribute to the observed spin rate drift; however better knowledge of Titan's gravitational moments and interior structure is required to estimate their effects."
    },
    "0807/0807.1477_arXiv.txt": {
        "abstract": "We present new multi-band imaging data in the optical ($BVRI$ and \\ha) and near infrared bands ($JHK$) of 15 candidate ring galaxies from the sample of \\citet{app87}. We use these data to obtain color composite images, global magnitudes and colors of both the ring galaxy and its companion(s), and radial profiles of intensity and colors. We find that only nine of the observed galaxies have multi-band morphologies expected for the classical collisional scenario of ring formation, indicating the high degree of contamination of the ring galaxy sample by galaxies without a clear ring morphology. The radial intensity profiles, obtained by masking the off-centered nucleus, peak at the position of the ring, with the profiles in the continuum bands broader than that in the \\ha\\ line. The images as well as the radial intensity and color profiles clearly demonstrate the existence of the pre-collisional stellar disk outside the star-forming ring, which is in general bluer than the disk internal to the ring. The stellar disk seems to have retained its size, with the disk outside the ring having a shorter exponential scale length as compared to the values expected in normal spiral galaxies of comparable masses. The rings in our sample of galaxies are found to be located preferentially at around half-way through the stellar disk. The most likely reason for this preference is bias against detecting rings when they are close to the center (they would be confused with the resonant rings), and at the edge of the disk the gas surface density may be below the critical density required for star formation. Most of the observed characteristics point to relatively recent collisions ($<80$~Myr ago) according to the N-body simulations of \\citet{ger96}. ",
        "introduction": "Ring galaxies are a class of objects whose optical appearance is dominated by a ring or a ring-like structure. In one of the earliest discussions of ring galaxies, \\citet{bur59} suggested that these objects may be the aftermath of a close collision between an elliptical and spiral galaxy. On the other hand, \\citet{fre74} suggested that the ring is probably the result of a collision between a spiral galaxy and an intergalactic cloud of \\hi. \\citet{the76,the77} systematically studied the basic observational properties of a sample of ring galaxies and suggested that rings are formed when an intruding galaxy passes nearly through the center of a normal disk galaxy. \\citet{lyn76} used numerical simulations and settled the issue regarding their origin. They demonstrated that ring galaxies are formed as a result of an on-axis collision between an intruder galaxy and a gas-rich disk-galaxy. In this scenario, the collision sets off an expanding ring wave, which in turn triggers star formation in its wake. Due to the expanding nature of the wave, the ring of star formation also advances successively to larger radii with time. This scenario has stood the test of the multi-band data those have become available since then (Appleton \\& Struck-Marcell 1996 and references therein).  Many of the predictions from theoretical models have been tested observationally in the Cartwheel, the prototype ring galaxy. The ring in this galaxy is expanding \\citep{fos77}, and is forming massive stars \\citep{hig95}. A radial color gradient was also noticed by \\citet{mar92}, which was found to be in agreement with the sequential ordering of stellar ages \\citep{kor01, vor01}. \\citet{app97} established the color gradient in some other ring galaxies. \\citet{mar95} carried out \\ha\\ imaging observations of eight ring galaxies and found that majority of the star-forming regions are located exclusively in the ring.  In the collisional scenario proposed by \\citet{lyn76}, a stellar density wave is set off as the collective response of the stars that were present in the pre-collisional disk of the target galaxy. However, we know very little about the underlying stellar disk of the target galaxy. It is generally believed that the observed rings, especially in the near infrared continuum bands, trace the location of the stellar density waves. The stellar disk outside the ring should carry important information on the nature of the pre-collisional disk. However, such a disk has not been traced even in the Cartwheel. \\citet{app97} tried to infer an outer disk using radial profiles of surface brightness and color in four galaxies, and found such a disk in two cases (IIHz4 and VIIZw466). Their images were not deep enough to study the azimuthal structure of the outer disks.  \\citet{kor01} have developed a method to estimate the contribution of the newly formed stars in different wavelength bands as the density wave expands in the target galaxy. The density wave in the underlying stellar disk can be recovered by subtracting the contribution of the newly formed stars. Availability of \\ha\\ and continuum images are fundamental in achieving this objective. The continuum images should reach at least 2 magnitudes deeper than the brightness of the ring in order to register the stellar disk on either side of the star-forming ring. With this goal in mind, we carried out new imaging observations of a sample of 15 candidate ring galaxies in the $BVRIJHK$ broad bands and in the emission line of \\ha. Some of the candidate galaxies may not have formed by the scenario proposed by \\citet{lyn76}, and these galaxies were included in our observing list with the hope that the new observations would allow us to filter out contaminating galaxies. The data and detailed surface photometric analysis of all sample galaxies are presented in this paper. In a forthcoming paper (Paper~II), we will discuss the results obtained from the extraction of the underlying stellar density wave in galaxies that are most likely formed by the classical collisional scenario.  Sample and the details of our observations are discussed in \\S~2. Morphological descriptions and one dimensional profiles of individual galaxies are presented in \\S~3. In \\S~4, we compare the photometric properties of ring galaxies with that of normal galaxies. Results obtained from the new data are discussed in the context of N-body simulations in \\S~5. Conclusions from our study are presented in \\S~6. Gray scale maps in the \\ha\\ and $B$-bands and one dimensional surface brightness and color profiles for each galaxy are presented in the Appendix. Throughout this paper, all distance scaling assumes a value for the Hubble Constant of 75 km\\,s$^{-1}$\\,Mpc$^{-1}$. ",
        "conclusions": "We have carried out an analysis of the optical and NIR images of 15 galaxies in the northern hemisphere from the sample of \\citet{app87}. The images are used to establish a variety of characteristics common to these systems. These characteristics are summarized below.  All the 15 galaxies studied show evidence for recent star formation in the form of detection of FIR and \\ha\\ emission. In general, the SFR derived using the FIR is higher than that derived using the \\ha, with the differences suggesting a median visual extinction of $A_v=2.0$~mag. Arp145 and Arp148 are the two galaxies where the observed FIR and \\ha\\ fluxes suggest effective visual extinction in excess of 4~mag. The median SFR (7~\\msunyr) and \\ha\\ equivalent width (44\\,\\AA) suggest that their current star formation activity is similar to that found in other kinds of interacting galaxies. The observed 20~cm fluxes of the sample galaxies seem to be having a higher thermal fraction as compared to that in normal spiral galaxies.  All the 15 galaxies studied show morphological signs of having suffered an interaction in the recent past. However, only nine of the galaxies have morphologies consistent with a ring-making collision of the type described by \\citet{lyn76}.  Our sample of ring galaxies shows the sequential color-color relation found in other ring galaxies by \\citet{mar92}. However, we found that the relation for ring galaxies is not very much different from that seen for the normal spiral galaxies --- and that the characteristic that distinguishes the ring galaxies from the normal galaxies in a color-color plot is that the ring galaxies occupy only the blue end of the relation for the normal spiral galaxies.  We were able to trace the underlying stellar disk well beyond the ring in all our sample galaxies. The outer extension of the stellar disk is in general symmetric around the ring, with noticeable asymmetries present in Arp141, Arp143, Arp147 and NGC2793. The disk colors in the external parts are redder than the ring, but are in general bluer than that at the ring-center.  The ring-center of our sample galaxies is fainter (at least by a magnitude in $B$), and bluer (by around 0.2 mag in $B-V$) as compared to the values at the disk centers of normal spiral galaxies. Recent star formation in the wave is expected to make the inner disk bluer (as is observed) and brighter (contrary to what is observed). Thus, the observed low surface brightness of the ring-center implies either considerable fraction of the inner disk stars have splashed out of the disk or that the pre-collisional galaxy was of low surface brightness.  The outer disk scale lengths of ring galaxies are found to be systematically smaller than that of normal galaxies of comparable mass. On the other hand, the disk radius $R_{B25}$, measured at the 25~\\magarc\\ on azimuthally averaged $B$-band profiles, seems unchanged due to collisions. Under the N-body simulations of \\citet{ger96} these results imply time elapsed since the collision less than 80~Myr.  The physical sizes of the ring in our sample varies from 2--12~kpc. However, the observed rings seem to be systematically located around half-way through the disk. This is most likely due to the difficulties in identifying a ring galaxy in its initial phase of evolution. The outer parts of the galaxies may not have sufficient gas surface densities to form stars."
    },
    "0807/0807.0647_arXiv.txt": {
        "abstract": "The Hubble constant, $H_0$, sets the scale of the size and age of the Universe and its determination from independent methods is still worthwhile to be investigated. In this article,  by using the Sunyaev-Zel`dovich effect and X-ray surface brightness data from 38 galaxy clusters observed by Bonamente {\\it{et al.}} (2006), we obtain a new estimate of $H_0$ in the context of  a flat $\\Lambda$CDM model. There is a degeneracy on the mass density parameter ($\\Omega_{m}$) which is broken by applying a joint analysis involving the baryon acoustic oscillations (BAO) as given by Sloan Digital Sky Survey (SDSS). This happens because the BAO signature does not depend on $H_0$. Our basic finding is that a joint analysis involving these tests yield $H_0= 0.765^{+0.035}_{-0.033}$ km s$^{-1}$ Mpc$^{-1}$ and $\\Omega_{m}=0.27^{+0.03}_{-0.02}$. Since the hypothesis of spherical geometry assumed by Bonamente  {\\it {et al.}} is questionable, we have also compared the above results to a recent work where a sample of triaxial galaxy clusters has been considered. ",
        "introduction": "Ten years ago, two different teams using SNe type Ia as standard candles announced that the expansion of the Universe is speeding up and not slowing as believed by the cosmology community\\cite{RiessPerl}. The authors interpreted their  results as implying  that the cosmological constant $\\Lambda$ is greater than zero, an assumption that is not only surviving but has  been strengthened  in the last decade by many independent astronomical observations\\cite{Astier,Sper07,LSS,BAO,Allen,Ages} (see also Peebles and Ratra\\cite{Ratra}, and Padmanabhan\\cite{PAD} for  reviews).  On the other hand, until now there is no consensus about the present value of the Hubble parameter even when the estimates of $H_0$ are limited to the Cepheid period-luminosity relation and the correlation of the luminosity at maximum light for SNe type Ia.  For instance, Riess {\\it et al.}\\cite{Riess05} using a hybrid sample composed by Cepheids in galaxies and recent type Ia SNe concluded that $H_0=73\\pm 6$ $km.s^{-1}.Mpc^{-1}$ while Sandage\\cite{Sand06} and coworkers based on data from the Hubble Space Telescope (HST) key project are claiming that $H_0=62\\pm 5$ $km.s^{-1}.Mpc^{-1}$. Therefore, is still worthwhile to improve the  direct local estimates of $H_0$, and, probably, more important, to investigate their concordance with independent methods. An interesting one is provided by the combination of Sunyaev-Zel'dovich Effect (SZE) and measurements of X-ray surface brightness from galaxy  clusters.  The so-called SZE is a small distortion on the Cosmic Microwave Background (CMB) spectrum provoked by the inverse Compton scattering of the CMB photons passing through a population of hot electrons\\cite{SunZel72,Itoh98}. Observing the temperature decrement of the CMB in the direction of galaxy clusters and the X-ray spectrum emitted by the hot electron gas pervading the cluster, it is possible to break the degeneracy between concentration and temperature\\cite{DEG}, and, therefore, to calculate  the angular diameter distance (ADD)\\cite{SilWhi78,Barttlet04,Caval79}. This technique for measuring distances is completely independent of other methods (as the one provided by the luminosity distance), and it can be used to measure distances at high redshifts directly. Review papers on this subject have been published by Birkinshaw\\cite{Birki99} and Carlstrom, Holder and Reese\\cite{Carlstrom02}. More recently,  such a technique has been applied for a fairly large number of clusters\\cite{Reese02,Jones05,DeFilippis05,Boname06,Roque06}.  In 2006, Bonamente and collaborators\\cite{Boname06} determined the ADD distance to 38 clusters of galaxies in the redshift range $0.14 \\leq z \\leq 0.89$ using  X-ray data from Chandra and Sunyaev-Zeldovich effect  data  data from the Owens Valley Radio Observatory and the Berkeley-Illinois-Maryland Association interferometric arrays. Assuming spherical symmetry, the cluster plasma and dark matter distributions were analyzed by using a hydrostatic equilibrium model accounting for radial variations in density, temperature and abundances. In addition, by fixing  the earlier cosmic concordance $\\Lambda$CDM model ($\\Omega_m=0.3$, $\\Omega_{\\Lambda}=0.7$), they obtained for the Hubble constant $H_0= 76.9^{+3.5 +10.0}_{-3.4 -8.0}\\, km. s^{-1}. Mpc^{-1}$ (statistical followed by systematic uncertainty at 1$\\sigma$ c. l.). The main advantage of this method for estimating $H_0$ is that it does not rely on the extragalactic distance ladder being fully independent of any local calibrator. In a point of fact, once the angular diameter distance has been measured a large amount of astrophysical and cosmological tests can be performed in order to constrain the relevant free parameters\\cite{Barttlet94,Rephaeli95,HKobayashi}.  It should be stressed, however, that in the above quoted determination of the Hubble constant, a specific cosmology was fixed from the very beginning. This is needed because the determination of the Hubble parameter from SZE/X-ray technique is endowed with a strong degeneracy with respect to the matter density parameter which can only be broken by considering additional cosmological tests. An interesting possibility is given by the signature of the  baryon acoustic oscillations (BAO) from the last scattering surface (LSS).  In 2005, Eisenstein {\\it{et al.}}\\cite{Eisenstein05} presented the large scale correlation function from the Sloan Digital Sky Survey (SDSS) showing clear evidence for the baryon acoustic peak at $100 h^{-1}$ Mpc scale, a result in good agreement with the analyses from CMB data\\cite{Sper07}. The breaking on the degeneracy between $H_0$ and $\\Omega_m$ by applying this BAO signature is possible because it does not depend on $H_0$ and is highly sensitive to the matter density parameter. The combination  SZE/X-ray with BAO has been recently discussed in the literature by Cunha, Marassi and Lima\\cite{cunha}.  In this article, we discuss the determinations of $H_0$ and $\\Omega_m$ by considering the same 38 clusters from Bonamente {\\it{et al.}} sample where spherical symmetry was assumed. Our basic findings follow from a joint analysis involving the data from SZE and X-ray surface brightness with the recent SDSS measurements of the baryon acoustic peak. The influence of  the intrinsic cluster geometry  on $H_{0}$ it will be quantified by comparing the results derived here with a recent determination of both parameters\\cite{cunha} based on De Filippis {\\it{et al.}} sample formed by  25 triaxial clusters\\cite{DeFilippis05}. ",
        "conclusions": "In this paper we have discussed a new determination of the Hubble constant based on the SZE/X-ray distance technique for a sample of 38 clusters as compiled by Bonamente {\\it et al.} assuming spherical symmetry\\cite{Boname06}. The degeneracy on the $\\Omega_{m}$ parameter was broken trough a joint analysis applying the baryon acoustic oscillation signature from the SDSS catalog. The Hubble constant was constrained to be $h = 0.765^{+0.035}_{-0.033}$ for $1\\sigma$ and $\\Omega_{m}=0.273^{+0.03}_{-0.02}$. These limits were derived assuming the flat $\\Lambda$CDM scenario and a spherical $\\beta$-model. The central $h$ value derived here is in perfect agreement with the Bonamente {\\it{et al.}} value, $h=0.769$ (assuming a hydrostatic equilibrium  and fixing the concordance model), as well as with others recent estimates coming from WMAP and Hubble Space telescope Key Project, where $h=0.73$. However, it differs slightly of a recent estimation  by Cunha, Marassi \\& Lima, $h=0.74$,  where  the SZE/X-Ray + BAO technique was applied for 24 triaxial galaxies clusters. In general ground, such results are suggesting that the combination of these three independent phenomena (SZE/X-ray and BAO) provides an interesting method to constrain the Hubble constant."
    },
    "0807/0807.2642_arXiv.txt": {
        "abstract": "The luminosities of SGRB host galaxies are anticorrelated with both the isotropic equivalent gamma ray energy and the gamma ray luminosity of the explosions.  Observational selection effects only strengthen the significance of this correlation.  The correlation may indicate that there are two physically distinct groups of SGRBs.  If so, it requires that the more luminous class of explosions be associated with the younger class of progenitors.  Alternatively, it could be due to a continuous distribution of burst and host properties.  As one possible explanation, we find that the effect of binary neutron star masses on inspiral time and energy reservoir produces a correlation of the appropriate sign, but does not automatically reproduce the correlation slope or the full range of SGRB energy scales.  Any future model of SGRB progenitors needs to reproduce this correlation. ",
        "introduction": "\\label{sec:intro} The host galaxies of short gamma ray bursts (SGRBs) exhibit a wide range of physical properties. SGRBs have been found not only in star-forming hosts, but also in elliptical galaxies with strong upper limits on their star formation rate (hereafter ``quiescent hosts'') (Prochaska et al 2006; Berger 2008). In contrast, long gamma ray bursts (LGRBs) are found exclusively in star forming galaxies (e.g., Le Floc'h et al 2003, Fruchter et al 2008, Savaglio et al 2009). Several lines of evidence now link LGRBs with the deaths of massive stars (Hjorth et al 2003, Stanek et al 2003, Fruchter et al 2006, Raskin et al 2008). The origin of the SGRBs is not firmly established, but suspicion centers on merging neutron star binaries, or other phenomena primarily associated with neutron stars (Belczynski et al 2006; Chapman, Priddey, \\& Tanvir 2008).  Hosts offer one of the best tools for unveiling the nature of the GRB progenitors, because the properties of the gamma ray emission and afterglow are decoupled from many important details of the progenitor objects under the fireball model of GRBs (e.g. Paczy\\'{n}ski \\& Rhoads 1993, Katz 1994, Piran 2005).  The situation is similar to supernovae, which divide into two groups, the type Ia and the core collapse events (which include types II, Ib, and Ic). Core collapse events are found only in star forming galaxies, and are associated with the deaths of stars having initial masses $\\ga 8 \\msun$. In contrast, type Ias are found in both quiescent and star forming hosts, and are thought to be powered by the nuclear explosion of a white dwarf whose mass is pushed above the Chandrasekhar limit by accretion.  However, despite having a characteristic progenitor mass set by fundamental physics, the Ia supernovae are not entirely homogeneous in their properties.  Those found in star forming hosts are both more frequent (per unit stellar mass) and more luminous than their cousins in quiescent galaxies (Hamuy et al 1995; Scannapieco \\& Bildsten 2005; Sullivan et al 2006).  In this {\\it Letter}, we look for similar trends linking the propertie of SGRBs and their hosts.  We find an anticorrelation between SGRB isotropic energy and host galaxy luminosity.  We find no obvious selection effects that could produce such an effect, nor do we find evidence for any similar effect in a larger control sample of LGRBs. ",
        "conclusions": "While our sample is small, the anticorrelation between SGRB $\\eiso$ and host luminosity is probably real, based on our Monte Carlo simulations.  The similar anticorrelation between brightness of type Ia supernovae and the stellar population ages of their host galaxies was first described on the basis of a similarly small sample (13 supernovae; Hamuy et al 1995). We now examine possible interpretations for this result.  Suppose, first, that two distinct mechanisms produce SGRBs.  Such a situation would arise naturally if two or more of the different proposed progenitors actually do produce SGRBs.  Combinations that have been explored include coalescence of double neutron star binaries and of NS-black hole binaries (Belczynski et al 2006; O'Shaughnessy et al 2008; Troja et al 2008), and combinations of NS-NS binaries with giant flares from magnetars (Chapman, Priddey, \\& Tanvir 2008). Additionally, a range of SGRB properties might arise from NS-NS binaries alone, provided those binaries evolved through a range of different binary interaction histories (Belczynski et al 2006; Salvaterra et al 2008).  If one mechanism is associated with stars younger than 1 Gyr, and the other with ancient populations (age $\\sim H_0^{-1}$), we can explain the properties of figures~\\ref{fig:eiso_absmag}a and \\ref{fig:ssfr_lgam}a. The fraction of old-progenitor SGRBs occurring in quiescent galaxies should match the fraction of stellar mass in quiescent galaxies.  The present, limited data shows two bursts in quiescent hosts, and a third similarly low-luminosity burst (070724) in a star-forming host galaxy. This matches well the fraction of stellar mass in elliptical galaxies at low redshift, which is 54\\% to 60\\% (Baldry et al 2004). Conversely, the fraction of more luminous SGRBs occurring in quiescent hosts is $0/9$ in the present sample (fig.~\\ref{fig:ssfr_lgam}a).  Under a two-population scenario, the mechanism responsible for the more luminous SGRBs should have a negligible rate for ages exceeding a few Gyr.  If instead there is a single, continuous distribution of SGRB and host properties, it suggests that the SGRB progenitor's properties vary systematically with either age or metallicity, which vary systematically along the Hubble sequence and have direct effects on stellar-scale physics.  If SGRBs are due to binary neutron star inspiral events, an anticorrelation of SGRB energy and the specific star formation rate of the host galaxy could come from the dependence of inspiral time and available energy on mass. Gravitational radiation gives an inspiral time of \\begin{equation} t_{inspiral} = {5 c^5 \\over 256 G^3} {a_0^4 \\over M_1 M_2 (M_1 + M_2)} ~ \\end{equation} (e.g., Landau \\& Lifschitz 1958), where $M_1$ and $M_2$ are the masses of the two bodies, $a$ is the semimajor axis of their orbit, and $G$ and $c$ are Newton's constant and the speed of light. Thus, if a star formation event yields a population of compact object binaries with a range of mass, we expect a range of inspiral times. The SGRBs occurring in quiescent hosts would be those with low masses and/or wide initial separations.  The available energy reservoir is $E \\sim G M_1 M_2 / R_{final}$, where $R_{final}$ is the characteristic size of the system when the GRB energy is released.  If the SGRB event produces a black hole, we expect $R_{final} \\sim R_{grav} \\propto (M_1+M_2)$, so that $E\\sim M_1 M_2 c^2 / (M_1 + M_2)$.  Then, for $M_1 \\approx M_2$, $t_{inspiral} \\sim a_0^4 E^{-3}$. Characteristic stellar ages for SGRB hosts range from $\\sim 10$ Gyr for quiescent hosts, to $\\sim s^{-1} \\sim 1$ Gyr for typical star forming SGRB hosts (see, e.g., Savaglio et al 2008). If this tenfold range of age corresponds to a tenfold range of $t_{inspiral}$, it would imply a range of $\\sim 2\\times$ in neutron star mass, which is comparable to the range of NS masses believed to exist in our Galaxy.  However, the corresponding $2\\times$ range of energy cannot explain the range of \\eiso\\ observed in figure~\\ref{fig:eiso_absmag}a.  A larger range in SGRB energetics might follow if $R_{final} \\sim R_{NS}$, especially for a neutron star equation of state that gives $d R_{ns} / dM < 0$. Moreover, $a_0$ may correlate with $M_1$ and $M_2$ in some complex way depending on on both binary star mass transfer and detailed supernova physics.  While a more thorough understanding of the SGRB mechanism is clearly required before we can fully model the relation between host and SGRB properties, the existence of such a relation will provide valuable clues to the nature of the short GRBs.  The correlation in figure~\\ref{fig:eiso_absmag}a can also provide a tool for estimating SGRB redshifts. $\\log(\\eiso)$ and $B_{abs}$ are related empirically by $\\log(\\eiso / \\erg) \\approx 50 + 0.725 (B_{abs} + 20.2)$. Defining ${\\cal D} = d_L / \\sqrt{1+z}$, so that $\\eiso = (4\\pi {\\cal D}^2) f$ and $B_{abs} = R_{AB} - 5 \\log[({\\cal D} / 10pc)]$, we can substitute for $\\eiso$ and $B_{abs}$ to obtain a distance estimate: \\begin{equation} \\log(\\De / \\hbox{Gpc}) = 0.012 + 0.129 (R_{AB}-20) - 0.177 \\log(f_{-6}) ~~, \\label{eq:dist_est} \\end{equation} where $f_{-6}$ is the SGRB fluence in units of $10^{-6} \\erg \\, \\cm^{-2}$. The conversion from $\\De$ to the estimated distance $d_{L,est}$ and redshift $z_{est}$  must of course follow the cosmology we used to derive equation~\\ref{eq:dist_est}. For our primary sample of 12 SGRBs, the rms residual in $\\log(\\De)$ is $0.17$ dex, and in redshift $z(\\De)$ is 0.147. The residuals from the fit substantially exceed the uncertainties expected from fluence and host magnitude measurements, implying that most of this scatter is intrinsic.  As an alternative, we can directly fit a relation of the form $\\log(z_{est}) = \\log(z_0) + a (R_{AB} - 20) + b \\log(f_{-6})$ by minimizing $\\Delta z_{rms} = \\langle (z_{est} - z_{obs})^2 / (N-1)\\rangle^{1/2}$. This approach gives $\\log(z_{est}) = -0.56 + 0.107 (R_{AB}-20) - 0.14 \\log(f_{-6})$, with $\\Delta z_{rms} =  0.142$. This redshift estimator is easier to apply, and is formally independent of the adopted cosmological parameters, though it is less directly linked to whatever physical mechanism links $\\eiso$ and $B_{abs}$.  Differences between the predicted redshifts from the two equations are all much smaller than the scatter in either.  Another possibility is that the external medium around SGRBs plays some role in their luminosity.  The interstellar medium is typically less dense in elliptical galaxies than in actively star forming galaxies.  However, this should only affect the observed flux from external shocks that GRB ejecta drive into the ambient medium.  External shocks are believed to power the afterglow emission, but not the GRB luminosity itself (e.g., Piran 2005).  Since our observed anticorrelation involves the SGRB $\\gamma$-ray flux itself, we disfavor this explanation.  To conclude, we have found an anticorrelation between the energy of a short gamma ray burst and the luminosity of its host galaxy.  Such an anticorrelation occurs with a probability of $\\sim 1\\%$ in simulations that account for observational selection effects, and so is probably real.  Its physical origin is unclear, though it is most likely due to a correlation between the age of an SGRB progenitor and the luminosity of the explosion.  If this correlation is a continuous distribution, it can provide approximate redshift estimates for SGRBs.  Time will tell whether this correlation holds for a larger sample, and whether it is a property of a single distribution, or if it instead reflects an underlying division of SGRBs into two physically distinct sets."
    },
    "0807/0807.1373.txt": {
        "abstract": "}}% %BeginExpansion \\begin{abstract}% %EndExpansion Magnetically driven astrophysical jets are related to accretion and involve toroidal magnetic field pressure inflating poloidal magnetic field flux surfaces. Examination of particle motion in combined gravitational and magnetic fields shows that these astrophysical jet toroidal and poloidal magnetic fields can be powered by the gravitational energy liberated by accreting dust grains that have become positively charged \\ by emitting photo-electrons. Because a dust grain experiences magnetic forces after becoming charged, but not before, charging can cause irreversible trapping of the grain so dust accretion is a consequence of charging. Furthermore, charging causes canonical angular momentum to replace mechanical angular momentum as the relevant constant of the motion. The resulting effective potential has three distinct classes of accreting particles distinguished by canonical angular momentum, namely (i) \\textquotedblleft cyclotron-orbit\\textquotedblright, (ii) \\textquotedblleft Speiser-orbit\\textquotedblright, and (iii) \\textquotedblleft zero canonical angular momentum\\textquotedblright\\ particles. Electrons and ions are of class (i) but depending on mass and initial orbit inclination, dust grains can be of any class. Light-weight dust grains develop class (i) orbits such that the grains are confined to nested poloidal flux surfaces, whereas grains with a critical weight such that they experience comparable gravitational and magnetic forces can develop class (ii) or class (iii)\\ orbits, respectively producing poloidal and toroidal field dynamos. %TCIMACRO{\\TeXButton{End abstract}{ ",
        "introduction": "Magnetohydrodynamically driven plasma jets having topology and dynamics analogous to astrophysical jets have been produced in laboratory experiments by %TCIMACRO{\\TeXButton{\\citet{Hsu2002,Bellan2005}}{\\citet{Hsu2002,Bellan2005}} }% %BeginExpansion \\citet{Hsu2002,Bellan2005} %EndExpansion and by %TCIMACRO{\\TeXButton{\\citet{Lebedev2005}}{\\citet{Lebedev2005}}}% %BeginExpansion \\citet{Lebedev2005}% %EndExpansion ; see discussion by %TCIMACRO{\\TeXButton{\\citet{Blackman2007}}{\\citet{Blackman2007}}}% %BeginExpansion \\citet{Blackman2007}% %EndExpansion . The feature of azimuthal symmetry, common to both the lab experiments and to real astrophysical jets, has important implications for the structure of the magnetic field. This is because an azimuthally symmetric magnetic field can be expressed using a cylindrical coordinate system $\\{r,\\phi,z\\}$ as$\\ $ \\begin{equation} \\mathbf{B=}\\frac{1}{2\\pi}\\left(  \\nabla\\psi\\times\\nabla\\phi+\\mu_{0}I\\nabla \\phi\\right)  \\ \\label{B general2}% \\end{equation} where the poloidal flux $\\psi(r,z,t)$ is defined by \\begin{equation} \\psi(r,z,t)=\\int_{0}^{r}2\\pi r^{\\prime}dr^{\\prime}B_{z}(r^{\\prime},z,t) \\label{psi def2}% \\end{equation} and the poloidal electric current $I(r,z)$ is defined by \\begin{equation} I(r,z,t)=\\int_{0}^{r}2\\pi r^{\\prime}dr^{\\prime}J_{z}(r^{\\prime},z,t). \\label{Idef2}% \\end{equation} The definition of $I(r,z)$ is consistent with Ampere's law for the toroidal field, since using $\\nabla\\phi=\\hat{\\phi}/r\\ $in Eq.\\ref{B general2} gives the toroidal magnetic field to be \\begin{equation} B_{\\phi}=\\frac{\\mu_{0}I}{2\\pi r}. \\label{Bphi}% \\end{equation}   Equations \\ref{B general2}-\\ref{Idef2} describe the magnetic field and electric current distribution of \\textit{any} axisymmetric magnetic field. Because astrophysical jets are azimuthally symmetric, their magnetic field must be of the form prescribed by Eqs.\\ref{B general2}- \\ref{Idef2} and, indeed, it is generally believed that astrophysical jets involve large-scale poloidal magnetic fields threading an accretion disk (e.g., see %TCIMACRO{\\TeXButton{\\citet{Livio2002,Ferreira2004}}{\\citet %{Livio2002,Ferreira2004}}}% %BeginExpansion \\citet{Livio2002,Ferreira2004}% %EndExpansion ) and in addition, toroidal magnetic fields. Application of Ampere's law to Eq.\\ref{B general2} shows that the poloidal and toroidal currents are respectively given by \\begin{equation} \\mathbf{J}_{pol}=\\frac{1}{2\\pi}\\nabla I\\times\\nabla\\phi\\label{Jpol2}% \\end{equation} and \\begin{equation} \\mathbf{J}_{tor}=\\ -\\frac{r^{2}\\nabla\\cdot\\left(  r^{-2}\\nabla\\psi\\right) }{2\\pi\\mu_{0}}\\ \\nabla\\phi\\label{Jtor}% \\end{equation} showing that poloidal magnetic fields are produced by a toroidal electric current and toroidal magnetic fields are produced by a poloidal current; toroidal vectors are those vectors in the $\\phi$ direction and poloidal vectors are those in any combination of the $r$ and $z$ directions. \\ Knowledge of the two stream-function quantities $I(r,z,t)$ and $\\psi(r,z,t)$ is thus necessary and sufficient to determine the complete vector magnetic field and the complete vector current density.  The term `magnetic axis' has traditionally been assigned different meanings in the respective contexts of astrophysics and laboratory toroidal magnetic confinement devices (e.g., tokamaks, reversed field pinches, or spheromaks). Specifically, a local maximum in $r$-$z$ space of $\\psi(r,z)$ is called a magnetic axis in the context of toroidal confinement devices whereas the $z$ symmetry axis of the magnetic field is called the magnetic axis in \\ the context of astrophysics. To avoid confusion, we will call the location of a maximum of $\\psi(r,z)$ the poloidal flux magnetic axis. Poloidal magnetic field lines follow level contours of $\\psi$ and so one can envision the projection of the magnetic field in the $r$-$z$ plane as being like a set of roads, each at a different altitude, encircling a mountain peak at a specific $r$-$z$ location which is the poloidal flux magnetic axis (also called an O-point). \\ Since a toroidal current at infinity is not physical, and since the net magnetic flux enclosed by a circle with infinite radius must vanish as field lines cannot go to infinity, $\\psi$ must vanish at infinity. Furthermore, mathematical regularity of physical quantities requires $\\psi$ to vanish on the $z$ axis %TCIMACRO{\\TeXButton{\\citep{Lewis1990}}{\\citep{Lewis1990}}}% %BeginExpansion \\citep{Lewis1990}% %EndExpansion . Thus, a non-trivial $\\psi$ can only be finite in the region $0<r<\\infty$, $-\\infty<z<\\infty.$ The simplest situation of physical interest is therefore where $\\psi$ has a single maximum in the $r$-$z$ plane. We will consider this situation, namely a single poloidal field magnetic axis with $\\psi(r,z)$ symmetric with respect to $z.$ This situation has been previously considered by %TCIMACRO{\\TeXButton{citet{Lovelace2002}}{\\citet{Lovelace2002}} }% %BeginExpansion \\citet{Lovelace2002} %EndExpansion and implies via Eq.\\ref{Jtor} that a toroidal current circulates in an accretion disk to produce the poloidal magnetic field \\begin{equation} \\mathbf{B}_{pol}=\\mathbf{\\ }\\frac{1}{2\\pi}\\nabla\\psi\\times\\nabla\\phi. \\label{Bpol def}% \\end{equation} An inescapable feature of this topology is that because $\\nabla\\psi=0$ at the maximum of $\\psi$, i.e., at the poloidal field magnetic axis, the poloidal magnetic field has a null on the poloidal field magnetic axis. In the $z=0$ plane, the poloidal flux $\\psi$ thus starts from zero at $r=0,$ increases to a maximum at the poloidal field magnetic axis, and then decays to zero as $r\\rightarrow\\infty.$ \\  We define $a\\ $to be the radius of the poloidal field magnetic axis. In addition, we define $\\left\\langle B_{z}\\right\\rangle $ to be the \\textit{spatially-averaged} axial magnetic field linked by the poloidal field magnetic axis and $\\psi_{0}$ to be the value of the poloidal magnetic flux at the poloidal field magnetic axis, so \\begin{equation} \\left\\langle B_{z}\\right\\rangle =\\frac{\\int_{0}^{a}dr2\\pi rB_{z}(r,0)}% {\\int_{0}^{a}dr2\\pi r}=\\frac{\\psi_{0}}{\\pi a^{2}}. \\label{Bbar}% \\end{equation} \\qquad\\qquad  The axial field $B_{z}=(2\\pi r)^{-1}\\partial\\psi/\\partial r$ reverses sign at $r=a$ and the radial field $B_{r}=-(2\\pi r)^{-1}\\partial\\psi/\\partial z$ reverses sign at $z=0.$ $\\ $An analytic representation for a physically realizable generic flux function satisfying all these properties is derived in Appendix \\ref{Generic Flux Function}. This generic flux function is \\begin{equation} \\psi(r,z)=\\ \\ \\frac{27\\left(  r/a\\right)  ^{2}\\psi_{0}}{8\\left(  \\left( \\frac{r}{a}+\\frac{1}{2}\\right)  ^{2}+\\left(  \\frac{z}{a}\\right)  ^{2}\\right) ^{3/2}} \\label{generic}% \\end{equation} and has the properties that (i)\\ $\\psi(r,z)$ has a maximum value of $\\psi_{0}$ at $r=a,$ $z=0,$ (ii) $\\psi\\sim r^{2}$ for $r\\ll a$ and $z=0,$ (iii) $\\psi\\sim r^{-1}$ for $r\\gg a,z$ and (iv) for $r\\ll a/2$ or $r\\gg a\\ $ and for $z\\gg a$ the contours of $\\psi$ are identical to the contours of the poloidal flux produced by a current loop located at $r=a/2,$ $z=0.$ This flux function thus encompasses simpler models which assume a uniform axial magnetic field $B_{z}% $; these simpler models would correspond to the $r,z\\ll a$ region here since in this region $\\psi\\sim r^{2}$ which corresponds to having a uniform axial magnetic field $B_{z}.$ This flux function could also\\ be used to describe the far-field of a dipole by assuming that $r,z\\gg a.$ Since any real axial magnetic field must always be generated by a toroidal current located at some finite radius, any real situation will have a poloidal flux function qualitatively similar to Eq.\\ref{generic}. The flux function prescribed in Eq.\\ref{generic} is similar in essence to the flux function used in Fig.1 of %TCIMACRO{\\TeXButton{citet{Lovelace2002}}{\\citet{Lovelace2002}}}% %BeginExpansion \\citet{Lovelace2002}% %EndExpansion .  Figure \\ref{GenericFluxFunction-new} plots $\\psi(r,z)$ as prescribed by Eq.\\ref{generic} and shows that $\\psi(r,z)$ has its maximum at the poloidal field magnetic axis $r\\ =a,$ $z=0$. This flux function corresponds to a smoothly varying toroidal current density prescribed by Eq.\\ref{Jtor} concentrated in the vicinity of $r=a,$ $z=0.$ Since for $z=0$ and small $r$, this function has the asymptotic dependence $\\psi\\simeq27\\psi_{0}\\left( r/a\\right)  ^{2}$, it corresponds to an approximately uniform axial magnetic field $B_{z}\\simeq27\\psi_{0}/\\pi a^{2}$ for $r,z\\ll a.$ The $r\\ll a$ inner-region $B_{z}$ is thus 27 times stronger than the average $B_{z}$ field between $0$ and $a.$ The total toroidal current $\\mathcal{I}_{\\phi}$ associated with the generic flux function given by Eq.\\ref{generic} is calculated in Appendix \\ref{current associated with flux function} using the integral form of Ampere's law and found to be \\begin{equation} \\mathcal{I}_{\\phi}=\\frac{27\\psi_{0}}{\\ \\pi a\\mu_{0}}\\ \\ . \\label{Iphi}% \\end{equation}   \\begin{figure}[ptb] \\caption{Plot of the normalized generic flux function $\\psi(r,z)/\\psi_{0}\\ $in coordinates normalized to the radius of the poloidal field magnetic axis, that is to the radial position of the maximum of $\\psi(r,z)\\,.$ Contours of iso-surfaces shown on top; these correspond to projection of poloidal magnetic field onto $r$-$z$ plane.}% \\plotone{f1.eps}\\label{GenericFluxFunction-new}% \\end{figure}  The laboratory jets involve the mutual interaction between poloidal and toroidal magnetic fields powered by laboratory capacitor banks. The jet acceleration mechanism results from the pressure of the toroidal magnetic field inflating flux surfaces associated with the poloidal magnetic field. The question arises as to what powers the toroidal and poloidal magnetic fields in an actual astrophysical situation. Existing models of astrophysical jets are based on the magnetohydrodynamic (MHD) approximation of plasma behavior and typically assume (i)\\ the poloidal field is pre-existing and (ii)\\ the toroidal field results from a rotating accretion disk twisting up this assumed primordial poloidal field. The purpose of this paper is to present an alternate model wherein it is postulated that the toroidal and poloidal field result instead from a non-MHD dusty plasma dynamo mechanism that converts the gravitational energy of infalling dust grains into an electrical power source that drives poloidal and toroidal electric currents creating the respective toroidal and poloidal fields. A brief outline of how infalling charged dust can drive poloidal currents has been presented in %TCIMACRO{\\TeXButton{\\citet{Bellan2007}}{\\citet{Bellan2007}}}% %BeginExpansion \\citet{Bellan2007}% %EndExpansion .  This model obviously requires existence of sufficient infalling dust to provide the jet power. Since the dust-to-gas mass ratio in the Interstellar Medium (ISM) is 1\\%, one might be tempted to argue that any jet driven by the proposed dust infall mechanism would be limited to having\\ less than 1\\% of the power available from infalling gas, a constraint that would contradict observations. However, in %TCIMACRO{\\TeXButton{\\citet{Bellan2008}}{\\citet{Bellan2008}} }% %BeginExpansion \\citet{Bellan2008} %EndExpansion (to be referred to as Paper I), we showed that the dust-to-gas mass ratio in a molecular cloud can be substantially enriched compared to the ISM\\ value (e.g., the dust to gas mass ratio in a molecular cloud could be enriched 20-fold compared to the 1\\% ISM value). This enrichment occurs because accreting dust slows down much more in proportion to its initial velocity than does accreting gas so that the density amplification resulting from dust slowing down is much greater than the corresponding density amplification of gas.  The condition for the toroidal magnetic field to inflate the poloidal magnetic field and create a jet can be expressed as \\begin{equation} \\frac{\\mu_{0}I}{\\psi}>\\lambda\\label{lambda threshold}% \\end{equation} where $\\lambda$ is a parameter of the order of the inverse characteristic linear dimension in the radial direction. The ratio $I/\\psi$ can be thought of as the ratio of the electric current flowing along a flux tube to the magnetic flux content of the flux tube and is proportional to the twist of the magnetic field. Equation \\ref{lambda threshold}, well-established in spheromak formation physics %TCIMACRO{\\TeXButton{\\citep %{Barnes1990,Jarboe1994,Geddes1998,Bellan2000,Hsu2005}}{\\citep %{Barnes1990,Jarboe1994,Geddes1998,Bellan2000,Hsu2005}}}% %BeginExpansion \\citep{Barnes1990,Jarboe1994,Geddes1998,Bellan2000,Hsu2005}% %EndExpansion , is essentially a statement that jet expansion (i.e., poloidal field inflation)\\ occurs when the toroidal magnetic field pressure force $\\sim B_{\\phi}^{2}A_{1}$ acting on area $A_{1}$ exceeds the restraining force $B_{z}^{2}A_{2}$ of the poloidal magnetic field `tension' acting on area $A_{2}$. Here $A_{1}$ and $A_{2}$ are not exactly the same because the toroidal and poloidal fields do not act over the same areas. \\ The equivalence between Eq. \\ref{lambda threshold} and the condition $B_{\\phi}^{2}>B_{z}% ^{2}A_{2}/A_{1}$ is seen by substituting $\\mu_{0}I=2\\pi aB_{\\phi}$ from Ampere's law and $\\psi\\sim B_{z}\\pi a^{2}$ in Eq.\\ref{lambda threshold}.  Paper I divided the regions of interest into successively smaller concentric regions and considered dust and gas behavior in the outermost regions. Simultaneous gas and dust accretion were considered and it was shown that the dust could be considered as a perturbation on the gas, so that the gas accretion problem could be solved first without considering dust and then the solution of this gas accretion problem could be used as an input for the dust accretion problem. Below is a listing showing which regions were considered in Paper I, which are considered in this paper, and which will be considered in a future paper; the nominal radii scales and star mass are from Table 3 in Paper I:  \\begin{description} \\item \\textit{ISM scale (considered in Paper I):} The outermost scale is that of the Interstellar Medium (ISM). The ISM\\ has a gas density $\\sim$10$^{7}$ m$^{-3}$, a dust-to-gas mass ratio of 1 percent, a gas temperature $T_{g}^{ISM}\\sim100$ K$,$ and is optically thin. The ISM is assumed to be spatially uniform and to bound a molecular cloud having radius $r_{edge}$ $\\sim10^{5}$ a.u.  \\item \\textit{Molecular cloud scale (considered in Paper I):}\\ The molecular cloud scale has much higher density than the ISM and is characterized by force balance between gas self-gravity\\ and gas pressure. The molecular cloud scale is sub-divided into a large, radially non-uniform low-density outer region and a small, approximately uniform, high-density inner core region. Clouds have a characteristic scale given by the Jeans length $r_{J}\\sim1.4\\times10^{4}$ a.u. The radial dependence of gas density is provided by the Bonnor-Ebert sphere solution which acts as the outer boundary of the Bondi accretion scale.  \\item \\textit{Bondi accretion scale (considered in Paper I):} The Bondi accretion scale is $\\sim r_{B}\\sim4.3\\times10^{3}$ a.u. which is sufficiently small that gas self-gravity no longer matters so equilibrium is instead obtained by force balance between gas pressure and the gravity of a central object assumed to be a star having mass $M\\sim0.4M_{\\odot}$ The Bondi scale is sub-divided into three concentric radial regions: an outermost region where the gas flow is subsonic, a critical transition radius at exactly $r_{B}$ where the flow is sonic, and an innermost region where the gas flow is free-falling and supersonic.  \\item \\textit{Collisionless dusty plasma scale (considered in this paper): } Free-falling dust grains collide with each other in one of the above scales and coagulate to form large-radius grains which are collisionless and optically thin. The optically thin dust absorbs UV photons from the star, photo-emits electrons and becomes electrically charged. The charged dust grains are subject to electromagnetic forces in addition to gravity. Motions of charged dust grains relative to electrons result in electric currents with associated poloidal and toroidal magnetic fields [see preliminary discussion in %TCIMACRO{\\TeXButton{\\citet{Bellan2007}}{\\citet{Bellan2007}}}% %BeginExpansion \\citet{Bellan2007}% %EndExpansion ]. \\ This region is assumed to have a scale of $10-10^{3}$ a.u.\\ \\ and corresponds to the scale of $a,$ the radius of the poloidal magnetic field axis. It is assumed that a distributed toroidal current peaked at a nominal radius $a$ is responsible for producing a poloidal field having the generic profile given in Fig.\\ref{GenericFluxFunction-new}.  \\item \\textit{Jet scale (to be considered in a future publication): }The electric currents interact with the magnetic fields to produce magnetohydrodynamic forces that drive astrophysical jets in a manner consistent with Eq.\\ref{lambda threshold} and analogous to that reported in \\ %TCIMACRO{\\TeXButton{\\citet{Hsu2002,Hsu2005}}{\\citet{Hsu2002,Hsu2005}} }% %BeginExpansion \\citet{Hsu2002,Hsu2005} %EndExpansion and %TCIMACRO{\\TeXButton{\\citet{Bellan2005}}{\\citet{Bellan2005}}}% %BeginExpansion \\citet{Bellan2005}% %EndExpansion . This region is assumed to have a scale $\\ll10^{3}$ a.u., possibly as small as a few a.u. and will involve a deformation of the generic poloidal field profile given in Fig.\\ref{GenericFluxFunction-new} because of the pressure of toroidal magnetic field inflating the poloidal flux surfaces. \\end{description} ",
        "conclusions": "We have shown that charging of collisionless dust grains incident upon a star causes the dust grain orbital dynamics to change from a relatively simple Kepler form to more complicated motion involving competition between magnetic and gravitational forces. This competition gives rise to five qualitatively different types of orbits. Two of these, the retrograde and prograde cometary orbits are just perturbations of Kepler cometary orbits. The orbit of a particle where magnetic forces overwhelm gravitational forces is just a Larmor (cyclotron) orbit and in this case the particle is constrained to remain within a poloidal Larmor radius of a poloidal flux surface in a manner similar to tokamak confinement. Particles where magnetic and gravitational forces are comparable can have two very different types of orbit depending on whether the incident particle is prograde or retrograde. Prograde particles of this latter type develop Speiser orbits; these orbits are paramagnetic with respect to the poloidal magnetic field and so can be the source of the poloidal magnetic field. Retrograde particles having comparable magnetic and gravitational forces can have a peculiar behavior whereby centrifugal force is eliminated with the result that the charged particle falls in towards the star along a spiral orbit. The accumulation of these \\textquotedblleft drain-hole\\textquotedblright\\ particles near the star provides a radial electric field oriented so as to drive the poloidal currents and toroidal magnetic fields of an astrophysical jet.  This paper showed the existence of these different types of orbits, how their orientation is suitable for generating the poloidal and toroidal magnetic fields associated with an accretion disk and astrophysical jet, and how questions of angular momentum conservation are inherently resolved. A future paper will investigate the quantitative values of dust grain parameters required to produce toroidal and poloidal fields in the accretion disk of a young stellar object.  Finally, we offer some remarks regarding the effect of deviations from axisymmetry. The model presented here assumed perfect magnetic field axisymmetry field whereas actual accretion disks are observed to have varying amounts of non-axisymmetry. This situation is  analogous to toroidal magnetic fusion devices such as tokamaks, reversed field pinches, and spheromaks all of which are modeled to first approximation as being axisymmetric, but in reality have  deviations from axisymmetry due to waves, turbulence, instability, and errors in machine construction. It is known from these devices that a modest breaking of symmetry  does not invalidate the results of the axisymmetric model, but rather weakens the conclusions, e.g., instead of cyclotron-orbiting particles being perfectly confined to the vicinity of a poloidal flux surface, when there is deviation from axisymmetry  cyclotron-orbiting particles can slowly wander away from  the  poloidal flux surface they started on. One would expect that deviations from axisymmetry in accretion disks  would cause a similar transport of cyclotron particles across poloidal flux surfaces. Because symmetry breaking  causes the canonical angular momentum of particles to change, it could be considered as being somewhat like a collision that changes the canonical angular momentum of each of two  particles involved in a collision while conserving the total canonical angular momenta. Hence,  deviations from axisymmetry would  cause a   jiggling   of the canonical angular momenta of individual particles so that particles  on the borderline between being drain-hole and cyclotron or on the borderline between being Speiser and cyclotron might spend part of the  time (i.e., between jiggles) being one type and part of the time  being the neighboring type. Similarly cyclotron particles that are on the borderline between being mirror-trapped and not mirror-trapped would, as they get kicked into and out of the mirror loss-cone,  spend part of the  time being mirror-trapped and part of the  time not being mirror-trapped. However, at any given time there would be a certain fraction of particles of each type, i.e., a certain fraction would be cyclotron, a certain fraction would be drain-hole, a certain fraction would be Speiser, and a certain fraction would be  cometary.  \\pagebreak  \\appendix   \\begin{center} \\textbf{APPENDICES} \\end{center}"
    },
    "0807/0807.1415.txt": {
        "abstract": " ",
        "introduction": "Many problems in modern particle physics and cosmology require setting up transport equations for quantum systems in out-of-equilibrium conditions, including electroweak baryogenesis~\\cite{EWBGrew}, leptogenesis~\\cite{buchmuller} and particle creation in the early universe~\\cite{TomGar}, just to mention a few.  Moreover, one often should be able to describe the evolution of coherent quantum correlations in the presence of decohering effects of the surroundings. This is true in particular for the case of the electroweak baryogensis (EWBG), where the problem is to reliably compute the fermionic out-of-equilibrium distribution functions in the neigbourhood of an expanding CP-violating phase transition wall. The EWBG problem can be divided to two regimes depending on the width of the phase transition wall in comparison to the mean free path of the fermions interacting with the wall. The case of thick wall has been addressed earlier in the semiclassical WKB approach in~\\cite{CJK,CJK2} (for earlier work see also~\\cite{Earlier}) and later in the context of CTP formalism in~\\cite{KPSW,gauge}. For a review see~\\cite{PSW}. In the thick wall limit the notion of localized particle distributions can be maintained and one can reduce the full quantum transport equations to local Boltzmann equations involving  CP-violating (and CP-even) force terms employing a well defined expansion in spatial gradients. In the thin wall limit the dominant source for the asymmetry comes from the quantum reflection processes which are inherently nonlocal and no consistent quantum field theoretical formalism exists for treating reflection phenomena together with decohering collisions. For early attempts to include collisions in a Dirac equation approach see~\\cite{hernandez,splicewall,rius}. One of the goals of this paper is to derive from field theory a density matrix formalism that can be used to solve this problem consistently. However, our methods can equally well be used to describe for exsample the coherent particle production or neutrino-oscillations in the early universe. Although our motivation comes mostly from cosmological applications and we will only consider fermionic fields, the formalism that we will develop is not restricted to solving only these problems. Instead, our generic approach to coherence within quantum field theory should be easily extended to the case of scalar fields and also to nonrelativistic problems. It should have a wide range of applications in generic problems where one is interested in quantitative description of quantum coherence in noisy backgrounds.  In this paper we shall consider only the noninteracting problem, but the formalism we develop is easily extendable to the case with interactions. Our main result is finding a phase space description for the quantum coherence in terms of new singular solutions in close analogy with the usual on-shell particle distributions. The basic objects of our study are the 2-point {\\em Wightmann} functions\\footnote{Note that our function $G^<$ does not contain an explicit minus sign often included to its definition in the literature.}: % \\beqa iG^<_{\\alpha\\beta ,ij}(u,v) &\\equiv& \\langle  \\bar \\psi_{\\beta ,j}(v){\\psi}_{\\alpha ,i}(u)\\rangle \\equiv {\\rm Tr}\\{\\hat \\rho \\, \\bar \\psi_{\\beta ,j}(v){\\psi}_{\\alpha ,i}(u) \\} \\nonumber \\\\ iG^>_{\\alpha\\beta ,ij}(u,v) &\\equiv& \\langle  {\\psi}_{\\alpha ,i}(u)\\bar \\psi_{\\beta ,j}(v)\\rangle \\equiv {\\rm Tr}\\{\\hat \\rho \\, {\\psi}_{\\alpha ,i}(u) \\bar \\psi_{\\beta ,j}(v)\\} \\,, \\label{G-less0} \\eeqa % which describe the most interesting properties of the out-of-equilibrium fermionic system. In section \\ref{propagatorth} we will first derive the standard form of these functions under the usual assumption of translational invariance both in space and time and in thermal equilibrium. We then generalize this derivation to the case where the translational invariance is lost either in time or in one of the spatial directions (denoted by $z$ hereafter). These studies are most easily done in a mixed representation, where the functions (\\ref{G-less0}) are Fourier transformed with respect to the relative coordinate $u-v$. In this representation the equation of motion for $G^<$ is found to separate into two sets of equations we call {\\em kinetic}, or evolution equations containing explicit space or time derivatives, and to algebraic (in the mean field limit) {\\em constraint} equations. How this division takes place depends on the special assumption on the spacetime symmetries.  Our most important observation is that giving up the translational invariance allows new solutions in the {\\em dynamical} phase space that carry information on the quantum coherence in the system. In the time-dependent, but spatially homogenous case considered in section \\ref{sect:tcase} this new class of solution is found at shell $k_0 = 0$ and in the planar symmetric static case, studied in section \\ref{sect:zcase}, at shell $k_z = 0$. (For a stationary problem, the position of the latter shell is shifted to $k_z = v_wk_0$, where $v_w$ is the velocity of the static frame.) These solutions are interpreted to describe the coherence between particles and antiparticles of opposite 3-momentas and same helicities on mass shells $k_0 =\\pm \\sqrt{\\vec k^2 + m^2}$ in the former case and between left and right moving states of same spin on shells $k_z =\\pm \\sqrt{k_0^2 - m^2}$ in the latter case. The new coherence solutions are present only in the dynamical functions $G^<$ and $G^>$. We will show that the spectral sum rule excludes these solutions from the spectral function ${\\cal A} = i(G^>+G^<)/2$, so that the coherence shells are not a part of the {\\em kinematical} phase space. This is as it should be, since an asymptotic state made out of pure coherence without the mixing mass-shell states does not make any physical sense. Moreover, the coherence solutions in $G^{<,>}$ are shown to be inconsistent both with the full translational invariance and with thermal equilibrium.  The phase space shell structure described above is entirely set by the algebraic constraint equations. At first sight this singular structure appears to render the kinetic equations to be of little use, but in the end the problem reveals an interesting connection to the measurement theory. In section \\ref{sect:dynamicaleqns} we show how physical, observable density matrix can be defined as a convolution of the singular phase space solution with a smooth phase space weight function that describes the existing external information on the system. This definition of physical density matrices is the second major result in this paper, along with our finding of the singular coherence shells in the phase space, and we will illustrate this principle with several examples. We show, for example, how a complete information of the momentum, the energy and spin of the state renders the quantum evolution to a trivial constant propagation of an eigenstate without coherence. Other nontrivial examples with quantum mixing include the usual Klein problem, reflection off a smooth phase transition wall and the definition and evolution of the particle number in a homogenous out-of-equlibrium system. In all these cases we are able to reproduce the known results in a way which underlines the appropriate choice of the weight function and the necessity of including the new coherence solutions.  In section \\ref{sect:flavourmixing} we extend our formalism to the case of mixing fields and show how the usual notion of the flavour mixing arises in the present context. Introducing flavour mixing through a Hermitian $N\\times N$ mass matrix leads to a very complicated shell structure in the phase space. Assuming a suffiently large $k_0$ and no mass-degeneracies, each mass shell solution separates into $N$ separate diagonal and $N(N-1)$ off-diagonal mass shells. Similarly, the number of free coherence functions at shell $k_z=0$ (in the static planar symmetric case) multiplies to $N$ and in addition $N(N-1)$ new coherence shells appear near, but not exactly at $k_z=0$. The role of each shell in the mixing phenomenon is described qualitatively. We show in particular how the usual density matrix equation for (flavour) mixing neutrino system emerges through a use of a weight function that encodes enough information to tell the direction of motion of the state, but not enough to collapse it into a singular mass-eigenstate. In section \\ref{sect:interaction} we outline how our formalism can be straightforwardly extended to the case with interations and finally, section \\ref{sect:conclusions} contains our conclusions and outlook.  % ",
        "conclusions": "\\label{sect:conclusions}  In this work we have derived quantum kinetic transport equations for fermionic systems including non-local quantum coherence. A crucial observation leading to our formalism was the finding that in cases where the full translational invariance is lost, the free fermionic 2-point correlation functions $G^{<,>}(k,X)$  have, in addition to the usual mass-shell solutions $k^2-|m|^2=0$, new solutions living on shells $k_0=0$ (homogenous case) or $k_z=0$ (planar symmetric static case). These solutions were identified as carrying the information about the quantum coherence between particle-antiparticle pairs of same helicities and opposite momenta in the former case, and between incoming and reflecting waves of equal spin in the latter case.  Another crucial element was the definition of a physical density matrix as a convolution of the singular phase space density matrix with a weight function encoding the amount of extrenous information (or the quantitative measure of the lack of it) about the state of the system. We illustrated the use of the formalism with several examples including reflection problems and definition of a particle number in the early universe during coherent particle production. We considered also the case of several mixing fermion fields and showed how the usual evolution equation for flavour mixing neutrino system arises from the singular phase space structures and an appropriate weight function. Finally we have outlined how our formalism can be extended to the case with interactions, with a slight technical reservation concerning the spatially varying problem.  Let us finally comment on our choice to limit the discussion to the mean field limit. This seems somewhat contradictory, since quantum effects become more important when the rate of change in the background gets larger in comparision with the wave length or the frequency of the probe.  Also, we got exactly the correct answers to our reflection calculations despite the mean field limit assumption. Understanding these apparent paradoxes begins from the observation that (at least for sufficiently smooth weight functions), the integrated evolution equations always have the same form as in the mean field limit, since all derivative corrections to them are reduced to vanishing surface terms. So, the only thing that {\\em does} get changed by gradients is the connetion between the averaged-out density matrix elements and the mass- and coherence shell distribution functions. Indeed, our particle numbers and fluxes are just mean field approximations to the full quantum system. However, these connections do become exact in the case of the Klein problem everywhere outside the potential step, and in the case of smooth wall at spatial infinity. This is of course why the formally mean field quantities in these cases give exact results for asymptotic currents. Similarly, in the coherent particle production case, our particle number and coherence functions provide an approximation to the full quantum phase space, that becomes exact when inflaton oscillation stops and mass becomes a constant.  One might wonder if these considerations render our results to be only of academic interest. This is obviously not so: first the complicated structure of the phase space is not optional; it {\\em is} there. One cannot just ignore the constraint equations and concentrate to the integrated form of the evolution equations. We have shown that in the mean field limit this structure is singular and allows a particularily transparent picture for separation to quantum coherence and mass-shell degrees of freedom. It is true that beyond mean field limit, the singular structure is lost. In the case of fermions this occurrs at the second order in gradients, while the first order can be computed within spectral limit and it gives rise to corrections that lead to the semiclassical effects discussed in refs.~\\cite{CJK,KPSW,PSW}. Even then the mean field limit can provide a good approximation to the phase space, capturing the most important features of the quantum evolution. Second, from the practical point of view, the singular shell structure for mass- and coherence shells is crucial when the formalism is extended to include interactions. Indeed, the entire success of the current approach relies on ones ability to find a spectral approximation to the dynamical and kinematical phase space of the system; only then can we compute the collision terms explicitly and describe the evolution of the coherence and particle numbers on these shells in a tractable manner.  As we have pointed out in many occasions, this paper is merely setting up the basic formalism which will be extended elsewhere to include decoherence~\\cite{HKR,HKR2} and then applied to various problems of interest in cosmology. This formalism will be crucial in particular to reliably compute the quantum reflection contribution to the baryon number production during the electroweak phase transition~\\cite{FarSha,hernandez,CKV}. It will also be possible to use it to study the effect of collisions on the coherent particle production~\\cite{HKR2}. We believe that the formalism could, and also will provide to be useful in other applications beyond the immediate application to the cosmology."
    },
    "0807/0807.2278_arXiv.txt": {
        "abstract": "We report on the results of observations of hard X-ray sources in the Galactic plane with the {\\em Chandra X-ray Observatory}.  The hard X-ray ``IGR'' sources were discovered by the {\\em INTEGRAL} satellite, and the goals of the {\\em Chandra} observations are to provide sub-arcsecond localizations to obtain optical and infrared counterparts and to provide constraints on their 0.3--10 keV spectra.  We obtained relatively short, $\\sim$5 ks, observations for 20 IGR sources and find a bright {\\em Chandra} source in {\\em INTEGRAL} error circles in 12 cases.  In 11 of these cases, a cross-correlation with optical and/or infrared source catalogs yields a counterpart, and the range of $J$-band magnitudes is 8.1--16.4.  Also, in 4 cases, the {\\em Chandra} X-ray spectra show evidence for absorbing material surrounding the compact object with a column density of local material in excess of $5\\times 10^{22}$~cm$^{-2}$.  We confirm that IGR~J00234+6141 is a Cataclysmic Variable and IGR~J14515--5542 is an Active Galactic Nucleus (AGN).  We also confirm that IGR~J06074+2205, IGR~J10101--5645, IGR~J11305--6256, and IGR~J17200--3116 are High Mass X-ray Binaries (HMXBs). Our results (along with follow-up optical spectroscopy reported elsewhere) indicate that IGR~J11435--6109 is an HMXB and IGR~J18259--0706 is an AGN. We find that IGR~J09026--4812, IGR~J18214--1318, and IGR~J18325--0756 may be HMXBs.  In cases where we do not find a {\\em Chandra} counterpart, the flux upper limits place interesting constraints on the luminosities of black hole and neutron star X-ray transients in quiescence. ",
        "introduction": "The hard X-ray imaging of the Galactic plane by the {\\em INTErnational Gamma-Ray Astrophysics Laboratory (INTEGRAL)} satellite is uncovering a large number of new or previously poorly studied ``IGR\" sources.  In the most recent comprehensive paper on {\\em INTEGRAL} sources, 500 sources had been detected by {\\em INTEGRAL} in the 20--40 keV band, including 214 IGR sources \\citep{bodaghee07}.  While {\\em INTEGRAL} provides hard X-ray images of unprecedented quality, it still only localizes sources to $1^{\\prime}$--$5^{\\prime}$, which is not nearly adequate for reliably finding optical or IR counterparts.  However, a relatively large fraction of the IGR sources are persistent, making it possible to detect the sources in the $\\sim$1--10 keV band with {\\em Chandra}, {\\em XMM-Newton}, or {\\em Swift}, vastly improving their localizations even with short observations with these satellites \\citep{tomsick06,rodriguez03,rtc08}.  The group of IGR sources that have been identified include 49 Active Galactic Nuclei (AGN), 32 High-Mass X-ray Binaries (HMXBs), 9 Low-Mass X-ray Binaries, 9 Cataclysmic Variables (CVs), and smaller numbers of other source types \\citep{bird06,bodaghee07}.  Probably the biggest surprise is the large number of HMXBs as well as the properties of these systems.  Many of the IGR HMXBs have large levels of intrinsic (i.e., local) absorption with $N_{\\rm H}\\sim 10^{23-24}$ cm$^{-2}$ \\citep[e.g.,][]{walter03,mg03,rodriguez03,combi04,walter04c,beckmann05,walter06}, and there is evidence in several cases that a strong stellar equatorial outflow is responsible for the absorption \\citep{fc04,chaty08}.  In addition to the high column densities, some members of the group of IGR HMXBs exhibit other extreme properties, including the high amplitude X-ray flaring of the Supergiant Fast X-ray Transients \\citep[SFXTs,][]{negueruela06,sguera06,wz07} and neutron stars with very long, $\\sim$1000--6000~s, spin periods \\citep{patel04,lutovinov05b}, which may, in at least some cases, be due to very high neutron star magnetic fields \\citep{patel07}.  In order to determine the nature of more of the unclassified IGR sources, we have been obtaining observations of IGR sources in the Galactic plane with the {\\em Chandra X-ray Observatory}.  We observed 4 sources during {\\em Chandra} observing cycle 6, obtained sub-arcsecond localizations for all 4, and found that 2 of these sources are CVs, 1 is an HMXB, and 1 is a likely HMXB \\citep{tomsick06,masetti06v,chaty08}. When we selected targets for our cycle 8 proposal (in 2006 March), we considered all the IGR sources that had been detected at that time (now, the number has nearly doubled, and we are observing many of the recently detected sources in cycle 9).  As we are interested in Galactic sources, we selected sources that are within $5^{\\circ}$ of the Galactic plane. We also eliminated sources from our list that were considered to be likely Active Galactic Nuclei (AGN), and this left us with a list of 20 targets.  Between 2006 December and 2008 February, we obtained cycle 8 {\\em Chandra} observations of these 20 IGR sources.  The goals of the observations are to provide a sub-arcsecond position for each source to facilitate an optical/IR identification and to obtain a soft X-ray spectrum.  In the following, we describe the {\\em Chandra} observations and our analysis of these data in \\S$2$.  For the cases where we obtained sub-arcsecond {\\em Chandra} positions, we report on the results of our search for optical and IR counterparts in \\S$3$.  Then, in \\S$4$, we describe the implications for each source and give our conclusions in \\S$5$. ",
        "conclusions": "During {\\em Chandra} observing cycles 6 and 8, we have obtained snapshot observations of 24 IGR sources located within $5^{\\circ}$ of the Galactic plane.  In 16 cases, a {\\em Chandra} counterpart is detected, providing the first sub-arcsecond X-ray positions for these sources.  Such an accurate position is required to obtain or confirm optical or IR counterparts in the crowded Galactic plane region, and a summary of the results for these 16 sources is given in Table~\\ref{tab:summary}.  The information in the table includes the position in Galactic co-ordinates, the $J$-band magnitude, the source type, the spectral type of the optical companion for Galactic sources or the AGN type for the one AGN in our sample, and an estimate of the contribution to the column density that is local to the system.  For this last quantity, $N_{\\rm H,local}$, we derive a range of possible levels of local absorption from the values given in Table~\\ref{tab:spectra} and from \\cite{tomsick06}, where the upper limit comes from assuming that all of the measured $N_{\\rm H}$ is local and the lower limit comes from subtracting the Galactic $N_{\\rm H}$ and two times the Galactic $N_{\\rm H_{2}}$ from the measured value.  If the lower limit is less than or consistent with zero, then there is no evidence for local absorption, and only the upper limit is quoted. For Galactic sources, a high level of local absorption provides evidence that a strong stellar wind is enveloping the compact object.  Considering the different types of sources with {\\em Chandra} counterparts in our sample, the HMXBs are the largest group, with 6 of the sources being spectroscopically confirmed HMXBs, and 4 of the sources being possible HMXBs.  At least 3 of the confirmed HMXBs are Be X-ray binaries (IGR~J06074+2205, IGR~J11435--6109, and IGR~J11305--6256), and the {\\em Chandra} spectra for the first two show evidence for significant local absorption ($N_{\\rm H,local}$ is in the range (4.6--9.7)$\\times 10^{22}$ and (10--20)$\\times 10^{22}$~cm$^{-2}$ for IGR~J06074+2205 and IGR~J11435--6109, respectively).  IGR~J16207--5129 has an early-type (i.e., high-mass) supergiant companion with no evidence for local aborption, and IGR~J10101--5645 and IGR~J17200--3116 have early-type companions, but better optical or IR spectra are required to determine the exact spectral type.  Our determination of which of the remaining sources are possible or likely HMXBs depend on several measurements that we report in this paper.  All 4 of these systems have hard X-ray spectra and none of the IR images indicate that they are extended sources. In addition, IGR~J09026--4812 is bright in the IR ($K_{\\rm s}\\sim 13$), but it shows relatively high optical/IR extinction, suggesting that it is relatively distant with a luminous companion.  A similar argument based on the optical/IR magnitudes holds for IGR~J18214--1318, IGR~J18325--0756, and IGR~J16195--4945.  In addition, the evidence that IGR~J18214--1318 and IGR~J18325--0756 are HMXBs is stronger because of the fact that $N_{\\rm H,local}$ is in the range (7--15)$\\times 10^{22}$~cm$^{-2}$ for the former and in the range (6--78)$\\times 10^{22}$~cm$^{-2}$ for the latter, possibly indicating a strong stellar wind.  Of the remaining 6 sources, 3 of them (IGR~J00234+6141, IGR~J16167--4957, and IGR~J17195--4100) are spectroscopically confirmed CVs.  All 3 of these sources have very low X-ray column densities and are likely to be relatively nearby. IGR~J14515--5542 and IGR~J18259--0706 are AGN.  While IGR~J14515--5542 has been spectroscopically classified as a Seyfert 2, our {\\em Chandra} spectrum confirms an earlier {\\em Swift} result that the X-ray column density is lower than would be expected from a Seyfert 2.  The final source with a {\\em Chandra} counterpart is IGR~J18256--1035. Our results confirm that this source has a hard X-ray spectrum, but it does not necessarily have any local absorption.  Also, it is in a region that is very crowded with IR sources, and better optical/IR images are necessary to determine its brightness.  Table~\\ref{tab:summary2} lists the remaining 8 sources without a clear {\\em Chandra} counterpart along with the upper limits on the X-ray flux from {\\em Chandra} as well as any information about the source type from the literature.  In 3 cases (IGR~J17285--2922, IGR~J17331--2406, and IGR~J18539+0727), there is evidence from earlier {\\em INTEGRAL} observations that these sources are black hole candidate (BHC) X-ray transients.  Such systems can be very faint in quiescence \\citep{garcia01,tomsick03_qbh}, so our {\\em Chandra} non-detections can be seen as further evidence for their BHC nature.  The non-detections for the likely HMXBs, IGR~J01363+6610 and IGR J17407--2808, are more surprising.  While many HMXBs show high levels of variability, flux upper limits in the range of $10^{-13}$ to $10^{-14}$ ergs~cm$^{-2}$~s$^{-1}$ are lower than expected, and may indicate that these systems are more distant or have atypical properties.  Although we do not know the source type for the final 3 sources (IGR~J17445--2747, IGR~J17507--2856, and IGR~J18193--2542), we note that all 3 sources have locations in the vicinity of the Galactic center.  This could indicate that they are part of bulge population of Low-Mass X-ray Binary (possibly BHC) transients."
    },
    "0807/0807.3754_arXiv.txt": {
        "abstract": " ",
        "introduction": "%  Gravitational wave (GW) astronomy is an inceptive branch of observational astronomy. There exist many aims for observing with GW detectors (LIGO, VIRGO, TAMA, GEO 600, e.t.c.), e.s. binary black hole - black hole (BH-BH), neutron star - black hole (NS-BH), white dwarf - black hole (WD-BH) systems. There are many evidences for the presence of Supermassive Black Holes in the centers of many galactic bulges. One is the detection of flares from tidally disrupted stars. In the disruption process GW are also emitted. \\noindent The coalescence of the black hole (BH) with the neutron star (NS) can probably also produce short GRBs with durations of milliseconds.  \\noindent Most results for the binary mergers obtained from relativistic numerical simulations were carried out with the purpose of investigating the emission of gravitational waves. For such a scenario the inspiral, merging and ringdown phases are typical. \\noindent As an example we can consider the evolution of a binary system of two nonspinning black holes of equal masses $M_{0}$, with an initial proper separation of approximately 16.6 $M_{0}$ (Pretorius 2005). Here we use c = G = 1. The binary is merged within approximately 1 orbit, leaving behind a black hole mass of $ M \\simeq 1.9 \\,M_{0}$. For $M_{0}= M_{\\odot}$ the time/space scales are equal to $250 \\,\\mu s$ and $148 \\,km$, respectively. Approximately 15\\% of the total scalar field energy does not collapse into black hole. The residual scalar field leaves the vicinity of the orbit quite rapidly, within of the order of the light crossing time of the orbit.  About 5 \\% of the total mass will be emitted as the gravitational waves during the final stages of the collision (ringdown phase) lasting only a few milliseconds.  Faber et al. (2005) have performed relativistic calculations of BH-NS mergers on the assumption that the BH is much more massive than the NS. Adiabatic evolution calculations for neutron stars with low compactness show that the neutron star typically disrupts completely within a few orbital periods. The majority of the mass transferred onto the black hole is accreted promptly; however a significant fraction ($\\simeq30$\\%) of the mass moves outward, some of which will become gravitationally unbound and ejected completely from the system. The remaining portion forms an accretion disk around the black hole, and could provide the energy source for short-duration gamma-ray bursts. As matter accretes onto the BH, it will excite quasinormal ringing modes that could in principle be detected both as gravitational waves and electromagnetic radiation. Gravitational wave emission calculations performed by Faber et al. (2005) showed also the {\\it \"chirping\"} signal, in which the binary separation decreases while the GW amplitude and frequency increases. This lasts until the onset of mass transfer from the neutron star onto the black hole. At this point we are coming across a much more rapid {\\it \"reverse chirp\"}, as the gravitational wave amplitude and frequency rapidly decrease while the NS is tidally disrupted.  \\noindent In recent years SWIFT has detected the afterglows of the short bursts GRB050724, GRB050813, and several others. These bursts together provide evidence for the fact that some short GRBs are occurring due to the merging of NS-NS binary systems (Gehrels \\& Leonard 2007).   In this paper we present the comparison of the numerical results and analytical predictions based on the detailed calculations for a test particle in a quasicircular orbit rotating about the BH. We report on the detection of high-frequency oscillations during short gamma-ray bursts. These phenomena might be evidence for to the coalescence of a BH-NS binary. ",
        "conclusions": "%  The detection of high-frequency oscillations in short gamma-ray bursts provides a new prospect for exploring the nature of this mysterious phenomenon. HFO with periods in the range of milliseconds and magnitudes of a few tens of percent of the burst luminosity may be related to the accretion of debris formed after the tidal disruption of a neutron star by a black hole in a coalescing binary system. For such a scenario, one would expect a few cycles of radiation at an orbital frequency {\\it chirping up and down}. This phenomenon indeed is seen in the bursts considered here.  \\newpage"
    },
    "0807/0807.4174_arXiv.txt": {
        "abstract": "We report the discovery of a giant Ly$\\alpha$ emitter (LAE) with a Spitzer/IRAC counterpart near the reionization epoch at $z=6.595$. The giant LAE is found from the extensive 1 deg$^2$ Subaru narrow-band survey for $z=6.6$ LAEs in the Subaru/XMM-Newton Deep Survey (SXDS) field, and subsequently identified by deep spectroscopy of Keck/DEIMOS and Magellan/IMACS. Among our 207 LAE candidates, this LAE is not only the brightest narrow-band object with $L({\\rm Ly}\\alpha) = 3.9 \\pm 0.2 \\times 10^{43}$ erg s$^{-1}$ in our survey volume of $10^6$ Mpc$^3$, but also a spatially extended Ly$\\alpha$ nebula with the largest isophotal area whose major axis is at least $\\simeq 3''$. This object is more likely to be a large Ly$\\alpha$ nebula with a size of $\\gtrsim 17$-kpc than to be a strongly-lensed galaxy by a foreground object. Our Keck spectrum with medium-high spectral and spatial resolutions suggests that the velocity width is $v_{\\rm FWHM}=251\\pm 21$ km s$^{-1}$, and that the line-center velocity changes by $\\simeq 60$ km s$^{-1}$ in a $10$-kpc range. The stellar mass and star-formation rate are estimated to be $0.9-5.0\\times 10^{10} M_\\odot$ and $>34$ $M_\\odot$yr$^{-1}$, respectively, from the combination of deep optical to infrared images of Subaru, UKIDSS-Ultra Deep Survey, and Spitzer/IRAC. Although the nature of this object is not yet clearly understood, this could be an important object for studying cooling clouds accreting onto a massive halo, or forming-massive galaxies with significant outflows contributing to cosmic reionization and metal enrichment of inter-galactic medium. ",
        "introduction": "\\label{sec:introduction}  Identifying the first stage of galaxy formation is one of the ultimate goals in astronomy today. Theoretical models predict that primordial gas accretes onto the center of halos via gravitational cooling with subsequent star-formation activity \\citep{fardal2001,yang2006}. These primordial galaxies make spatially extended Ly$\\alpha$ nebulae caused by hydrogen cooling, and it is suggested that high-$z$ extended Ly$\\alpha$ sources, or Ly$\\alpha$ blobs, are candidates for primordial galaxies (e.g. \\citealt{matsuda2004,saito2006,nilsson2006,smith2007}). Ly$\\alpha$ blobs are found mostly at $z\\simeq 2-3$, and have angular extents of $\\simeq 5-16$ arcsec with total Ly$\\alpha$ luminosities ranging from $\\simeq 6\\times 10^{42}$ to $10^{44}$ erg s$^{-1}$ \\citep{matsuda2004}. The most prominent Ly$\\alpha$ nebulae known to date are blobs 1 and 2 found by \\citet{steidel2000}, which extend over $\\gtrsim 100$ kpc with $L(Ly\\alpha)\\simeq 10^{44}$ erg s$^{-1}$. Although Ly$\\alpha$ blobs are candidates for galaxies with gas inflow of cooling accretion, it is also suggested that Ly$\\alpha$ blobs can be produced by intensive starbursts associated with significant outflows (e.g., \\citealt{taniguchi2000,wilman2005}), by a hidden AGN (e.g., \\citealt{haiman2001}), or by both of them (e.g., \\citealt{dey2005}; Yang et al. in preparation). \\begin{figure} \\epsscale{1.1} \\plotone{f1.eps} \\caption{ Optical to infrared images of Himiko. North is up and east is to the left. We display $10''\\times10''$ images at $BVRi'z'$ and $NB$ ($NB921$) bands from Subaru/SXDS, at $JHK$ bands from UKIDSS-UDS DR3, and $3.6-24\\mu$m bands from Spitzer/SpUDS. We show intensity contours in $NB921$ and $z'$ images. The black contours denote 1 $\\sigma$ level of sky fluctuation. The yellow contours represent (2, 3, 5, 10, 15) and (2,3,4) $\\sigma$ levels of sky fluctuations in $NB921$ and $z'$ images, respectively. We also plot the position of the DEIMOS slit by the red box. The dispersion direction towards red spectrum is shown by the magenta arrow. The green arrows point to the $position$ 1 and 2 that are probable peaks in the $z'$ image. \\label{fig:z7blob_images}} \\end{figure} In fact, the infrared-submm and X-ray observations indicate that \\citeauthor{steidel2000}'s blobs 1 and 2 would be powered by a heavily obscured starburst (\\citealt{geach2007,matsuda2007}; see also \\citealt{chapman2004}) and an AGN \\citep{basu-zych2004}, respectively. \\citet{matsuda2006} claim that all of their spectroscopically-identified Ly$\\alpha$ blobs are likely to be the sites of massive galaxy formation because of their large line widths of $v_{\\rm FHWM}\\gtrsim 500$ km s$^{-1}$. It is also well known that such bright large Ly$\\alpha$ nebulae are associated with radio-loud AGN (e.g. \\citealt{mccarthy1987,vanojik1997,reuland2003,barrio2008,smith2009}) or radio-quiet quasars (\\citealt{weidinger2005,hennawi2008}). In this way, extended Ly$\\alpha$ nebulae shed light not only on primordial galaxies but also on massive-galaxy formation and AGN activity.   \\begin{figure} \\epsscale{1.1} \\plotone{f2.eps} \\caption{ Composite pseudo-color image of Himiko. The RGB colors are assigned to $3.6\\mu$m , $z'$, and $NB921$ images, respectively. North is up and east is to the left. The image size is $5''\\times5''$. The white bar at the bottom right represents the length of one arcsecond. The brightest peak with a bluish white color corresponds to $position$ 1. The $position$ 2 is located $1.1$ arcsec west of the $position$ 1. \\label{fig:z7blob_composite}} \\end{figure}   So far, extended Ly$\\alpha$ nebulae have been found only at $z=2-5$ with the majority at $z\\simeq 2-3$ (e.g., \\citealt{saito2008}). Due to this current observational limit, it is difficult to identify primordial galaxies as well as to study the early stage of massive galaxy formation. At $z\\simeq 2-3$, the mean metallicity of the inter-stellar medium is already as high as $Z=0.1 Z_\\odot$ (\\citealt{sadat2001}; see also \\citealt{pettini1997}). It is predicted that the fraction of primordial galaxies to metal enriched galaxies would be quite low at $z\\simeq 2-3$ (\\citealt{scannapieco2003}), and that the fraction of population III to population II star-formation rate (SFR) rises with increasing redshift \\citep{trac2007}. On the other hand, the importance of the early stage of massive-galaxy formation has been recognized by the downsizing behavior of stellar-mass assembly \\citep{cowie1996}. A massive population at $z=2-3$ selected from a distant-red galaxy sample is old, $2-3$ Gyr, and their mean formation redshift is estimated to be $z\\gtrsim 5$ (\\citealt{labbe2005}; see also \\citealt{kriek2006}). It is implied that galaxies at the massive end would have a very high specific star-formation rate (SSFR) at $z\\gtrsim 4$ \\citep{drory2008}, and that a major active star-formation in massive galaxies probably takes place at $z\\gtrsim 4$. Thus, it is important to study extended Ly$\\alpha$ nebulae at a redshift greater than the current-observational limit, especially at the reionization epoch of $z\\simeq 6-11$ \\citep{fan2006,komatsu2008}. This epoch is also the today's observational limit of galaxy studies \\citep{iye2006,stark2007,ota2008,bouwens2008}. Moreover, such bright Ly$\\alpha$ sources can be a good laboratory for understanding reionization (e.g. \\citealt{kashikawa2006,dijkstra2007b,mcquinn2007,kobayashi2007}) and metal enrichment (e.g. \\citealt{martin2002,bouche2007}) of inter-galactic medium (IGM).                        In this paper, we report our discovery of an extended Ly$\\alpha$ nebula, which we named Himiko \\footnote{ Himiko is a name of legendary queen in ancient Japan. } , near the reionization epoch at $z=6.595$. We describe the photometric identification and spectroscopic confirmation of this object in \\S \\ref{sec:discovery}, and present detailed properties such as kinematics and stellar population in \\S \\ref{sec:giant_LAE}. We discuss the nature of this object and prospects of future observations in \\S \\ref{sec:discussion}. Throughout this paper, magnitudes are in the AB system. We adopt $(h,\\Omega_m,\\Omega_\\Lambda,n_s,\\sigma_8)=(0.7,0.3,0.7,1.0,0.8)$.     \\begin{figure} \\epsscale{0.88} \\plotone{f3.eps} \\caption{ Surface brightness (SB) profiles of our object. Open squares show a SB profile in $NB921$ image, and filled circles denote that of a continuum-subtracted (Ly$\\alpha$) image. The dotted line is the Gaussian profile representing the point-spread function of our image with an $FWHM=0''.8$. \\label{fig:radial_profile_z7blob}} \\end{figure}     \\vspace{2cm} ",
        "conclusions": "\\label{sec:discussion}  \\subsection{Comparisons with Ly$\\alpha$ Blobs at $\\lowercase{z}\\sim 3$} \\label{sec:comparisons}  We compare properties of our giant LAE with those of Ly$\\alpha$ blobs found at $z\\simeq 3$. Because objects experience more severe surface-brightness dimming at $z=6.6$ than at $z\\simeq 3$, careful comparisons are needed. We produce simulated $NB921$ images of blob 1 of \\citet{steidel2000} and blob 28 of \\citet{matsuda2004} redshifted to $z=6.595$ based on the narrow-band images of \\citet{matsuda2004}. We choose the blob 1 and 28, because they have the brightest Ly$\\alpha$ luminosity and the highest surface brightness, respectively, in \\citeauthor{matsuda2004}'s catalog. We carry out the simulations in the same manner as \\citet{saito2006}, but with an improved random noise whose amplitude exactly matches to those of real $NB921$ image. Figure \\ref{fig:z7blob_simu} presents the simulated $NB921$ images, together with the original narrow-band images of \\citet{matsuda2004}. The simulated image of blob 1 indicates that our observations would miss $\\simeq 95$\\% of blob 1's $L({\\rm Ly}\\alpha)$, and that no diffuse Ly$\\alpha$ nebula could be identified. The simulated blob 1 has ($NB921$[total], $A_{\\rm iso}$, $\\left <SB\\right >$) $=$ ( $25.5$ mag, $0.8$ arcsec$^2$, $26.7$ mag arcsec$^{-2}$). We plot this simulated object in Figure \\ref{fig:mag_isoarea_NB921LAE}. The simulated blob 1 is indistinguishable from the cloud of normal LAEs in the planes of $A_{\\rm iso}$ vs. $NB921$ and $A_{\\rm iso}$ vs. $\\left< SB \\right>$. On the other hand, the simulated blob 28 shows ($NB921$[total], $A_{\\rm iso}$, $\\left <SB\\right >$) $=$ ( $24.1$ mag, $3.7$ arcsec$^2$, $26.3$ mag arcsec$^{-2}$). This simulated object is recognizable in Figure \\ref{fig:mag_isoarea_NB921LAE}. However, this simulated object is more similar to the other LAEs with the second and third largest $A_{\\rm iso}$ than to our object. There exist remarkable differences in $A_{\\rm iso}$ and $\\left <SB\\right >$ (or $NB921$ magnitude) between our object and the simulated objects. In other words, the large $A_{\\rm iso}$ and the high $\\left <SB\\right >$ of our object at $z=6.6$ cannot be realized even for blob 28 whose very high surface brightness would minimize the effect of cosmological surface brightness dimming. It implies that our object might be a population that has not yet been identified at $z\\simeq 3$.  \\begin{figure} \\epsscale{1.2} \\plotone{f9.eps} \\caption{ Original narrow-band images of blob 1 and 28 at $z=3.1$ from \\citet{matsuda2004} (top panels) and simulated narrow-band images at $z=6.6$ (bottom panels). The size of each panel is $25''\\times25''$. The intensity contours are presented with yellow lines. The contours represent (2, 3, 5, 10, 15) $\\sigma$ levels of sky fluctuations, where the $2\\sigma$ sky fluctuations correspond to $28.0$ and $26.8$ mag arcsec$^{-2}$ in the original images \\citep{matsuda2004} and the simulated images, respectively. Note that the absolute values of contour levels in the simulated images are the same as the yellow contours in the $NB921$ image of Figure \\ref{fig:z7blob_images}. In the original images, blob 1 is the large diffuse extended source, while blob 28 is the compact source at the center. Note that the original and simulated images include foreground and background sources neighboring the blobs in the $25''\\times25''$ areas. \\label{fig:z7blob_simu}} \\end{figure}    Although the Ly$\\alpha$ brightness and profile of our object seem different from those of Ly$\\alpha$ blobs at $z\\simeq 3$, the other properties of our object show both similarities and differences. The line width of our object is comparable to some of the extended LAEs presented in \\citet{matsuda2006} in the plane of FWHM(Ly$\\alpha$) vs. isophotal area, although the threshold of isophotal area is different from ours. On the other hand, there are no objects with such a luminous Ly$\\alpha$ line in the range of $v_{\\rm FWHM}=100-900$ km s$^{-1}$ in the catalog of \\citet{saito2008}. The velocity width of our object is $251$ km s$^{-1}$ which is significantly smaller than that of \\citeauthor{steidel2000}'s blob 1 with a broad-velocity width, $\\simeq 1000-1500$ km s$^{-1}$ (\\citealt{ohyama2003}; see also \\citealt{bower2004}). The velocity width of our object is closer to the one of blob 28 ($v_{\\rm FWHM}=362$ km s$^{-1}$; Y. Matsuda in preparation). The stellar masses of the sub-components shaping \\citeauthor{steidel2000}'s blobs 1 and 2 range from $10^{10}$ to $10^{11}M_\\odot$ \\citep{uchimoto2008}, which is comparable to ours. Similarly, \\citet{smith2008} report that the stellar mass of their Ly$\\alpha$ blob at $z=2.83$ is as massive as $3-4\\times 10^{11} M_\\odot$. These stellar masses are comparable to ours within an order of magnitude. The estimated number density of our object is only $1.2\\times 10^{-6}$ comoving Mpc$^{-3}$, which is lower than those of $z\\sim 3$ Ly$\\alpha$ blobs found by \\citet{matsuda2004} ($3\\times 10^{-4}$ Mpc$^{-3}$), \\citet{saito2006} ($1\\times 10^{-5}$ Mpc$^{-3}$), and Yang et al. in preparation ($3\\times 10^{-6}$ Mpc$^{-3}$). This number density is also lower than the upper limit of number density of \\citeauthor{steidel2000}'s two Ly$\\alpha$ blobs ($2\\times 10^{-5}$ Mpc$^{-3}$), which are estimated from the number of objects (2) and \\citeauthor{matsuda2004}'s survey volume ($1.3\\times 10^{5}$ Mpc$^{3}$). Although these number densities depend on the criteria of sample selections and observing fields, our object at $z=6.6$ would be as rare as (or even rarer than) these Ly$\\alpha$ blobs at $z\\sim 3$. Regarding the environment, our giant LAE resides in a high density region of LAEs (\\S \\ref{sec:photometric_identification}). \\citet{matsuda2004} find that, at $z=3.1$, the distribution of their Ly$\\alpha$ blobs trace the dense regions of LAEs. The similar spatial correlation between our extended LAE and compact LAEs is also seen at $z=6.6$.   We also compare our object with Ly$\\alpha$ nebulae around high-$z$ radio galaxies (HzRGs; \\citealt{miley2008}). The Ly$\\alpha$ luminosity of our object is comparable to those of HzRGs ($\\sim 10^{43}-10^{44}$ erg s$^{-1}$; \\citealt{mccarthy1991}). On the other hand, our object has a much smaller FWHM line width of Ly$\\alpha$, $251\\pm21$ km s$^{-1}$, than HzRGs ($\\gtrsim 1000$ km s$^{-1}$; \\citealt{mccarthy1996}). However, recent studies have found an HzRG with a strikingly relaxed dynamics whose Ly$\\alpha$ FWHM line width is $<300$ km s$^{-1}$ \\citep{villar-martin2007}. It would be difficult to distinguish between our object and such a dynamically-relaxed HzRG powering a Ly$\\alpha$ halo based on dynamics, although our object has a radio flux fainter than that of a typical radio-loud AGN (Section \\ref{sec:agn}).        \\subsection{Inferred Kinematics} \\label{sec:inferred_kinematics}   If this object forms a single virialized system whose possible two $z'$-band components would be just bright {\\sc Hii} regions in a disk (see \\S \\ref{sec:nature}), the dynamical mass from the rotation is naively estimated to be $M_{\\rm rot} \\sin(i) \\simeq 1\\times 10^9 M_\\odot$ by $M_{\\rm rot} \\sin(i)= ([v_c \\sin(i)]^2 r )/G$ with $v_c \\simeq \\Delta v/2 = 30$ km s$^{-1}$ and $r= D/2 = 5$ kpc (\\S \\ref{sec:kinematics}), where $i$ and $G$ are the inclination of a rotating disk and the gravitational constant. On the other hand, from the 1-dimensional velocity dispersion of $\\sigma_{v}=v_{\\rm FWHM}/2.35=107$ km s$^{-1}$ in a half size of the extraction aperture, $R\\sim 6/2=3$ kpc (\\S \\ref{sec:kinematics}), we obtain the mass of random motion of $M_{\\rm rand}=4\\times 10^{10} M_\\odot$ via $M_{\\rm rand} = (5/3) (3 \\sigma_{v}^{2}) R/G$, assuming a uniform sphere. If the blue half of Ly$\\alpha$ is absorbed by the external IGM, the velocity dispersion and mass are $\\simeq 214$ km s$^{-1}$ and $M_{\\rm rand}=2\\times 10^{11} M_\\odot$, respectively. Thus, $M_{\\rm rand}$ is 1-2 order(s) of magnitude larger than $M_{\\rm rot}$ in a reasonable range of inclination ($i=10-90^\\circ$). If we consider the smaller radius for the $M_{\\rm rand}$ estimate ($R= 3$ kpc) than that for the $M_{\\rm rot}$ estimate ($r= 5$ kpc), the difference of these dynamical masses defined in a common radius becomes even larger. If Ly$\\alpha$ reflects dynamics, our giant LAE would be a system more dominated by random motion than rotation. Note that these estimates of dynamical masses depend on the size of spectrum extraction aperture with the uncertainties of seeing smearing, and that these results are only true under the assumptions of the single-virialized system and no significant effects of resonant scattering.               If this giant LAE is an outflow object whose Ly$\\alpha$ emission is produced by shock heating (cf. \\citealt{dijkstra2008}), the dynamical time scale required to form the extended Ly$\\alpha$ nebula is $\\simeq 7\\times 10^7$ yr, where we assume the size of the major axis ($\\simeq 17$ kpc) and the typical velocity width ($v_{\\rm FWHM}=251$ km s$^{-1}$). Since this time scale is as long as the cosmic time between $z=6.595$ and $7.1$, it would start making an ionized-bubble since $z=7.1$ in this shock heating scenario.        \\subsection{Nature of the giant LAE} \\label{sec:nature}  The nature of the Ly$\\alpha$ nebula of our object is not yet clearly understood within the currently available observational data. There are five possible explanations: (1) halo gas photoionized by a hidden AGN, (2) clouds of {\\sc Hii} regions in a single virialized galaxy, (3) cooling gas accreting onto a massive dark halo associated with an initial onset of starburst at the halo center, (4) merging bright LAEs with clouds of {\\sc Hii} regions, and (5) outflowing gas excited by shocks or UV radiation from starbursts and/or mergers. There is a chance of (1), although no positive evidence of an AGN is found. Our object shows the lack of {\\sc Nv} line, no line broadening, and no detections in X-ray, MIR, submm, and radio bands (Figure \\ref{fig:z7blob_multibandsed}). But, these constraints are not strong enough to discard the possibility of a hidden AGN. The case of (2) seems surprising, because it means that such a large galaxy exists in a very early epoch of $z=6.6$. This galaxy would have a size of $\\gtrsim 17$ kpc and possibly two large star-forming regions in a disk seen in our $z'$ band image (\\S \\ref{sec:gravitational_lensing}). However, there is a chance to explain this large galaxy in the framework of Cold Dark Matter (CDM) models. We estimate properties of the most massive dark halo whose number density is the same as our object at $z=6.6$. Based on the analytic CDM model of \\citet{sheth1999}, we find that the dark halo has a radius of 47 kpc and a mass of $1\\times 10^{12} M_\\odot$ with a circular velocity of 380 km s$^{-1}$. All of these values of the dark halo are significantly larger than those of our giant LAE measured via radiation, indicating that the single-galaxy picture can be compatible with the CDM model in terms of halo properties. But it is not obvious that such a big virialized baryonic system at $z=6.6$ can be reproduced in the scheme of CDM model. The case of (3) is possible with the potentially large $EW_0$ of $57-402$\\AA. Moreover, our DEIMOS spectrum presents a possible weak red peak at $\\sim 9245$\\AA\\ (Figure \\ref{fig:blob_2dspec_1dspec}) which is redshifted from the main Ly$\\alpha$ peak ($9233$\\AA) by $\\sim 400$ km s$^{-1}$. This possible red peak may be similar to the one produced by IGM infalling into a collapsing cloud, which is claimed by \\citet{dijkstra2006a,dijkstra2006b}. On the other hand, the comparable SFR values from Ly$\\alpha$ and UV continuum (\\S \\ref{sec:stellar_population}) may not prefer this scenario, since the Ly$\\alpha$ luminosity could be explained solely by normal star-forming activities that the UV continuum indicates. However, this argument is not strong due to the underestimation of SFR from Ly$\\alpha$ luminosity. In fact, the intrinsic Ly$\\alpha$ luminosity could be brighter by a factor of 20 if our giant LAE has diffuse Ly$\\alpha$ components similar to \\citeauthor{steidel2000}'s blob 1 (\\S \\ref{sec:comparisons}). The explanation of (4) would be reasonable, given the possible existence of two peaks in $z'$ band with a separation of $1''.1$ (6.0 kpc). The $NB921$-image profile can be reproduced by two exponential disks with a reasonably small half-light radius of $R_{\\rm hl}=0''.3$ corresponding to $R_{\\rm hl}=1.6$ kpc (\\S \\ref{sec:gravitational_lensing}; cf. \\citealt{simard1999}). The merger would induce star-formation activities, and could produce the bright Ly$\\alpha$-line and UV-continuum emission. Figure \\ref{fig:z7blob_specdiff} shows that velocity widths at $position$ 1 and 2 are similar, implying that dynamical masses of these components would be comparable. We may be witnessing the site of a major merger near the reionization epoch. On the other hand, our object shows a Ly$\\alpha$ nebula potentially larger than the isophotal scale with the major axis of $3''.1$ (17 kpc; \\S \\ref{sec:photometric_identification}). If the Ly$\\alpha$ nebula really extends beyond the $2\\sigma$-level isophotal area, it becomes difficult to explain the Ly$\\alpha$ morphology with a profile of two merging LAEs. The (5) case seems plausible, since our object has a relatively high SFR and a large stellar-mass with possible multiple components in $z'$ band (\\S \\ref{sec:gravitational_lensing}). In either case of (4) or (5), our object would be a massive galaxy in formation with significant star-formation contributing to cosmic reionization (cf. \\citealt{iliev2006}) and/or with outflows for the metal enrichment of IGM \\citep{bouche2007}. Since our object has the small velocity offset ($\\Delta v=60$km s$^{-1}$) and the line width ($v_{\\rm FWHM}=251$ km s$^{-1}$), the dynamics of merger or outflow would have to be well collimated to the direction perpendicular to the line of sight.   The angular size of the Ly$\\alpha$ nebula is $\\gtrsim 17$ proper kpc (\\S \\ref{sec:gravitational_lensing}), which is comparable to the diameter of the stellar disk of the present-day Milky Way. It is impressive, if we consider that the age of the Universe at $z=6.595$ is only 6\\% of the one of the present-day Universe. Moreover, such an extended Ly$\\alpha$ source is very rare in the cosmological volume only with the number density of $1.2\\times 10^{-6}$ comoving Mpc$^{-3}$ at $z=6.6$ (\\S \\ref{sec:spectroscopic_confirmation}). If our selection of large Ly$\\alpha$ nebula does not miss a significant fraction of massive galaxies at this early epoch ($z=6.6$), our object could be an ancestor of a bright-cluster or cD galaxy, and should be a good laboratory of massive-galaxy formation near the reionization epoch.                      \\subsection{Future Prospects} \\label{sec:future}  In section 4.3, we have found that the currently available data do not provide a clear answer to the question about the nature of this object. It is obvious that deeper NIR and infrared images of this object can be taken with Hubble and Spitzer Space Telescopes to constrain SFR, dust extinction, and stellar age, which will trace back through the star-formation history of this object (e.g. \\citealt{yan2006,eyles2007}). These deep images are critically important to give stronger constraints on the SED of our object (Figure \\ref{fig:z7blob_sed}) that characterizes SFR and dust extinction, and to identify Ly$\\alpha$ photons that are not originated from star-formation activities but from others, such as cold accretion. The high-resolution NIR image of Hubble Space Telescope is useful to investigate the possibilities of mergers and outflows. A $J$-band spectrum will test the existence of AGN with a relatively strong {\\sc Civ} emission line. We estimate an expected {\\sc Civ} flux to be $\\sim 2$ and $4 \\times 10^{-17}$ erg s$^{-1}$ cm$^{-2}$ for Seyfert II and QSO, assuming the ratios of $f_{\\rm CIV}/f_{\\rm Ly\\alpha}\\sim 0.2$, and $f_{\\rm CIV}/f_{\\rm Ly\\alpha}\\sim 0.5$, respectively \\citep{mccarthy1993}. These flux limits are achieved by deep NIR spectroscopy with 8m-class telescopes. Moreover, we can characterize star-formation activities and metal enrichment in our Ly$\\alpha$ nebula by deep submillimeter and millimeter observations with Atacama Large Millimeter/Submillimeter Array (ALMA) which will start the operation in 2012 preceded by the early-science operation. ALMA observations will allow us to investigate emission from dust and molecular-clouds in our object. A detection of spatially-extended metal line from the Ly$\\alpha$ nebula could reject the possibility of cooling accretion of primordial gas, and provide an independent probe of star-formation properties.   We estimate an expected intensity of dust emission in $850\\mu$m, $S(850\\mu{\\rm m})$, and a flux of molecular CO(6-5) line, $S({\\rm CO})$, assuming typical parameters of local starbursts. We start the calculations from our lower-limit of SFR, $34 M_\\odot$yr$^{-1}$. The far-infrared luminosity, $L({\\rm FIR})$, is calculated from $SFR [M_\\odot {\\rm yr}^{-1}]=1.7\\times 10^{-10} L({\\rm FIR}) [L_\\odot]$ \\citep{kennicutt1998}. We obtain $L({\\rm FIR})=2.0\\times 10^{11} L_\\odot$. Assuming the modified blackbody radiation with a dust emissivity index of $\\beta=1.3$ and a dust temperature of $T_{\\rm dust}=35.6$ K \\citep{dunne2000}, we estimate the $850\\mu$m dust emission from $L({\\rm FIR})$ to be $S(850\\mu{\\rm m})=0.28$ mJy (see, e.g., \\citealt{ouchi1999}). This moderately bright $850\\mu$m emission is expected because of the negative k-correction \\citep{blain2002}. Again from the $L({\\rm FIR})$ value with the same modified blackbody radiation, the dust mass is $M_{\\rm dust}=6.0\\times 10^{7} M_\\odot$ via the relation presented in \\citet{debreuck2003}. The mass of molecular hydrogen is $M(H_2)=3.0\\times 10^{9} M_\\odot$ which is calculated with the relation of $M({\\rm H_2})/M_{\\rm dust}=50$ ($\\sim 25-75$; \\citealt{seaquist2004,young1996}). Finally, the flux of molecular CO(6-5) line is $S({\\rm CO})=0.1$ Jy km s$^{-1}$ with the assumptions of an H$_2$-to-CO conversion factor of 0.8 \\citep{downes1998}, and a line ratio of ${\\rm CO}(6-5)/{\\rm CO}(1-0)=0.5$ \\citep{bayet2006}. In summary, we expect $S(850\\mu{\\rm m})\\gtrsim 0.28$ mJy and $S({\\rm CO})\\gtrsim 0.1$ Jy km s$^{-1}$, considering that our SFR is the lower limit ($>34 M_\\odot$yr$^{-1}$). We use ALMA Sensitivity Calculator \\footnote{ http://www.eso.org/sci/facilities/alma/observing/tools/etc/index.html }, and estimate the on-source integration time to be $\\lesssim 0.2$ and $\\lesssim 3$ hours for $5\\sigma$ detections of a $850\\mu{\\rm m}$ continuum and a CO(6-5) line, respectively. Here we assume the large beam size of $1''$ in a compact configuration of 50 12m-arrays for the point-source detection, the band width of 16 GHz, and the CO-line width of $250$ km s$^{-1}$ \\citep{nishiyama2001} with a 50 km s$^{-1}$ spectral resolution. Either of an $850\\mu$m-thermal continuum or a molecular-CO(6-5) line may be detected in reasonable observing time under the assumptions of local starbursts. If our object does not have dust or molecular gas as much as the local starbursts, a deficit of dust or molecular-line emission would be identified by ALMA observations. In either case, dust and molecular-gas properties of our object could be characterized in a few years."
    },
    "0807/0807.1567_arXiv.txt": {
        "abstract": " ",
        "introduction": "Most inflation models predict a nearly Gaussian distribution of the primordial curvature perturbation. Deviations from an exactly Gaussian distribution are characterized by a dimensionless parameter $f_{NL}$ \\cite{Bartolo:2004if}. In the case of single field inflation model $f_{NL}\\sim {\\cal O}(n_s-1)$ \\cite{Maldacena:2002vr}, which is constrained by WMAP ($n_s=0.960_{-0.013}^{+0.014}$) \\cite{Komatsu:2008hk} to be much less than unity. A Gaussian distribution of the primordial curvature perturbation is still consistent with WMAP five-year data \\cite{Komatsu:2008hk}: \\e -9<f_{NL}^{local}<111 \\quad \\hbox{and} \\quad -151<f_{NL}^{equil}<253 \\quad (95\\% \\hbox{CL}), \\q where ``local\" and ``equil\" denote the shapes of the non-Gaussianity. In \\cite{Yadav:2007yy} the authors reported that a positive large non-Gaussianity \\e 27<f_{NL}^{local}<147 \\q is detected at $95\\%$ C.L.. Planck is expected to bring the uncertainty of $f_{NL}^{local}$ to be less than 5 \\cite{Komatsu:2001rj}. If a large value of $f_{NL}$ is confirmed by the forthcoming cosmological observations, the simplest model of inflation is ruled out and some very important new physics of the early Universe will be showed up.   In general a large number of light scalar fields are expected in the theories beyond the standard model, such as string theory. The consistent perturbative superstring theory can only live in ten-dimensional spacetime. To connect string theory with experiments, string theory must be compactified on some six-dimensional manifold and many dynamical moduli fields emerges in four dimensions. The typical number of moduli fields is $N\\sim {\\cal O}(10^2-10^3)$. One can expect that the expectation values of some of these scalar fields are displaced from the minimum of their potential due to the quantum fluctuations during inflation. Usually these scalar fields are subdominant during inflation and their fluctuations are initially of isocurvature type. After the end of inflation they are supposed to completely decay into thermalized radiation before primordial nucleosynthesis and thus the isocurvature perturbations generated by them are converted to be a final adiabatic perturbations. These scalar fields are called curvatons.  The curvaton mechanism to generate an initially adiabatic perturbation deep in the radiation era is proposed in \\cite{Linde:1996gt,Enqvist:2001zp,Lyth:2001nq,Moroi:2001ct}. The primordial curvature perturbation in curvaton model with single curvaton has been discussed in \\cite{Linde:1996gt,Enqvist:2001zp,Lyth:2001nq,Moroi:2001ct,Lyth:2002my,Bartolo:2003jx,Malik:2006pm,Sasaki:2006kq}. If many curvatons make contributions to the primordial density perturbation, the calculation becomes much more complicated \\cite{Choi:2007fya,Assadullahi:2007uw,Valiviita:2008zb}. Since curvaton model can give a large positive local-type non-Gaussianity, recently some topics related to curvaton model are discussed in \\cite{Huang:2008ze,Ichikawa:2008iq,Multamaki:2008yv,Suyama:2008nt,Beltran:2008ei,Li:2008jn,Li:2008fm,Huang:2008qf}.    In single-curvaton model $f_{NL}$ is inverse proportional to the fraction of curvaton energy density in the energy budget at the epoch of curvaton decay. Smaller the energy density of curvaton, larger the non-Gaussianity. Since the curvaton mass is smaller than the Hubble parameter $H_*$ during inflation, or equivalently its Compton wavelength is large compared to the curvature radius of the de Sitter space $H_*^{-1}$, the gravitational effects play a crucial role on the behavior of curvaton field in such a scenario. The typical energy density of curvaton field in such a background is roughly $H_*^4$, which leads to an upper bound on $f_{NL}$ \\cite{Huang:2008ze}: $f_{NL}<522\\cdot r^{1\\over 4}$ (up to an order one coefficient), where $r$ is the tensor-scalar ratio.   Since multiplicity of scalar fields is generally expected, we focus on the multi-curvaton scenario in this paper. Here we consider a special case in which all of curvatons have the same decay rate and their masses are larger than the decay rate. In order to simplify the calculation of the primordial curvature perturbation, we assume all of curvatons suddenly decay into radiation at the same time. We give a name, ``N-vaton\", to this scenario. As a realistic N-vaton model, the axions in the KKLT compactification of Type IIB string theory are suggested to be curvatons. Based on the random matrix theory, the mass spectrum of axion obeys the Marcenko-Pastur law and a rich phenomenology of this model is shown up.   Our paper is organized as follows. In Sec. 2, we use the $\\delta \\cn$ formalism \\cite{Starobinsky:1986fxa,Sasaki:1995aw,Lyth:2005fi} to calculate the primordial curvature perturbation on large scales in N-vaton model. The various bounds on $f_{NL}$ are discussed in Sec. 3. In Sec. 4 we consider a more general case where the curvature perturbation generated by inflaton cannot be ignored and we find that the spectral index of primordial power spectrum can be red-tilted naturally. In Sec. 5, we propose a realistic N-vaton model in which the axions in the KKLT compactification of Type IIB string theory are taken as curvatons. At the end we give some discussions on N-vaton model in Sec. 6.   \\setcounter{equation}{0} ",
        "conclusions": "In this paper we explicitly calculate the primordial curvature perturbation in N-vaton model. Multiplicity of light scalar fields is generic in the theories going beyond standard model. Even though the total energy density of these light scalar fields is subdominant during inflation, the perturbation produced by them can dominate the density perturbation on large scale. We also suggest a realistic N-vaton model in which the axions in the KKLT compactification of Type IIB string theory are taken as curvatons, and a rich phenomenology is shown up. If a large local-type non-Gaussianity is confirmed by the forthcoming experiments, it can shed a light on these light scalar fields.   In order to fit the spectral index from WMAP data, the inflaton fluctuation is still required to play a significant role in the total primordial power spectrum. Generally the tensor-scalar ratio is required to be not smaller than $10^{-5}$ if $f_{NL}^{local}>10$. Many inflation models constructed in string theory, such as brane inflation \\cite{Dvali:1998pa,Kachru:2003sx}, happen in a quite low energy scale with $r \\sim 10^{-10}$ which is too small to generate a large non-Gaussianity in curvaton/N-vaton scenario. How to construct an inflation model with $r\\sim {\\cal O}(10^{-5}-10^{-3})$ and $\\eta\\sim -0.1$ is still an open question.    In general, the curvaton decay rate mediated by particles of mass $M_X$ is expected to be of order $m^3/M_X^2$. So it is natural to assume that the different curvatons have different decay rates and decay at different time, in particular for the case where they have different masses. The authors in \\cite{Assadullahi:2007uw} gave a concrete example to show that a large non-Gaussianity can be obtained even when the curvatons dominate the total energy density at the time of decays in the case of two curvatons. However if there are hundreds or thousands of curvatons, whether this enhancement of the non-Gaussianity is generic or not is still an open question. It is worth studying this problem in the future.       \\vspace{.5cm}  \\noindent {\\bf Acknowledgments}  We would like to thank Daniel Chung, Hironobu Kihara, Han-Tao Lu, K.~P.~Yogendran and Yu-Feng Zhou for useful discussions.    \\newpage  \\appendix  \\setcounter{equation}{0}"
    },
    "0807/0807.3338_arXiv.txt": {
        "abstract": "Using deep {\\it Chandra} observations of M84 we study the energetics of the interaction between the black hole and the interstellar medium of this early-type galaxy. We perform a detailed two dimensional reconstruction of the properties of the X-ray emitting gas using a constrained Voronoi tessellation method, identifying the mean trends and carrying out the fluctuation analysis of the thermodynamical properties of the hot ISM.  In addition to the $PV$ work associated with the bubble expansion, we identify and measure the wave energy associated with the mildly supersonic bubble expansion. We show that, depending on the age of the cavity and the associated wave, the waves can have a substantial contribution to the total energy release from the AGN. The energy dissipated in the waves tends to be concentrated near the center of M84 and in the direction perpendicular to the bubble outflow, possibly due to the interference of the waves generated by the expansion of northern and southern bubbles.  We also find direct evidence for the escape of radio plasma from the ISM of the host galaxy into the intergalactic medium. ",
        "introduction": "The detailed investigations of the balance between the heating and cooling of the hot interstellar medium (ISM) and intracluster medium (ICM) have been made possible due to the key observations made by {\\it Chandra} and XMM-{\\it Newton} and the successes in numerical modeling (see McNamara \\& Nulsen 2007 for a review).  The commonly accepted paradigm states that the feedback mechanism is ultimately linked to the activity of the supermassive black holes located at the bottom of the potential wells in systems with cool cores.\\\\ \\indent A standard way of estimating the energy released by the AGN is to measure the $PV$ work associated with X-ray cavities (bubbles) inflated by the AGN.  This is an indirect measurement as the bubbles are typically ``observed'' as depressions in X-ray emissivity.  However, the task of estimating the AGN energy injected into the ICM can be accomplished by assuming that such X-ray cavities are in pressure balance with the ICM. One can then infer the total energy contained in the cavities that is available for doing mechanical work (i.e., enthalpy) from $H=\\gamma/(\\gamma -1)PV$, where $\\gamma$ is the adiabatic index, $P$ is the pressure and $V$ is the bubble volume.  Such an observational estimate is further complicated by the fact that the effective adiabatic index $\\gamma$ of the material inside the bubbles is not known.  That is, it is not known if the bubbles are predominantly thermal, in which case $\\gamma = 5/3$ (e.g., Mazzotta et al. 2002 shows that both the thermal and non-thermal models fit the data in the case of MKW 3s galaxy cluster) or non-thermal, in which case $\\gamma = 4/3$ (e.g., Sanders \\& Fabian 2007 in the case of the Perseus cluster). The content of radio lobes and X-ray cavities may depend on such factors as the initial jet composition ($e^{+}-e^{-}$ or $p-e^{-}$), the efficacy of entrainment of colder thermal ICM gas or the magnetic pressure support inside the bubbles (Dunn et al. 2006 and references therein).  Furthermore, the energy content obtained this way may be a lower limit to the actual total energy released by the AGN if the pressure balance assumption does not hold.  An extreme example of this effect is seen in the simulations of very high Mach number jets interacting with the ICM (Binney, Bibi \\& Omma 2007). \\\\ \\indent Moreover, X-ray observations show that the standard method of inferring AGN energies leads to the bubble energy {\\it increasing} in the statistical sense with the distance from the center of the gravitational potential (Diehl et al. 2008).  This apparently counter intuitive result may be the result of the erroneous assumption that the bubbles are ``born'' in pressure equilibrium with the surrounding ICM or ISM.\\\\ \\indent If the bubbles are overpressured with respect to the ISM, then one would need the internal bubble pressure to correctly estimate the energy inside the bubble.  Moreover, if they were overpressured in the past, then this overpressure should have created waves that have carried some portion of the energy away from the cavities. The standard method of estimating the AGN energy based on the pressure balance and cavity size does not include such contributions to the total energy balance.  In addition to increasing the overall energy budget, such waves provide a more gentle and spatially-distributed heating.\\\\ \\indent The above effects clearly demonstrate the need for alternative or supplementary measurements of the energy injected by the AGN. In this Paper we report on the long time exposure observation of the elliptical galaxy M84 in the Virgo cluster. We focus on the interaction of the AGN with the ISM and the energy content in the thermodynamical fluctuations generated by the the AGN outburst and argue that they are waves. An approach to measure the wave energy that is similar to the one described in this Paper, has been considered by Sanders \\& Fabian (2007) in the case of the observations of the Perseus cluster and in the numerical simulations of wave dissipation (Ruszkowski et al. 2004). ",
        "conclusions": "We have analyzed a deep {\\it Chandra} observation of the AGN feedback in the Virgo elliptical galaxy M84. We have applied a constrained Voronoi tessellation binning method to {\\it Chandra} data. Our main results are: (1) the AGN outflow is mildly supersonic, (2) the non-thermal plasma from the AGN-inflated cavities appears to be porous and the non-thermal particles escape from the cavities, (3) co-centric density perturbations (weak shock waves) are present and their morphology is qualitatively consistent with the results of numerical simulations, (4) we have estimated the mechanical energy in the waves and found that it may contribute substantially to the overall mechanical energy delivered by the AGN."
    },
    "0807/0807.5108_arXiv.txt": {
        "abstract": "\\ We present observations of the neutron star X-ray binary and relativistic jet source Circinus X-1 made at 4.8 and 8.6 GHz with the Australia Telescope Compact Array during a time interval of almost 10 years. The system shows significant variations in the morphology and brightness of the radio features on all timescales from days to years. Using the time delay between the successive brightening of the different components of the radio emission we were able to provide further evidence for the relativistic nature of the arcsec scale outflow, with an apparent velocity $\\beta_{app} \\geq$ 12. No compelling evidence for an evolution of the orientation of the jet axis was found.  We also place an upper limit on the proper motion of the system which is consistent with previous optical studies. Besides the previously reported radio flares close to the orbital phase 0.0 (interpreted as enhanced accretion at periastron passage), we also identified outbursts with similar properties near the orbital phase 0.5. The global spectral index revealed a preferentially steep spectrum over the entire period of monitoring with a mean value and standard deviation $\\alpha$=-0.9 $\\pm$ 0.6 ($F_{\\nu} \\propto \\nu^{\\alpha}$), which became significantly flatter during the outbursts. Polarization was detected in one third of the epochs and in one case Faraday rotation close to the core of the system was measured. ",
        "introduction": "Circinus X-1 is one of the most exotic X-ray binary systems. The detection of type I X-ray bursts \\citep*{Ten86a,Ten86b} and twin kHz quasi-periodic oscillations in the X-ray power spectra \\citep{Bou06} strongly suggests that the compact object in the system is a neutron star. The nature of the companion star is more debatable: some evidence is pointing towards a low-mass object \\citep{Joh99}, while other towards a high-mass one \\citep*{Mur80,Jon07}. If the companion star is a supergiant, as recent evidence seems to suggest \\citep{Jon07}, then this is at odds with the low magnetic field that must be present in the system (as inferred from type I X-ray bursts observations, for instance). Thus Circinus X-1 appears to be a very exotic system, harbouring a young neutron star with a low magnetic field.  Neutron star X-ray binaries (XRBs) are classified according to their X-ray spectral and timing properties as Z and Atoll sources. The Z class comprises seven objects accreting near the Eddington limit. All of them have been detected in the radio band and show variable emission at cm wavelengths. The Atoll type form the largest class of neutron star XRBs and have X-ray properties similar to black hole XRBs. Due to their systematic lower radio fluxes, one order of magnitude less radio loud than Z sources, only five of them have been detected in the radio band. The behaviour of Circinus X-1 is puzzling, exhibiting similarities with Z sources (e.g. \\citealt*{Shi99}), but at times showing characteristics reminiscent of Atoll sources \\citep{Oos95}. The orbit of the binary system probably has a relatively high eccentricity (e $\\simeq$ 0.4-0.9) that produces variations in the accretion rate on the compact object \\citep*{Mur80,Nic80,Tau99,Jon07}. This offers an explanation for the periodic flares (P$\\simeq$16.6 d) observed in X-ray \\citep{Kal76}, infrared \\citep{Gla78,Gla94} and radio \\citep{Whe77,Hay78}, which are interpreted as enhanced accretion near the periastron passage.  Circinus X-1 is associated with an arcmin scale synchrotron nebula \\citep{Hay86,Ste93} which is probably powered by the jet originating close to the binary system \\citep{Hei02,Tud06}. The jet is observed in radio at arcmin as well as arcsec scales and was recently detected in X-ray band by Chandra (\\citealt{Hei07}; Soleri et al. 2008). Observations at cm wavelengths offered evidence for the presence of a relativistic outflow aligned very close to the line of sight \\citep{Fen04}. However, X-ray \\citep{Sch08,Iar08} and optical \\citep{Jon07} spectroscopy seem to favour a significantly higher angle between the jet axis and the line of sight.  Besides the flux modulation at the orbital period, Circinus X-1 also exhibits `secular' changes, documented in X-ray and radio. The ASM/RXTE (All Sky Monitor/Rossi X-ray Timing Explorer) light curve (2-10 keV) shows that after a quasi-constant X-ray output, the flux decreased starting from 2000 up to now by one order of magnitude (Fig. 1, left, bottom panel). The picture is even more dramatic in radio. In the 1970s-1980s the radio flares reached up to 1 Jy (e.g. \\citealt{Whe77,Hay78}), then they were detected at mJy level only \\citep{Ste91,Fen97,Fen98}, until recently when a relatively increased radio activity from Circinus X-1 was noticed \\citep*{Fen05,Nic07, Del07}. This offered the opportunity to obtain new VLBI (Very Large Baseline Interferometry) observations of the system within the framework of the first Southern e-VLBI (electronic VLBI) experiment \\citep{Phi07}, after almost 25 years since the last successful VLBI observations \\citep{Pre83}.  Circinus X-1 complex is very close to the supernova remnant SNR G321.9-0.3 on the plane of the sky. Optical observations \\cite{Mig02} ruled out the possibility that the two objects are physically associated as it was assumed for a long time \\citep{Cla75}. ",
        "conclusions": "We have presented ATCA radio data of the XRB Circinus X-1 at 4.8 and 8.6 GHz covering a period of time of almost 10 years. The radio maps reveal a complex structure of the radio emission with variations in the morphology and brightness of the different compact emitting regions (sometimes seen as far as 10 arcsec from the core, corresponding to a physical scale of 0.05 $r$ pc, with $r$ the distance to Circinus X-1 expressed in kpc) at timescales of days. The orientation of the radio emission is on a NW-SE direction in the plane of the sky and remains unchanged during the period covered by our observations. This suggests that if the jet is indeed precessing, as proposed by \\cite{Iar08}, then its precession period should be much larger then the time span of our data set.  We were able to confirm the claim of \\cite{Fen97} for a radio flare not associated with the orbital phase 0.0. Further evidence was found that such outbursts do occur near the orbital phase 0.5 (within a region of $\\pm$ 0.1 in phase). Their properties are similar to those of the outbursts commonly observed near the orbital phase 0.0. We suggested that the wind accretion from the massive companion might be responsible for producing them.  \\cite{Fen04} showed that the outflow in the arcsec scale jet of Circinus X-1 is relativistic. Reanalyzing the data in some more detail and determining the time delay between the successive brightening of the various compact radio emitting regions, it was found that indeed the outflow seems to be relativistic with an apparent velocity $\\beta_{app} \\geq$ 12.  A study of the proper motion of the system was carried out. The errors in determining the positions of the core (dominated by the errors associated with the coordinates of the phase-referencing calibrator) allowed only an upper limit to be estimated, namely 50 mas yr$^{-1}$($v_{system} <$ 240 $r$ km s$^{-1}$, where $r$ is the distance to Circinus X-1 expressed in kpc). This is however a conservative value and future observations together with improvements in the systematic errors of the calibrators will pull down the limit and offer a better constraint.  The global spectral index of Circinus X-1 (between 4.8 and 8.6 GHz) had a mean value and standard deviation of $\\alpha=-0.9 \\pm 0.6$ $(F_{\\nu} \\propto \\nu^{\\alpha})$ during the monitoring period. The spectrum was flattening when flares were active, either close to phase 0.0 or 0.5.  Polarization emission was detected in one third of the epochs and in one epoch Faraday rotation was observed at a level of -175 $\\pm$ 30 rad m$^{-1}$. The nature of the Faraday screen is unknown. Circumstantial evidence points towards an intrinsic origin, related to the immediate vicinity of Circinus X-1 itself, however this has to be tested in detail in future observations. In the jet the electric field vectors seem to be oriented preferentially along its axis tentatively suggesting the presence of shocks at the interface between the ejected matter and the surrounding environment."
    },
    "0807/0807.0535.txt": {
        "abstract": "Baryon Acoustic Oscillations (BAO) %,subtle features in the large scale %distribution of galaxies produced by acoustic waves in the photon-baryon %plasma during the early stages of the universe, provide a ``standard ruler'' of known physical length, %easy to %interpret observationally, making it one of the most promising probes of the nature of dark energy. %approaches to obtain an accurate and unbiased determination of %the equation of state of dark energy. The detection of BAO as an excess of power in the galaxy distribution at a certain scale requires measuring galaxy positions and redshifts. ``Transversal'' (or ``angular'') BAO measure the angular size of this scale projected in the sky and provide information about the angular distance. ``Line-of-sight'' (or ``radial'') BAO require very precise redshifts, but provide a direct measurement of the Hubble parameter at different redshifts, a more sensitive probe of dark energy. %the matter-energy content of %the universe. %It is implied in the literature that ``transversal'' BAO can and should be %studied using photometric redshifts, which are much more ``cheap'' in terms of observing time %than spectroscopic redshifts, and that the precision afforded by the latter should be reserved %for ``line-of-sight'' BAO. The main goal of this paper is to show that %, contrary to usual assumptions, it is possible to obtain photometric redshifts with enough precision ($\\sigma_z$) to measure BAO along the line of sight. There is a fundamental limitation as to how much one can improve the BAO measurement by reducing  $\\sigma_z$. We show that $\\sigma_z\\sim 0.003 (1+z)$ is sufficient: a much better precision will produce an oversampling of the BAO peak without a significant improvement on its detection, while a much worse precision will result in the effective loss of the radial information. This precision in redshift can be achieved for bright, red galaxies, featuring a prominent $4000$~\\AA\\  break, by using a filter system comprising about 40 filters, each with a width close to $100$~\\AA, covering the wavelength range from $\\sim4000$~\\AA\\  to $\\sim8000$~\\AA, supplemented by two broad-band filters similar to the SDSS $u$ and $z$ bands. We describe the practical implementation of this idea, a new galaxy survey project, PAU\\footnote{Physics of the Accelerating Universe (PAU): {\\tt http://www.ice.cat/pau}}, to be carried out with a telescope/camera combination with an {\\it etendue} about 20~m$^2$~deg$^2$, equivalent to a 2~m telescope equipped with a 6~deg$^{2}$-FoV camera, %with about $500$~Mpixels, and covering 8000~sq.~deg.\\ in the sky in four years. We expect to measure positions and redshifts for over 14 million red, early-type galaxies with $L>L_\\star$ and $i_{AB}\\lesssim 22.5$ in the redshift interval $0.1<z<0.9$, with a precision $\\sigma_z < 0.003(1+z)$. This population has a number density $n\\gtrsim 10^{-3}$ $\\textrm{Mpc}^{-3}\\textrm{h}^{3}$ galaxies within the 9~$\\textrm{(Gpc/h)}^{3}$ volume to be sampled by our survey, ensuring that the error in the determination of the BAO scale is not limited by shot-noise. By itself, such a survey will deliver precisions of order $5\\%$ in the dark-energy equation of state parameter $w$, if assumed constant, and can determine its time derivative when combined with future CMB measurements. In addition, PAU will yield high-quality redshift and low-resolution spectroscopy for hundreds of millions of other galaxies, including a very significant high-redshift population. The data set produced by this survey will have a unique legacy value, allowing a wide range of astrophysical studies. % ",
        "introduction": "\\label{sec:intro}  Physical Cosmology has recently entered the precision era. This transition has been propelled by the gathering, over the past decade, of unprecedented high-precision data sets for several cosmological observables. The combined analysis of the Cosmic Microwave Background (CMB) anisotropies (e.g., \\citet{wmap3ext, wmap5ext, boomerang, cbi, acbar, vsa}) with distance-scale measurements at increasingly higher redshifts (e.g., \\citet{SNe1, astier06, SNe2}) and probes of large-scale structure \\citep{2df, sdss1, sdss2, Tegmark06, hutsi-2} yields a remarkably consistent picture: a spatially flat universe that has started a phase of acceleration of the expansion at the present epoch. From the observations gathered so far this acceleration is consistent with the effect of a cosmological constant, but it may also be caused by the presence of a dynamical energy component with negative pressure, now termed Dark Energy (DE), or might also point to a fundamental modification of our description of gravity. The answer to what is the exact cause is likely to have profound implications for cosmology and for particle physics.  Two recent collective reports, one by the US Dark Energy Task Force %\\citep[DETF][]{detf}, (DETF)\\ \\citep{detf}, convened by NASA, NSF and DOE, and another by the European ESA-ESO Working Group on Fundamental Cosmology \\citep{esa-eso}, have identified the most promising observational strategies to characterize DE properties in the near future. These reports concluded that the method based on measurements of Baryon Acoustic Oscillations (BAO) from galaxy redshift surveys is less likely to be limited by systematic uncertainties than other methods that are proposed. %, although photometric calibration needs to be sufficiently accurate. It appears that, while recognizing the need for a combined strategy involving two or more independent techniques, BAO measurements can substantially contribute to increase the accuracy on the DE equation of state.  Baryon Acoustic Oscillations are produced by acoustic waves in the photon-baryon plasma generated by primordial perturbations \\citep{eh98}. At recombination ($z\\sim 1100$), the photons decouple from the baryons and start to free stream, whereas the pressure waves stall. As a result, baryons accumulate at a fixed distance from the original overdensity. This distance is equal to the sound horizon length at the decoupling time, $r_{BAO}$. The result is a peak in the mass correlation function at the corresponding scale. First detections of this excess were recently reported, at a significance of about three standard-deviations, both in spectroscopic~\\citep{eisenstein,2dFb,hutsi} and photometric~\\citep{nikhil,blake} galaxy redshift surveys.  The comoving BAO scale is accurately determined by CMB observations ($r_{BAO}=146.8 \\pm 1.8$~Mpc  for a flat $\\Lambda CDM$ Universe~\\citep{Hinshaw08}), and constitutes a ``standard ruler'' of known physical length. The existence of this natural standard ruler, measurable at different redshifts, makes it possible to probe the expansion history of the Universe, and thereby dark energy properties and the Universe geometry (see, e.g., \\citet{seo,Glazebrook} and references therein). This motivates the present efforts to measure BAO (e.g., \\citet{adept}, \\citet{des}, \\citet{hetdex}, \\citet{panstarrs}, \\citet{space}, \\citet{wfmos}, \\citet{wigglez}).  Broad-band photometric galaxy surveys can measure the angular scale of $r_{BAO}$ in several redshift shells, thereby determining $d_{A}(z)/r_{BAO}$, where $d_{A}(z)$ is the angular distance to the shell at redshift $z$. If galaxy redshifts can be determined precisely enough, the BAO scale can also be measured along the line of sight, providing a direct measurement of the instantaneous expansion rate, the Hubble parameter (or actually of $H(z)\\, r_{BAO}$), at different redshifts. This quantity is more sensitive to the matter-energy contents of the universe compared to the integrated quantity $d_{A}(z)$. The direct determination of $H(z)$ distinguishes the BAO method from other methods. In addition, since systematic errors affect the radial and tangential measurements in different ways, the consistency between the measured values of $H(z)$ and $d_{A}(z)$ offers a test of the results.  As a rule of thumb, in order to get the same sensitivity to the dark-energy parameters, a galaxy redshift survey capable of exploiting the information along the line of sight needs to cover only $\\sim$ 10\\% of the volume covered by a comparable survey that detects the scale in the transverse direction only~\\citep{briddle}. When covering a similar volume, precise enough redshift measurements can provide substantially tighter constraints on the DE parameters.  Large volumes have to be surveyed in order to reach the statistical accuracy needed to obtain relevant constraints on dark-energy parameters. Enough galaxies must be observed to reduce the shot noise well below the irreducible component due to sampling variance (see section~\\ref{sec:power}). The usefulness of the correlation along the line of sight favors spectroscopic redshift surveys that obtain very accurate redshifts, but the need for a large volume favors photometric redshifts that can reach down to fainter galaxies.  The intrinsic comoving width of the peak in the mass correlation function is about 15~Mpc/h, due mostly to Silk damping~\\citep{Silk}. This sets a requirement for the redshift error of order $\\sigma(z)=0.003(1+z)$, corresponding to 15~Mpc/h along the line of sight at $z=0.5$. A much better precision will result in oversampling of the peak without a substantial improvement on its detection, while worse precision will, of course, result in the effective loss of the information in the radial modes \\citep{seo}. Note also that in the presence of substantial redshift errors, the error distribution needs to be known and accurately corrected for when inferring the BAO scale.  It has usually been assumed in the literature that photometric redshifts are not precise enough to measure ``line-of-sight'' BAO \\citep{seo,briddle}. While this is true for broad-band photometric surveys, here we examine how one can reach the required redshift precision with narrow-band photometry. We find (see section~\\ref{sec:survey}) that redshifts of luminous red galaxies can be measured with a precision $\\sigma(z)\\sim 0.003(1+z)$ using a photometric system of $40$ filters of $\\sim 100$~\\AA, continuously covering the spectral range from $\\sim 4000$ to $\\sim 8000$~\\AA , plus two additional broad-band filters similar to the $u$ and $z$ bands.  We describe the practical implementation of this idea, a photometric galaxy redshift survey called PAU (Physics of the Accelerating Universe). PAU will measure positions and redshifts for over 14 million luminous red galaxies over 8000 deg$^{2}$ in the sky, in the range $0.1<z<0.9$ (comprising a volume of 9~(Gpc/h)$^{3}$), and with an expected photometric redshift precision $\\sigma(z)\\lesssim 0.003(1+z)$. This redshift precision makes it possible to measure radial BAO with minimal loss of information. The PAU survey can be carried out in a four year observing program at a dedicated telescope with an effective {\\it etendue} $\\sim 20$~m$^2$deg$^2$.  The outline of the paper is as follows. Section~\\ref{sec:req} discusses the scientific requirements for such a survey. An optimization of the survey parameters follows in section~\\ref{sec:survey}, while section~\\ref{sec:base} presents a possible baseline design for the instrument. The science capabilities of the survey are given in section~\\ref{sec:science}. Finally, section~\\ref{sec:summary} contains the summary and conclusions. ",
        "conclusions": "\\label{sec:summary} % In 1998, the discovery of the accelerated expansion of the universe changed completely our understanding of the universe and its components. Ten years on, the quest to understand what causes the acceleration continues. Along the way, the Baryon Acoustic Oscillation (BAO) technique has been identified as a systematically robust, yet statistically powerful, probe of dark energy properties. In particular, measuring the BAO feature along the line of sight as a funcion of redshift provides a direct determination of $H(z)$, which, in turns, depends on the amount and characteristics of dark energy. However, such a measurement requires a very precise determination of galaxy redshifts.  We have presented here a novel approach to photometric redshift determination which allows the measurement of the BAO feature along the line of sight in an efficient way, by using a set of about 40 narrow-band (FWHM$\\approx$~100~\\AA) filters. The approach complements (for BAO and for other science) spectroscopic surveys, which typically measure much more precisely the spectra of a much reduced sample of galaxies.  Because of the intrinsic width of the peak in the galaxy-galaxy correlation function of about 15~Mpc/h, there is a fundamental limitation as to how much one can improve the BAO measurement by reducing the photo-z errors. A redshift precision of order $\\sigma(z)\\approx 0.003(1+z)$, corresponding to 15~Mpc/h along the line of sight at $z=0.5$, is about optimal for this measurement. Redshift space distortions, biasing and non-linear effects produce distortions that can be comparable to the effect of this photo-z precision.  Simulations show that both the target galaxy density ($n\\sim 10^{-3}$(h/Mpc)$^3$) and precision in redshift ($\\sigma_z/(1+z) \\sim 0.003$) can be achieved with the proposed system. These simulations indicate that PAU by itself %(no priors) can determine the equation of state of the dark energy assumed constant ($w$) to about 5\\%, while when the PAU data is combined with expected supernova and CMB data samples, a sizable increase in the Dark Energy Task Force figure-of-merit (inverse of the area of the error ellipse in the $w_0$--$w_a$ plane) by about a factor three is achieved, making the PAU very competitive when compared to other planned ground-based BAO surveys, photometric or spectroscopic.  The survey will produce a unique data set with low-resolution spectroscopy in the optical wavelengths for all objects in the northern sky up to $m_{AB}=23$--$23.5\\ \\textrm{arcsec}^{-2}$ (five sigma). A survey like PAU, producing such a catalogue, will have enormous legacy value and will be extremely useful for many areas of astrophysics, with contibutions that are different from, and complementary to, those a spectroscopic survey can deliver. %"
    },
    "0807/0807.2943_arXiv.txt": {
        "abstract": "We investigate the physical conditions of the interstellar medium and stellar components in the regions of the southern Galactic star forming complexes associated with \\mbox{IRAS~10049-5657} and \\mbox{IRAS~10031-5632}. These regions have been mapped simultaneously in two far infrared bands ($\\lambda_{eff}$ $\\sim$ 150 \\& 210 $\\mu$m), with $\\sim 1'$ angular resolution using the TIFR 1-m balloon borne telescope. Spatial distribution of the temperature of cool dust and optical depth at 200 $\\mu$m have been obtained taking advantage of the similar beams in the two bands. The HIRES processed IRAS maps at 12, 25, 60 and 100 $\\mu$m have been used for comparison. Using the 2MASS near infrared sources, we find the stellar populations of the embedded young clusters. A rich cluster of OB stars is seen in the \\mbox{IRAS~10049-5657} region. The fits to the stellar density radial profile of the cluster associated with \\mbox{IRAS~10049-5657} has been explored with inverse radius profile as well as the King's profile; the cluster radius is obtained to be $\\sim2$ pc. The source in the cluster closest to the IRAS peak is IRA-7 which lies above the zero age main sequence curve of spectral type O5 in the colour-magnitude diagram. Unlike \\mbox{IRAS~10049-5657}, a small cluster comprising of a few deeply embedded sources is seen at the location of \\mbox{IRAS~10031-5632}. Self consistent radiative transfer modelling aimed at extracting important physical and geometrical details of the two IRAS sources show that the best fit models are in good agreement with the observed spectral energy distributions. The geometric details of the associated cloud and optical depths ($\\tau_{100}$) have been estimated. A uniform density distribution of dust and gas is implied for both the sources. In addition, the infrared ionic fine-structure line emission from gas has been modelled for both these regions and compared with data from IRAS-Low Resolution Spectrometer. For \\mbox{IRAS~10049-5657}, the observed and modelled luminosities for most lines agree to within a factor of four while for \\mbox{IRAS~10031-5632}, we find a discrepancy of a factor of 100 and it is likely that some basic assumptions of the model are not valid in this case. ",
        "introduction": "HII regions are  astronomical sources that represent early stages of deeply embedded high mass (O or early B) stars. Their study can provide vital information about high mass star formation as well as their interaction with the parent molecular cloud. Being deeply embedded in interstellar cloud including the dust component, almost all of their energy is absorbed and re-emitted in the infrared wavebands. \\mbox{IRAS~10049-5657} (G282.0-1.2) and \\mbox{IRAS~10031-5632} (G281.6-1.0) are Galactic star forming regions in the southern sky, which are generally less-studied. Radio measurements indicate that \\mbox{IRAS~10049-5657} is an extended HII region (Hill, 1968; Manchester, 1969).  A number of distance estimates to \\mbox{IRAS~10049-5657} can be found in the literature, ranging from 5.1 to 7.1 kpc. Here we use the distance of 6.3 kpc estimated by Caswell \\& Haynes (1987) based on radio recombination line measurements (for R$_{\\odot}$ = 8.5 kpc and Galactic rotation velocity = 220 km s$^{-1}$ at R$_{\\odot}$). The distance to \\mbox{IRAS~10031-5632} is estimated to be 3.7 kpc (Caswell \\& Haynes, 1987).  Recently, it has been concluded that \\mbox{IRAS~10049-5657} complex harbours a very massive OB star cluster (Bik et al. 2005; Hanson et al., 2003). Bik et al. (2005) have carried out near-infrared K-band spectroscopy of few members of this cluster and find two very massive (O3-O4) stars here. \\mbox{IRAS~10049-5657} has been studied as a part of surveys for the search of emission lines including masers. Formaldehyde absorption has been detected towards this source at 4.8 GHz (Whiteoak and Gardner, 1974) and 14.5 GHz (Gardner and Whiteoak, 1984). Whiteoak et al (1982) detected CO (1-0) line emission from this source using the 4-m radio telescope of CSIRO. Search for methanol transition (Peng and Whiteoak, 1992), methanol maser (Schutte et al, 1993) and OH maser (Cohen et al, 1995) close to this source have led to negative results. \\mbox{IRAS~10049-5657} has been imaged by Puchalla et al. (2002) at 42 GHz using the Mobile Anisotropy Telescope on Cerro Toco (MAT/TOCO) and an integrated flux of $22.2\\pm2.1$ Jy within a 0.3$^{\\circ}$ circular beam has been obtained by them. Both these regions, \\mbox{IRAS~10049-5657} and \\mbox{IRAS~10031-5632} have been studied as a part of the Parkes-MIT-NRAO (PMN) survey (Kuchar and Clark, 1997). They find the peak radio flux density of \\mbox{IRAS~10049-5657} and \\mbox{IRAS~10031-5632} at 5 GHz to be 25.7 Jy/beam and 1.0 Jy/beam, respectively, where the beam is $\\sim 4\\arcmin.9$. CS~(2-1) line emission has been observed close to both these regions (Bronfman et al., 1996) using the SEST telescope. The IRAS-LRS (Low Resolution Spectrometer) spectrum of \\mbox{IRAS~10031-5632} shows strong [Ne~II] emission line at 12.8 $\\mu$m (de Muizon et al., 1990; Simpson \\& Rubin, 1990), a relatively weak [S~III] emission line at 18.7 $\\mu$m as well as emission in the Unidentified Infrared Bands (UIBs) at 7.7, 8.6 and 11.3 $\\micron$ (Zavagno et al., 1992; de Muizon et al., 1990). Water vapor maser (Braz et al. 1989) as well as methanol maser (Schutte et al. 1993) have not been found close to \\mbox{IRAS~10031-5632}.  The less known southern Galactic massive star forming regions are being studied under a long-term program which involves observing these sources in far infrared (Vig et al. 2007, Ojha et al. 2002, Karnik et al. 2001, Ghosh et al. 2000, Verma et al. 1994). In this paper, we present a systematic study of the star forming regions associated with \\mbox{IRAS~10049-5657} and \\mbox{IRAS~10031-5632}. The star forming region associated with \\mbox{IRAS~10049-5657} is believed to harbour a cluster of very massive stars. \\mbox{IRAS~10031-5632} is, on the other hand, a young star forming region with few members belonging to the cluster. We have carried out an infrared study of these southern Galactic HII regions in detail with the aim of understanding the energetics, physical sizes, the spatial distribution of interstellar dust and its temperature as well as the associated young clusters. Section 2 describes observations and other data sets used. Section 3 describes the results and in Section 4, radiative transfer modelling of these sources is presented. A comprehensive discussion of these sources is carried out in Section 5 and a brief summary is presented in Section 6. ",
        "conclusions": "\\subsection{IRAS~10049-5657}  Using the FIR map of \\mbox{IRAS~10049-5657} (Fig.~\\ref{FIRmap_a}), we see an extension of the dust emission towards north-east. This is also seen in the 60 and 100 $\\micron$ emission IRAS-HIRES maps of this region, shown in Fig.~\\ref{HIRES_a}. The IRAS-HIRES mid infrared emission from warm dust, however, shows an extension towards the west. This western extension is also clearly seen in the MSX maps. The flux densities from the TIFR maps at 150 and 210 $\\micron$ have been used to compute the mass of dust using the formulation of Hildebrand (1983) and Sandell (2000). For a temperature of 30 K, the dust mass obtained is  $\\sim21$ $M_{\\odot}$. As can be seen from Fig.~\\ref{SUMSS_a}, a north-east extension seen also in the cold dust emission is observed for the ionised gas also. The total flux of 42.4 Jy at 843 MHz can be compared with the flux density of 40 Jy (beam $\\sim14'$) at 1410 MHz (Manchester, 1969).  From the CM (J-H vs J) diagram of the 2MASS sources (Fig.~\\ref{cmcc_a} left) within the cluster radius ($\\sim65''$), we see that there are 11 sources lying above the ZAMS curve of O9. These 11 sources are designated as IRA-1, ... , IRA-11 and a list of their positions and magnitudes is given in Table~\\ref{2MASS_parm_a}. However, it is important to note that the spectral types inferred from CM diagram are the earliest possible spectral type (upper limits) when the stars have infrared excess. Sources with infrared excess can be found from the CC (J-H vs H-K) diagram. In the CC diagram (see Fig.~\\ref{cmcc_a}), 8 `asterisks' lie in the band occupied by reddened ZAMS stars. One source, IRA-7 (shown as a triangle lying in the `T' region), shows an infrared excess while IRA-10 lies to the left of the reddening band of the ZAMS objects (drawn from the top of the main sequence branch) in the CC diagram. The upper left part of the CC diagram is not an allowed region for young stellar objects (Lada \\& Adams, 1992). IRA-10 is faint in J band ($16.9\\pm0.2$ mag) with larger errors though it is relatively brighter in the H ($14.2\\pm0.02$ mag) and K$_s$ ($13.2\\pm0.04$) bands as compared to the J band. One possibility to explain the near infrared colours of IRA-10 as well as its position in the CC diagram is that IRA-10 could comprise two or more unresolved sources. The sources lying near the unreddened main sequence in the CC diagram are possibly foreground sources not associated with the cluster. A small but significant fraction of cluster stars lie outside and to the right of the reddening band of the ZAMS stars. These are mostly young stellar objects (YSOs) with intrinsic colour excess. Fourteen objects lie in this infrared excess zone, i.e. in the `T' and `P') regions. By dereddening the stars (on the CC diagram) that fall within the reddening vectors encompassing the main sequence stars, we find the visual extinction (A$_V$) towards each star. The individual extinction values range from $A_V \\sim 0 - 17$ magnitudes. From the CM diagram, we find that the extinction values of most `asterisks' lie at about $A_V\\sim 15 - 18$ magnitudes.  Figure~\\ref{2mass_a} shows the 2MASS K$_s$-band image of the region around \\mbox{IRAS~10049-5647} in grayscale. The 2MASS image of this region shows diffuse emission apart from the sources (earlier than spectral type O9) clustered together. It is interesting to note the clustering of these sources close to the IRAS peak and the distribution of the other such stars in the NE-SW direction which is the direction of extension of the ionised gas as well as the cold dust. The morphology of emission in the UIB as well as the emission from warm dust in MSX bands also shows extension along the south-west direction. The brightest infrared source among the selected 2MASS sources is IRA-11. IRA-7 is closest ($\\sim 2.7''$) to the IRAS peak, is of spectral type earlier than O5 and has an infrared excess. This is consistent with the spectral type determined by Bik et al. (2005) for IRA-7 (referred to as 10049nr411 in their paper) using high resolution K-band spectra. They find that it is of spectral type O3-O4. Recent multi-epoch radial velocity measurements by Apai et al. (2007) show large amplitude variations in radial velocities pointing towards IRA-7 being a massive binary system (50 and 20 M$_{\\odot}$). Bik et al. (2005) have also carried out spectroscopy of IRA-3 (10049nr324) and find it to be of spectral type O3-O4/O5-O6. This is consistent with its ZAMS spectral type, O5-O6, obtained from the CM diagram. It is interesting to note that the `asterisk' closest to the radio peak is IRA-8. The CM diagram indicates IRA-8 to be of ZAMS spectral type O9-O6.  We have also cross-correlated the MSX PSC sources with those from 2MASS PSC which lie within the cluster radius. There are 2 such MSX PSC sources namely G282.0341-01.1810 and G282.0176-01.1793 (hereafter M1 and M2, respectively; listed in Table~\\ref{msx}). M1 coincides ($1.3''$) with an infrared excess source, 2MASS-J10064115-5712377 (J$\\sim16.0\\pm0.2$, H$\\sim14.9\\pm0.2$, K$_s\\sim13.5\\pm0.3$). M2 coincides ($2''$) with a source detected only in K$_s$ band, 2MASS-J10063592-5711563 (K$_s\\sim14.1\\pm0.2$ mag). From the 2MASS-MSX colour-colour diagram ($F_{21}/F_8$ vs $F_8/F_K$; Lumsden et al. 2002, their Fig.~9), we find that these sources lie in the region generally covered by compact H II regions and massive young stellar objects. We can, therefore, conclude that these are young stars associated with the cluster.  \\subsection{IRAS~10031-5632}  The dust emission around \\mbox{IRAS~10031-5632} is compact (unresolved) at 150 $\\mu$m, as well as at the IRAS-HIRES wavebands. \\mbox{IRAS~10031-5632} is not resolved in the SUMSS radio map also. However, \\mbox{IRAS~10031-5632} is barely resolved in the 210 $\\mu$m map. Deconvolving the beam from 210 $\\mu$m image gives us $\\sim0.9'$ as an estimate of its size. The dust mass obtained using the flux density at 210 $\\micron$ from the TIFR map is $\\sim7$ $M_{\\odot}$ for a temperature of 30 K. However, from the radiative transfer modelling, we obtain a larger dust mass of 24 $M_{\\odot}$. The total radio flux density is 1.1 Jy. Considering that the radio flux is from a single ionising ZAMS star, we use the formulation of Schraml \\& Mezger (1969) and Panagia (1973) to estimate its spectral type. We find it to be of ZAMS spectral type O9-O8.5. This can be compared with the ZAMS spectral type of O9 obtained from the FIR luminosity.  Unlike the case of \\mbox{IRAS~10049-5657}, we do not see a rich cluster with a large number of stars around \\mbox{IRAS~10031-5632}. Rather a small group comprising of a few stars is seen close to this region. There are seven sources (labelled as IRB-1 to IRB-7) clustered close to the FIR (IRAS) peak; of which two, IRB-1 and IRB-5, are detected in all the three (J, H and K$_s$) bands. The 2MASS designation and flux details of these sources are given Table~\\ref{2MASS_parm_b}. From the CM diagram (H-K vs K; Fig.~\\ref{cmcc_b} left), it is observed that IRB-1 is of ZAMS spectral type earlier than O5. It is to be noted that this spectral type of IRB-1 obtained from NIR study is the earlier possible spectral type (upper limit) as it shows an infrared excess. IRB-5 is however of later spectral type, B3-B2. It lies among a well-defined group of stellar sources which have lower extinction values (A$_V\\sim5$ mag). From the CC diagram (J-H vs H-K), it is found to lie near the main-sequence curve. It is therefore possible that IRB-5 is not associated with the cluster but is a foreground source. The other five sources are detected in H and K$_s$ or only in K$_s$ bands. The source closest ($\\sim 1.5''$) to the IRAS peak, IRB-3, is detected in H ($15.1\\pm0.2$ mag) and K$_s$ ($12.8\\pm0.1$ mag) bands but not detected in the J band. From the CM diagram, it is found that IRB-3 is a heavily extincted early spectral type (O9-O6) ZAMS star. It is located at the peak of 843 MHz radio emission. The other four sources IRB-2, IRB-4, IRB-6 and IRB-7 are detected only in K$_s$ band. This implies that these are deeply embedded objects. Thus, among the small number of cluster members, the majority of them are deeply embedded indicating that the star formation here is in an early stage. The 2MASS image of the region around \\mbox{IRAS~10031-5632} is shown in Fig.~\\ref{2mass_b}. The sources lying above the ZAMS curve of spectral type O9 are shown as asterisk symbols while the `asterisk' having an infrared excess is shown as a solid triangle. The open squares represent sources detected either in H and K$_s$ bands or only in K$_s$ band. The other sources are represented with plus-symbols. The seven sources discussed above are labelled as ``IRB\" in the figure. It is interesting to note the presence of young objects (detected in H and K$_s$ or only K$_s$ bands) clustered near the IRAS peak.  In the \\mbox{IRAS~10031-5632} region, we find that there is 1 MSX PSC source within $1'$ radius centered on the IRAS peak, G281.5857-00.9706 (designated as M3 and listed in Table~\\ref{msx}). M3 could possibly be associated with either IRB-3 ($\\sim2.8''$) or IRB-2 ($\\sim2.8''$). In both the cases, we obtain the 2MASS-MSX colours ($F_{21}/F_8$ and $F_8/F_K$, Lumsden et al. 2002) and observe that M3 lies in the general region covered by compact H II regions."
    },
    "0807/0807.0341_arXiv.txt": {
        "abstract": "This article investigates the non-linear evolution of cosmological perturbations on sub-Hubble scales in order to evaluate the unavoidable deviations from Gaussianity that arise from the non-linear dynamics. It shows that the dominant contribution to modes coupling in the cosmic microwave background temperature anisotropies on small angular scales is driven by the sub-Hubble non-linear evolution of the dark matter component. The perturbation equations, involving in particular the first moments of the Boltzmann equation for photons, are integrated up to second order in perturbations. An analytical analysis of the solutions gives a physical understanding of the result as well as an estimation of its order of magnitude. This allows to quantify the expected deviation from Gaussianity of the cosmic microwave background temperature anisotropy and, in particular, to compute its bispectrum on small angular scales. Restricting to equilateral configurations, we show that the non-linear evolution accounts for a contribution that would be equivalent to a constant primordial non-Gaussianity of order $\\fNL\\sim25$ on scales ranging approximately from $\\ell\\sim1000$ to $\\ell\\sim3000$. ",
        "introduction": "\\label{sec1}  The cosmic microwave background (CMB) offers a unique window on the physics of the early Universe, and in particular on inflationary models. The angular power spectrum of the CMB anisotropies has been extensively used to set constraints on the shape of the inflationary potentials; see e.g. Ref.~\\cite{Komatsu:2008hk}. The statistical properties of the temperature anisotropies and polarisation depend both on the inflationary period during which they were created and on the physics at play after Hubble-radius crossing and during the recombination. At linear order in metric perturbations, those latter physical processes amount to affect the metric perturbations by a multiplicative transfer function. The characteristic features observed in the temperature anisotropy spectrum originate from the development of acoustic oscillations that this transfer function encodes.  The overall amplitude of the metric perturbation and its scale dependence are however determined by the inflationary phase.  At linear order, the calculation of the transfer function - and hence the detailed shape of the temperature power spectra - for generic inflationary models requires the identification of the relevant degrees of freedom during inflation (see e.g. Refs.~\\cite{1980PhRvD..22.1882B,1992PhR...215..203M,1996PThPh..95...71S}), as well as a full resolution of the dynamics up to recombination time. All these aspects are now fully understood (see e.g. Refs.~\\cite{PeterUzanBook,BernardeauBook} and references therein).  At this level of description, the metric perturbations are linearised so that the non-linear couplings that are inherently present in the Einstein equations are ignored. Therefore models that predict Gaussian initial metric fluctuations  are expected to induce cosmic fields with Gaussian statistical properties. This is a priori the case for generic models of inflation. It has to be contrasted to models with active topological defects - such as cosmic strings - that have soon been recognised as a source of large non-Gaussianities~\\cite{1994PhRvD..49..692P,Gangui:1996cg,Durrer:2000gi,Perivolaropoulos:1992gy}. The current data however clearly favor only mild non-Gaussianities although those might be larger than those induced by pure gravity couplings. This is not the case for single field slow-roll inflation for which it has been unambiguously shown in Ref.~\\cite{2003JHEP...05..013M} that it can produce only very weak non-Gaussian signals, that are bound to be overridden by the gravity induced couplings. It has however been realised that some models of inflation might produce significant deviation from Gaussianity in the context of multiple-field inflation~\\cite{2002astro.ph..6039K, 2002PhRvD..65j3505B, 2002PhRvD..66j3506B, 2003PhRvD..67l1301B, 2002PhLB..524....5L, 2004PhR...402..103B, 2004JHEP...04..006B,Bernardeau:2006tf,2004PhRvD..70d3533B} or with non-standard kinetic terms~\\cite{2004PhRvD..70l3505A}. The question of the observation of primordial non-Gaussianities is largely open.\\\\  In general however, primordial deviations from Gaussianity are in competition with the couplings induced during the non-linear evolution of the cosmic fields. It has triggered general studies aiming at characterising the bispectrum to be expected in the observation of the cosmic microwave background temperature anisotropies and polarisations whether it arises from inflation or from subsequent effects.  This task is multi-fold. It requires a proper identification of the mode couplings (at the quantum level) during the inflationary phase - so taking into account the usual gauge freedom - as well as a second order treatment of the post-inflationary evolution. While the former has been set on firm ground~\\cite{2003JHEP...05..013M,2005PhRvD..72d3514W}, the latter issue is still largely unexplored. This article proposes both numerical and analytical insights into it.  Hereafter, we assume that on super-Hubble scales, the only significant scalar perturbations are adiabatic and that they obey a nearly Gaussian statistics~\\footnote{Note that the choice of $\\zeta$ as the primordial field is not unique and one could have chosen the Bardeen potential. With such a choice however the $\\fNL^{\\zeta}$ incorporates only the inflation dependent couplings - $\\fNL^\\zeta$ is proportional to the slow roll parameter in single field inflation for instance. The other coupling terms induce by the change of variable can be incorporated into $\\mT^{(2)}$; see Eq.~(\\ref{fnltdef}) below.}. To be more precise, they are described in Fourier space by a single variable $\\zeta_{0}(\\vk)$, $\\vk$ being a comoving wave-number, that satisfies \\begin{equation}\\label{zetaP} \\langle \\zeta_{0}(\\vk_{1})\\zeta_{0}(\\vk_{2})\\rangle =\\Dirac(\\vk_{1}+\\vk_{2})\\ P_{\\zeta}(k_{1}) \\end{equation} and \\begin{widetext} \\begin{eqnarray}\\label{zetaBispec} \\langle\\zeta_{0}(\\vk_{1})\\zeta_{0}(\\vk_{2})\\zeta_{0}(\\vk_{3})\\rangle= \\hspace{0cm}2\\Dirac(\\vk_{1}+\\vk_{2}+\\vk_{3})\\ \\fNL^{\\zeta}(\\vk_{1},\\vk_{2}) P_{\\zeta}(k_{1})P_{\\zeta}(k_{2})+\\sym \\ , \\end{eqnarray} where ``$\\sym$'' stands for the two other terms obtained by permutation of the wave-numbers. This defines the primordial power spectrum $P_\\zeta(k)$ and the primordial mode coupling amplitude~\\footnote{This is the expression for the bispectrum obtained assuming $\\zeta$ could be expanded as $\\zeta=\\zeta_{G}+\\fNL^\\zeta\\zeta_{G}\\zeta_{G}$ where $\\zeta_{G}$ is assumed to obey Gaussian statistics. This is not however a valid description when the bispectrum originates from multiple-field couplings or from quantum calculation. The formal expression (\\ref{zetaBispec}) is always valid though; see Refs.~\\cite{2004PhRvD..69f3520B,2005PhRvD..72d3514W}.} $\\fNL^{\\zeta}$. Considering an observable quantity $\\theta$ related to the perturbation variables, the effect of evolution can generically be recapped~\\footnote{Things are actually slightly more complicated since usually observables cannot be decomposed into 3D Fourier modes. The functions $\\mTt^{(1)}(k)$ and $\\mTt^{(2)}(\\vk_{1},\\vk_{2})$ should then be thought as projection operators. This is in particular the case for temperature anisotropies and polarisations. This does not affect however the general point we want to make in this introduction.} as \\begin{eqnarray} \\label{FormalExp} \\theta(\\vk)&=&\\mTt^{(1)}(k)\\, \\zeta_{0}(\\vk) +\\int\\frac{\\dd^3\\vk_{1}\\dd^3\\vk_{2}}{(2\\pi)^{3/2}}\\,\\Dirac(\\vk\\!-\\!\\vk_{1}\\!-\\!\\vk_{2}) \\mTt^{(2)}(\\vk_{1},\\vk_{2})\\, \\zeta_{0}(\\vk_{1}) \\zeta_{0}(\\vk_{2}) +\\dots\\ , \\end{eqnarray} where $\\mTt^{(1)}(k)$ is the linear transfer function and $\\mTt^{(2)}(\\vk_{1},\\vk_{2})$ is the second order transfer function. $\\theta$ can be thought as being e.g. the observed temperature anisotropies, but it could also stand for the CMB polarisation, or even cosmic shear surveys. When computing the bispectrum of $\\theta$ there will be a contribution from the mode couplings induced by the second order transfer function $\\mT^{(2)}$ and the possible initial non-Gaussianities, \\begin{eqnarray}\\label{FormalBispec} \\langle \\theta(\\vk_{1})\\theta(\\vk_{2})\\theta(\\vk_{3})\\rangle&=& 2\\Dirac(\\vk_{1}+\\vk_{2}+\\vk_{3})\\left[\\fNL^{\\zeta}(\\vk_{1},\\vk_{2})+ \\fNL^{\\theta}(\\vk_{1},\\vk_{2})\\right]\\mTt^{(1)}(k_{1}) \\mTt^{(1)}(k_{2}) \\mTt^{(1)}(k_{3})\\ P_{\\zeta}(k_{1})P_{\\zeta}(k_{2})\\nonumber\\\\ &&\\qquad\\qquad+\\sym \\end{eqnarray} \\end{widetext} where $\\fNL^{\\theta}$ is related to the second order transfer function by \\begin{equation}\\label{fnltdef} \\mTt^{(2)}(\\vk_{1},\\vk_{2})\\equiv \\fNL^{\\theta}(\\vk_{1},\\vk_{2})\\,\\mTt^{(1)}(|\\vk_{1}+\\vk_{2}|)\\ . \\end{equation} The full derivation of the details of $\\mTt^{(2)}$ is a fantastic task. It requires an understanding of the metric fluctuations behaviour at second order, from radiation dominated super-Hubble scales to matter dominated era at sub-Hubble scale, as well as a comprehension of the physics of recombination - through the Boltzmann equation - at a similar order. Such a task has been undertaken by several authors~\\footnote{Early derivations are to be found in Ref.~\\cite{Maartens:1998xg,2006JCAP...06..024B,2007JCAP...01...19B}. A more rigorous and comprehensive calculation -- including a proper derivation of the Boltzmann coupling terms and taking into account the polarisation effects -- is to be found in Refs.~\\cite{Pitrou:2007jy,pitrounew}.} and the multitude of effects at play needs to be sorted out. So, the goal of this article is not to provide an end to end calculation of $\\mTt^{(2)}$, but to show that on small scales one can extract the dominant terms in order to get an insight into this physics at second order.  Modes coupling due to gravitational clustering is, by far, not a novel subject. It can be traced back to early works by Peebles \\cite{1980lssu.book.....P} where the function $\\mT^{(2)}$ for the non-linear sub-Hubble evolution of cold dark matter field (CDM) during a matter dominated era was derived. General modes coupling effects, within the same regime, has been extensively studied in the eighties and nineties where a whole corpus of results has been obtained (see e.g. Ref.~\\cite{2002PhR...367....1B} for an exhaustive review). On sub-Hubble scales, the second order mode coupling function for the gravitational potential reads \\begin{eqnarray}\\label{F2Peeb} && \\fNL^{\\Phi}(\\vk_{1},\\vk_{2}) =\\frac{k_{1}^2\\,k_{2}^2}{\\frac32\\,H^2\\,a^2\\,\\vert \\vk_{1}+\\vk_{2}\\vert^2}\\nonumber\\\\ &&      \\quad\\times\\left(\\frac{5}{7}+\\frac{1}{2} \\frac{\\vk_{1}\\cdot\\vk_{2}}{k_{1}^2} +\\frac{1}{2}\\frac{\\vk_{1}\\cdot\\vk_{2}}{k_{2}^2} +\\frac{2}{7}\\frac{(\\vk_{1}\\cdot\\vk_{2})^2}{k_{1}^2k_{2}^2}\\right) \\end{eqnarray} in the particular case of an Einstein-de Sitter universe (here $a$ is the scale factor and $H$ the Hubble parameter). This well-established result proved to be useful for observational cosmology. The angular modulation it exhibits has indeed been observed in actual data sets; see e.g. Ref.~\\cite{2001ApJ...546..652S}.  The fact that on sub-Hubble scales, that is $k^2\\gg H^2\\,a^2$, the non-Gaussianity is driven by the non-linearities of the CDM sector, and that they can start developing even before equality is one of the leading ideas of the present study. Indeed temperature anisotropies on small angular scales -- i.e. beyond the first acoustic peak -- mostly trace the gravitational potential~\\footnote{At least to some extent as we shall see in the course of this paper.} long after it has entered a sub-Hubble evolution and already during the matter dominated era. It is then natural to expect that the temperature anisotropies should be substantially determined by a form close to that of Eq.~(\\ref{F2Peeb}).  The goal of this paper is to evaluate how close we are from the behaviour~(\\ref{F2Peeb}) depending on scales, to which extent the temperature anisotropies trace this form and finally to estimate the amplitude of the temperature bispectrum on small angular scales. In this work two approaches will be compared: a full numerical integration of the second-order equations presented in \\S~\\ref{sec2}, where the main approximation lies in the modelisation of the Compton scattering collision term at second order, see Eq.~(\\ref{e.Hyp}), and an approximate analytical resolution discussed in \\S~\\ref{sec3a}.  The bottom line of our analysis is that on small angular scales, the density perturbation of the cold dark matter starts to dominate the Poisson equation so at the time of decoupling we can assume that the system is split in (1) the evolution of CDM and (2) the evolution of the photons-baryons plasma which develops acoustic oscillations in the gravitational potential determined by the CDM component. As we shall demonstrate, at second order the dominant term of the temperature fluctuations is driven by the second order gravitational potential. Our approximation requires to consider a regime in which the Silk damping is efficient, that is wave-modes larger than the damping scales, hence corresponding to multipoles roughly larger that 2000. This picture will be shown to be in agreement with the numerical estimation (see \\S~\\ref{sec-num-an}). We then proceed in \\S~\\ref{sec3} by a computation of the bispectrum in which we show that, for equilateral configurations, the non-linear dynamics has an amplitude equivalent to that of a primordial non-Gaussianity with constant $\\fNL$ of order 25. A back-of-the-envelop argument allows us to understand the magnitude of this number. ",
        "conclusions": "\\label{sec5}  This article investigates the non-Gaussianity that arises in the CMB temperature anisotropies due to the post-inflationary non-linear dynamics during the radiation and matter dominated era. More specifically it aims at identifying the leading mechanisms and leading terms that determine the shape of the CMB bispectrum on small angular scales.\\\\  The driving idea that we have pursued throughout the paper is that at small angular scales the second order CMB anisotropies trace the second order gravitational potential as it is shaped by the CDM component during its sub-Hubble evolution. To give support to this picture, we have developed both analytical insights into the joint evolution of the density potentials and the temperature fluctuations and numerical tools.  We have thus solved numerically the joint evolution equations of the cosmic fluids up to second order. We have been able to check that at the time of decoupling, the second order potential indeed traces its expected shape. This conclusion is summarised and illustrated on Fig.~\\ref{Phi2Tight}. At this stage, and as long as one restricts these results to the tight coupling regime, no approximations have been made. The accuracy with which the $k$-dependence of the matter dominated mode coupling kernel, i.e. Eq.~(\\ref{F2Peeb}), is recovered is truly remarkable. We stress that this is due to the fact that for the physics at work at small scales, i.e. $k/k_\\EQ\\gg1$, the density perturbations of the CDM component start to dominate the Poisson equation much before equality. This implies that the non-linearities developed by the CDM can be transferred very efficiently to the gravitational potential even before the beginning of the matter era.  Determining exactly how this mode coupling kernel is actually transferred to the source term of the CMB anisotropies relies on further numerical integrations through the recombination era. At this stage, the {\\it only approximation} we make concerns the Compton scattering collision term entering the Boltzmann equation for radiation at second order. We did not use its full second order expression but we argue that it can be reduced to its formal first order form, namely to Eq.~(\\ref{e.Hyp}). Such an assumption is clearly valid in the tight coupling regime. Actually this is the only term appearing in the collision term in the baryons rest-frame as long as tight coupling at first order is efficient. We then argue that when the coupling drops, Silk damping effects effectively suppress all other contributions.  This leads to the behaviour depicted on Fig.~\\ref{SW2Fig} for the main source term of the temperature anisotropies. In particular $\\Theta_{SW}^{(2)}$, the monopole of the second order source term, is found to be attracted toward a non-vanishing and non-oscillatory term, $-R\\Phi^{(2)}$, where we recall that $R$ is the baryon to radiation ratio, while the dipole contribution, and thus the Doppler effect, vanishes. As a result the main contribution to the CMB temperature anisotropies at second order is found to be directly proportional to the second order gravitational potential. It has to be noted that the efficiency with which the second order term converges to this form is considerably accelerated by the Silk damping effects which efficiently suppress the oscillatory parts of the solution. This results are clearly illustrated on Fig.~\\ref{SW2Fig}. We argued from order of magnitude arguments that, because the damping scale is typically of order $15k_\\EQ$, our description shall be valid for $l\\gtrsim2400$. This observation is the basis of the main result of this paper. Actually,as Fig.~\\ref{figbispec} tends to show, it seems that this description may be valid at lower multipoles.  Finally we explore the consequence on the CMB bispectrum. For obvious reasons we use the small angle approximation to perform the numerical integrations. The bispectrum for equilateral configurations is illustrated on Fig.~\\ref{figbispec}. We show that for these configurations, its amplitude corresponds to what a primordial non-Gaussian potential of $\\fNL^{\\Phi}$ of order 25 would have given (also for an equilateral configuration). As shown in the text, this number can easily be recovered from back-of-the-envelop calculations. The first lesson that can be drawn from this result is that it gives a signal larger than what a model with a primordial $\\fNL$ of order unity would give! The second lesson is that the $l$-dependence of the bispectrum is expected to be different from the one induced by primordial mode couplings. It is expected to have a specific shape as encoded in the CDM kernel expression.\\\\  In conclusion, this work offers a breakthrough insight into the physics of CMB in the non-linear regime and on small angular scales. It identifies what is, as we argued, the main small scale contribution of the bispectrum, hence filling the gap with the standard results that have been obtained in the weakly non-linear regime of gravitational clustering of dark matter. We did not check this result against a  (yet non-existing) full second order Boltzmann code, and this is probably desirable, but we argue that, given the amplitude of the effects, all other contributions will be subdominant.  With such a large signal, detection of this bispectrum should be easily within reach of future CMB experiments!\\\\  {\\bf Acknowledgements:} We thank J. Martin-Garcia for his help in using the tensorial perturbation calculus package xPert~\\cite{xact} that was used to derive the second order expressions of this paper. We also thank G. Faye, Y. Mellier, S. Prunet, D. Spergel and N. Aghanim for many discussions."
    },
    "0807/0807.4754_arXiv.txt": {
        "abstract": "We report on a new mechanism for heat conduction in the neutron star crust. We find that collective modes of superfluid neutron matter, called superfluid phonons (sPhs), can influence heat conduction in magnetized neutron stars. They can dominate the heat conduction transverse to magnetic field when the magnetic field $B \\gsim 10^{13}$ G. At density $\\rho \\simeq 10^{12}-10^{14} $ g/cm$^3$ the conductivity due to sPhs is significantly larger than that due to lattice phonons and is comparable to electron conductivity when temperature $\\simeq 10^8$ K. This new mode of heat conduction can limit the surface anisotropy in highly magnetized neutron stars.  Cooling curves of magnetized neutron stars with and without superfluid heat conduction could show observationally discernible differences. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0807/0807.3342_arXiv.txt": {
        "abstract": "{ Silicate carbon stars show the 10~\\micron\\ silicate emission, despite their carbon-rich photospheres. They are considered to have circumbinary or circum-companion disks, which serve as a reservoir of oxygen-rich material shed by mass loss in the past. } {We present $N$-band spectro-interferometric observations of the silicate carbon star BM~Gem using MIDI at the Very Large Telescope Interferometer (VLTI). Our aim is to probe the spatial distribution of oxygen-rich dust with high spatial resolution. } { BM~Gem was observed with VLTI/MIDI at 44--62~m baselines using the UT2-UT3 and UT3-UT4 baseline configurations. } { The $N$-band visibilities observed for BM~Gem show a steep decrease from 8 to $\\sim$10~\\micron\\ and a gradual increase longward of $\\sim$10~\\micron, reflecting the optically thin silicate emission feature emanating from sub-micron-sized amorphous silicate grains. The differential phases obtained at baselines of $\\sim$44--46~m show significant non-zero values ($\\sim \\! -70$\\degr) in the central part of the silicate emission feature between $\\sim$9 and 11~\\micron, revealing a photocenter shift and the asymmetric nature of the silicate emitting region. The observed $N$-band visibilities and differential phases can be fairly explained by a simple geometrical model in which the unresolved star is surrounded by a ring with azimuthal brightness modulation. The best-fit model is characterized by a broad ring ($\\sim$70~mas across at 10~\\micron) with a bright region which is offset from the unresolved star by $\\sim$20~mas at a position angle of $\\sim$280\\degr. This model can be interpreted as a system with a circum-companion disk and is consistent with the spectroscopic signatures of an accretion disk around an unseen companion recently discovered in the violet spectrum of BM~Gem. } {} ",
        "introduction": "\\label{sect_intro}  ``Silicate carbon stars'' are characterized by oxygen-rich circumstellar material as evident from the prominent silicate emission, despite their carbon-rich photospheres (Little-Marenin \\cite{little-marenin86}; Willems \\& de Jong \\cite{willems86}). The detection of masers from O-bearing molecules such as OH and \\HOH\\ from some of the silicate carbon stars confirms the presence of oxygen-rich circumstellar gas (e.g., Nakada et al. \\cite{nakada87}; Engels \\cite{engels94}; Szymczak et al. \\cite{szymczak01} and references therein). Moreover, optical spectroscopic studies (e.g., Lloyd-Evans \\cite{lloyd-evans90}; Chan \\cite{chan93}; Ohnaka \\& Tsuji \\cite{ohnaka99}) reveal that silicate carbon stars have \\cisotp\\ ratios as low as $\\sim$4--5 (classified as ``J-type''), which is difficult to explain by the present stellar evolution theory. Currently, silicate carbon stars are considered to have an unseen, low-luminosity companion. The oxygen-rich material which was shed by mass loss when the primary star was an M giant can be stored in a circumbinary disk (Morris \\cite{morris87}; Lloyd-Evans \\cite{lloyd-evans90}) or in a circumstellar disk around the companion (``circum-companion disk'', Yamamura et al. \\cite{yamamura00}), even after the primary star becomes a carbon star. In the latter case, Yamamura et al. (\\cite{yamamura00}) suggest that the primary star drives an outflow from the circum-companion disk, and the silicate emission emanates from this outflow.  Thanks to the recent progress of high-angular resolution observation techniques, there is growing observational evidence for the presence of circumbinary or circum-companion disks around silicate carbon stars. Our $N$-band spectro-interferometric observations of IRAS08002-3803 with the MID-infrared Interferometric instrument (MIDI) at ESO's Very Large Telescope Interferometer (VLTI) spatially resolved the dusty environment of a silicate carbon star for the first time (Ohnaka et al. \\cite{ohnaka06}), and our radiative transfer modeling suggests the presence of an optically thick circumbinary disk in which amorphous silicate and a second grain species (amorphous carbon, large silicate grains, or metallic iron) coexist. The presence of a circumbinary disk with multiple grain populations is also suggested for another silicate carbon star IRAS18006-3213 by Deroo et al. (\\cite{deroo07}). The spatial distributions of the 22 GHz \\HOH\\ maser emission toward two silicate carbon stars, V778~Cyg and EU~And, have recently been resolved by Szczerba et al. (\\cite{szczerba06}), Engels (priv. comm.), and Ohnaka \\& Boboltz (\\cite{ohnaka_boboltz08}) using MERLIN and the Very Long Baseline Array (VLBA). The observed spatio-kinematic structures are consistent with the presence of circum-companion disks. However, the formation mechanisms of these disks are little understood, and the possible link between the disk formation and the anomalously low \\cisotp\\ ratios is by no means clear, either.   Direct detection of a companion toward a silicate carbon star is hampered by the huge luminosity contrast between the primary carbon star ($10^3$--$10^4$~\\LSOL) and the companion ($\\la$1~\\LSOL), as well as small angular separations (e.g., Engels \\& Leinert \\cite{engels_leinert94}). However, Izumiura (\\cite{izumiura03}) and Izumiura et al. (\\cite{izumiura08}) have overcome the first problem by taking advantage of the fact that the violet flux of a carbon star (primary star) is heavily suppressed due to molecular and atomic absorption.  Using the High Dispersion Spectrograph (HDS) on the Subaru Telescope, they discovered continuum emission below 4000~\\AA\\ and Balmer lines showing P~Cygni profiles with an outflow velocity of $\\ga$400~\\KMS\\ toward the silicate carbon star BM~Gem. These spectroscopic signatures strongly suggest the presence of an accretion disk around an unseen companion, although it could not be spatially resolved, and its nature---whether it is a main-sequence star or a white dwarf---remains to be determined. Comparison between the observationally derived accretion luminosity ($\\sim$0.2~\\LSOL) and the theoretical prediction from the Bondi-Hoyle-type accretion resulted in estimated binary separations of $\\sim$3--60~AU for a main-sequence star companion and $\\sim$100--500~AU for a white dwarf companion. These values translate into angular separations of 2.5--50~mas and 83--420~mas for a distance of 1.2~kpc adopted for BM~Gem by Izumiura et al. (\\cite{izumiura08}).  Given these estimated separations, it is challenging to resolve the circum-companion disk toward BM~Gem with a single-dish telescope. On the other hand, infrared interferometry provides us with a unique opportunity to achieve sufficiently high spatial resolution. In this paper, we present the results of mid-infrared spectro-interferometric observations of BM~Gem with VLTI/MIDI to probe the presence of the circum-companion disk with dust thermal emission. ",
        "conclusions": "\\label{sect_concl}   Our VLTI/MIDI observations of the silicate carbon star BM~Gem have revealed that the $N$-band visibilities are characterized by a steep decrease from 8 to $\\sim$10~\\micron\\ and a gradual increase longward of 10~\\micron, reflecting the optically thin 10~\\micron\\ emission feature emanating from sub-micron-sized amorphous silicate grains. Furthermore, we have detected significant non-zero differential phases ($\\sim \\! -70$\\degr) between $\\sim$9 and 11~\\micron, which reveals the asymmetric nature of the silicate emitting region. The wavelength dependence of the observed $N$-band visibilities and differential phases is remarkably different from that previously observed toward two silicate carbon stars, IRAS08002-3803 and IRAS18006-3213. This is most likely to reflect the differences in mineralogy, optical thickness, and geometry (circumbinary or circum-companion) of the disks.   Our simple geometrical modeling of the observed visibilities and differential phases suggests the presence of a bright region which is offset from the central star by 13--34~mas (16--41~AU at 1.2~kpc) at a position angle of $280 \\pm 40$\\degr. This bright region can be interpreted as a circum-companion disk. The derived offset from the central star makes a white dwarf companion unlikely (thus leaving a main-sequence star companion more favorable), because a white dwarf companion at these separations would produce too much accretion luminosity compared to that observed.  However, the lack of detailed information on the geometry of the silicate dust distribution makes this circum-companion disk interpretation not yet definitive.  For example, we cannot yet entirely exclude the possibility that the asymmetric-disk model represents a circumbinary disk with an asymmetric structure. MIDI observations with shorter baselines are indispensable for putting tighter constraints on the size and shape of the silicate emitting region. Also, high-resolution imaging at violet/blue wavelengths is important for direct detection of the circum-companion accretion disk.  Combination of the violet/blue high-resolution imaging and MIDI observations will provide definitive confirmation that silicate dust is stored in a circum-companion disk.      \\begin{acknowledgement} We thank the ESO VLTI team on Paranal and in Garching and the MIDI team for carrying out the observations and making the data reduction software publicly available. We also thank the referee, H.~van~Winckel, for his constructive comments.    \\end{acknowledgement}"
    },
    "0807/0807.3497_arXiv.txt": {
        "abstract": "We model the high-energy emission that results from  the interaction of relativistic particles with photons and matter in the cluster of galaxies Abell 3376. The presence of relativistic particles is inferred from the recently found radio relics in this cluster, being one of the most prominent examples of double opposite, giant ringlike radio structures. Assuming that diffusive shock acceleration takes place in the cluster regions where radio relics are observed, we calculate the spectral energy distribution resulting from the most relevant non-thermal processes, which are synchrotron radiation, inverse Compton scattering, relativistic Bremsstrahlung, and inelastic proton-proton collisions. In the context of our model, the major radiative component at high energies is inverse Compton scattering, which could reach luminosities $L\\sim 9\\times 10^{41}$~erg~s$^{-1}$ in the energy range between $\\sim 1$~MeV and 10~TeV. Hadronic interactions would yield a minor contribution to the overall non-thermal emission, but would dominate at ultra-high energies. The cluster Abell~3376 might be detectable at gamma-rays by HESS, GLAST satellite and future planned Cherenkov arrays. ",
        "introduction": "\\label{s_intro}   Clusters of galaxies are characterized by thermal emission in the $2$ to $10$~keV soft X-ray region ($L_{\\rm X} \\sim 10^{43} - 10^{45}\\;\\rm{erg\\;s^{-1}}$) as a result of Bremsstrahlung radiation of the hot intracluster gas ($kT\\sim 5-8$~keV). However, the detection of radiation at radio wavelengths \\citep{Feretti04}, as well as in the extreme ultraviolet ($0.07-0.4$~keV; EUV) \\citep{Bowyer04} and at hard X-rays ($20-80$~keV; HXR) \\citep{fuscofemiano04}, indicates the presence of non-thermal activity in these systems. There is general consensus that radio emission is produced by synchrotron radiation originated from the interaction of relativistic electrons with the cluster magnetic field ($B \\sim \\mu{\\rm G}$) \\citep{GovoniFeretti04}. On the other hand, EUV radiation might be produced either by a cooler thermal component ($kT\\sim 2$~keV) (\\citealt{Lieu96}; \\citealt{Mittaz98}) or by inverse Compton (IC) scattering of cosmic microwave background (CMB) photons by the same population of relativistic electrons responsible for the radio emission (\\citealt{EnsslinBiermann98}; \\citealt{Lieu99}). The latter process might also explain the HXR emission \\citep{FuscoFemiano99}, although this possibility is still matter of debate.  Regarding the radio emission, whose non-thermal origin is firmly established, clusters of galaxies present two different sources of diffuse large-scale synchrotron emission, known as `radio halos' and `radio relics'. The former are located at the centre of clusters and are characterized by unpolarized radio emission with structures that roughly resemble that shown by the X-ray emission, whereas the latter are polarized radio sources more irregularly shaped found at the periphery of clusters (e.g. \\citealt{FerettiGiovannini08}). Radio relics have been observed in several clusters. The most extended and powerful sources of this class have been detected in clusters with central radio halos, such as Coma \\citep{Giovannini91} and Abell clusters A2163 \\citep{Feretti01}, A2255 \\citep{Feretti97}, A2256 \\citep{Rottgering94} and A2744 \\citep{Govoni01}. Only few clusters present double opposite relics, being the most prominent examples those found in A3667 \\citep{Rottgering97}, and A3376 \\citet{Bagchi06}. The large radio flux recently detected from the relics of A3376 is an evidence of that relativistic leptons are present in those regions. For this reason, in the present work we will be concerned with the non-thermal processes taking place in the outskirts of the cluster Abell 3376 where the two giant and almost symmetric radio relics are located. Nevertheless, the model that we will present in this paper is also valid for clusters which present only one relic. In the case of Abell 3376 the radio flux is more than one order the magnitude larger than in other radio emitting clusters, like Abell 521 and Abell 2163.   There is increasing evidence that the radio emitting particles in relics are accelerated by accretion and merger shocks generated in the intracluster medium (ICM) during cosmological large-scale structure formation. Because of the electron short radiative lifetimes, radio emission can be efficiently produced close to the location of the shock waves. Thus, the origin of radio relics can be explained if they are considered to trace the position of these very large shocks (\\citealt{Ensslin98}; \\citealt{Hoeft04}). According to the standard diffusive shock acceleration theory, shock waves in the presence of even modest and turbulent magnetic fields are sites of efficient acceleration of charged particles \\citep{Drury83}. Another possibility may be adiabatic compression in these environmental shock waves \\citep{EnsslinGopal01}.  Based on the inferred population of high energy particles involved in the previously mentioned non-thermal processes, it is natural to expect gamma-ray emission from galaxy clusters. This emission could be generated from neutral pion decay in cosmic ray collisions in the ICM \\citep{VolkAha96}, or by IC scattering that involves CMB photons and relativistic electrons \\citep{AtoyanVolk00}. Hydrodynamical cosmological simulations are a useful tool to estimate the properties of the cosmic ray populations in galaxy clusters and their effects on thermal cluster observables, as well as to predict the radiation emitted by these large virialized objects at $\\gamma$-ray energies (\\citealt{Keshet03}, \\citealt{Pfrommer07}, \\citealt{Pfrommer08}). A different numerical approach has been used by \\citet{BerringtonDermer03} to investigate the temporal evolution of particle and photon spectra resulting from non-thermal processes at the shock fronts formed in merging clusters of galaxies. All these works point to clusters in general as sources of $\\gamma$-ray emission observable with GLAST satellite.  These theoretical results are consistent with the fact that no $\\gamma$-ray emission has been detected so far with the current observational facilities, as indicated by the lack of correlations between unresolved EGRET $\\gamma$-ray sources and nearby X-ray bright galaxy clusters \\citep{Reimer03}. Only marginal evidence for emission from A1758 within the location error contours of the source 3 EG J1337+5029 was reported by \\citet{Fegan05}. Perseus and Abell 2029 galaxy clusters were observed by \\citet{Perkins06} with the Whipple 10 m Cherenkov telescope. They find no evidence of point source or extended $\\gamma$-ray emission in the TeV energy range. In addition, recent observations of the clusters Coma and Abell 496 made by \\citet{Domainko07} with the array of Chrerenkov telescopes HESS have not detected significant signals of $\\gamma$-ray emission in exposure times of $\\sim 10-20$ hours.  In the present work, we estimate the high energy emission of the recently detected radio relics in the nearby cluster Abell 3376. Their radio power at $1.4$~GHz, together with the assumption of an equipartition magnetic field, constrain effectively our model which is specific of this particular cluster. We assume a diffusive shock acceleration mechanism, and include different non-thermal processes that may be acting in the radio structures. We predict that the $\\gamma$-ray photons created in this cluster should be detected in the near future by GLAST, as well as by the new generation of Cherenkov arrays in the southern hemisphere, such as HESS II.   The paper is organized as follows. Section \\ref{s_cluster} describes the main features of the cluster A3376. Section \\ref{s_accel} presents the acceleration and loss mechanisms that affect the content of relativistic particles in the ICM. The estimates of production of $\\gamma$-rays and lower energy radiation are given in Section \\ref{s_emission}. Finally, we summarize our results and present our concluding remarks in Section \\ref{s_prospects}. ",
        "conclusions": "\\label{s_prospects}  We have presented a model for the emission produced by the nearby cluster Abell 3376 in a wide range of energies, from radio to gamma-rays, considering the contribution of different non-thermal radiative processes. This cluster presents strong evidence of merger activity, characterized by large-scale ring-shaped synchrotron radio structures, that are identified with radio relics \\citep{Bagchi06}. These kind of radio structures detected in several clusters of galaxies suggest a rich content of accelerated particles in the ICM, with energies as high as  $\\sim 10^{14}$ and $\\sim 10^{18}$ eV for electrons and protons, respectively. Taking into account that radio relics are tracers of merger shocks, we assume that a diffusive shock acceleration mechanism acting at the wave shocks gives place to such a population of relativistic particles. These particles are subsequently cooled by different leptonic and hadronic radiative processes.  We have considered syncrotron radiation, IC scattering and relativistic Bremsstrahlung that affect relativistic primary electrons. Relativistic protons are involved in $pp$-interactions with the thermal protons of the ICM. This process give raise to a population of secondary pairs, $e^{\\pm}$, that are cooled by the same processes that affect the population of primary electrons. The parameters involved in these non-thermal processes, such as the magnetic field and normalization of the energy distribution of accelerated particles, have been estimated from the constraints established by the observed radio power in the relics and the assumption of equipartition. X-ray observations allows us to extract information about the volume of these acceleration sources. On the other hand, results from a simulated cluster have provided the typical gas density of the relics and the velocity of ths shock.  We find that particle energy distributions reach a stationary regime by the lifetime of the relics ($\\sim 1$~Gyr). The steady energy distribution of particles are computed and the SEDs of the radiation produced at the relic position derived. In our calculations, we have adopted different ratios ($a = 0, 1, 100$) for the relativistic proton/electron energy densities. At larger energies, gamma-ray emission is mainly produced by $\\pi^0$-decay but with a considerably smaller luminosity, even for the case where a large energy density of protons is considered ($L_{pp} \\sim 4\\times 10^{39}$ erg s$^{-1}$). This is consistent with the fact that no cluster has been yet observed at these high energies.  The proximity of the cluster studied in this paper and its high content of relativistic particles in the radio-relics make Abell 3376 an interesting potential target for the investigation of gamma-ray emission in this type of objects. This source might be detected at $\\gamma$-rays by GLAST satellite and by HESS array with a reasonable exposure time. In addition, future Cherenkov telescopes in the southern hemisphere, like HESS II, could easily detect the high-energy emission from Abell 3376."
    },
    "0807/0807.3204_arXiv.txt": {
        "abstract": "In this paper we will discuss the current status of coherent integration with the Navy Prototype Optical Interferometer (NPOI)\\cite{armstrong:1998}. Coherent integration relies on being able to phase reference interferometric measurements, which in turn relies on making measurements at multiple wavelengths. We first discuss the generalized group-delay approach, then the meaning of the resulting complex visibilities and then demonstrate how coherent integration can be used to perform very precision measurement of stellar properties. For example, we demonstrate how we can measure the diameter of a star to a precision of one part in 350, and measure properties of binary stars. The complex phase is particularly attractive as a data product because it is not biased in the same way as visibility amplitudes. ",
        "introduction": "At the 2004 and 2006 SPIE meetings we demonstrated how to coherently integrate NPOI\\cite{armstrong:1998} data by fitting a model of the fringes to the raw data and using the fitting parameters to coherently integrate the complex visibilities from individual frames \\cite{jorgensen:2004, jorgensen:2006}. This discussion was also published in a separate paper \\cite{jorgensen:2007}. In this paper we take a different approach to fringe tracking which is an extension of the conventional group-delay approach, but including the atmospheric dispersion as well as time-dependence in the group-delay calculation. This is discussed in Section~\\ref{section_group_delay}. We then discuss the meaning of the resulting coherently integrated complex visibilities, including the phase (Section~\\ref{section_coherent_phase}), and the amplitude (Section~\\ref{section_coherent_amplitude}). We then show two examples of how to use complex visibilities for scientific measurements, including an precise diameter measurement (Section~\\ref{section_diameter}) and precise measurements on a binary star (Section~\\ref{section_binary}), before concluding.  The Navy Prototype Optical Interferometer (NPOI) \\cite{armstrong:1998}, located in Flagstaff, AZ, measures fringes by scanning the optical path-difference across several fringe periods (usually between one and eight) and using photon-counting avalanche photo diodes to detect the photons and bin them according to fringe phase at up to 32 different wavelengths in three different spectrographs. This produces a regular array every 2 ms with fringe phase along one dimension and wavelength along the other dimension, making it easy to process the arrays with Fourier-based techniques. For technical reasons the NPOI data has been restricted to 16 wavelength channels in two spectrographs in recent years, and those are the data we will show here. It is important however that NPOI measures many channels (which we will define as approximately ten or more) because it is this multi-wavelength capability which makes coherent integration possible and interesting. Coherent integration consists of two primary steps, as outline in Figure~\\ref{figure_ci_overview}. The first step uses the data to estimate the fringe phase whereas the second step uses the estimate of the fringe phase together with the data to produce coherently integrated visibilities, as outline in the next two sections.  \\begin{figure} \\includegraphics[width=\\linewidth]{ciProcess.eps} \\caption{\\label{figure_ci_overview} An overview of the coherent integration procedure. It consists of two primary steps. The first is the determination of the fringe phase, and the second is the coherent integration using the measured fringe phase.} \\end{figure} ",
        "conclusions": "\\noindent In this paper we have outlined the current approach to coherent integration at the NPOI. The meaning of the visibility phase is now well-understood, the visibility amplitudes can be calibrated, and the resulting coherently integrated visibilities can be used to make high-precision scientific measurements, as illustrated by the measurement of the diameter - and it's variation with wavelength - of $\\nu\\,\\text{Ophiuchus}$."
    },
    "0807/0807.3817_arXiv.txt": {
        "abstract": "We have measured  the  projected rotational velocities (\\vsi) in a number of symbiotic stars and  M giants using high resolution spectroscopic observations. On the basis  of our measurements and data from the literature, we  compare the rotation of mass-donors in symbiotics  with \\vsi\\ of field giants and  find  that: (1) the K giants in S-type symbiotics rotate at \\vsi$\\: >\\!4.5$~\\kms, which is 2-4 times faster than the field K giants; (2) the M giants in S-type symbiotics rotate on average 1.5 times faster than the field M giants. Statistical tests show that these differences are highly significant -- p-value $< 10^{-3}$  in the spectral type bins K2III-K5III, M0III-M6III, and M2III-M5III; (3) our new observations of  D'-type symbiotics also confirm that they are fast rotators.  As a result of the rapid rotation, the cool giants in symbiotics should have 3-30 times larger mass loss rates. Our results suggest also that bipolar ejections in symbiotics seem to happen in objects where the mass donors rotate faster than the orbital period.  All spectra used in our series of  papers can be obtained upon request from the authors. ",
        "introduction": "The Symbiotic stars (SSs)  are interacting binaries consisting of a cool giant and a hot companion. In most cases the hot component is a white dwarf (WD), but it can also be a main sequence star, hot subdwarf, or neutron star (see Kenyon 1986, Corradi et al. 2003). The symbiotic phase represents a late stage in stellar evolution and about 200 such objects are known (Belczy{\\' n}ski et al. 2000). On the basis of their IR properties, SSs have been classified into stellar continuum  (S) and dusty (D or D') types (Allen 1982). The D--type systems contain Mira variables as mass donors. D'-type symbiotics are characterized by an earlier spectral type giant (F-K) and lower dust temperatures. Soker (2002) has predicted theoretically that the cool companions in symbiotic systems are likely to rotate much faster than isolated cool giants or those in wide binary systems.  This is the third in a series of papers exploring the rotational velocities of the mass donating (cool) components of SSs. In this paper we present \\vsi\\ measurements and expand our sample of SSs. We then perform comparative analyses and  explore theoretical predictions that the mass donors in S-type symbiotics are faster rotators compared with  field giants. ",
        "conclusions": "In this paper we report measurements of the projected rotational velocity (\\vsi) of the cool giants in  9 symbiotic stars and 28 field giants, using high resolution spectroscopic observations. Collecting  together all available data (from Paper I,II,III  as well as from the literature), we compare the rotation of the mass donors in symbiotic stars with the rotation  of field  giants of similar spectral types and find  that:  (1) The symbiotic  K giants included in our survey rotate on average more than twice as fast as the field K giants.  (2) The  M giants in symbiotics are more rapid rotators  than most of the field M giants.  (3) For  the D'-type (yellow) symbiotics, 5 out of the six southern D'-type SSs are the fastest or among the fastest rotators in their spectral classes.  (4) As a result of the rapid rotation SSs should have on average $\\sim\\!$10 times larger mass loss rates than normal red giants.  (5)  We find suggestions that in the jet-ejecting symbiotics the mass donors  rotate faster than the orbital periods.  This is the first observational investigation that clearly confirms theoretical predictions that the mass donors in symbiotics are fast rotators."
    },
    "0807/0807.1486_arXiv.txt": {
        "abstract": "We present multi-color light curves for the W UMa-type eclipsing binary TU Boo for two epochs separated by 22 years. An analysis of the O-C diagram indicates the earlier observations took place right in the middle of a major period change, thus allowing for a unique study on mass transfer and period changes in this W UMa-type system. We compute model fits to our light curves, along with the only other published set, using the Wilson-Devinney program, and find temporally correlated changes in the size of the secondary component with anomalies in the O-C diagram. We investigate the cause of these changes and find support for the existence of rapid, large-scale mass transfer between the components. We postulate that this interaction allows them to maintain nearly equal surface temperatures despite having achieved only marginal contact. We also find support for the evolutionary scenario in which TU Boo has undergone a mass ratio reversal in the past due to large-scale mass transfer so that what is presently the secondary component of TU Boo is in an advanced evolutionary state, oversized due to a helium-enriched core, with a total system age of $\\geq$ 10 Gyr. ",
        "introduction": "The eclipsing binary TU Boo, [14$^{h}$ 04$^{m}$ 59$^{s}$, +30$\\degr$ 00$\\arcmin$ 00$\\arcsec$, m$_{V}$ $\\approx$ 12], is a short period W UMa-type system, (P $\\approx$ 0.324 days), initially discovered by \\citet{Guthnick26} and found to have a spectral type of G3 by \\citet{Schwassmann47}. It has  only been subjected to a single published analysis in the past 50 years by \\citet{Niarchos96}, who modeled the system in 1996 based on B and V light curves obtained photoelectrically in 1982. They obtained a slightly over-contact solution for the system with q = 0.498, and noted period changes and other interesting aspects of the system.  The classical theory of W UMa systems comes from the seminal work of \\citet{Lucy68} in which he proposed they were zero-age contact binaries with a common convective envelope. Although the thermodynamic arguments are sound, the formation scenario encounters many observational problems. W UMa systems are virtually non-existant in young clusters, while quite populous in open clusters with ages exceeding 4-5 Gyrs, as well as globular clusters \\citep{Kaluzny93,Rucinski98,Rucinski00}. As well, numerical simulations of binary formation from protostellar clouds favor early fragmentation and formation of detached systems, and seem to preclude the possibility of fission from a rapidly rotating protostar \\citep{Boss93,Bonnell01}. In contrast to Lucy, a new general theory has recently arisen, \\citep{Stepien06,Eker07}, which postulates that W UMa-type systems start out as detached systems with P $\\approx$ 2 days and significantly different component masses. They lose angular momentum via magnetized stellar winds over several Gyr, during which time the initially more massive component evolves to terminal age main-sequence (TAMS). At this time the more massive component starts to rapidly transfer mass to the less massive, and a contact system is formed. Rapid mass transfer continues past the point of mass ratio reversal, so that what was the more massive component becomes the less massive, and visa versa. This continues until a mass ratio of $\\sim$0.5 is reached, at which point a tentative equilibrium is reached and we find a typical W UMa-type system in marginal contact. Over the next several Gyr, the system undergoes thermal oscillations, although overall conservative mass transfer from the now more evolved secondary, oversized due to helium enrichment, to the primary occurs. Eventually the now more massive primary evolves to TAMS, and the stars coalesce into a single, rapidly rotating star. ",
        "conclusions": "We have complied a comprehensive O-C diagram for the past $\\sim$80 years showing major period changes in the TU Boo system, with photometric data obtained at the time of a major period change. We postulate that these changes are be caused by mass transfer between the components due to thermal oscillations resulting from marginal contact. We interpret the 1984 O-C spike, temporally correlated with the increase in the size of the secondary, as evidence for rapid mass-exchange between components in this, and thus likely other W UMa-type systems. This may very well explain how these systems can maintain near-equal surface temperatures despite such marginal degrees of contact. The well-constrained scale of the system due to physical considerations and absolute magnitude calculations validates the results for absolute masses and radii of the components. Along with the observed current mass transfer rate, this strongly supports the theory that secondary stars in W UMa systems are evolved and oversized due to helium-enriched cores."
    },
    "0807/0807.0449_arXiv.txt": {
        "abstract": "An analytic solution for the accretion of ultra-hard perfect fluid onto a moving Kerr-Newman black hole is found. This solution is a generalization of the previously known solution by Petrich, Shapiro and Teukolsky for a Kerr black hole. We show that the found solution is applicable for the case of a non-extreme black hole, however it cannot describe the accretion onto an extreme black hole due to violation of the test fluid approximation. We also present a stationary solution for a massless scalar field in the metric of a Kerr-Newman naked singularity. ",
        "introduction": "% The only known three-dimensional exact solution for the accretion flow onto the Kerr black hole is the analytical solution by Petrich, Shapiro and Teukolsky \\cite{PetShapTeu}. This solution describes the stationary accretion of a perfect fluid with the ultra-hard equation of state, $p=\\rho$, with $p$ being the pressure and $\\rho$ being the energy density, onto a moving Kerr black hole (see also \\cite{shapiro89,psst89,AbrahamsShapiro90,KarasMucha}). Here we generalize this solution to the case of a moving Kerr-Newman black hole. For the Kerr-Newman metric with naked singularity, we present the stationary solution for a massless scalar field.  The problem of a steady-state accretion onto a (moving) black hole can be formulated as follows. Consider a black hole moving through a fluid with a given equation of state, $p=p(\\rho)$ \\footnote{Alternatively, one can choose a reference frame in which the black hole is at rest and the fluid has a non-zero velocity at the infinity.}. Usually it is assumed that the back-reaction of fluid to the metric is negligible, in which case the problem is being solved in the {\\sl test fluid approximation}. It is also assumed that the black hole mass changes in time slowly enough, such that the steady-state accretion is established (so called {\\sl quasi-stationary} process). The goal is to find a stationary solution to the equations of motion for the flowing fluid in the gravitational field of the black hole.  The Fig.~\\ref{arrows} illustrates an example of this solution for the case of fluid passing the Schwarzschild black hole. It is unlikely that the Kerr-Newman black hole (or a naked singularity) is found in astrophysical context, as well as the accreting fluid with the ultra-hard equation of state. However, the theoretical study of these questions can be useful for better understanding both the principal General Relativity problems, e.g. the approaching to the extreme black hole state, and the real matter accretion onto the astrophysical compact objects.  The velocity of relativistic perfect fluid, $u^\\mu$, in the absence of vorticity can be expressed as the gradient of the scalar potential $\\psi$, see, e.~g.~\\cite{LLfluid}, \\begin{equation} hu_\\mu=\\psi_{,\\mu}, \\label{potential} \\end{equation} where $h=d\\rho/dn=(p+\\rho)/n$ is the fluid enthalpy properly normalized, $h=(\\psi^{,\\alpha}\\psi_{,\\alpha})^{1/2}$, $n$ is the particle number density. When the equation of state for the perfect fluid has a simple form, $p=p(\\rho)$, then the number density can be expressed in terms of the energy density and the pressure: \\begin{equation} \\label{n1} \\frac{n(\\rho)}{n_\\infty}= \\exp\\left[\\,\\,\\int\\limits_{\\rho_{\\infty}}^{\\rho} \\frac{d\\rho'}{\\rho'+p(\\rho')}\\right], \\label{numberdens} \\end{equation} where $n_\\infty$ and $\\rho_\\infty$ are correspondingly the number density and the energy density at the infinity. The above equation can be seen as a formal definition of the ``number density'' for a perfect fluid with an arbitrary equation of state $p=p(\\rho)$.  \\begin{figure}[t] \\begin{center} \\includegraphics[width=0.45\\textwidth]{arrows.eps} \\end{center} \\caption{\\label{arrows} The three-velocity field of the fluid, ${\\bf v}$ for the case of Schwarzschild black hole, and we take ${\\bf v}_{\\infty}=0.5$} \\end{figure}  A specific feature of a fluid with the ultra-hard equation of state is that the resulting equation for the scalar potential is linear \\cite{moncrief}. For $p=\\rho$, from Eq.~(\\ref{numberdens}) one easily obtains, $\\rho\\propto n^2$. As a result, the number density conservation, $(nu^\\alpha)_{;\\alpha}=0$, gives the Klein-Gordon equation for a massless scalar field, $\\Box\\psi=0$. This reveals a way to obtain an exact solution of the accretion problem in the Kerr metric by separation of variables and decomposition in spherical harmonic series \\cite{PetShapTeu}. Our study of the perfect fluids and scalar fields in the Kerr-Newman metric closely follows the analysis of Petrich {\\sl et al.} \\cite{PetShapTeu}.  The introduction of the potential $\\psi$, Eq.~(\\ref{potential}), provides a way to solve a problem of hydrodynamics dealing with an equation for a scalar function. This is not surprising, since any scalar field $\\psi$ is fully equivalent to the hydrodynamical description, provided that the vector $\\psi^{,\\mu}$ is timelike. For example, the canonical massless scalar field $\\psi$ with the energy-momentum tensor, \\begin{equation} T_{\\mu\\nu}=\\psi_{,\\mu}\\psi_{,\\nu}- \\frac{1}{2}g_{\\mu\\nu}g^{\\rho\\sigma}\\psi_{,\\rho}\\psi_{,\\sigma}. \\label{energymom} \\end{equation} is equivalent to a perfect fluid with the ultra-hard equation of state, provided the following identifications are taken into account: \\begin{equation} \\label{ident} u_\\mu \\equiv  \\frac{\\psi_{,\\mu}}{\\sqrt{\\psi_{,\\mu}\\psi^{,\\mu}}}, \\quad p=\\rho\\equiv \\frac{1}{2}\\psi_{,\\mu}\\psi^{,\\mu}. \\end{equation} Therefore, in what follows we do not make a distinction between massless scalar field and ultra-hard perfect fluid, as far as vector $\\psi^{,\\mu}$ is timelike. For example, the problem of accretion of a perfect fluid can be viewed as accretion of a scalar field, with a specific form of the boundary condition at the infinity, $\\psi\\to \\dot{\\psi}_\\infty t$. Here $\\dot{\\psi}_\\infty$ is expressed in terms of the fluid density at the infinity, $\\rho_\\infty$, as follows, $\\dot{\\psi}_\\infty=\\left(2\\rho_\\infty\\right)^{1/2}$. However, if a solution for the scalar field contains spacelike $\\psi^{,\\mu}$ somewhere then such a solution can be identified with no perfect fluid analogue.  The paper is organized as follows. Accretion of ultra-hard perfect fluid (massless scalar field) onto the moving Kerr-Newman black hole is considered in Sec.~\\ref{sec fluid}. The basic equation of motion for the scalar potential is derived in Sec.~\\ref{sec general}; in Sec.~\\ref{sec moving} we solve this equation for the case of the moving non-extreme black hole; in Sec.~\\ref{sec extreme} we analyze in detail the case of the extreme black hole. In Sec.~\\ref{sec sf} we find a stationary solution for the scalar field in the Kerr-Newman singularity. In Sec.~\\ref{sec conclusion} we briefly describe the results of the work. ",
        "conclusions": "% \\label{sec conclusion} We found an analytic solution for the stationary accretion of ultra-hard perfect fluid onto the moving Kerr-Newman black hole. The presented solution is a generalization of Petrich, Teukolsky and Shapiro solution for the accretion onto the moving Kerr black hole. Our solution describes the velocity, and the density distribution of the perfect fluid with the equation of state $p(\\rho)=\\rho$ in space, when the Kerr-Newman black hole moves through the fluid. The solution is unequally determined in terms of the scalar potential, the exact solution for which is given by Eq.~(\\ref{PsiFinBlackHole}). One can think of the found solution as of one to the problem for a canonical massless scalar field in the metric of a moving Kerr-Newman black hole, since these two description are equivalent for the considered problem. In this case the solution for the scalar field coincides with the scalar potential for the fluid, Eq.~(\\ref{PsiFinBlackHole}).  Our results are only valid in the test fluid approximation, put differently, the back reaction of the accreting fluid is ignored. In the case of a non-extreme Kerr-Newman black hole one can always find such a range of parameters that the back reaction can indeed be safely ignored. While this is not the case for the extreme black hole. We showed that for the extreme Kerr-Newman black hole the test fluid approximation is always violated, since the energy density of ultra-hard fluid diverges at the event horizon. A similar behavior of the energy density for the perfect fluids with sound velocity $c_s\\geq1$ occurs for the extreme Reissner-Nordstr\\\"om black hole \\cite{bcde08a}. Apparently, if the black hole is almost extreme, $M^2-Q^2-a^2\\to +0$, the value of the energy density at the infinity must be tuned considerably, to avoid the violation of the test fluid approximation. Thus to solve correctly the problem of accretion in the case of an (almost) extreme black hole the back reaction of the fluid must be taken into account. This is in accord with \\cite{hod02,hod08}, where it was suggested that the back reaction is important for a description of a scalar particle absorption with a large angular momentum by a near extreme black hole. We will investigate this problem in a separate work \\cite{bcde08b}.  We also presented an analytic solution for the massless scalar field in the metric of the Kerr-Newman naked singularity, i.~e. when $M^2<Q^2+a^2$. This solution describes the scalar field distribution around the naked singularity, in general with non-zero influx. The found solution contains some pathologies which are the consequences of the causality violation in the vicinity of the Kerr-Newman space-time with a naked singularity. For the Kerr or the Reissner-Nordstr\\\"om naked singularity the presented solution is well behaved in the whole spacetime, except for the point of the physical singularity."
    },
    "0807/0807.0213_arXiv.txt": {
        "abstract": "This work deals with defect structures in models described by scalar fields. The investigations focus on generalized models, with the kinetic term modified to allow for a diversity of possibilities. We develop a new framework, in which we search for first-order differential equations which solve the equations of motion. The main issue concerns the introduction of a new function, which works like the superpotential usually considered in the standard situation. We investigate the problem in the general case, with an arbitrary number of fields, and we present several explicit examples in the case of a single real scalar field. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0807/0807.0814_arXiv.txt": {
        "abstract": " ",
        "introduction": "There are two approaches to study chemical evolution. The straightforward approach relies on the knowledge of all nucleosynthetic sources, their elemental yields, their occurrences in the interstellar medium (ISM), and their exchange of material with the ISM. Based on this knowledge, both the average and the scatter for the abundance of any element in the ISM can be calculated as functions of time. Unfortunately, our knowledge of nucleosynthetic sources is rather incomplete. So this forward approach cannot be applied to all elements. On the other hand, the elemental abundances in stars generally reflect the composition of the ISM from which they formed. Given a large collection of data on stellar abundances, it is possible to infer the characteristics of some nucleosynthetic sources from these data with the help of some basic understanding of stellar models. This reverse approach works best at early times when only a small number of potential sources could have contributed to the ISM. This paper presents a combination of forward and reverse approaches to study chemical evolution of heavy elements in the early Galaxy. The goal is to illustrate what can be learned about stellar sources from such studies.  Here ``early'' means the first Gyr or so of Galactic history. During this epoch, only short-lived massive stars could have had time to evolve and contribute their nucleosynthetic products to the ISM. As $\\sim 1/3$ of the solar Fe inventory was contributed by core-collapse supernovae (SNe, mostly Type II) associated with massive stars over a period of $\\sim 10$~Gyr, ``early times'' also correspond to $<1/30$ of the solar Fe abundance in the ISM. In the notation $[{\\rm Fe/H}]\\equiv\\log({\\rm Fe/H})-\\log({\\rm Fe/H})_\\odot$, such times correspond to metal-poor ISM with $[{\\rm Fe/H}]<-1.5$. Three groups of heavy elements will be discussed: (1) the elements from C through Zn with mass numbers $A<70$, (2) the elements from Sr through Ag with $A\\sim 88$--110, and (3) the elements with $A\\sim 130$ through the actinides attributed to the rapid neutron capture process (the $r$-process). The solar abundances of the elements with $A\\sim 88$--209 received important contributions from the slow neutron capture process (the $s$-process). However, as the major $s$-processing occurs in long-lived (with lifetimes of $>1$~Gyr) stars of a few $M_\\odot$, the $s$-process contributions to the ISM at $[{\\rm Fe/H}]<-1.5$ can be ignored to good approximation. Likewise, contributions to the Fe group elements from Type Ia SNe associated with white dwarfs evolved from long-lived low-mass stars in binaries can also be ignored at $[{\\rm Fe/H}]<-1.5$. ",
        "conclusions": "\\label{sec-con} As summarized in Table~\\ref{tab}, considerations of chemical evolution of heavy elements in the early Galaxy based on stellar observations and current understanding of stellar models show that (1) Fe core-collapse SNe from progenitors of $>11\\,M_\\odot$ are the major source for the elements from C through Zn in the early Galaxy; (2) O-Ne-Mg core-collapse SNe from progenitors of $\\sim 8$--$11\\,M_\\odot$ are the major source for the heavy $r$-nuclei with $A>130$; and (3) as both Fe and O-Ne-Mg core-collapse SNe can produce protoneutron stars that emit neutrinos to drive winds, both kinds of SNe can produce the elements from Sr through Ag by CPRs during the $\\alpha$-process. The neutrino-driven winds in some Fe core-collapse SNe (e.g., those producing heavier protoneutron stars) may even produce $r$-nuclei with $A\\sim 130$ (corresponding to the elements Te, I, and Xe), but stellar observations cannot help us to identify the sources for these nuclei as they are inaccessible to spectroscopic studies.  \\begin{table}[h] \\begin{center} \\caption{Stellar Sources for Heavy Elements}\\label{tab} \\begin{tabular}{lcc} \\hline & $>11\\,M_\\odot^a$ & $\\sim 8$--$11\\,M_\\odot^b$\\\\ \\hline C to Zn ($A<70$) & yes & no\\\\ Sr to Ag ($A\\sim 88$--110)& yes & yes\\\\ $r$-nuclei ($A\\sim 130$) & maybe & ?\\\\ $r$-nuclei ($A>130$) & no & yes\\\\ \\hline \\end{tabular} \\medskip\\\\ $^a$Progenitors of Fe core-collapse SNe.\\\\ $^b$Progenitors of O-Ne-Mg core-collapse SNe.\\\\ \\end{center} \\end{table}  \\begin{figure}[h] \\begin{center} \\includegraphics[scale=0.3, angle=270]{fig5a.eps} \\includegraphics[scale=0.3, angle=270]{fig5b.eps} \\caption{(a) Data on HD~122563 (squares with error bars linked by line segments; \\citealt{honda}) compared with  the solar ``$r$''-pattern translated to pass through the Eu data (curves labeled as such). Squares with downward arrows indicate upper limits. The abundance pattern of the heavy $r$-elements [Ba ($Z=56$) and above] in HD~122563 shown in (a) exhibits substantial differences from the solar ``$r$''-pattern, especially for Ce and Pr. In addition, HD~122563 has much larger proportions of the elements from Sr through Ag relative to the heavy $r$-nuclei as compared to the solar ``$r$''-pattern. (b) Comparison of the data on HD~122563 (squares with error bars linked by line segments; \\citealt{honda}) with those on CS~22892--052 (curves labeled as such; \\citealt{sneden}) normalized to the same $\\log\\epsilon({\\rm Y})$ as for HD~122563.} \\label{fig-hl} \\end{center} \\end{figure}  Based on the above attribution, one may try to identify two templates for the overall production of heavy elements by Fe and O-Ne-Mg core-collapse SNe, respectively. The squares with error bars in Fig.~\\ref{fig-hl}a show the data on the elements from Sr through Ir in HD~122563 \\citep{honda}. These data are linked by line segments to guide the eye. The solar ``$r$''-pattern translated to pass through the Eu data is shown as the curves labeled so in Fig.~\\ref{fig-hl}a. It can be seen that the data on the heavy $r$-elements Ce, Pr, Nd, and Sm lie significantly above the translated solar ``$r$''-pattern while those on the elements from Sr through Ag lie far above it. This suggests that the source responsible for the abundances in HD~122563 is not a major contributor to the heavy $r$-elements but mainly produces the elements from Sr through Ag. Consequently, the overall abundance pattern in this star including the elements from C through Zn may be taken as representative of the yields of Fe core-collapse SNe.  For convenience, this pattern will be referred to as the $L$-pattern (dominated by lighter nuclei). As Fe core-collapse SNe are the major source for Fe in the early Galaxy, the Fe abundance in a metal-poor star can be used along with the $L$-pattern to identify the absolute contributions from such SNe to all the elements in this star.  As shown in Figs.~\\ref{fig-decouple1} and \\ref{fig-decouple2}, a number of metal-poor stars exhibit a highly regular abundance pattern of the heavy $r$-elements that is essentially identical to the solar $r$-pattern. The data on one such star, CS~22892--052, which are shifted to pass through the Y data for HD~122563 for comparison with the latter star, are shown as the curves labeled so in Fig.~\\ref{fig-hl}b. The Fe abundance in CS~22892--052 indicates that Fe core-collapse SNe contributed very little to its abundances of Sr and heavier elements. Figure~\\ref{fig-hl}b shows that if the source responsible for the abundances of these elements in CS~22892--052 produces the elements from Sr through Ag at a similar level to an Fe core-collapse SN, then this source must be the major source for the heavy $r$-elements. Consequently, the abundance pattern of Sr and heavier elements in CS~22892--052 can be taken as representative of the yields of O-Ne-Mg core-collapse SNe. For convenience, this pattern will be referred to as the $H$-pattern (dominated by heavier nuclei). As O-Ne-Mg core-collapse SNe are the major source for the heavy $r$-element Eu, the Eu abundance in a metal-poor star can be used along with the $H$-pattern to identify the absolute contributions from such SNe to all the elements in this star.  Based on the above discussion, the abundance of element E in a metal-poor star can be determined as \\citep{qw07} \\begin{equation} \\left(\\frac{\\rm E}{\\rm H}\\right)= \\left(\\frac{\\rm E}{\\rm Fe}\\right)_L\\left(\\frac{\\rm Fe}{\\rm H}\\right)+ \\left(\\frac{\\rm E}{\\rm Eu}\\right)_H\\left(\\frac{\\rm Eu}{\\rm H}\\right), \\label{eq-hl} \\end{equation} where (E/Fe)$_L$ and (E/Eu)$_H$ are the yield ratios representing the $L$-pattern and the $H$-pattern taken from the data on HD~122563 \\citep{honda} and CS~22892--052 \\citep{sneden}, respectively. As an example, the abundances of Sr for a large number of metal-poor stars are calculated from Eq.~(\\ref{eq-hl}) using the observed Fe and Eu abundances and compared with the data in terms of $\\Delta\\log\\epsilon({\\rm Sr})\\equiv \\log\\epsilon({\\rm Sr})_{\\rm cal}-\\log\\epsilon({\\rm Sr})_{\\rm obs}$ in Fig.~\\ref{fig-cpr}a. It can be seen that the calculated Sr abundances are within 0.3 dex of the data for the majority of the stars. With the limited data available for the elements Nb, Mo, Ru, Rh, Pd, and Ag, even better agreement between the calculated and observed abundances is obtained for these elements as shown in Fig.~\\ref{fig-cpr}b. In conclusion, the two-component model based on Eq.~(\\ref{eq-hl}) appears to provide a very good description of the abundances in metal-poor stars. It remains to be seen whether self-consistent ab initio models of stellar evolution and nucleosynthesis can show that the $L$-pattern and the $H$-pattern are indeed characteristic of the yields of Fe and O-Ne-Mg core-collapse SNe, respectively.  \\begin{figure}[h] \\begin{center} \\includegraphics[scale=0.3, angle=270]{fig6a.eps} \\includegraphics[scale=0.3, angle=270]{fig6b.eps} \\caption{Comparison of the two-component model based on Eq.~(\\ref{eq-hl}) with the data for a large sample of stars. (a) The difference between the calculated abundance of Sr and the observed value is shown in terms of $\\Delta\\log\\epsilon({\\rm Sr})\\equiv \\log\\epsilon_{\\rm cal}({\\rm Sr})-\\log\\epsilon_{\\rm obs}({\\rm Sr})$ as a function of [Fe/H]. The calculation uses only the observed Fe and Eu abundances. The symbols represent the data sets as follows: squares for \\citet{johnson}, pluses for \\citet{aoki}, circles for \\citet{barklem}, crosses for \\citet{otsuki}, and asterisks for \\citet{hill} and \\citet{inese}. (b) Comparison between the model and data on Nb (squares), Mo (pluses), Ru (circles), Rh (crosses), Pd (asterisks), and Ag (diamonds). The limited data available are taken from \\citet{crawford,hill,cowan,johnson}; and \\citet{inese}. It can be seen from these examples that the elemental abundances in a metal-poor star are rather well estimated from the model using only the Fe and Eu abundances to identify the contributions from two different sources to the star.} \\label{fig-cpr} \\end{center} \\end{figure}"
    },
    "0807/0807.0657.txt": {
        "abstract": "We present X-ray and infrared observations of the X-ray source CXOGC~J174536.1-285638. Previous observations suggest that this source may be an accreting binary with a high-mass donor (HMXB) or a colliding wind binary (CWB). Based on the {\\it Chandra} and {\\it XMM-Newton} light curve, we have found an apparent $189 \\pm 6$ day periodicity with better than 99.997\\% confidence. We discuss several possible causes of this periodicity, including both orbital and superorbital interpretations. We explore in detail the possibility that the X-ray modulation is related to an orbital period and discuss the implications for two scenarios; one in which the variability is caused by obscuration of the X-ray source by a stellar wind, and the other in which it is caused by an eclipse of the X-ray source. We find that in the first case, CXOGC~J174536.1-285638 is consistent with both CWB and HMXB interpretations, but in the second, CXOGC~J174536.1-285638 is more likely a HMXB. ",
        "introduction": "{\\it Chandra} observations of the Galactic Center (GC) have revealed a large new population of low luminosity X-ray sources with $L_X[D/8kpc]^2 \\sim 10^{31} - 10^{35} erg~s^{-1}$ \\citep{mun03}. In addition, the Swift and INTEGRAL missions have recently revealed a new population of highly-absorbed X-ray sources, believed to be high-mass X-ray binaries \\citep[HMXBs; e.g.][]{beck05,bod07,neg07}. In 2005, we identified an infrared (IR) star as the first spectroscopically confirmed IR counterpart to the low luminosity {\\it Chandra} source CXOGC~J174536.1-285638 \\citep[hereafter Paper I]{mikles06}. Based on the X-ray and IR spectra and the X-ray to IR luminosity ratio, we showed that the source is most likely a massive star in a binary system. The source shows strong HeI (2.114-$\\mu$m), Brackett-$\\delta$ (1.945-$\\mu$m), and Brackett-$\\gamma$ (2.166-$\\mu$m) emission lines, typical of both accretion-powered binaries and CWBs. Additionally, we observe Brackett series, HeI, HeII, CIII, and NIII line emission. P~Cygni profiles are visible in several HeII lines, suggesting wind activity around a massive star. The X-ray spectrum of this source is particularly intriguing, having prominent Fe-XXV emission centered at 6.7~keV with an equivalent width of 2.2~keV. This is one of the highest equivalent width Fe-XXV lines ever seen (Paper I).  Our initial IR observations were aimed toward the discovery of a short ($<1$~d) period in the CXOGC~J174536.1-2856 binary. We use IR spectroscopy to search for variations in CXOGC~J174536.1-285638's IR emission features. We analyze {\\it Chandra} and {\\it XMM-Newton} archival data to search for X-ray variability over short and long baselines. From the combined X-ray light curve, we find a period of $189 \\pm 6$ days. We discuss CXOGC~J174536.1-2856's variability in the IR and X-ray, and examine the implications of a 189~d period for the nature of the source. In \\S 2, we summarize our IR and X-ray observations and analysis, detailing both our IR radial velocity study and X-ray period analysis. In \\S 3, we discuss the IR and X-ray variability in CXOGC~J174536.1-2856, specifically exploring an orbital period interpretation of the identified X-ray period. ",
        "conclusions": "We have searched for evidence of periodic variability in the IR spectra and long-term X-ray light-curve of the GC X-ray source CXOGC~J174536.1-285638. We find no evidence of IR variability on short ($<3~d$) timescales or between the 2005 and 2006 spectra. We compare the IR line ratios Br-$\\gamma$/HeI and Br-$\\gamma$/HeII in CXOGC~J174536.1-285638 to known HMXBs and CWBs and find the relative emission line strengths to be more consistent with an HMXB. We have identified an apparent $189 \\pm 6$ d period in the CXOGC~J174536.1-285638 X-ray light curve. We find no evidence of periodic X-ray variability at timescales less than 189~d. Using a Monte Carlo simulation, we test the significance of the 189~d period detection; despite our fairly sparse time sampling, we find this period is significant with a confidence level greater than 99.997\\%. We explore several interpretations of the X-ray modulation.  It is plausible, if the source is a HMXB, that the periodic modulation is superorbital in nature and related to a precessing accretion disk, in which case, further observations are required to determine the orbital period of the system and thus the nature of the system components. If the source is a HMXB and the 189~d period is superorbital, then we expect to find a shorter orbital period. This putative orbital periodicity is not necessarily observable in the IR as in this scenario, the IR emission is dominated by a single bright source. If the orbital period is detectable in the X-ray, targeted observations with a sensitive detector over a time interval of $\\sim 1-2$ weeks during the high flux stage are required to ensure sufficient counts to test for variability.  We also explore an orbital period interpretation and summarize scenarios for this in Table \\ref{tbl-7}. If the observed period is orbital in nature, and the X-ray modulation is caused by obscuration of the X-ray source due to a dense wind, then CXOGC~J174536.1-285638 is consistent with both CWB and HMXB interpretations. The further constraint of the X-ray luminosity is consistent with a massive ($M_{OB} > 80 M_\\odot$) donor with a neutron star companion. If X-ray modulation is caused by an eclipse, the mass ratio is low and CXOGC~J174536.1-285638 is more consistent with an HMXB interpretation. If the 189~d period is orbital, we may be able to identify the source nature by obtaining long term photometric observations in the IR. Also, targeted IR follow-up spectroscopy to cover multiple phases of the source period will allow us to search for a relationship between the X-ray and IR variability in this system. In the low flux phase, additional IR spectroscopic line features (e.g., absorption, P~Cygni variation) may become apparent that can help us discern the nature of the stellar components.  Recently, \\citet{hyodo08} reported the discovery of an early-type, Galactic Center source which appears to have many characteristics in common with CXOGC~J174536.1-285638. The source, CXOGC J174645.3-281546, has an unusually strong FeXXV line ($\\sim$1~keV), shows X-ray variability of a factor of $\\sim 2$ on a $\\sim$year timescale, and appears to have a high-mass star as its likely IR counterpart.  As in CXOGC~J174536.1-285638, its X-ray to IR luminosity ratio is $\\sim 10^{-4}$.  These intriguing similarities in X-ray spectral appearance, variability timescale, and luminosity lead us to suggest that it would be interesting in future observations to study this source in concert with CXOGC~J174536.1-285638.  Although there are only two sources with these properties known at present, it is possible that they could ultimately define a new (sub)class of early-type Galactic sources with strong FeXXV emission."
    },
    "0807/0807.2093_arXiv.txt": {
        "abstract": "The observation by the {\\it Swift} X-ray Telescope of the Fe K$\\alpha_1,\\alpha_2$ doublet during a large flare on the RS CVn binary system II~Peg represents one of only two firm detections to date of photospheric Fe~K$\\alpha$ from a star other than our Sun.  We present models of the Fe~K$\\alpha$ equivalent widths reported in the literature for the II~Peg observations and show that they are most probably due to fluorescence following inner shell photoionisation of quasi-neutral Fe by the flare X-rays. Our models constrain the maximum height of flare the to 0.15~R$_*$ assuming solar abundances for the photospheric material, and 0.1~R$_*$ and 0.06~R$_*$ assuming depleted photospheric abundances ([M/H]~=~-0.2 and [M/H]~=~-0.4, respectively). Accounting for an extended loop geometry has the effect of increasing the estimated flare heights by a factor of $\\sim$3. These predictions are consistent with those derived using results of flaring loop models, which are also used to estimate the flaring loop properties and energetics. From loop models we estimate a flare loop height of 0.13~R$_*$, plasma density of $\\sim 4\\times10^{12}$~cm$^{-3}$ and emitting volume of $\\sim 6\\times10^{30}$~cm$^{3}$. Our estimates for the flare dimensions and density allow us to estimate the conductive energy losses to $E_{cond} \\leq 2\\times10^{36}$~erg, consistent with upper limits previously obtained in the literature.  Finally, we estimate the average energy output of this large flare to be $\\sim10^{33}$~erg~sec$^{-1}$, or 1/10th of the stellar bolometric luminosity. ",
        "introduction": "\\label{s:intro}  \\begin{figure*} \\begin{minipage}{17.cm} \\epsscale{0.32}\\plotone{f1a.eps} \\epsscale{0.32}\\plotone{f1b.eps} \\epsscale{0.32}\\plotone{f1c.eps}  \\end{minipage} \\caption[]{Example Fe~K$\\alpha$ equivalent widths as a function of flare height and inclination angle for three different photospheric metallicities for the 1a segment of the flare, as defined by O07. The observed value and its error is indicated by the horizontal black dashed line and teh shaded area. Left: Models calculated at solar abundances (Grevesse \\& Sauval, 1998); Central: Models calculated with [M/H]~=~-0.2; Right: Models calculated with [M/H]~=~-0.4.} \\label{f:f1} \\end{figure*}   A fraction of the X-rays emitted during stellar flares or by hot stellar coronae are directed downwards towards the stellar photosphere. Here they interact with the photospheric gas and are reprocessed through scattering and photoionisation events. The downward cascade following inner shell photoionisation of quasi-neutral gas in the stellar photosphere produces characteristic fluorescent emission from astrophysically abundant species. While the reprocessed X-ray spectra reach the observer at low flux levels compared to that of the flares/coronae, fluorescent emission lines, can be significantly stronger and can be detected against the spectra of the flare/corona itself.  The intensity or equivalent width of the fluorescent lines detected at Earth depends uniquely on the details of the fluorescing spectrum, the photospheric abundance of the photoionised species and the geometry of the system (flare height and inclination angle; e.g. see fig. 2 of Testa et al., 2008), making them potentially powerful diagnostics. A number of theoretical studies exist which provide a framework for the interpretation of the strong iron fluorescent emission in the solar context (e.g. Tomblin, 1972; Bai, 1979), and more recently for arbitrarily photoionised slabs (Kallman et al., 2004) and stellar photospheres (Drake, Ercolano and Swartz, 2008, DES08). Drake \\& Ercolano (2008a,b) have also recently showed that fluorescent emission from neon and oxygen should also be detectable in the solar spectrum with current and future instrumentation. These may provide an independent measure of the photospheric solar oxygen and neon abundance, elements that are central to the solution of the ``solar oxygen crisis'' (e.g. Ayres et al.\\ 2006, Socas-Navarro 2007, Basu \\& Antia 2008).  Observationally, the $2s-1p$~6.4~KeV Fe K$\\alpha$ doublet from low ionisation stages of iron has often been detected in solar spectra (e.g. Neupert et al., 1967; Doschek et al., 1971; Fedelman et al., 1980; Tanaka et al., 1984; Parmar et al. 1984; Zarro et al., 1992), from disks around pre-main sequence stars (e.g. Tsujimoto et al., 2005; Favata et al., 2005), and from the photospheres of the single G-type giant HR~9024 (Testa et al., 2007) and of the RS CVn binary system II~Peg (Osten et al., 2007, hereafter O07).  The detection of the 6.4~keV fluorescent iron line in the {\\it Swift} X-ray Telescope spectrum of II~Peg taken during a large flare (O07) is particularly interesting as it represents the first of only two detections (with HR~9024, Testa et al., 2007) of photospheric fluorescence emission in stars other than the Sun. O07, however, assigned the excitation mechanism to electron impact ionization of photospheric Fe, rather than photoionisation. In this paper we show by means of Monte Carlo fluorescence modeling that the Fe~K$\\alpha$ equivalent widths reported by O07 are perfectly consistent with photoionisation induced fluorescence. We also show that our predictions of the flaring loop scale height agree with those from flaring loop models (Reale, 2007) and discuss other implications on flaring plasma density, volume and energetics.  In Section~2 we briefly describe our Monte Carlo model, and in Section~3 we present and discuss our results. Section~4 deals with the flaring loop properties and energetics from hydrodynamic models. Our final conclusions are given in Section~5. ",
        "conclusions": "We have presented a set of Monte Carlo calculations able to reproduce the Fe~K$\\alpha$ equivalent widths from the {\\it Swift} X-ray Telescope observations of II~Peg during a 'superflare' (Osten et al., 2007, O07). Our models show that the data are consistent with the Fe~K$\\alpha$ emission being produced by fluorescence following K-shell photoionisation of quasi-neutral iron in the stellar photosphere. This contrasts with the interpretation of O07, who favoured collisional ionisation by non-thermal electrons to produce the Fe~K$\\alpha$ emission. They argued that the normal incidence flare X-ray penetration depth required in the photosphere to obtain the observed equivalent widths is similar to the $\\tau=1$ Compton scattering depth, and consequently any fluorescent photons produced would not easily escape.  However, this assessment was based on a simple analytical formula appropriate for optically thin cases in which only a small fraction of the incident X-ray flux undergoes photoabsorption or scattering (e.g. Liedahl, 1999; Krolik \\& Kallman, 1987).  The semi-infinite photospheric case lies outside the range of applicability of this formula.  Furthermore, incident angles on the photosphere range from $\\sim 0$--$90^\\circ$ and path lengths for escape can therefore be much smaller than gas penetration depths (by a factor equal to the inverse cosine of these angles).  Our stochastic treatment includes Compton scattering and shows that the observed equivalent widths can be produced by fluorescence.  We also note that the impact excitation mechanism is a very low efficiency process and requires a large amount of energy in the form of accelerated electrons.  This was already noted by, e.g., Parmar (1984) in reference to solar Fe fluorescence and by Ballantyne (2003), and was discussed in reference to a large flare on II~Peg by Testa et al. (2008).  The derived flare loop heights are $h$~=~0.15~R$_*$, assuming solar photospheric abundances, $h$~=~0.1~R$_*$ assuming photospheric abundances depleted by [M/H]~=~-0.2 and $h$~=~0.06~R$_*$ assuming [M/H]~=~-0.4. These value are in good agreement with the predictions from our alternative analysis based on hydrodynamic models (e.g. Reale, 2007) which yield $h \\approx 0.13$~R$_*$. We estimate flaring loop properties and energetics using scaling laws based on the hydrodynamic loop modeling of Reale (2007) and obtain a plasma density of $4\\times10^{12}$~cm$^{-3}$ and volume of $6\\times10^{30}$~cm$^{3}$. Using the derived values for the flare dimensions and densities we obtain a more stringent upper limit for the conductive energy loss, $E_{cond}$, than previously possible, setting $E_{cond} \\leq 2\\times10^{36}$~erg. Considering the duration of the flare and the radiative and conductive losses, the energy output of the flare is estimated at $\\sim$10$^{33}$~erg sec$^{-1}$, which is approximately one tenth of the stellar bolometric luminosity.  While we cannot rule out a contribution to the observed Fe~K$\\alpha$ flux from electron impact with nonthermal electrons, as proposed by O07, such contribution is not necessary to explain the available data."
    },
    "0807/0807.3007_arXiv.txt": {
        "abstract": "We study the properties of powerful X-ray flares from 161 pre-main sequence (PMS) stars observed with the Chandra X-ray Observatory in the Orion Nebula region.  Relationships between flare properties, protoplanetary disks and accretion are examined in detail to test models of star-disk interactions at the inner edge of the accretion disks.  Previous studies had found no differences in flaring between diskfree and accreting systems other than a small overall diminution of X-ray luminosity in accreting systems.   The most important finding is that X-ray coronal extents in fast-rotating diskfree stars can significantly exceed the Keplerian corotation radius, whereas X-ray loop sizes in disky and accreting systems do not exceed the corotation radius.  This is consistent with models of star-disk magnetic interaction where the inner disk truncates and confines the PMS stellar magnetosphere.   We also find two differences between flares in accreting and diskfree PMS stars.   First, a subclass of super-hot flares with peak plasma temperatures exceeding 100~MK are preferentially present in accreting systems.   Second, we tentatively find that accreting stars produce flares with shorter durations.  Both results may be consequences of the distortion and destabilization of the stellar magnetosphere by the interacting disk.  Finally, we find no evidence that any flare types, even slow-rise flat-top flares are produced in star-disk magnetic loops.  All are consistent with enhanced solar long-duration events with both footprints anchored in the stellar surface. ",
        "introduction": "}    A broad consensus has emerged in the past decade concerning the structure and astrophysics of pre-main sequence (PMS) stars.  The angular momentum inherited from the collapsing interstellar material is mostly distributed into a rapidly rotating protostar, a Keplerian protoplanetary disk, and frequently the orbital motion of multiple stellar components. The young star interacts with its disk in a complex fashion with accretion and ejection of collimated outflow.  It is widely believed that strong magnetic fields generated within the young star mediate this star-disk interaction, truncating the disk and funneling accretion onto limited portions of the stellar surface.  These issues are developed in a variety of reviews \\citep[e.g.,][]{Hartmann98, Bouvier07a, Bouvier07b}.    It has been difficult, however, to elucidate in detail the nature of the PMS stellar magnetospheres and their role in facilitating star-disk interactions and accretion.  Well-developed analytical theory assumes that a dipolar field extends to several stellar radii where the Keplerian orbits and stellar fields are in co-rotation, interacts with the disk magnetic field,  and guides both accretion onto the surface and into a high-velocity collimated outflow \\citep[e.g.,][]{Shu94, Lovelace95}.  However, considerable evidence has emerged that the magnetic fields on the surfaces of PMS stars are concentrated in complex multipolar active regions similar to those on the Sun, rather than a simple dipolar.  This emerges from photometric, Doppler imaging, circular polarization and spectroscopy of Zeeman-sensitive lines from PMS photospheres \\citep[e.g.,][]{Daou06,Yang07, Johns-Krull07, Donati07}.  Recently, theoretical efforts have begun to calculate the resulting complex PMS magnetosphere and accretion process \\citep{Jardine06, Gregory06a, Long07}.    X-ray studies are a potentially useful tool for investigating these issues.  It has long been known that late-type stars exhibit their highest levels of X-ray emission, arising mostly from violent magnetic reconnection events, during their PMS phase \\citep[e.g., reviews by][]{Feigelson99, Feigelson07}.  Recent X-ray surveys of nearby PMS stellar populations give detailed insights into PMS magnetic flaring; these include the Chandra Orion Ultradeep Project \\citep[COUP,][]{Getman05} and the XMM/Newton Extended Survey of Taurus \\citep[XEST,][]{Gudel07}. Both astrophysical study of the properties of individual flares, and statistical study of many flares, from the COUP and XEST observations reveal that most events are similar to solar magnetic flaring but with X-ray luminosities enhanced $10^3-10^5$ fold in intensity \\citep[e.g.][]{Favata05, Wolk05, Flaccomio05, Stassun06, Maggio07, Caramazza07, Arzner07, Stelzer07, Franciosini07}.    The decay of PMS X-ray flares on timescales of $10^3-10^5$~s have proved particularly amenable to astrophysical modeling.   One favored model, developed by \\citet{Reale97} and extensively applied to solar and stellar X-ray flares, considers the X-rays produced by thermal plasma at $T \\sim 10^7$~K confined in a cylindrical loop cooling by radiation and conduction but subject to possible reheating by later magnetic reconnection events. Using such models, inferences can be made concerning magnetic fields responsible for the flare, including loop length and magnetic field strength.  For the most powerful PMS flares, inferred loop lengths reach $4-20$~R$_\\star$, comparable to the expected inner edges of protoplanetary disks \\citep{Favata05}\\footnote{Recall, however, that the vast majority of weaker PMS flares arise in magnetic loops no larger than the star \\citep[e.g.][]{Imanishi03,Wolk05, Franciosini07}.}. Thus, X-ray flaring potentially probes the region of star-disk interaction. Links are also emerging between other features of PMS X-ray emission, such as the statistical suppression of flares in accreting systems and the rotational modulation of X-rays, and models of multipolar PMS magnetospheres \\citep{Gregory06, Gregory07}.    The present study extends the analysis of powerful flares from the COUP made by \\citet{Favata05}.  They modeled 32 flares using a traditional technique of time-resolved spectroscopy. In \\citet[][Paper I]{Getman08} and \\citet{Getman06}, we introduce a more sensitive data analysis method based on adaptively smoothed median energies which permits modeling of 216 COUP flares.  This is a sufficiently large sample that permits new investigation of possible links between flaring, PMS magnetospheres and protoplanetary disks. Here, we present observational evidence that X-ray emitting coronal structures are in fact truncated by inner disks around the Keplerian corotation radius, just as predicted by PMS theoretical models outlined above.  We also report a variety of other results, both positive and negative, linking magnetic flare properties to the presence or absence of disks.    The paper is organized as follows.  Preliminary work and recovery of some established results are given in \\S \\ref{disk_section}. The main results of our study appear in the following three sections: the absence of strong links between some flare properties and disks (\\S \\ref{flare_prop_section}); clear evidence that PMS magnetospheres do not extend beyond disk inner edges (\\S \\ref{truncation_section}); and possible relations between super-hot flares,  accretion and non-dipolar magnetic fields (\\S \\ref{superhot_section} and \\ref{magnetic_fields.section}). Discussion follows in section \\S \\ref{discussion_section}. Readers are encouraged to consult Paper I for details on the selection, modeling, and properties of the 216 COUP flares discussed here. ",
        "conclusions": "}    We examine here empirical relationships between magnetic reconnection flares, protoplanetary disks and accretion. The current work provides detailed flare modeling of a much larger dataset than previously available: we consider 216 bright X-ray flares from 161 PMS stars observed during the 13-day COUP exposure of the Orion Nebula.  Our sample is larger because, as described in Paper I, we use data analysis techniques that permit modeling of fainter flares than feasible with traditional parametric modeling of variable $Chandra$ ACIS spectra.  We thus have opportunity to uncover more subtle relationships between flaring, disks and accretion than were previously possible.  The main results of our study are as follows:    1. Perhaps with the exception of the $Chandra$ study of NGC~2264, where CTTS are seen to be more variable than WTTS \\citep{Flaccomio06}, past studies using smaller samples have not found differences in flare statistics or properties as a function of PMS evolutionary state or mass: accreting Class II systems appeared to flare like non-accreting Class III systems \\citep[e.g.][]{Stelzer00,  Wolk05, Favata05, Stelzer07}. We confirm here that Class II and Class III systems produce X-ray flares with indistinguishable morphologies. Typical fast-rise slow-decay flares, composite step and double flares, and the unusual slow-rise flat-top flares are present in both classes. We have no clear evidence for distinctive flare types emerging from star-disk magnetic field lines. Accretion variations producing optical band variations are not associated with X-ray flares \\citep{Stassun06}. Accretion and magnetic flaring thus appear to be unrelated, even though the theory of PMS accretion involves magnetic truncation of disks and magnetic funneling of disk material onto the stellar surface.  The COUP PMS flares are consistent with solar-type magnetic structures with both footpoints anchored in the stellar surface.  Many may represent analogies of a much less powerful class of solar flares known as long duration events which produce giant coronal arches and X-ray streamers.    2. A distinct difference is found in the distribution of loop sizes. We find that X-ray coronal extents in fast rotating Class~III stars sometimes exceeds the Keplerian corotation radius, whereas X-ray magnetospheres in Class~II stars appear to be truncated at the inner edge of accreting protoplanetary disks. This directly supports theoretical models of magnetically mediated accretion, magnetic star-disk rotational coupling, and  disk confinement of PMS magnetospheres.    3. A related, but statistically less secure, result is that flares from highly accreting Class~II stars have somewhat shorter durations and weaker total X-ray energies than Class~III flares. This might reflect the destabilization of magnetic arcades in accreting systems, and could account for the reduction of time-averaged X-ray luminosities in Class II compared to Class III populations noted in previous studies.    4. A subclass of super-hot flares with peak plasma temperatures greater than 100~MK is noted for the first time.  These are inconsistent with formation in normal dipolar magnetic loops attached to the stellar surface, and appear preferentially in accreting Class~II systems.  They may reflect compression and distortion of the large scale magnetospheric topology by star-disk magnetic interactions, as predicted by recent theoretical calculations."
    },
    "0807/0807.1658_arXiv.txt": {
        "abstract": "In May 2008 the soft gamma-ray repeater \\src\\ resumed its bursting activity after nearly a decade of quiescence. After detection of a bright burst, \\swift\\ pointed its X-ray telescope in the direction of the source in less than five hours and followed it for over five weeks. In this paper we present an analysis of the data from these \\swift\\ observations and an \\xmm\\ one performed when \\src\\ was still in a quiescent state. The analysis of the bursts detected with \\swift/BAT shows that their temporal and spectral properties are similar to those found in previous observations of \\src\\ and other soft gamma-ray repeaters. The maximum peak luminosity of the bursts was $\\sim$$2\\times10^{41}$ \\lum. Our data show that the outburst was accompanied by a fast flux enhancement and by a hardening of the spectrum with respect to the persistent emission. ",
        "introduction": "The soft gamma-ray repeater (SGR) \\src\\ is likely to host a `magnetar', i.e. an isolated neutron star believed to have an extremely strong magnetic field ($B\\sim 10^{14}$--$10^{15}$ G) powering their bright X-ray emission and peculiar bursting activity \\citep[e.g.][]{mereghetti08}. Several magnetars, including \\src, have been observed to emit short bursts ($<$1 s) in the hard X/soft gamma-ray band, with characteristic peak luminosities of the order of $\\sim$$10^{39}$--$10^{41}$ \\lum. Besides short bursts, SGRs are known to emit intermediate and giant flares, with typical durations of 0.5--500 s, during which luminosities up to $\\sim$$10^{47}$ \\lum\\ can be achieved.\\\\ \\indent \\src\\ was discovered in 1998 by the \\emph{Compton Gamma Ray Observatory} because of the intense bursts it emitted at the time \\citep{kouveliotou98_1627}. These bursts, more than a hundred in six weeks, were also observed by other X-ray satellites \\citep{hurley99_1627,woods99,feroci98,smith99,mazets99}. Soon after the discovery of the bursts, the persistent X-ray emission of this SGR was detected by \\sax\\ at a luminosity of $\\sim$$10^{35}$ \\lum\\ \\citep[assuming a distance to the source of 11 kpc;][]{corbel99}. The quiescent spectrum was well modelled by an absorbed power law ($N_{\\rm H}\\approx8\\times10^{22}$ cm$^{-2}$ and photon index $\\Gamma\\approx2.5$; \\citealt{woods99}). No further bursting activity has been reported since then, but several X-ray satellites observed the X-ray persistent counterpart of \\src\\ in the past ten years \\citep{kouveliotou03,mereghetti06}. Since its discovery, this persistent emission showed a slow luminosity decay, from about $10^{35}$ to $10^{33}$ \\lum, the lowest value ever observed for an SGR, and a spectral softening from $\\Gamma\\approx2$ to 4 \\citep{kouveliotou03,mereghetti06}. The post-burst cooling trend seen in X-rays is peculiar among SGRs; rather it resembles the behaviour of transient anomalous X-ray pulsars (AXPs), a sub-class of the magnetar family.\\\\ \\indent Here we report on the last observation of \\src\\ performed at the end of the ten year long stretch of quiescence and on the burst re-activation of the source on 28 May 2008 \\citep{palmer08,golenetskii08,woods08atel1549}. ",
        "conclusions": "\\label{disc} The recent, spectacular re-activation of \\src, following a quiescent interval of nearly a decade, triggered the BAT instrument on board \\swift\\ on 28 May 2008. Tens of bursts were observed, with fluxes exceeding the underlying continuum by a factor $>$$10^5$. The bursts achieved a maximum luminosity of $\\sim$$10^{41}$ \\lum\\ and had a duration of $<$0.5 s, typical of the bursts detected in SGRs. Thanks to the rapid response of \\swift, the source was repeatedly observed with XRT in the days following the `burst forest' emission, leading to the earliest post-burst observations ever obtained for this SGR. In fact, at the time of the previous active period, the persistent emission was observed only one month after the first burst detection \\citep{woods99}.\\\\ \\indent With respect to the last pre-burst \\xmm\\ observation, the source was detected in May--June 2008 at a much larger flux level (see Figure~\\ref{history}) and with a considerably harder spectrum. A serendipitous \\cxo\\ observation performed only 20 hours before the detection of the bursting activity provides a 3\\,$\\sigma$ upper limit on the absorbed flux of $5 \\times 10^{-12}$ \\flux\\ \\citep[0.5--10 keV band;][]{tiengo08atel1559}, showing that most of the flux enhancement occurred in less than a day. These facts indicate a significant phase transition marked by the burst activation. The correlated spectral hardening/flux increase is in line with what is observed in the long term evolution of other magnetars \\citep[e.g.][]{mereghetti08}, and expected in models in which the non thermal X-ray emission is due to resonant up-scattering by magnetospheric currents \\citep{tlk02}. \\\\ \\indent The early flux decay of the source is shown in Figure~\\ref{comp}, where, for comparison, we also plot the flux of the decay that followed the past bursting activity of \\src\\ and those of two AXPs, CXOU\\,J164710.2--455216 \\citep{icd07} and 1E\\,2259+586 \\citep{woods04}. The \\src\\ data taken after more than two days from the May 2008 trigger are well fit by a power-law decay (index $\\sim-0.2$), but the XRT points at earlier times shows a marked excess over this trend, indicating a very steep initial decay. This behaviour closely resembles that of 1E\\,2259+586: in that case, after the June 2002 bursts active phase the source flux showed a double component decay, with a steep component that decayed rapidly during the first $\\sim$2 days, followed by a slower year-long decay phase \\citep{woods04}. Interestingly, the phase of steep flux decay (and harder spectra) was associated with a long-lasting period of bursting activity. This is consistent with what was observed in CXOU\\,J164710.2--455216 and \\src: in the former case, in which no steep/prompt decay was observed, the bursting activity was already over at the time of the first observation of the persistent flux. In the latter case, after the activation of June 1998 only a shallow decay was monitored $>$60 days from the first bursts. The light-curve of the recent \\src\\ decay reveals both, the steep and shallow phase, and the last burst was detected by Konus/\\emph{Wind} between the first and the second XRT pointing \\citep{golenetskii08}. The new observations support the presence of two distinct time-scales (a short, $\\sim$1 day, and a longer one, $\\sim$month) in the flux decay following the outburst of a magnetar. In this respect, we notice that while a steep decay seems to point toward a magnetospheric effect (for instance following current dissipation or other forms of activity), the decay index of the shallow phase ranges from about $-0.6$ (in the case of \\src, 1998) to $-0.2$ (for 1E\\,2259+586 and \\src, 2008) and is roughly compatible with crustal cooling (considering the uncertainties in the theoretical models, see \\citealt{eichler06} and references therein). The energetics and relative importance of the steep/shallow decay phases, nevertheless, greatly vary from source to source. Indeed, in some cases this requires a powering mechanism for the tail emission much more energetic than the bursts energy deposition. This might be associated with magnetospheric current dissipation or crustal cooling following an impulsive heat deposition. \\begin{figure} \\resizebox{\\hsize}{!}{\\includegraphics[angle=-90]{fig4.eps}} \\caption{\\label{comp} Comparison among flux decays of \\src\\ (this work, \\citealt{woods08atel1564}, \\citealt{mereghetti06}), CXOU\\,J164710.2--455216 (\\citealt[][]{icd07}; filled circles)  and 1E\\,2259+586 (\\citealt[][]{woods04}; filled triangles) following a period of bursts emission and/or glitch. In the case of \\src\\ we report the available data for both the 2008 (filled squares) and 1998 (open squares) activation periods. The solid lines represent the power-law best-fits (see Section \\ref{disc} for more details).} \\end{figure}\\\\ \\indent For several bursts we had enough counts to perform a spectral analysis. The parameter values derived from the 2BB fits are in agreement with the results of a comprehensive analysis of a `burst forest' emitted by SGR\\,1900+14 \\citep{israel08short}, which showed that the bursts populate almost homogeneously all temperatures between $\\sim$2 and $\\sim$12 keV, with a bimodal distribution behaviour and a sharp edge in the \\mbox{$kT$--$R^2$} plane. This supports the idea of two distinct emitting regions, a cold and larger one and a hot and smaller one which in turn may be associated to the escaping regions of two populations of photons, from the O- and E- polarisation modes. Interestingly, the bright bursts detected from \\src\\ lie within the cloud of the SGR\\,1900+14 bursts (Figure~\\ref{2bb}) and the luminosities of the two blackbody components are in agreement with the relation shown in Figure~6 of \\citet{israel08short}, suggesting that short bursts form a continuum in terms of spectral properties, duration and fluence. Finally, we notice that in the scenario proposed by \\citet{israel08short}, the luminosity of the hot blackbody is not expected to exceed the magnetic Eddington luminosity \\citep{thompson95}. In the case of \\src, where the maximum luminosity observed by the BAT for the hot blackbody is $\\sim$$10^{41}$ \\lum, this translates into a lower limit for the magnetic field of $B>1.8\\times10^{14}$ G."
    },
    "0807/0807.1308_arXiv.txt": {
        "abstract": "We have carried out a survey of X-ray emission from stars with giant planets, combining both archival and targeted surveys. Over 230 stars have been currently identified as possessing planets, and roughly a third of these have been detected in X-rays. We carry out detailed statistical analysis on a volume limited sample of main sequence star systems with detected planets, comparing subsamples of stars that have close-in planets with stars that have more distant planets.  This analysis reveals strong evidence that stars with close-in giant planets are on average more X-ray active by a factor $\\approx 4$ than those with planets that are more distant. This result persists for various sample selections. We find that even after accounting for observational sample bias, a significant residual difference still remains.  This observational result is consistent with the hypothesis that giant planets in close proximity to the primary stars influences the stellar magnetic activity. ",
        "introduction": "\\label{s:intro}  After centuries of ignorance on planetary systems beyond our own, we are now in an era when stars are being routinely identified as possessing planets. Since the discovery by Mayor \\& Queloz (1995) of a Jupiter-type planet (mass $0.44~M_J$) orbiting close in to the G2 star 51\\,Peg ($0.05$~AU), almost three hundred new Extrasolar Giant Planets (EGPs) have been found.\\footnote{ The current state of this rapidly advancing field is summarized by the International Astronomical Union's Working Group on Extrasolar Planets at {\\tt http://www.ciw.edu/boss/IAU/div3/wgesp/} and at the Exoplanet Encyclopedia at {\\tt http://exoplanet.eu/} }  One of the surprising results so far have been the detection of planets with orbital radii, $\\scrA$, much smaller than seen in the Solar System, with values as low as $\\scrA\\approx$0.03~AU.  At such small separations, it is likely that these giant planets will have measurable tidal ($\\propto \\scrA^{-3}$) or magnetic ($\\propto \\scrA^{-2}$ for large $\\scrA$) effects on the primary stars (Cuntz, Saar, \\& Musielak 2000; Saar, Cuntz, Shkolnik 2004).  Both these effects increase non-linearly for small values of $\\scrA$, leading to potentially dramatic disruptions of the stellar environment for stars with close-in giant planets. How these disruptions affect coronal activity remains an open question.  It is possible that activity will be enhanced because of the magnetic interactions between star and planet.  Tidal bulges will also have an effect on the stability and energetics of the chromosphere. Hitherto, analyses of variables that control stellar activity levels for single stars have focused on rotation rate, mass, age, evolutionary state, and to some extent metallicity, as factors determining coronal activity.  Now we are presented with another possibility, viz., the existence of close-in EGPs may also be a controlling parameter.  It is well known that binary stars with close stellar companions are generally more active than single stars (see e.g., Pye et al.\\ 1994).  Stars with close-in EGPs may be of the same type, though it is unclear whether similar mechanisms of activity enhancements hold at both extremes of this family.  A preliminary search for planet-induced stellar activity enhancement was carried out by Bastian, Dulk \\& Leblanc (2000) in the radio and by Saar \\& Cuntz (2001) in the optical.  But due to insufficient sensitivity, neither study succeeded in uncovering clear evidence for this phenomenon.  However, a sensitive search for activity (Shkolnik et al.\\ 2003; see also Cuntz \\& Shkolnik 2002) detected enhanced stellar emission in the Ca~II~HK lines from the chromosphere of HD\\,179949, in phase with the orbit of its close-in giant planet ($P_{\\rm orb} = 3.09$ days, $M \\sin i = 0.93 M_J$, $\\scrA = 0.045$ AU). The enhancement is clearly planet-related because it is not in phase with the stellar rotation rate $P_{\\rm rot} = 8 - 9$ days (Tinney et al.\\ 2001).  Because the Ca~II~HK enhancement  is seen only on the hemisphere facing the planet, a magnetic interaction is preferred over a tidal one.  Since a unipolar inductor model (viz., the Io--Jupiter interaction; Zarka et al.\\ 2001) does not fit the data as well (Saar et al.\\ 2004), this appears to be the first observational evidence for an exoplanetary magnetosphere.  These observations also provide an opportunity to probe the high-energy particle environment near EGPs since the interaction strength depends on the magnitudes of the stellar and planetary magnetic fields $B_*$ and $B_P$ (Saar et al.\\ 2004).  Note however that these emissions appear to be not phase locked, as is expected if the stellar magnetic fields are varying (McIvor, Jardine, \\& Holzwarth 2006; Cranmer \\& Saar 2006).  The peak strength of the HK emission flux enhancement found is $\\approx$4\\% (Shkolnik et al.\\ 2004).  Such low-level enhancements are difficult to detect and study in any detail.  However, if the emission scales like other manifestations of stellar activity, coronal enhancements should be much larger: e.g., Ayres et al.\\ (1995) found surface fluxes $F_X/F_{bol} \\propto (F_{\\rm C\\,IV}/F_{bol})^3$. Follow-up observations of HD\\,179949 in X-rays (\\chandra/ACIS; Saar et al.\\ 2006) show significant spectral and temporal variability phased to the planetary orbit, but there is some residual ambiguity due to the poorly constrained stellar rotational period ($P_{\\rm rot} \\approx 7-10$ d) -- some fraction of the variation may be due to changes in the underlying stellar activity.  This picture has been reinforced by recent observations of similar activity on $\\upsilon$~And and a confirmation of the activity synchronization on HD\\,179949 (Shkolnik et al.\\ 2008).  There are also indications of the influence of a planet on stellar activity in $\\tau$~Boo and HD\\,189733, though the results are inconclusive. The newer observations also show that the possible influence of the giant planet on stellar activity is complex, intermittent, and prone to phase shifts (Saar et al.\\ 2006, Shkolnik et al.\\ 2008, Lanza 2008).  In \\S\\ref{s:data}, we detail the X-ray data available for stars with giant plants.  In \\S\\ref{s:analysis}, we carry out statistical searches for trends in the data as a function of $\\scrA$, and show evidence of a significant deviation between extremal subsamples. In these analyses, we take into account the large numbers of censored data (that is, undetected stars) and properly include their effect. We discuss these results in \\S\\ref{s:discuss} with careful attention to the biases in the sample.  We summarize our results in \\S\\ref{s:summary}. ",
        "conclusions": "\\label{s:discuss}  \\subsection{Sample Bias \\label{s:bias}}  The sample of candidate stars for which planet searches are conducted is subject to some subtle biases.  Some of these biases have the effect of masking the signature of planet induced activity and are difficult to quantify since the planet detection processes are numerous, the programs are still incomplete, and the rates of false positives and false negatives are unknown.  We may however determine the approximate extent of these biases by studying the ensemble properties of the sample of stars with detected planets. Here we describe these biases, and consider their effect on our ability to detect intrinsic trends in X-ray activity. Furthermore, because the set of EGPs is dominated by those identified with the radial velocity method, we shall limit our sample to those stars which have been verified to have planetary systems by this method, and will thus consider only those biases introduced by that method.  \\begin{enumerate}  \\item {\\bf Spectral Homogeneity} Most of the stars for which planets are searched for are solar like. This is advantageous to our study since the stars considered here are relatively homogeneous; therefore effects that arise due to changes in the convective turnover timescales leading to changes in the nature of the magnetic dynamo at the high- and low-mass ends of the coronal sequence, or due to the changing evolutionary states of the systems may generally be ignored. We further homogenize the sample by limiting it to main sequence stars within 60~pc (additional types of filtering to further homogenize the samples has no effect on our results; see \\S\\ref{s:activity}).  Because the sample of stars is homogeneous, we do not expect any correlations between the stellar radii $R_*$ and $\\scrA$ to be present, and indeed we find that the data are consistent with these expectations (Table~\\ref{t:acorrs}).  We have also carried out a full analysis of the extremal subsamples for various subsets of the full dataset (see Figure~\\ref{f:lohiratmany}) and show that our results are robust to such selections. We thus conclude that in our sample there is no bias present due to stellar type or size effects.  \\item {\\bf Stellar Distance} For a given planetary mass, the detectability of planets decreases as the distance of the planet from the star ($\\scrA$) increases, since the radial velocity amplitude decreases.  Since detectability in general decreases with distance to the star ($d_*$), we expect to find more stars with small $\\scrA$ at larger $d_*$.  In the sample of stars with detected planets, we therefore expect that $\\scrA$ and $d_*$ are anti-correlated.  Indeed, we find a slight negative correlation between $\\scrA$ and $d_*$, though it is weak (see Table~\\ref{t:acorrs}). This bias exists simply as a result of the limitations of measurement statistics and is independent of any stellar activity effects.  Thus, if we assume that X-ray properties of stars are independent of the distance limit, the only effect on the X-ray data will be to have a larger fraction of higher upper limits at smaller $\\scrA$ (due to the larger numbers of more distant stars).\\footnote{ In our correlation analyses we have tested for the effect that these higher upper limits may have by using uniform distributions (bounded above by the value of the upper limit) in both log space and normal space to describe the censored data and find no qualitative difference between the two cases. } Thus, stellar distance should have no effect on the ensemble properties of the {\\sl X-ray} luminosities of stars from such a sample.  In any case, we carry out most of our analysis with the volume limited sample, which is explicitly designed to remove any such effect.  \\item {\\bf Selecting for reduced activity} The process of planet detection via the radial velocity method (RV; e.g., Butler, et al.\\ 1996) is limited by the amount of intrinsic RV jitter that may be caused by stellar activity (Saar \\& Donahue 1997, Saar et al.\\ 1998; see also Baluev 2008), and hence planets are preferentially detected when they are close-in, massive, and the primary star is not active.  Because radial velocity signatures are hard to detect around active stars, there is a tendency to select candidate stars for planet search {\\sl against} activity. Thus, {\\sl a priori} there is an expectation that the sample of stars with detected planets would have lower activity levels than field stars of the same type.  But as shown in \\S\\ref{s:samplestat}, the X-ray emission from these stars is by and large consistent with field star emission levels.  In any case, in order to avoid introducing accidental biases by comparing the sample of stars with EGPs to those without detected EGPs (note that a lack of detection does not preclude the existence of planets around the star -- even if the star is part of the planet search program, there may be a planet that is less massive or more distant than the sensitivity that has been reached), we have chosen to compare the extrema of a single distribution, viz., the sample of stars known to possess giant planets.  Thus any global selection effect that may exist in the sample towards lesser activity will apply equally to both subsamples and its effects are irrelevant for this study.  Also note that in some types of planet detection methods, persistent stellar activity fixed to an active longitude can mimic the signature of a planet (Lanza et al.\\ 2008). This results in false detections of planets at the stellar orbital period.  However, this has an effect on photocentric methods and have no effect on radial velocity or transit methods currently in use.  \\item {\\bf Intra-sample trend in inhibition of activity} If stellar activity inhibits planet detection, then it will inhibit it selectively.  Planets at large $\\scrA$ are preferentially detected around stars with weak activity, since high activity may mask the RV signal of the distant planet.  Thus, in a sample of stars selected based on the existence of planets, those stars with distant planets are presumptively less active.  This bias thus produces the same signature in activity trends that we search for (see \\S\\ref{s:farout}), and thus interferes with our study.  How can this bias be quantified?  Consider that the cause of the bias is the excess jitter in velocity induced by magnetic activity, which is related to the energy deposited in the corona, which in turn is tracked by the ratio\\footnote{ $\\frac{L_X}{L_{bol}}$ is preferred for this bias measurement because it is surface area independent. } of X-ray to bolometric luminosity $\\frac{L_X}{L_{Bol}}$.  This jitter serves to mask the RV signal, which is due to wobble from gravitational effects and is therefore inversely proportional to the orbital semi-major axis of the planet. Thus, the bias will manifest itself as a positive correlation between the formally independent parameters $\\frac{L_X}{L_{Bol}}$ and $\\frac{1}{\\scrA}$.  Note that this correlation is not necessarily linear because it is dependent on the observation process, the cumulative observation time, etc.  Any existing correlation also does not imply causality, and indeed the velocities that are generated by these two processes are quite different.  Nevertheless, we expect that stars with large values of $\\frac{L_X}{L_{Bol}}$ will primarily be present in our sample provided they are also accompanied by small values of $\\scrA$.  Thus, a measurement of the correlation between these two variables serves as a measurement of the sample bias.  For the sample of main sequence stars within 60~pc that have EGPs detected via the RV method, we find that \\begin{equation} \\label{e:vel2bias} \\log_{10}\\frac{L_X}{L_{Bol}} = (-5.65 \\pm 0.04) + (0.22 \\pm 0.05) \\cdot \\log_{10}\\frac{1}{\\scrA} \\,, \\end{equation} where $\\scrA$ is in units of [AU]. Here the errors are derived via Monte Carlo bootstrapping on the X-ray flux measurement errors.  Thus, we find that there does exist a weak, but statistically significant, correlation between $\\frac{L_X}{L_{Bol}}$ and $\\frac{1}{\\scrA}$. This trend is shown in Figure~\\ref{f:vel2}.  {\\sl Conservatively assuming that all of this correlation is attributable to the inherent bias}, we can then estimate its effect on the observed activity trend.  That is, for each of the stars in the close-in and distant samples (in \\S\\ref{s:farout}), given $\\scrA$ and $L_{Bol}$ we can compute a predicted $L_X(\\scrA,L_{Bol})$ that can then be used in the place of the measured $L_X$ to carry out the same analysis.  The result of this calculation is shown in Figure~\\ref{f:velohi}, which shows that the close-in and distant samples differ by $\\approx{1}\\sigma$, and account for only a factor of $\\sim 2$ of the {\\sl observed} difference between the two samples. It is important to note that the magnitude of this bias will vary for every subsample, and that the magnitude and direction of the bias will be case specific.  \\begin{figure} \\centerline{\\includegraphics[width=4.5in,angle=0]{f7.eps}}\t% \\caption{The distribution of $\\log\\frac{L_X}{L_{Bol}}$ with $\\log\\frac{1}{\\scrA}$. The X-ray detections are marked with `x's, with error bars denoted with vertical lines; the X-ray non-detections are shown with downward arrows.  Also shown are the best-fit regression line (solid curve) and the envelope bound of the bootstrapped regression curves (dotted lines).  Since more active stars have higher radial velocity jitter, we expect fewer planets to be detected around stars at high $\\frac{L_X}{L_{Bol}}$ unless they also have high $\\frac{1}{\\scrA}$; this correlation tracks an important sample bias (see text). This trend is removed from the rest of the analysis. \\label{f:vel2}} \\end{figure}  \\begin{figure} \\centerline{\\includegraphics[width=4.5in,angle=0]{f8.eps}}\t% \\caption{Same as Figure~\\ref{f:lohi}, but using $L_X(\\scrA,L_{Bol})$ estimated from the regression analysis on $\\frac{L_X}{L_{Bol}}$ with $\\frac{1}{\\scrA}$ (Equation~\\ref{e:vel2bias}). The two distributions differ by $\\approx{2}\\sigma$. \\label{f:velohi}} \\end{figure}  Note that $\\frac{L_X}{L_{Bol}}$ is the preferred proxy for activity because it is surface area independent. However, because there are no $R_*$ biases in the sample, $L_X$ and $\\frac{L_X}{L_{Bol}}$ are similarly correlated with $\\scrA$.  Thus, a priori, we expect that accounting for the selective inhibition bias of Equation~\\ref{e:vel2bias} in the $\\scrA$ vs $L_X$ dataset should result in a complete removal of any luminosity difference between the extremal samples, i.e., that $\\frac{<L_X({\\rm close-in}>}{<L_X{\\rm distant}>}\\approx1$. That this is manifestly not so (Figure~\\ref{f:lohirat}) is an indication that the differences in the distribution of $L_X$ for the extremal samples is systematic, and an inherent property of the observed sample.  Furthermore, the bias measurement is conservative, i.e., we include some of the desired signal within the estimate of the bias, and thus we {\\sl overestimate} the magnitude of the bias.  Because this bias is removed from the observed ratios of the mean luminosities, our results provide a conservative estimate of the true magnitude of the effect of planet-induced activity.  \\item {\\bf Intrinsic sample variance} Other parameters known to affect the level of stellar activity are age, rotation, the Rossby number, metallicity, temporal variability, etc.  Stellar ages are not known with sufficient accuracy for us to consider them here, and it is possible that there exists a trend where $\\scrA$ is correlated with stellar age.  But note that such hidden variables do not affect the results derived here, and simply point to a more complex explanation of the connection between close-in planets and stellar activity.  It is also as yet unclear what type of effect a massive Jupiter-type planet will have on these parameters.  For instance, Rotational synchronization by tidal interactions may be thought to increase $v\\sin{i}$ and thereby coronal activity for close-in EGPs compared to distant EGPs.  However, not only is this as likely to spin-down a fast rotating star as to spin-up a slow-rotator, but also the timescales for the synchronization of the stellar rotation period with the planetary orbital period are very large to begin with, and increase as $\\scrA^6$ for distant planets (see e.g., Drake et al.\\ 1998 and references therein), implying that most stars in the sample are {\\sl not} synchronized to the orbits of their close-in giant planets. We therefore conclude that both theoretical and observational prejudice points to evenly scattered values for these stellar parameters in our sample.  \\end{enumerate}  \\subsection{Activity Enhancement} \\label{s:activity}  We use the parameterization of the bias described above (see \\S\\ref{s:bias}) to determine the residual signal, i.e., the magnitude of the excess X-ray activity in the close-in subsample compared to the distant subsample that can be attributed to the effect of EGPs.  In order to simplify the calculation, we consider the ratio of the average luminosities of the close-in and distant subsamples calculated during the Monte Carlo simulation described above (see analysis in \\S\\ref{s:farout}).  The data from Figure~\\ref{f:lohi} are shown as the solid histogram in Figure~\\ref{f:lohirat}, and confirm that the close-in subsample is on average more X-ray luminous by a factor $4^{<7}_{>2.8}$, where the bounds on the number indicate the extent of the asymmetrical $1\\sigma$-equivalent errors. After correcting for the sample bias (by dividing the original ratio by the ratio of the luminosities as predicted by Equation~\\ref{e:vel2bias}), we find that the bias-corrected luminosity ratio decreases, but is $>1$ at a significance of $0.95$.  This constitutes a definite detection of activity enhancement, by a factor $2.3^{<4}_{>1.3}$ (dashed histogram in Figure~\\ref{f:lohirat}). This remaining enhancement can be attributed to the effect of close-in giant planets on their parent stars. We therefore conclude that the close-in and distant samples indeed show differing X-ray activity levels and that a factor $\\sim 2$ can be attributed to the effect of close-in EGPs. Note that the existence of this residual enhancement does not in itself allow us to conclude that there is a causal connection between the closeness of giant planets to their primary stars and the activity levels on those stars.  However, in conjunction with the Ca\\,II~HK enhancements observed by Shkolnik et al.\\ 2003), our results suggest that such a connection may be present.  \\begin{figure} \\centerline{\\includegraphics[width=4.5in,angle=0]{f9.eps}}\t% \\caption{Distributions of $<L_X{\\rm (close-in)}>/<L_X{\\rm (distant)}>$. The distribution of the ratios of simulated $L_X$ for the close-in ($\\scrA<0.15$ AU) and distant ($\\scrA>1.5$ AU) samples is shown as the solid histogram.  The same ratios, modified to account for the sample bias (see \\S\\ref{s:bias}) are shown with dashed lines. The means and standard deviations of the two distributions are denoted with vertical and horizontal lines respectively (solid and dashed, for the two distributions).  The data show that the two extremal samples differ by a factor $\\approx{4}$ (between 2.8 to 7 times), and after sample bias is accounted for, by a factor $\\approx{2}$ (between 1.3 and 4 times). \\label{f:lohirat}} \\end{figure}  As a further check on the sensitivity of our analysis to the samples used, we compute the same ratio as in Figure~\\ref{f:lohirat} for a number of different subsamples. The resulting distributions of the ratios of mean luminosity for close-in versus distant subsamples are shown in Figure~\\ref{f:lohiratmany}; in all cases the close-in subsamples have larger $<L_X>$ than the distant subsamples. %   \\begin{figure} \\centerline{\\includegraphics[width=4.5in,angle=0]{f10.eps}}\t% \\caption{Effect of variations of sample composition on the distribution of $<L_X{\\rm (close-in)}>/<L_X{\\rm (distant)}>$.  The ratios of $<L_X(\\scrA<0.15)>/<L_X(\\scrA>1.5)>$ are compared for different representative subsets of the full sample of stars with detected planets. In all plots, the baseline comparison from Figure~\\ref{f:lohirat}, for the main-sequence sample within 60~pc, is shown as the dashed histogram.  The solid histograms are: {\\sl (a)} all X-ray detections, {\\sl (b)} the full sample of all stars with EGPs detected via the RV method, {\\sl (c)} the nearby main-sequence sample excluding dM stars, and {\\sl (d)} the full sample as in (b), excluding giants. In all cases, the bias correction (see \\S\\ref{s:bias}) is done separately for each sample.  In all cases, the evidence is strong that stars with close-in giant planets are more active than those with distant planets. \\label{f:lohiratmany}} \\end{figure}   Note that the above result is dependent on the bounds chosen for $\\scrA$ for both subsamples. Clearly, the contrast between the two extremal sets would be heightened if the range of $\\scrA$ defining the two subsamples is shrunk further, and vice versa.  This must also be balanced by the decreased robustness of the result due to the consequent dearth of X-ray detections in the subsamples. A rigorous estimate of the ``best'' split between the close-in and the distant samples is not feasible because of the large number of stars that remain undetected in X-rays and the need to compute the bias factor (Equation~\\ref{e:vel2bias}) separately for each case. Nevertheless, as we have shown above (Figure~\\ref{f:lohirun}) there exists a persistent difference in the mean luminosity between the close-in and distant subsamples, and it favors the suggestion that stars with close-in planets tend to be more active. The best way to improve this result is to obtain X-ray detections of more stars at the extreme ranges of $\\scrA$.   \\subsection{Dependence of the X-ray excess on $\\scrA$} \\label{s:binary}  For tidally interacting close binaries in a sample that included RS\\,CVn stars, Schrijver \\& Zwaan (1991) found that the X-ray surface flux $F_X \\propto \\sqrt{a_*}$, where $a_*$ is the distance to the cooler component. Comparing the average X-ray luminosities for close-in and distant subsamples that we derive above (see \\S\\S\\ref{s:farout},\\ref{s:activity}), after sample biases have been corrected for, we find that \\begin{equation} <F_{X_{\\rm corr}}> ~~ \\propto ~~  <\\scrA>^{-0.35 \\pm 0.17} \\,, \\end{equation} which is a smaller dependence than that found by Schrijver \\& Zwaan, but the large error bar prevents a definitive conclusion.  It is well known that binary stars are generally more X-ray active than single stars of the same type and rotation rate (see, e.g., Zaqarashvili, Javakishvili, \\& Belvedere 2002 and references therein).  Numerous mechanisms have been proposed to account for this so-called ``overactivity,'' generally based on tidal and magnetic interactions.  These studies were extended to the case of stars with EGPs and brown dwarfs by Cuntz et al.\\ (2000). They suggested that the energy generation due to tidal interactions is proportional to the gravitational perturbation \\begin{mathletters} \\begin{equation} \\frac{\\Delta g_*}{g_*} = \\frac{M_P}{M_*} \\frac{2 R_*^3}{(\\scrA-R_*)^3} \\end{equation} where $M_*$ is the mass of the star and $M_P$ that of the EGP, and the height of the tidal bulge, \\begin{eqnarray} \\nonumber h_{tide} &\\propto& \\frac{\\Delta g_*}{g_*} R_* \\\\ \\nonumber &\\propto& \\frac{M_P}{M_*} R_*^4 \\scrA^{-3} \\\\ &\\propto& \\scrA^{-3} \\,. \\end{eqnarray} \\end{mathletters} Saar et al.\\ (2004) estimated the energy released via reconnection during an interaction of the planetary magnetosphere with the stellar magnetic field, \\begin{eqnarray} \\nonumber F_{int} &\\propto& B_* B_P v_{rel} \\scrA^{-n} \\\\ &\\propto& \\scrA^{-n} \\,, \\end{eqnarray} where $B_*$ and $B_P$ are the stellar and planetary magnetic fields, $v_{rel}$ is the relative velocity between them that produces the shear in the magnetic fields and leads to the reconnection.  Here, $n=3$ very close to the star and $n=2$ farther away (in the ``Parker spiral''). Thus, if the magnitude of the enhancement and the stellar magnetic field can be measured for these systems, then the planetary magnetic environment can be investigated.  Since our measurements are an ensemble average, i.e., they do not take into account variations in mass, magnetic fields, and relative velocities, we cannot directly verify the above models.  Nevertheless, we can obtain a rough estimate for the variation of the ``excess'' luminosity with $\\scrA$ by assuming that \\begin{equation} \\delta{L_X} \\propto \\scrA^x \\end{equation} where $\\delta{L_X}$ is the difference between the actual luminosity and a luminosity unaffected by a close-in EGP.  For the former, we use the values simulated during the Monte Carlo analysis of \\S\\ref{s:farout}, and approximate the latter with the average $L_X$ derived for the distant sample.  Fitting straight lines to $\\log{\\delta{L_X}}$ vs $\\log{\\scrA}$, and accounting for the bias in the same manner as described above in \\S\\ref{s:activity}, we find that $x\\approx{-1}$. The error on this value is however very large ($\\sim 100\\%$). Therefore, we conclude that while the data are qualitatively consistent with the scenarios proposed by Cuntz et al.\\ (2000), a reliable test of the theory is not feasible without more observations.     We have searched for X-ray emission from a sample of 230 stars with known giant planets with a view to characterizing the effect of the EGPs on the parent star. We first find that the overall sample of stars with known EGPs is similar in gross X-ray properties to field star samples (see \\S\\ref{s:samplestat}), and thus provides a representative sample of X-ray stars. We carry out a careful search for statistical trends with various parameters (see \\S\\ref{s:adepend}) and find some evidence that the activity levels of stars with close-in giant planets is higher than for stars with planets located further out, though the correlations are contradictory and inconclusive. We then carry out a more powerful test by analyzing in detail two extremal subsamples (see \\S\\ref{s:farout}) where we compare the X-ray emission from nearby (distance$<$60~pc) main-sequence stars with close-in EGPs against similar stars with much more distant EGPs.  For the sake of definiteness, we choose a close-in sample where the orbital semi-major axis $\\scrA<0.15$ AU, and a distant sample where $\\scrA>1.5$ AU (reducing the separation between the samples increases the number of stars considered, but reduces the contrast between the samples).  We verify that our adopte ranges of $\\scrA$ are not special by varying the limiting ranges of $\\scrA$ over which the subsamples are defined, and find that invariably the close-in samples have X-ray luminosities higher than that of the distant sample.  The above result must be understood in the context of selection biases in our sample of stars.  In \\S\\ref{s:bias}, we demonstrate that observational biases account for about half of the observed differences seen in the data.  After these biases are accounted for, we find that the close-in sample is more active by a factor of $2_{>1.3}^{<4}$ on average.  This result holds even when the full data are filtered with different selection criteria.  The robustness of the result is limited by the large number of systems yet to be detected in X-rays.  New observations by {\\sl Chandra} or {\\sl XMM-Newton} would consequently improve the statistics and would constrain the magnitude of this effect with better precision.  We note that a simple model of the X-ray emission enhancement suggests an interaction strength proportional to the product of the stellar and planetary magnetic fields, $B_*\\,B_P$ at their point of interaction.  This suggests that if $B_*$ is known or can be estimated, $B_P$ for exoplanets can potentially be studied. Since the point of interaction between $B_P$ and $B_*$ depends in part on the strength of the stellar wind, close-in wind properties can also potentially be probed by such observations."
    },
    "0807/0807.3282_arXiv.txt": {
        "abstract": "{} {The present work is aimed at studying the distribution of galaxies of different types and luminosities along different structural scaling relations to see whether massive and dwarf ellipticals have been shaped by the same formation process.} {This exercise is here done by comparing the distribution of Virgo cluster massive and dwarf ellipticals and star forming galaxies along the B band effective surface brightness and effective radius vs. absolute magnitude relations and the Kormendy relation to the predictions of models tracing the effects of ram-pressure stripping on disc galaxies entering the cluster environment and galaxy harassment.} {Dwarf ellipticals might have been formed from low luminosity, late-type spirals that recently entered into the cluster and lost their gas because of a ram-pressure stripping event, stopping their activity of star formation. The perturbations induced by the abrupt decrease of the star formation activity are sufficient to modify the structural properties of disc galaxies into those of dwarf ellipticals. Galaxy harassment induce a truncation of the disc and generally an increase of the effective surface brightness of the perturbed galaxies. The lack of dynamical simulations of perturbed galaxies spanning a wide range in luminosity prevents us to drive any firm conclusion on a possible harassment-induced origin of the low surface brightness dwarf elliptical galaxy population inhabiting the Virgo cluster.} {Although the observed scaling relations are consistent with the idea that the distribution of elliptical galaxies along the mentioned scaling relation is just due to a gradual variation with luminosity of the Sersic index $n$, the comparison with models indicates that dwarf ellipticals might have been formed by a totally different process than giant ellipticals. } ",
        "introduction": "Scaling relations are often used as major constraints for models of galaxy formation and evolution. They can be used to trace the contribution of the different stellar components to the total luminosity of galaxies, as in the case of the color magnitude relation (Visvanathan \\& Sandage 1977; Bower et al. 1992 for early-type galaxies, Tully et al. 1982, Gavazzi et al. 1996 for late-type galaxies) or to study the relationship between kinematical, structural and stellar population properties of galaxies as in the case of the Tully-Fisher relation for spirals (Tully \\& Fisher 1977) and the fundamental plane for ellipticals (Dressler et al. 1987; Djorgovski \\& Davis 1987). Scaling relations have been used also to compare the total gas (Boselli et al. 2002), the present or past star formation activity (Boselli et al. 2001) and the metallicity (Zaritsky et al. 1994; Bender et al. 1993) of galaxies with their luminosity or mass. The study of these different scaling relations has been crucial for showing the role of mass in the formation of galaxies (Gavazzi et al. 1996; Boselli et al. 2001), a result now generally called downsizing effect, which is a new major constraint for hierarchical models of galaxy evolution (De Lucia et al. 2006).\\\\ The surface brightness vs. absolute magnitude and the Kormendy relation (effective surface brightness vs. effective radius; Kormendy 1985) have been often used in the literature to study the process governing the formation of early-type galaxies. Although limited to structural properties, these two relations can be easily determined for large samples of galaxies including both giant and dwarf ellipticals. High resolution spectroscopic data, necessary for measuring velocity dispersions needed in the construction of the fundamental plane, are still relatively rare for low luminosity systems. \\\\ The study of the B band surface brightness vs. absolute magnitude relation and of the Kormendy relation has originally shown a strong, apparent dichotomy in the behavior of dwarf and giant ellipticals. While in dwarfs the effective or central surface brightness increases with luminosity, the opposite trend is seen in giants (e.g. Ferguson \\& Binggeli 1994; Graham \\& Guzman 2003). An opposite trend between giants and dwarfs has been also observed in the Kormendy relation (e.g. Kormendy 1985; Capaccioli et al. 1992). This surprising result has been originally interpreted as a clear indication that dwarf ellipticals are not the low luminosity extension of giants but rather an independent class of objects. The recent study of Graham \\& Guzman (2003) based on HST data has however shown that this dichotomy is only apparent since it is due to a gradual steepening of the central radial profile with luminosity: the radial light profile of ellipticals is characterized by a Sersic law, with an index $n$ increasing from $\\sim$ 1 in dwarfs ($M_B$ $\\sim$ -13) to $\\sim$ 4 in brighter galaxies until the detection of core formation ($M_B$ $\\sim$ -20.5), this last probably related to the presence of a massive black hole (Faber et al. 1997; C\\^ot\\'e et al. 2006). The different behavior of E and dE in the surface brightness vs. absolute magnitude relation and in the Kormendy relation is just due to a gradual increase of $n$. Graham \\& Guzman (2003) thus concluded that dE appear to be the low luminosity extension of massive ellipticals and suggested that {\\it \"the mechanisms of how dE and E galaxies collapsed to form stars are therefore expected to be similar\"}.\\\\ Several recent observational evidences and simulations, however, seem to indicate that local group dwarf spheroidals (Mayer et al. 2006), Virgo cluster dwarf ellipticals (Barazza et al. 2002; Conselice et al. 2003; van Zee et al. 2004; Mastropietro et al. 2005; Lisker et al. 2006a,b; 2007, 2008, Lisker \\& Han 2008; Michielsen et al. 2008) or generally dwarf spheroidals in other clusters such as Coma (Smith et al. 2008), Perseus (Penny \\& Conselice 2008) or in the SDSS (Haines et al. 2007) might be late-type galaxies recently perturbed by an hostile environment through ram-pressure stripping or galaxy harassment. In a recent work based on a multifrequency analysis of dwarf ellipticals in the Virgo cluster combined with multizone model of the chemical and photometric evolution of galaxies Boselli et al. (2008) have shown how several observational properties of dE are consistent with those of star forming discs that recently entered the cluster environment loosing on a short time scale ($\\leq$ 150 Myr) their gas reservoir and thus abruptly stopped their star formation activity becoming quiescent systems. It is thus natural to wonder whether the classical scaling relations such as the surface brightness vs. absolute magnitude relation and the Kormendy relation of dwarf ellipticals could be produced by environmental effects or are still a clear indication that dE have been formed by the same process than giant ellipticals, that the most recent models of hierarchical galaxy formation indicate as major merging (De Lucia et al. 2006). \\\\ This work is the continuation of what presented in Boselli et al. (2008), where most off the specific aspects of the comparison between models and observations have been presented. The surface brightness luminosity relation has been already discussed, although in a different form, in Boselli et al. (2008). Here we consider it in its general form ($<\\mu_e(B)>$ vs. $M_B$) and combined with the effective radius $R_e(B)$ vs. absolute magnitude relation and the Kormendy relation. The aim of the present paper is that of comparing our observations to ram-pressure stripping model predictions and to the dynamical simulations of harassed galaxies in clusters. In particular we want to investigate whether the $<\\mu_e(B)>$ vs. $M_B$, $R_e(B)$ vs. $M_B$ and the Kormendy relations can be univocally taken, as claimed by Graham \\& Guzman (2003), as a clear indication that dE have been shaped through the same formation process than giant ellipticals. The comparison of our model predictions with observations indeed indicate that dE might have been formed through a totally different evolutionary process than their giant counterparts and not results from a hierarchical assembly of matter as indicated by cosmological models.  \\begin{figure*} \\centering \\includegraphics[width=12cm,angle=0]{grahamE.eps} \\caption{The effective surface brightness $<\\mu_e(B)>$ (upper left) and radius $R_e(B)$ (lower left) vs. B band absolute magnitude relations and the Kormendy relation ($<\\mu_e(B)>$ vs. radius $R_e(B)$) (upper right) for early-type galaxies: red, open circles are for ellipticals and lenticulars (E-S0-S0a), magenta asterisks for nucleated dwarf ellipticals (dE,N), orange, filled dots for dE. Sizes of the symbols have been chosen according to the light concentration index: large for $C_{31}(B)$ $>$5; medium for 3.5 $<$ $C_{31}(B)$ $\\leq$5 and small for $C_{31}(B)$ $\\leq$3.5. The horizontal dashed line gives the VCC surface brightness detection limit (25.3 mag arcsec$^{-2}$). Typical error bars for giants and dwarfs are given. } \\label{grahamE} \\end{figure*}   \\begin{figure*} \\centering \\includegraphics[width=12cm,angle=0]{grahamS.eps} \\caption{The effective surface brightness $<\\mu_e(B)>$ (upper left) and radius $R_e(B)$ (lower left) vs. B band absolute magnitude relations and the Kormendy relation ($<\\mu_e(B)>$ vs. radius $R_e(B)$) (upper right) for late-type galaxies (blue crosses for Sa-Sc, green open triangles for Scd-Im-BCD). Sizes of the symbols have been chosen according to the light concentration index: large for $C_{31}(B)$ $>$5; medium for 3.5 $<$ $C_{31}(B)$ $\\leq$5 and small for $C_{31}(B)$ $\\leq$3.5.} \\label{grahamS} \\end{figure*} ",
        "conclusions": "The analysis presented so far is totally consistent with the idea that the origin of the different behavior of dwarf and massive early-type galaxies along the B band effective surface brightness and effective radius vs. absolute magnitude relations and along the Kormendy relation often described in the literature is deeply related to the shape of their radial light profiles, as suggested by Graham \\& Guzman (2003). We indeed observe that the effective surface brightness decreases (and the effective radius increases) with luminosity only in early-type galaxies with high concentration indices or high Sersic indices $n$, while it increases (it is almost constant) whenever $C_{31}(B)$ $<$ 5 or $n$ $<$ 2.5. Massive ($M_B$ $\\leq$ -20) ellipticals and lenticulars all have $C_{31}(B)$ $\\geq$ 5, while quiescent dwarfs ($M_B$ $\\geq$ -18) have $C_{31}(B)$ $\\sim$ 3 consistent with a roughly exponential profile.\\\\ In late-type galaxies the effective surface brightness and the effective radius monotonically increase with luminosity, but with flatter slopes and larger scatters than in low luminosity and massive ellipticals respectively. They all have $C_{31}(B)$ $\\sim$ $\\leq$ 6, with $C_{31}(B)$ $\\sim$ 3 in $\\sim$ 70\\% of the massive objects. \\\\ The comparison with models have shown that the distribution of quiescent galaxies along the aforementioned scaling relations cannot be univocally taken as a proof of a similar formation process for dwarf and giants, as stated by Graham \\& Guzman (2003). The lack of dynamical simulations for the whole range in luminosity of the disc galaxy population prevents us to say whether the scaling relations observed in E-dE galaxies here discussed can be the result of the harassment suffered by spirals entered into the cluster.\\\\ The observed relations between the effective surface brightness, effective radius and absolute magnitude of low luminosity ellipticals are consistent with their formation after a recent and rapid decrease of the star formation activity of low luminosity spirals recently entered into the cluster medium and swept of their gas by a ram-pressure stripping event, as claimed by Boselli et al. (2008). This evolutionary path is thus significantly different than that followed by massive ellipticals for which hierarchical models of galaxy formation predict they formed most of their stars at early epochs ($z$ $\\geq$ 2) and later assembled through major merging events (De Lucia et al 2006). A different evolution of massive and dwarf ellipticals is also confirmed by the analysis of their spectro-photometric properties and of their [$\\alpha$/Fe] ratios, both indicating a rapid formation of the bulk of the stellar population in the early universe in massive ellipticals while a rather continuous and moderate star formation activity for dwarf systems (Thomas et al. 2005; Renzini 2006).\\\\ As a general conclusion we can say that the present work is a further indication that the interpretation of single scaling relations in terms of evolution is not always straightforward but rather needs to be included in a complete and coherent study combining multifrequency observations with model predictions."
    },
    "0807/0807.4141_arXiv.txt": {
        "abstract": "We measure the sizes of 82 massive ($M\\geq10^{11} M_{\\odot}$) galaxies at $1.7\\leq z\\leq 3$ utilizing deep HST NICMOS data taken in the GOODS North and South fields. Our sample is almost an order of magnitude larger than  previous studies at these redshifts, providing the first statistical study of massive galaxy sizes at $z > 2$, confirming the extreme compactness of these galaxies.  We split our sample into disk-like ($n\\leq2$) and spheroid-like ($n>2$) galaxies based on their S\\'{e}rsic indices, and find that at a given stellar mass disk-like galaxies at $z\\sim2.3$ are a factor of 2.6$\\pm$0.3 smaller than present day equal mass systems, and spheroid--like galaxies at the same redshifts are 4.3$\\pm$0.7 smaller than comparatively massive elliptical galaxies today.  At $z>2$ our results are compatible with both a leveling off, or a mild evolution in size. Furthermore, the high density ($\\sim$2$\\times$10$^{10}$M$_{\\odot}$kpc$^{-3}$) of massive galaxies at these redshifts, which are similar to present day globular clusters, possibly makes any further evolution in sizes beyond z=3 unlikely. ",
        "introduction": "One of the most exciting discoveries in extragalactic astronomy in the last few years is that massive ($M\\geq10^{11} M_{\\odot}$) galaxies at $z>1$ were extremely compact (Daddi et al. 2005; Trujillo et al. 2006b, 2007; Longhetti et al. 2007), particularly those with the lowest estimated star formation rates (Zirm et al. 2007; Toft et al. 2007; Cimatti et al. 2008; van Dokkum et al. 2008; P\\'{e}rez-Gonz\\'{a}lez et al. 2008). Since only a very few dense and massive objects are found at z$\\sim$0 (Bernardi et al. 2006; cf. with none at r$_{e} < $1 kpc) it is clear that significant growth in the sizes of these galaxies has occurred during cosmic history.  Within the current galaxy formation paradigm, the origin of these galaxies can be described by the collapse and merging of dark matter haloes. Models suggest that at very early times galaxies contain large amounts of cold gas, resulting in efficient starbursts (e.g., Khochfar \\& Silk 2006). As star formation occurs, the gas in these galaxies becomes heated due to various feedback processes (e.g., Granato et al. 2004, Menci et al 2006), leading to reduced star formation rates, creating compact and massive remnants. Observationally, we know that at $z < 2$ there are few gas rich mergers in massive galaxies based on structural analyses (e.g. Conselice et al. 2003; Conselice 2006; Conselice et al. 2008) preventing significant starbursts and new star populations from forming. Since at these lower redshifts the amount of available gas has decreased, ``dry'' mergers are expected to be the dominant mechanism for size and stellar mass growth (Ciotti \\& van Albada 2001, Dominguez-Tenreiro et al. 2006; Boylan-Kolchin et al 2006), although other processes are possible (Naab et al. 2007; Pipino \\& Matteucci 2008).  One of the ways to trace this evolution is through measuring the sizes of galaxies through time.  At $z < 2$, the size evolution of the most massive galaxies has been well characterized with large samples of objects. Recently, Trujillo et al. (2007) using $\\sim$800 sources found that, at a given stellar mass, disk--like objects at z$\\sim$1.5 were a factor of two smaller than their present-day counterparts. For spheroid--like objects the evolution is even stronger. These spheroidal objects are a factor of four smaller at $z \\sim 1.5$ compared with similar mass modern ellipticals. This evolution is also in qualitative agreement with hierarchical semi-analytical model predictions which find a factor of $1.5-3$ evolution in size since that redshift (e.g., Khochfar \\& Silk 2006) .  At $z>$2, however, our knowledge of the size evolution of the most massive objects is much more scarce. There are only a few attempts to explore this issue  using small samples of massive galaxies at z$\\sim$2.5 (Zirm et al. 2007; Toft et al. 2007; van Dokkum et al. 2008). This is due to the intrinsic scarcity of distant massive galaxies, and the relative small sizes of previous deep NIR imaging surveys. With the aim of substantially increase our knowledge of the size evolution of massive galaxies in the redshift interval $1.7\\leq z\\leq 3$, we have imaged a sample of 82 very massive galaxies in the GOODS North and South fields  within the H-band filter as part of the GOODS NICMOS Survey (Conselice et al. 2008, in prep).  To allow a comparison with both the local SDSS stellar mass-size relations, and the results obtained at lower redshifts ($z<2$), we split our sample according to light concentration using the S\\'{e}rsic index $n$ to separate disk-like galaxies from more concentrated spheroid-like systems. We find that both types continue, and perhaps level off, in their size evolution at $z > 2$.   We assume the following cosmology throughout: H$_{0}$=70 km s$^{-1}$Mpc$^{-1}$, $\\Omega_{\\lambda}$ = 0.7, and $\\Omega_{\\rm m}$ = 0.3, and use AB magnitude units ",
        "conclusions": "As has been demonstrated in previous work, some massive galaxies at $z < 2$ grow in size by up to an order of magnitude (e.g., Trujillo et al. 2007).  An interesting question is whether these massive galaxies become progressive smaller at higher redshifts, containing possibly even smaller sizes at $z>2$. Our results provide the first statistical sample in which to answer this question.  As seen in Fig. 2, the objects in our sample are compatible with the idea that the size evolution reaches a plateau beyond $z=2$. To shed some light on this question we compute the stellar mass density of our galaxies and compare these to the densest collection of stars in the local Universe -- globular clusters. This comparison is an interesting one, since globular clusters are also expected to be form either very early, or more recently as a result of mergers of gas clouds during galaxy collisions.   A typical spheroid--like galaxy in our sample at z$\\sim$2.75 has a stellar mass of $\\sim$2$\\times$10$^{11}$M$_{\\odot}$, and a size of r$_e$$\\sim$1 kpc. The stellar density for this object, assuming spherical symmetry, is $\\rho=(0.5M)/(4/3\\pi r_{e}^{3})$$\\sim$2.4$\\times$10$^{10}$M$_{\\odot}{\\rm kpc}^{-3}$. A disk--like galaxy at z$\\sim$2.75 has a typical mass of $\\sim$2$\\times$10$^{11}$M$_{\\odot}$ and size r$_e$$\\sim$2 kpc. Assuming a disk symmetry, the stellar mass density within these disk-like systems is $\\rho=(0.5M)/(\\pi r_{e}^{2}h)$$\\sim$2.6$\\times$10$^{10}$M$_{\\odot}{\\rm kpc}^{-3}$, where we have used h$\\sim$0.3 kpc. In both cases the stellar mass densities are similar. A typical globular cluster (r$_e$=10 pc and M$\\simeq$10$^5$M$_{\\odot}$) has a density of $\\sim$1.2$\\times$10$^{10}$M$_{\\odot}$kpc$^{-3}$. This is remarkably similar to our massive galaxies at $z > 2$, and reveals that these high--z galaxies may in principle have an origin similar to globular clusters. These high densities also suggest that their stellar mass densities likely do not become much larger at high redshifts ($z > 3$). A massive galaxy at $z > 2$ must also have formed very quickly, and consequently these high stellar densities could reflect the high gas densities in the primeval Universe. The compactness of our objects, and their similar densities to globulars, is consistent with a scenario whereby more massive haloes start collapsing earlier and drag along a large amount of baryonic matter that later forms into stars.   If, as suggested by the high density of our galaxies, the size evolution is stopped or diminished at $z>2$, this perhaps reveals a different evolutionary mechanism for massive galaxies at $z<2$. A faster size evolution is in agreement with theoretical models, which predict that the amount of gas involved in galaxy mergers decreases with lower redshifts. A lower amount of gas results in a more efficient size growth, as the energy of the collision is not dissipated into the formation of new stars (e.g., Khochfar \\& Silk 2006).  We thank S\\'{e}bastien Foucaud and Asa F. L. Bluck for instructive discussions, and Leonel Guti\\'{e}rrez and In\\'{e}s Flores Cacho for their valuable help with computing issues.  We also thank the other members of the GNS team, particularly Emanuele Daddi and Casey Papovich for their participation in the various aspects of the survey. The GNS is supported by NASA grant HST-GO-11082.  \\renewcommand{\\arraystretch}{1.2} \\begin{center} \\begin{table*}  \\caption{Size evolution of massive (M$_{\\star} > 10^{11}$ M$_{\\sun}$) galaxies at 1.7$<$z$<$3} \\hspace{1.6in} \\begin{tabular}{ccc} \\hline Redshift Range & n$<$2 &  n$>$2   \\\\ &  $<$r$_e$/r$_{e,SDSS}$$>$($\\pm$1$\\sigma$) &  $<$r$_e$/r$_{e,SDSS}$$>$($\\pm$1$\\sigma$) \\\\ \\hline 0.1 (SDSS)& 1 \t\t& 1           \\\\ 1.7-2.0\t  & 0.44(0.07)  & 0.23(0.04)  \\\\ 2.0-2.5\t  & 0.39(0.04)  & 0.23(0.04)  \\\\ 2.5-3.0\t  & 0.31(0.03)  & 0.14(0.03)  \\\\  \\hline \\label{data1} \\end{tabular} \\end{table*} \\end{center} \\renewcommand{\\arraystretch}{1}  \\renewcommand{\\arraystretch}{1.2} \\begin{center} \\begin{table*}  \\caption{Fit to $\\sim \\alpha(1+z)^{\\beta}$ to Trujillo et al. (2007), our data and Figure 2} \\hspace{2in} \\begin{tabular}{ccc} \\hline Disk--like galaxies \\\\ \\hline \\hline Redshift & $\\alpha$($\\pm$1$\\sigma$) &  $\\beta$($\\pm$1$\\sigma$)\\\\ \\hline 0.0-2.0\t  & 1.08(0.02)  & -0.78(0.04)  \\\\ 1.7-3.0\t  & 1.85(0.28)  & -1.34(0.59)  \\\\ 0.0-3.0\t  & 1.08(0.01)  & -0.82(0.03)  \\\\  \\hline \\hline Spheroid--like galaxies\\\\ \\hline \\hline Redshift & $\\alpha$($\\pm$1$\\sigma$) &  $\\beta$($\\pm$1$\\sigma$)\\\\ \\hline 0.0-2.0\t  & 1.16(0.01)  & -1.51(0.04)  \\\\ 1.7-3.0\t  & 1.42(0.29)  & -1.66(0.92)  \\\\ 0.0-3.0\t  & 1.15(0.01)  & -1.48(0.04)  \\\\  \\hline \\label{data2} \\end{tabular} \\end{table*} \\end{center} \\renewcommand{\\arraystretch}{1}"
    },
    "0807/0807.3885_arXiv.txt": {
        "abstract": "We present an interferometric and single dish study of small organic species toward Comets C/1995 O1 (Hale-Bopp) and C/2002 T7 (LINEAR) using the BIMA interferometer at 3 mm and the ARO 12m telescope at 2 mm.  For Comet Hale-Bopp, both the single-dish and interferometer observations of CH$_3$OH indicate an excitation temperature of 105$\\pm$5 K and an average production rate ratio $Q(CH_3OH)/Q(H_2O)$$\\sim$1.3\\% at $\\sim$1 AU.  Additionally, the aperture synthesis observations of CH$_3$OH suggest a distribution well described by a spherical outflow and no evidence of significant extended emission. Single-dish observations of CH$_3$CN in Comet Hale-Bopp indicate an excitation temperature of 200$\\pm$10 K and a production rate ratio of $Q(CH_3CN)/Q(H_2O)$$\\sim$0.017\\% at $\\sim$1 AU. The non-detection of a previously claimed transition of cometary (CH$_2$OH)$_2$ toward Comet Hale-Bopp with the 12m telescope indicates a compact distribution of emission, D$<$9$''$ ($<$8500 km). For the single-dish observations of Comet T7 LINEAR, we find an excitation temperature of CH$_3$OH of 35$\\pm$5 K and a CH$_3$OH production rate ratio of $Q(CH_3OH)/Q(H_2O)$$\\sim$1.5\\% at $\\sim$0.3 AU. Our data support current chemical models that CH$_3$OH, CH$_3$CN and (CH$_2$OH)$_2$ are parent nuclear species distributed into the coma via direct sublimation off cometary ices from the nucleus with no evidence of significant production in the outer coma. ",
        "introduction": "Over the last several years, complementary observations of comets with both single-dish telescopes and interferometric arrays, have shed new light on the abundance, production rate, distribution and formation of molecules in cometary comae.  Recently, Milam et al.\\ (2006) employed this method to investigate the origins of cometary formaldehyde (H$_2$CO) by comparing observations with both the Arizona Radio Observatory (ARO) 12m\\footnote{The Kitt Peak 12m telescope is currently operated by Arizona Radio Observatory (ARO), Steward Observatory, University of Arizona.} single-dish telescope and the Berkeley-Illinois-Maryland Association (BIMA) array\\footnote{Operated by the University of California, Berkeley, the University of Illinois, and the University of Maryland with support from the National Science Foundation.}. From this investigation toward Comet C/1995 O1 (Hale-Bopp), the authors clearly showed that the distribution and abundance of H$_2$CO must be coming from a source other than the comet nucleus. The extended source of H$_2$CO may be due to an enhanced production of this molecule, ($\\sim$4$\\times$ larger) from grains consisting of silicates and organic material (Milam et al.\\ 2006), or possibly from the photodissociation of a larger parent species such as methanol (CH$_3$OH) (Hudson, Moore \\& Cook 2005).  These complementary observations, however, are very difficult to obtain, as most of the molecular observations toward comets are performed using solely single-dish telescopes or interferometric arrays, depending on the telescope available at the time as the comet approaches the inner solar system and perigee.  Multiple facility collaborations for cometary observations are necessary if we are to further understand the formation, excitation and distribution of organic cometary material.  Recent cometary models of the abundance of organic material suggest Oort Cloud comets may have had an origin in the giant planet region, where they were thermally and chemically processed, before being ejected into the Oort Cloud (Weissman 1999).  Furthermore, observations comparing long- and short-period comets suggest distinct classes - organic ``rich'', ``normal'', and ``depleted'' (Bonev et al.\\ 2008; Mumma et al.\\ 2003).  However, the current sample size is too limited to accurately characterize each class of comets.  It should also be noted that many of the comets described as organic ``rich'' have undergone breakup or fragmentation events, thus exposing more internal, unprocessed material than an object just entering the inner solar system (Bonev et al.\\ 2008). So, the question persists, how does the classification of organic ``rich'', ``normal'' or ``depleted'' hold with regards to comets that are for example, dust rich versus dust poor, having undergone a recent breakup, or bear high abundances of hypervolatile species?  The organic species examined in this study, methanol (CH$_3$OH), methyl cyanide (CH$_3$CN) and ethylene glycol ([CH$_2$OH]$_2$) are believed to be parent species directly sublimated off the cometary ices in the nucleus.  Since the first reported single-dish detections of cometary CH$_3$OH towards Comet Brorsen-Metcalf (1989~o) and Comet Austin (1989~c1) (Colom et al.\\ 1990), CH$_3$OH remains the most abundant large ($>$5 atoms) organic (containing a C-atom) molecule measured in cometary ices. In a recent review by Charnley \\& Rodgers (2008), CH$_3$OH has a measured abundance range relative to H$_2$O between 1-4\\%.  However, aperture synthesis observations of cometary CH$_3$OH are very rare and no high spatial resolution data of this molecule currently exists. Similarly, single-dish observations of CH$_3$CN and (CH$_2$OH)$_2$ also suggest a purely nuclear origin (Biver et al.\\ 2002, Crovisier et al.\\ 2004). CH$_3$CN, as well as HC$_3$N, may be parent species that would increase the abundance of the CN radical as well as C$_2$ in the coma. Yet, current observations do not preclude the possibility of CH$_3$CN having its own parent species (Bockel{\\'e}e-Morvan et al.\\ 1987).  In addition, the high abundance of (CH$_2$OH)$_2$ observed in Comet Hale-Bopp at high spatial resolution also strongly suggests a nuclear origin.  To date, complementary observations of the same transitions of either CH$_3$CN or (CH$_2$OH)$_2$ with another facility were not available to follow-up on the claimed distribution of these species in the coma and their possible formation pathways.  In this paper, we report on complementary observations of cometary CH$_3$OH, CH$_3$CN  and (CH$_2$OH)$_2$ taken with the BIMA array near Hat Creek, CA and the ARO 12m telescope on Kitt Peak, AZ toward Comets C/1995 O1 (Hale-Bopp) and C/2002 T7 (LINEAR).  From these observations, we continued to investigate the abundance, production rate, distribution and formation of molecules in comets by directly comparing the measurements of these organic species with different telescopes.  Our data supports the hypothesis that CH$_3$OH, CH$_3$CN and (CH$_2$OH)$_2$ are primarily parent species, generated by direct sublimation of cometary ices contained in the nucleus. ",
        "conclusions": "We presented an interferometric and single dish study of cometary molecules toward comets C/1995 O1 (Hale-Bopp) and C/2002 T7 (LINEAR) using the Berkeley-Illinois-Maryland Association (BIMA) interferometer at 3 mm and the Arizona Radio Observatory (ARO) 12m telescope at 2 mm.  The overall conclusions from our analysis of these data are:  \\begin{enumerate}  \\item The CH$_3$OH Hale-Bopp data are well fit by an excitation temperature of T$_{ex}$=105$\\pm$5 K and a total beam-averaged column density of N$_{T}$=$5.7\\pm0.7\\times10^{14}$~cm$^{-2}$.  Assuming a nuclear origin for CH$_3$OH and using a Monte Carlo model, we find an average CH$_3$OH production rate from both instruments of $Q(CH_3OH)$$\\sim$1.3$\\times$10$^{29}$ s$^{-1}$ or $Q(CH_3OH)$/$Q(H_2O)$$\\sim$1.3\\% at $\\sim$ 1 AU.  \\item The CH$_3$OH ARO 12m LINEAR data centered around 157.26 GHz are best fit by an excitation temperature of T$_{ex}$=35$\\pm$5 K and a total beam-averaged column density of N$_{T}$=2.2$\\pm$0.3$\\times10^{13}$~cm$^{-2}$ and thus, a production rate $Q(CH_3OH)$$\\sim$2.0$\\times$10$^{27}$ s$^{-1}$ or $Q(CH_3OH)$/$Q(H_2O)$=1.5\\% at 0.3 AU.  \\item  From the combination of the single-dish and aperture synthesis observations of CH$_3$OH, we find the distribution of CH$_3$OH toward Comet Hale-Bopp is well described by a spherical outflow with an increase in column density closer to the cometary nucleus.  The data presented in the array images show no evidence of significant enhanced production of CH$_3$OH in the extended coma or from jets as any extended emission is completely resolved out by the array.  \\item The CH$_3$CN 12m Hale-Bopp data is best fit by an excitation temperature, T$_{ex}$=200$\\pm$10 K and a total beam-averaged column density, N$_{T}$=2.6$\\pm$0.3$\\times10^{12}$~cm$^{-2}$. Assuming a nuclear origin for CH$_3$CN and using a Monte Carlo model, we find an average CH$_3$CN production rate of $Q(CH_3CN)$$\\sim$1.7$\\times$10$^{27}$ s$^{-1}$ or $Q(CH_3CN)$/$Q(H_2O)$$\\sim$0.017\\% at $\\sim$ 1 AU.  A comparison between the single-dish observations from the ARO 12m and the IRAM 30m of CH$_3$CN suggest that the ARO observations are beam diluted.  The excitation of the transitions observed of CH$_3$CN takes place at a smaller radii where the density is higher and suggest a nuclear origin of CH$_3$CN.  \\item The non-detection of a previously claimed transition of cometary (CH$_2$OH)$_2$ toward Comet Hale-Bopp with the ARO 12m telescope indicates a compact distribution of emission on the order of $<$9$''$ ($<$8500 km).  This supports the hypothesis that the cometary production of (CH$_2$OH)$_2$ is direct sublimation off cometary ices from the nucleus.   We thank D. K. Yeomans for ephemerides assistance and G. Engargiola, T. Helfer, W. Hoffman, R. L. Plambeck, and M. W. Pound for invaluable technical contributions. We also thank an anonymous referee for a favorable review of this work and whose comments and suggestions provided additional clarity to this manuscript.  This material is based on work supported by the National Aeronautics and Space Administration through the NASA Astrobiology Institute under Cooperative Agreement CAN-02-OSS-02 issued through the Office of Space Science. S. N. M. would like to thank the Phoenix Chapter of ARCS, specifically Mrs. Scott L. Libby, Jr. endowment, for partial funding. This work was partially funded by NASA NAG5-4292, NAG5-4080, NAG5-8708, and NGT5-0083; NSF AST 96-13998, AST96-13999, AST96-13716, AST96-15608, and AST99-81363; and the Universities of Illinois, Maryland, and California, Berkeley. M. W. was funded by NSF AST 9625360, AST 9796263, and AST 0098583 and NASA NAG5-4349.   \\end{enumerate}"
    },
    "0807/0807.0138_arXiv.txt": {
        "abstract": "% Properties of Lyman Break Galaxies at $z\\sim5$ are presented and are compared with those at $z\\sim 2-3$ LBGs to see evolution of LBGs from $z\\sim5$ to $z \\sim3-2$. The LBGs at $z\\sim5$ are tend to have smaller stellar masses, but they may show large star formation rates and specific star formation rates. ",
        "introduction": "We have been studying properties of Lyman Break Galaxies (LBGs) at $z\\sim5$, $\\sim 1$ Gyr after the Big Bang and before $z\\sim3$. In this contribution, we present our recent results obtained for $z\\sim5$ LBGs and possible evolution from $z\\sim5$ LBGs  to $z \\sim 2-3$ LBGs. Our sample at $z\\sim5$ was made in a field including the Hubble Deep Field North (HDF-N) by using the Subaru SCam through $V$-, $I_C$-, and $z'$-filters. Since in and around the HDF-N field many spectroscopic data were available, the color criteria to select the $z\\sim 5$ LBGs can be reliably examined. A sample of $z\\sim5$ LBGs was chosen from galaxies brighter than $z^{\\prime}=26.5$ mag. A number of the sample is $\\sim 600$ in an area of $\\sim500$ arcmin$^2$. ",
        "conclusions": ""
    },
    "0807/0807.0624_arXiv.txt": {
        "abstract": "We present a new Monte Carlo Markov Chain algorithm for CMB analysis in the low signal-to-noise regime. This method builds on and complements the previously described CMB Gibbs sampler, and effectively solves the low signal-to-noise inefficiency problem of the direct Gibbs sampler. The new algorithm is a simple Metropolis-Hastings sampler with a general proposal rule for the power spectrum, $C_{\\ell}$, followed by a particular deterministic rescaling operation of the sky signal, $\\mathbf{s}$. The acceptance probability for this joint move depends on the sky map only through the difference of $\\chi^2$'s between the original and proposed sky sample, which is close to unity in the low signal-to-noise regime. The algorithm is completed by alternating this move with a standard Gibbs move. Together, these two proposals constitute a computationally efficient algorithm for mapping out the full joint CMB posterior, both in the high and low signal-to-noise regimes. ",
        "introduction": "Since the detection of anisotropy in the Cosmic Microwave Background (CMB; Smoot et al.\\ 1992), there has been an emphasis on likelihood or Bayesian methods for the inference of cosmological parameters and their error bars, or more generally, their confidence intervals. CMB analysis is most suitably addressed in a Bayesian, as opposed to frequentist, framework, simply because the observed microwave sky is interpreted as a single realization of a spatial random process.  Early measurements of the CMB were limited to signal to noise ratios of order unity at relatively low angular scales, where direct evaluation of the likelihood for the power spectrum or cosmological parameters is possible. However, the ${\\cal O}(N^{3})$ scaling of computational expense with pixel number $N$ prohibits direct likelihood evaluation for current and future CMB observations. Motivated by the scientific potential of CMB data with increasingly high spatial resolution, yet beset with systematics including partial sky coverage and foregrounds, an iterative method of sampling from the Bayes posterior, using a special case of Markov Chain Monte Carlo (MCMC) known as Gibbs sampling, was introduced by \\citep{jewell:2002,jewell:2004}. The method was later independently discovered and applied to COBE data \\citep{wandelt:2004}, numerically extended to high-resolution on the sphere \\citep{eriksen:2004}, applied to analysis of the WMAP \\citep{bennett:2003, hinshaw:2007, page:2007} data \\citep{odwyer:2004, eriksen:2007a, eriksen:2007b}, as well as generalized to include inference of foreground model parameters \\citep{eriksen:2008a, eriksen:2008b}.  While Gibbs sampling provably converges to the Bayes posterior over the entire range of angular scales probed by the data, the run-time required to generate enough independent samples at the low signal-to-noise, small angular scale regime was found to be prohibitive \\citep{eriksen:2004}. The reason for this is that typical variations in the power spectrum from one sample to the next are determined by cosmic variance alone, whereas the posterior itself is given by both cosmic variance and noise.  This results in a long correlation length in the sequence of spectra in the low signal to noise regime, thus requiring a very long run time to generate a sufficient number of independent samples.  In this paper we generalize the original Gibbs sampling algorithm to include a new type of MCMC step alternating with standard Gibbs sampling, which solves this problem of slow probabilistic convergence in the low signal to noise regime.  This method therefore makes possible an exact Bayesian approach to CMB analysis over the entire range of angular scales probed by current and future experiments.  The paper is organized as follows. We first review the CMB Gibbs sampler, and describe the associated numerical difficulties in analysis at small angular scales.  We then introduce the new MCMC step to the Markov chain, designed specifically to allow large variations in the high-$\\ell$ CMB spectrum, precisely where the signal to noise is $\\le 1$. We derive the required Metropolis-Hastings acceptance probability correctness in Appendix \\ref{app:proof}, and numerically demonstrate the method in Section \\ref{sec:simulations}, for both temperature and polarization. Finally, we summarize and conclude in Section \\ref{sec:conclusions}. ",
        "conclusions": "\\label{sec:conclusions}  We have presented a new MCMC algorithm for the high-L, low signal to noise limit of the joint posterior which solves the slow probabilistic convergence of the traditional Gibbs sampler in this regime.  This in principle allows sampling over the joint posterior $p(C_{l}, \\Bs | \\Bd)$ over the entire range of angular scales probed by current and future CMB experiments. The limiting computational burden is now entirely in the map-making step of Gibbs sampling, for which the cost per Gibbs iteration now scales with the expense of multiplication by the inverse noise matrix $\\BN^{-1}$.  Assuming pixel uncorrelated (but scan weighted) noise as a good approximation at small angular scales, the cost of an $\\BN^{-1}$ multiplcation is that of a forward and inverse spherical harmonic transform, or ${\\cal O}(\\ell_{\\textrm{max}}^{3})$.  Future work will attempt to push the generalized Gibbs + MCMC sampling scheme presented here to smaller angular scales, ultimately limited by the degree to which we can compute harmonic transforms."
    },
    "0807/0807.2635_arXiv.txt": {
        "abstract": "Strong incompressible three-dimensional magnetohydrodynamic turbulence is investigated by means of high resolution direct numerical simulations. The simulations show that the configuration space is characterized by regions of positive and negative cross-helicity, corresponding to highly aligned or anti-aligned velocity and magnetic field fluctuations, even when the average cross-helicity is zero. To elucidate the role of cross-helicity, the spectra and structure of turbulence are obtained in `imbalanced' regions where cross-helicity is non-zero. When averaged over regions of positive and negative cross-helicity, the result is consistent with the simulations of balanced turbulence. An analytical explanation for the obtained results is proposed. \\pacs{52.35.Ra} ",
        "introduction": " ",
        "conclusions": ""
    },
    "0807/0807.2545_arXiv.txt": {
        "abstract": "{ We present a 3-dimensional study of the local  ($\\leq 100 \\; h^{-1}$ kpc) environment of Sy1, Sy2 and Bright IRAS Galaxies. For this purpose we use three galaxy samples (Sy1, Sy2, BIRG) located at high galactic latitudes as well as three control sample of non-active galaxies having the same morphological, redshift and diameter size distributions as the corresponding Seyfert or BIRG sample. Using the CfA2 and SSRS galaxy catalogues as well as our own spectroscopic observations, we find that the fraction of BIRGs with a close neighbor is significantly higher than that of their control sample. We also find that Sy2 galaxies demonstrate the same behaviour with BIRG galaxies but not with Sy1s which do not show any excess of companions with respect to their control sample galaxies. An additional analysis of the relation between FIR colors and activity type of the BIRG's shows a significant difference between the colors of strongly-interacting and non-interacting starbursts and a resemblance between the colors of non-interacting starbursts and Sy2s. Our results support an evolutionary scenario leading from Starbursting to a Sy2 and finally to an unobscured Sy1 galaxy, where close interactions play the role of the triggering mechanism. ",
        "introduction": "Despite the numerous studies conducted over the last decade investigating the relation among interacting galaxies, starbursting and nuclear activity (eg. Hernandez-Toledo et al. 2001;  Ho 2005, Gonz\\'alez et al. 2008), the correlation between these physical processes remain uncertain. However, there is evidence that AGN galaxies host a post-starburst stellar population (eg. Gonz\\'alez Delgado 2001, Kaviraj 2008, Muller Sanchez 2008) while Kauffmann et al. (2003) showed that the fraction of post-starburst stars increases with AGN emission. Proving such a relation would simultaneously solve the problem of the AGN triggering mechanism. Interactions would be the main cause of such activities, being starbursting and/or the feeding of a central black hole. Indeed some studies seem to conclude that there is an evolutionary sequence from starburst to type 2 and then to type 1 AGN galaxies (e.g. Oliva et al. 1999, Storchi-Bergmann et al. 2001, Krongold et al. 2002, Chatzichristou 2002). In addition, Kim, Ho and Im (2005) using the [OII] emission line as a SFR indicator, reach the conclusion that type 2 are the precursors of type 1 quasars supporting the previous claims. Finally, a study by Ya-Wen Tang et al. (2008) tracing interactions by HI imaging, showed that 94\\% of their Seyfert sample exhibit HI disturbances in contrast with their inactive counterpart (15\\%). The above results raise doubts about the simplest version of the unification scheme of AGNs. ",
        "conclusions": "\\begin{figure}[t] \\resizebox{\\hsize}{!}{\\includegraphics[clip=true]{colorsbw.ps}} \\caption{ \\footnotesize FIR color-color diagram: $\\alpha(60,25)$ versus $\\alpha(25,12)$ : Starbursts are represented by circles, Sy2s by triangles, Liners by squares and normal galaxies by crosses. Highly interacting galaxies are represented by filled shapes, weakly interacting by open dotted shapes and non-interacting by open shapes.} \\end{figure}   Our results can be accommodated in a simple evolutionary scenario, starting with an interaction, and ending in a Sy1. First, close interactions would drive molecular clouds towards the central area, creating a circumnuclear starburst. Then, material could fall even further into the innermost regions of the galaxy, feeding the black hole, and giving birth to an AGN which at first cannot be observed due to obscuration. At this stage only a starburst would be observed. As starburst activity relaxes and obscuration decreases, a Sy2 nucleus would be revealed (still obscured by the molecular clouds from all viewing angles). As a final stage, a Sy1 would appear. In this case, the molecular clouds, initially in a spheroidal distribution, could flatten and form a ``torus'' (as in the unification scheme for Seyferts). As more material is accreted, it is possible that the AGN strengthens driving away most of the obscuring clouds, and leaving a ``naked'' Sy1 nucleus.  In order to better understand the role of interactions in driving starburst and nuclear activity, we are in the process of completing a study of AGN and starburst manifestations in the nearest neighbors of the active galaxies in our samples. Preliminary results show that more than 70\\% of the close companions of both type of Seyfert galaxies show some kind of activity. Most companions are HII galaxies and only a few are Seyferts. Furthermore, the fraction O[III]/Hb is significally higher in all Sy2 companions indicating a higher ionization. The correlation between high ionization and galaxy interactions is relatively well established for HII galaxies (eg. Woods, Geller \\& Barton 2007), although Alonso et al. (2007) presents such a connection also for AGNs."
    },
    "0807/0807.3891_arXiv.txt": {
        "abstract": "The IceCube neutrino observatory, under construction at the South Pole, consists of three sub-detectors: a km-scale array of digital optical modules deployed deep in the ice, the AMANDA neutrino telescope and the surface array IceTop. We summarize results from searches for cosmic neutrinos with the AMANDA telescope and review expected sensitivities for IceCube at various installation phases. Reliability and robustness of installation at the South Pole has been demonstrated during the past four successful construction seasons. The 40 installed IceCube strings are working well. We are developing detailed plans for the final construction of IceCube, including extensions optimized for low and high energy. We describe the IceCube Deep Core project which will extend the low energy response of IceCube. ",
        "introduction": "\\label{sec:intro}  IceCube, the largest neutrino telescope in history is rapidly becoming  a reality after four successful construction seasons. IceCube searches primarily for neutrinos in the GeV-PeV energy region expected to be produced at acceleration sites of high energy cosmic rays \\cite{physics} \\cite{ic_sensi}. The IceCube design \\cite{ic_design} employs 80 uniformly spaced strings, each one equipped with 60 digital optical modules (DOMs) \\cite{dom}. Among the strings of this \"in-ice\" array, 40 have already been deployed at depths between 1.5 and 2.5~km in the icecap of the South Pole. Moreover, an additional installation in the center-bottom part of the detector of an inner core array (the so-called IceCube Deep Core) composed of additional strings has been financed and recently approved. The AMANDA-II neutrino telescope \\cite{amanda} - the predecessor of IceCube - is maintained in operation. AMANDA-II is completely embedded in the IceCube instrumented volume, and the hits registered by AMANDA and IceCube are merged together into combined events \\cite{ag}. The third fundamental component of the neutrino observatory is also under construction: the IceTop air shower detector \\cite{icetop}. Combining IceTop data with data from the deep-ice array provides a unique opportunity to study cosmic ray composition in the region of the \"knee\" extending earlier measurements performed using the combination of the SPASE and AMANDA detectors \\cite{icetop_p} \\cite{Bai:2007pu}. The PMT pulses of the IceCube DOMs are first converted into digital waveforms. Every time the DOM  triggers, the waveform is digitized, read out and time stamped. The digitized PMT pulses from different DOMs are sorted into a time-ordered stream. The trigger is accomplished entirely via software. Single majority triggering (SMT) is applied in the in-ice array and in the IceTop array. For the in-ice, it requires the coincidence of hits in eight or more DOMs within a time window of $5 \\mu$sec and for IceTop six or more DOMs hit in a time window of the same size. Details about the data acquisition can be found in \\cite{kael}. There are two primary detection channels considered for the study of high energy neutrinos: {\\it track-like events} from upward through-going muons induced by charged-current $\\nu_{\\mu}$ interactions and {\\it cascade-like events} from charged-current (CC) interactions of $\\nu_e$, $\\nu_{\\tau}$ and neutral-current interactions of neutrinos of all flavor. A third channel concerns composite events  in which tracks and cascades from high energy CC interactions are observed together \\cite{tau}. Three {\\it observables} are available for each detection channel: the time of an event which is known with a relative time resolution \\cite{time_res}, the incoming direction characterized by the angular resolution \\cite{ic_sensi} and finally the estimated energy of the event characterized by the energy resolution. Average values of these observables for the track-like and cascade-like detection channels are reported in Table~\\ref{obs}.  \\begin{table} \\begin{tabular*}{0.45\\textwidth}{@{\\extracolsep{\\fill}}  l|cc  } \\hline\\hline Track-like Events & IceCube & AMANDA \\\\ \\hline Rel. Time Resolution [nsec] & 2 & 5-7\\\\ Angular Resolution  & $<1^\\circ$ & $2^\\circ-3^\\circ$\\\\ Energy Resolution (log$_{10}E$) & 0.3-0.4 & 0.3-0.4 \\\\ Field of view & $2\\pi$  & $2\\pi$\\\\ \\hline Cascade-like Events & &\\\\ \\hline Rel. Time Resolution [nsec] & 2 & 5-7\\\\ Energy Resolution ($log_{10}E$)& 0.18 & 0.18\\\\ Field of view & $4\\pi$ & $4\\pi$\\\\ \\hline\\hline \\end{tabular*} \\caption{Best values obtained for AMANDA and expected for IceCube observables.} \\label{obs} \\end{table}  The first level processing sorts events by detection channel or by specific low level analysis channel. Successive higher level processing are progressively more specific for various physics analyses. For example, track-like events are selected for the search of neutrinos from point-like sources \\cite{ps_5y}, from Gamma-Ray-Bursts (GRBs) \\cite{grb_mu}, from annihilation of WIMPs in the Sun \\cite{wimp-mu} and from a diffuse flux \\cite{diffuse_mu}. Cascades-like events  are considered for the search of a diffuse neutrino flux \\cite{diffuse-oxana} and as an additional channel for neutrinos from GRB explosions \\cite{grb_diffuse}. The data taking of the IceCube observatory started in 2006 with 9 strings (IC-9) and continued in 2007 with 22 strings (IC-22). The new physics run with 40 strings started in April 2008 and will continue until March 2009. The data taking demonstrated excellent stability and for the first time, operated successfully during drilling operations. As a result, the overall live time of IC-22 has been maximized. The procedures for installation, drilling operation and detector operation have been consolidated during the four deployment seasons, delivering a stable detector that performs as expected or even better without significant problems. ",
        "conclusions": ""
    },
    "0807/0807.4155_arXiv.txt": {
        "abstract": "We present near and mid-infrared observations of the pulsar-wind nebula (PWN) B0540-69.3 and its associated supernova remnant made with the {\\it Spitzer Space Telescope}. We report detections of the PWN with all four IRAC bands, the 24 $\\mu$m band of MIPS, and the Infrared Spectrograph (IRS).  We find no evidence of IR emission from the X-ray/radio shell surrounding the PWN resulting from the forward shock of the supernova blast wave. The flux of the PWN itself is dominated by synchrotron emission at shorter (IRAC) wavelengths, with a warm dust component longward of 20 $\\mu$m. We show that this dust continuum can be explained by a small amount ($\\sim 1-3 \\times 10^{-3} \\msun$) of dust at a temperature of $\\sim 50-65$ K, heated by the shock wave generated by the PWN being driven into the inner edge of the ejecta. This is evidently dust synthesized in the supernova.  We also report the detection of several lines in the spectrum of the PWN, and present kinematic information about the PWN as determined from these lines. Kinematics are consistent with previous optical studies of this object. Line strengths are also broadly consistent with what one expects from optical line strengths. We find that lines arise from slow ($\\sim 20$ km s$^{-1}$) shocks driven into oxygen-rich clumps in the shell swept-up by an iron-nickel bubble, which have a density contrast of $\\sim 100-200$ relative to the bulk of the ejecta, and that faster shocks ($\\sim 250$ km s$^{-1}$) in the hydrogen envelope are required to heat dust grains to observed temperatures. We infer from estimates of heavy-element ejecta abundances that the progenitor star was likely in the range of 20-25 $M_\\odot$. ",
        "introduction": "\\label{intro}  Many core-collapse supernovae (SNe) leave behind a neutron star as a compact remnant.  Some of these neutron stars are active pulsars which inflate a bubble of relativistic particles and magnetic fields confined by the ejecta or interstellar medium (ISM), known as a pulsar-wind nebula.  The combination of a shell supernova remnant (SNR) and associated pulsar-wind nebula can allow the investigation of various issues of importance in supernova and pulsar physics, including pulsar kicks, ejecta structure and composition, and particle acceleration at relativistic shocks.  Pulsar-wind nebulae serve as calorimeters for pulsar spindown energy loss, and as test systems to study the behavior of relativistic shocks where the pulsar wind is thermalized. We know of few cases of a ``normal'' radio and X-ray shell supernova remnant containing an active pulsar and synchrotron nebula.  Probably the best known such ideal case is the Large Magellanic Cloud remnant B0540-69.3 (or ``0540'' for short). 0540 is also one of a highly exclusive group of ``oxygen-rich'' SNRs, a group that includes Cas A, Puppis A, G292+1.8, 1E0102-72.3, and N132D.  Theoretical studies of PWNe have either concentrated on the gross evolution, assuming a homogeneous nebula \\citep{rees74,pacini73,reynolds84} or the detailed spatial structure, neglecting evolution \\citep{kennel84}.  Since the advent of the new generation of X-ray observatories, the study of PWNe has accelerated, with the identification of many new objects and more detailed information on known ones (see Gaensler \\& Slane 2006 for a recent review).  \\cite{chevalier05} modeled PWNe for different assumptions about the ejecta profiles into which they expand, to relate properties of supernovae to those of the PWNe.  PWNe produce extremely broad-band spectral-energy distributions (SEDs), well described in various frequency regimes with power laws. Most PWNe are observed in radio and X-rays; only a few are detected optically (here as in many other ways the Crab Nebula is an exception), and almost nothing is known about infrared or ultraviolet spectra.  Typical radio spectra are featureless, and are well described by power-laws with spectral indices $\\alpha < 0.3$ ($S_\\nu \\propto \\nu^{-\\alpha}$), with X-ray indices steeper by 0.5 -- 1.3 (see data in Chevalier 2005).  Since simple models of synchrotron losses predict a steepening of exactly 0.5, they lack some essential physics, which may be constrained if the complete spectrum is known.  Galactic PWNe are all found close to the Galactic plane, where they suffer from extinction in optical and UV and confusion in IR.  Filling in the SED between radio and X-rays can best be done with a high-latitude object. For this reason as for many others, 0540 is an interesting target.  0540 was first catalogued as a radio source of unknown nature, a minor feature on a 408 MHz map of the 30 Dor region made with the Molonglo telescope \\citep{lemarne68}.  \\cite{mathewson73} first classified it as a supernova remnant on the basis of its steep radio spectrum, although their optical survey did not detect it.  Early reports associated 0540 with the H$\\alpha$ emission nebula N 158A \\citep{henize56}, though that object is $3'$ from the centroid of the early radio positions (which could be localized to better than $10''$).  The absence of strong H$\\alpha$ emission from 0540 further demonstrates that the association with N 158A is erroneous. Subsequent radio observations \\citep{milne80} gave an improved spectral index of $-0.44$, typical for a shell supernova remnant.  The first indication of something unusual was the X-ray detection \\citep{long79} with the {\\sl Einstein} Observatory, in which 0540 was the third brightest X-ray remnant in the LMC.  The X-ray spectrum was shown to be featureless by \\cite{clark82} with the {\\sl Einstein} Solid-State Spectrometer.  The first optical detection was reported by Mathewson et al.~(1980), motivated by pre-publication reports of the observations of Clark et al~(1982).  Mathewson et al.~did not see H$\\alpha$ but instead a spectacular ring in [O III] of $8''$ diameter, with a smaller ring in fainter [N II] emission ($4''$ diameter) and no appreciable Balmer emission. In addition to classifying 0540 as an ``oxygen-rich'' SNR, \\cite{mathewson80} also reported spectroscopic observations indicating expansion speeds of order 1500 km s$^{-1}$. The discovery of the 50 ms X-ray pulsar \\citep{seward84} and optical synchrotron nebula \\citep{chanan84} added to the complexity and interest of the system.  The optical emission was shown definitively to be synchrotron by the discovery of polarization \\citep{chanan90}. The pulsar spindown timescale $P/2 {\\dot P}$ is about 1660 yr \\citep{seward84}, somewhat longer than the kinematic age estimate resulting from dividing the radius ($4'' = 1$ pc at our assumed distance of 50 kpc) by the expansion speed of about 1500 km s$^{-1}$, which yields a value of $\\sim 700$ yr.  The pulsar spindown luminosity is $1.5 \\times 10^{38}$ erg s$^{-1}$.  \\cite{reynolds85} modeled 0540 with the formalism of \\cite{reynolds84}, with the pulsar driving an accelerating synchrotron nebula into the inner edge of expanding ejecta.  At that time, there were no more than hints of extended structure that could be identified with the outer blast wave.  \\cite{reynolds85} found that the current radio, optical, and X-ray observations could be explained without requiring extreme values for the pulsar initial energy or other parameters.  He deduced an initial pulsar period of about $30 \\pm 8$ ms, that is, relatively slow, and concluded that the true age of 0540 was between 800 and 1100 yr, somewhat longer than the kinematic age due to the pulsar-driven acceleration.  Up to this time, all observations were consistent with 0540 being a standard Crab-like remnant (i.e., a nonthermal center-brightened radio and X-ray nebula surrounding a pulsar), except for the hint of larger-scale structure from radio images and from X-ray observations \\citep{seward84}.  Definitive information on the structure came from higher-resolution radio observations with the Australia Telescope \\citep{manchester93} which showed a clear radio shell with diameter about $65''$ surrounding a radio nebula with size (about $5''$ FWHM) comparable to the bright X-ray nebula and [O III] ring.  The shell has a radio spectral index $\\alpha$ of about $-0.4$, while the central nebula has $\\alpha = -0.25$.  At this point it was clear that 0540 is even more Crab-like than the Crab, as it possesses a clear outer blast wave interacting with surrounding material, so that we could be sure that the interior PWN is interacting with the inner SN ejecta as in Reynolds \\& Chevalier (1984).  X-ray emission from the blast wave was confirmed with {\\it Chandra} observations \\citep{hwang01}; the emission is brightest in the W and SW, like the radio.  Spectral fits indicated abundances typical of the LMC, with a temperature of order 4 keV (for a Sedov blast wave model) and ionization timescale $\\tau \\equiv n_e t = 3.7 \\times 10^{10}$ cm$^{-3}$ s, though spectral differences are apparent in different regions, and a hard component may be called for.  The most thorough optical spectroscopic study to date was reported by \\cite{kirshner89}.  They confirmed the high velocities (FWZI $\\sim 2800$ km s$^{-1}$), and reported weak H$\\alpha$ emission.  The average centroid of SNR lines (as opposed to narrower lines from a nearby H II region) was shifted by $+370$ km s$^{-1}$.  No [Ne III] was reported ($< 1.5$ \\% of [O III]); they concluded that this was a real abundance deficit rather than a temperature or density effect.  A detailed study by \\cite{serafimovich04}, focusing on the optical nonthermal continuum, revised the reddening and optical slope to give a power-law index in the optical of $\\alpha_o = -1.07.$ Recent observations by \\cite{morse06} report the discovery of faint [O III] emission extending to a radius of $8''$, with a velocity of 1650 km s$^{-1}$. They find the centroid of this velocity component to be the same as that of the LMC, so that a large peculiar velocity of the system is not required.  \\cite{chevalier05} modeled 0540, along with several other PWNe, with the goal of learning more about the SN explosion.  He obtained several results for a simple dynamical model of a PWN expanding into ejecta of various density profiles driven by a pulsar of given power. He interpreted 0540 as the result of a SN Ib/c, an exploding Wolf-Rayet star, with the prediction of a lack of significant emission from hydrogen. However, recent optical observations by \\citet{serafimovich04} and \\citet{morse06} have detected hydrogen. In light of this, it is now believed \\citep{chevalier06} that 0540 is the result of a type IIP supernova.  The infrared observations of 0540, which was detected by the {\\it Infrared Space Observatory} (ISO) \\citep{gallant99}, presented in this paper promise to advance our understanding on several fronts. The outline of our paper is as follows: In section 2, we describe the observations and data reduction, and results are given in section 3. In section 4.1, we discuss a general picture of the PWN, and sections 4.2 and 4.3 discuss in detail the line emission and dust continuum emission, respectively. In section 4.4, we discuss the origin of the O-rich clumps, whose existence we posit in section 4.2. Section 5 serves as a summary of our findings. ",
        "conclusions": "\\label{disc}  \\subsection{General Picture}  We aim at a self-consistent, semi-quantitative picture of the PWN that accounts for the presence of lines (optical and IR), the extent of the synchrotron nebula, and the source of the [O III] emission at $8''$ radius. We find it useful to first point out some contrasts between 0540 and the most widely-known object of its class, the Crab Nebula. Although 0540 has been referred to as ``The Crab's Twin,'' the two differ in some important ways. The most obvious difference is the lack of an outer shell in the Crab, while 0540's $30''$ shell has been seen in both radio and X-ray observations. For the purposes of this paper, however, the important differences lie in the PWN. In the Crab, the size of the nebula decreases with increasing frequency, so that the radio nebula is larger than the optical, which is larger than the X-ray, etc. In 0540, the synchrotron nebula is approximately identical in extent throughout all wavelengths, around $5''$. The other fundamental difference is the presence in 0540 of emission located beyond the synchrotron nebula (the [O III] halo). There is nothing like this seen in the Crab, where the radio synchrotron emission extends to the outer boundary of anything known to be associated with the nebula.  In modeling the Crab Nebula, \\cite{sankrit97} considered two models, one a pure shock model and the other a pure photoionization model to explain the optical emission. They concluded that shocks from an expanding shell were more likely. In the case of 0540, however, a pure shock model cannot reproduce the [O III] extended emission. We therefore propose an extension to their models that incorporates both a global shock {\\it and} photoionization. The specifics of our model will be described more fully in the sections below, but our general picture of the nebula is as follows.  It is based on the dynamical picture of Chevalier (2005; C05).  Approximately a millenium ago, a star exploded via the core-collapse mechanism, leaving behind a pulsar, and sending a shock wave out into the interstellar medium. The outer boundary of this forward shock is now about 8 pc (angular distance of about $30''$) from the pulsar, and the reverse shock into the ejecta is somewhere between $10'' - 30''$, having not yet reached back to the inner ejecta. The pulsar has since formed a pulsar-wind nebula, which itself is driving a shock into the inner edge of the surrounding ejecta, which are in free expansion. The shock wave heats the inner ejecta and sweeps them into a thin shell. Since the shell of material is being continuously injected with energy from the pulsar, it is accelerating and overtaking less dense material as it expands. The shock speed relative to upstream material, however, reaches a maximum and then begins to drop since the free-expansion speed of the ejecta material is also higher at larger radii. There is no reason, however, to expect the ejecta to be completely homogeneous. The $^{56}$Ni synthesized in the explosion will have heated the central ejecta by radioactive decay, causing them to expand in an ``iron-nickel bubble'' \\citep{li93}, and compressing intermediate-mass ejecta into a denser surrounding shell.  We propose that the PWN shock has reached a radius of about 1.2 pc from the pulsar, which corresponds to a size of $\\sim 5''$, the size of the nebula as determined by X-ray observations. The layer of shocked ejecta is geometrically thin, bounded on the inside by a contact discontinuity separating it from the the PWN proper, which is the shocked pulsar wind. The shock has already encountered and propagated through the low-density iron-nickel bubble and its surrounding shell. That shell is likely to be highly clumpy \\citep{basko94}; shocks driven into the clumps of heavy-element ejecta will be slow. Finally, at a sub-arcsecond radius we expect the inward-facing pulsar wind shock where the relativistic pulsar wind is thermalized. Interior to the shock driven into the ejecta, emission in optical/IR is both thermal and non-thermal, with the dominant component being synchrotron continuum emission from the relativistic electrons.  However, multiple emission lines are clearly detected from dense clumps and filaments of thermal gas.  In addition to this, we identify a rising continuum in the mid-infrared above the synchrotron continuum that we interpret as a small amount of warm dust, collisionally heated by electrons heated by the shock. Most lines seen in optical and infrared then come from dense clumps of ejecta, where the shock wave has slowed significantly and become highly radiative.  What remains is to explain the faint [O III] emission seen at $8''$. We propose that this is material that is still in free expansion, i.e. unshocked, that has been photoionized by ultraviolet photons emitted from the shockwave. The source of photoionization is two-fold; ultraviolet photons from the synchrotron nebula and those produced in fast radiative shocks both contribute appreciable amounts of ionizing radiation. We show below that to within a factor of 2, there are enough ionizing photons produced to account for the [O III] halo at $8''$.  We have included, in Figure~\\ref{cartoon}, a cartoon sketch of this picture, which will be further discussed in the following sections. A factor of a few is all we expect to be able to accomplish in modeling the nebula, due to the large uncertainties involved.  These uncertainties include, but are not limited to; nature of the progenitor star (which affects the post-explosion density distribution of the ejecta), heavy element abundances in the ejecta, degree of clumping of the ejecta, etc. We have endeavored in the following sections to point out places where uncertainties arise, and where possible, to assign quantitative values to them.  \\subsubsection{PWN Model}  C05 discusses a model, based on a thin-shell approximation, for a pulsar wind nebula interacting with an inner supernova ejecta density profile. We have used this model along with our observations to determine various quantities about 0540, including how much hot gas should be present. Observable quantities for the pulsar include period ($P$), period derivative ($\\dot P$), luminosity ($\\dot E$) and for the nebula, size ($R$) and shell velocity ($V_{sh}$). While the quantities for the pulsar are fairly well established by previous observations, those for the PWN are much more uncertain. Previous optical studies of the remnant \\citep{mathewson80, kirshner89}, as well as radio observations \\citep{dickel02} interpreted the PWN as a bubble of radius $\\sim 4''$, and the optical observations gave expansion velocities less than 1400 km s$^{-1}$. Based on {\\it Chandra} observations, \\citet{petre07} concluded the nebula was slightly larger, with a radius of $\\sim 5''$. We shall adopt $5''$ (1.2 pc) as an estimate of the location of the ejecta shock.  However, \\citet{morse06} detected faint [O III] emission in images extending out to a radius of $\\sim 8''$. Based on similarities to the Crab Nebula, they interpreted this [O III] halo as being the outer edge of the shock from the pulsar wind overtaking the slower moving ejecta. Here we present an alternative interpretation of this [O III] halo emission. We propose that it is undecelerated ejecta that have been photoionized, rather than shock-ionized. The FWZI of the [O III] emission from Morse et al. was 3300 km s$^{-1}$, which, given the extent of 1.8 pc and our interpretation of this as undecelerated ejecta, provides the remnant age of 1140 years. While this is somewhat longer than the favored model of \\citet{reynolds85}, it is at least reasonable given other age estimates made by previous studies of the object. Photoionization calculations are discussed in Appendix A.  As a first attempt to model the observations, we considered the case of a spherically symmetric shock wave driven by the energy input from the pulsar expanding into a medium with density profile described by $\\rho_{SN} = At^{-3}(r/t)^{-m}$. We considered different values of the parameter $m$, as dynamical mixing between the ejecta and surrounding medium would produce a complicated density structure. The swept-up mass does not exceed $1 \\msun$ in this model. Although this model did a reasonably good job at producing shock speeds high enough to account for the necessary dust grain heating rate, a spherically symmetric model does not adequately reproduce line radiation observed in both optical and IR. A slow shock into dense material is required to explain these lines, and the spherical model cannot account for this, since presence of lines requires a departure from the overall homogeneous density profile. We present the spherically symmetric calculations in Appendix B. A more robust model is required to explain both the slow shocks required for lines and the faster shocks required for dust emission. We will return to this picture at the end of the following section, but we must first describe our line observations in detail.  \\subsection{Lines} \\label{lines}  Eight emission lines are detected in the Spitzer spectrum.  They provide constraints on the density and temperature of the emitting gas, and perhaps more importantly on the elemental abundances.  They complement the optical spectra published by \\citet{kirshner89} (K89), \\citet{morse06} (M06) and \\citet{serafimovich05} (S05). We first summarize the implications of the optical spectra, then consider shock wave models for the combined optical and IR emission.  Several temperature estimates are available from the optical spectra. The [O III] line ratio I(4363)/I(5007) gives temperatures of about 24,000 K according to S05 or 34,000 K (K89).  According to the CHIANTI database \\citep{landi06}, the ratio given by K89 corresponds to 50,000 K, while that given in M06 implies 24,000 K.  K89 also find temperatures $>$30,000 K from the [O II] I(7325)/I(3727) ratio and $<$10,000 K from [S II] I(4072)/I(6723).  The [S II] ratio of M06 implies $T$ = 14,000 K.  Assuming a temperature of 10,000 K, S05 find a density of 1400 -- 4300 $\\rm cm^{-3}$, and at 14,000 K the range would be 1700 -- 5000 $\\rm cm^{-3}$.  The differences among the various temperature estimates may result partly from different reddening corrections and different slit positions, but it is clear that the [O III] lines are formed in hotter gas than the [S II] lines. The Spitzer data include only one pair of lines from a single ion, [Fe II] I(17.9$\\mu$)/I(26.0$\\mu$), which is constrained to be larger than 1.13.  Again using CHIANTI, this requires a density above about 5000 $\\rm cm^{-3}$ and a temperature above 4000 K.  However, the ratio depends upon the deblending of the [Fe II] and [O IV] lines at 26$\\mu$m, and the uncertainty may be larger than the formal value. The density contrast between the optically emitting material and the mean post-shock density from the global model indicates that as in the Crab Nebula, optically emitting material is concentrated in dense knots and/or filaments.  The next step in interpreting the spectra and constructing models is to estimate the relative importance of photoionization and shock heating.  In the Crab nebula, photoionization dominates, though shocks are important for the UV lines produced at higher temperatures and for compressing the gas to increase the optical emissivity \\citep{sankrit97}.  In the oxygen-rich SNRs, such as N132D and 1E0102-7219, shock heating dominates \\citep{blair00}.  0540 shows both synchrotron emission reminiscent of the Crab and extreme heavy element enhancement.  A pure photoionization model with strongly enhanced abundances and the observed density gives too low a temperature to account for the [O II], [O III] and [S II] line ratios, while shock models cool so rapidly that they produce little [O I] or [S II] unless they produce no [O III] at all. Therefore, it seems likely that a model of a shock including the PWN ionizing radiation is needed.  We have computed models with the shock model described in \\citet{blair00} illuminated by the power law continuum described by \\citet{serafimovich04}.  Briefly, the code is similar to that of \\citet{raymond79} and \\citet{cox85}, but it has been modified to describe SNR ejecta with little or no hydrogen.  The most important difference is that the cooling rate is enormous, so that the electron temperature is well below the ion temperature in the hotter parts of the flow.  The model is similar to those of \\citet{itoh81} and \\citet{sutherland95}.  Unlike those models, we do not include the photoionization precursor of the shock, because the ionizing emission from the shock is considerably weaker than the ambient synchrotron radiation.  In comparison with the spectra of Cas A, N123D and 1E0102-7219, shock models have the problems that no single shock model produces the observed range of ionization states, and that they tend to predict too much emission in the O I 7774 \\AA\\/ recombination line unless the cooling region is somewhat arbitrarily truncated (Itoh 1988).  However, they do predict reasonable relative intensities from the UV to the near IR for O III and O II. Below, we attribute the truncation to mixing with hotter, lower density gas.  We assume a 20 $\\rm km~s^{-1}$ shock with a pre-shock density of 30 $\\rm cm^{-3}$, which produces a density of around 5000 $\\rm cm^{-3}$ where the [S II] lines are formed.  The elemental abundances are O: Ne: Mg: Si: S: Ar: Ca: Fe = 1: 0.2: 0.1: 0.1: 0.1: 0.1: 0.1: 0.1 by number.  H, He and N are not included in the model, because it seems likely that the lines from these elements arise in some other gas, like either the quasi-stationary flocculi or the outer shell of ejecta seen as very fast knots in Cas A \\citep{kirshner77,fesen01}.  The normalization of the power law flux assumes that the shocked gas is 1 pc from the center of the PWN.  The list of caveats is long.  There is undoubtedly a range of shock speeds and pre-shock densities.  The shocked gas is unlikely to be a uniform mixture of the various elements, and large variations in the composition among different clumps, as observed in Cas A, are likely. There may well be a significant contribution from unshocked photoionized gas for some lines \\citep[e.g.,][]{blair89}, as we shall argue below for the [O III] halo.  The shock models are plane parallel, with the power law illumination incident from the PWN, while the X-rays are more likely to illuminate the shocked gas from behind.  The models terminate somewhat arbitrarily at 250 K because of numerical limitations.  This will affect the IR lines and the O I recombination line at 7774 \\AA .  Also, as a compromise between energy resolution and energy range of the ionizing radiation, the power law only extends to 2 keV.  This means that the inner shell ionization and Auger ionization of S and Fe is not included.  Finally, the atomic data in the code are somewhat out of date and need to be updated. Nevertheless, the code gives a reasonable idea of the relative line intensities.  To compare this model with our observed IR spectra, we must place the Spitzer spectrum on the same scale as the optical spectra.  We normalize the IR lines to [O III] 5007 = 100 by dividing the Spitzer intensities by 4 times the [O III] 5007 intensity given by M06.  The factor of 4 is meant to account for the fact that the $2''$ slit used by M06 covers only about 1/4 of the remnant.  This is obviously not a very accurate correction, but it is probably good to a factor of 2. Since the relative fluxes of many optical lines differ by a factor of 2 between M06 and K89, this is unfortunately the best we can do until an optical spectrum of the entire remnant becomes available.  The result is shown in Table 3.  Overall, the agreement is astonishingly good for such a simple model.  Several of the low ionization lines, [O I], [Ne II] and [Si II] are underpredicted, though the [Si II] line could be increased simply by increasing the silicon abundance.  The ratio of the [S IV] to [S III] IR lines is too low, but inclusion of the harder part of the power law spectrum would improve that.  Inclusion of the harder X-rays would also increase the intensity of the [Fe VII] line, though a lower pre-shock density or a higher shock speed would have the same effect. The oxygen column density of the model is only about $10^{14}~\\rm cm^{-2}$, and the thickness of the emitting region is only $6 \\times 10^{11}$ cm.  If the thickness were large enough to allow the remaining O$^+$ to recombine, the predicted O I recombination line, which is comfortably lower than the weakest detected lines, would increase to about 4 times the apparent detection limit of K89.  The agreement would improve if the argon abundance were cut in half, but otherwise the abundances appear to match the observations.  The shock model shown in Table 2 produces $1.1 \\times 10^{-14} ~\\rm erg~cm^{-2}~s^{-1}$ in the [Ne III] 15.5 micron emission line, so the flux shown in Table 1 would come from a region with surface area of $2 \\times 10^{38}$ cm$^{2}$. This area is roughly equal to the area of a sphere of $5''$ radius but the emission could come from many smaller volumes with a total filling fraction of a few percent. Heavy-element ejecta enter these slow shocks at a rate of $\\sim 0.01 \\msun$ yr$^{-1}$.  We conclude that, as in the Crab \\citep{sankrit97}, the observations can be explained by shocks that heat and compress the gas in the radiation field of the PWN.  The shock heating seems needed to reach the high temperatures seen in some line ratios and to provide the high densities observed, while the photoionization heating strengthens the low and moderate ionization lines.  Oxygen is about ten times as abundant as the other elements.  The shock speed and pre-shock density are not very well constrained, but a shock as fast as 80 $\\rm km~s^{-1}$ requires a low pre-shock density to match the observed density, and that in turn implies a very high pre-shock ionization state, overly strong [O IV] and overly weak [O II] emission.  In the last 100 years alone, about 1 $\\msun$ of heavy-element ejecta have been shocked, more than the total mass of the swept-up ejecta in the global model described in Appendix B. This casts some doubt on the validity of the global, spherically symmetric model, where density within freely expanding ejecta was assumed to be a smooth power-law. It is possible that the innermost ejecta have been swept-up by an iron-nickel bubble, as inferred for SN 1987A by \\cite{li93} and modeled by \\cite{basko94} and \\cite{wang05} (see also brief discussion in C05). We explored the possibility that the global shock could be contained within the shell swept-up by the iron-nickel bubble. In this one-dimensional picture, the shock passed through the inner, low-density region in $\\sim 50$ years, and has since been contained within the high-density ($n \\sim 30$ cm$^{-3}$) bubble wall. We varied parameters of the model until the shock speed in the bubble was approximately 20 km s$^{-1}$, as required by line models. We find, however, that the mass flux of material entering the shock throughout the remnant's entire lifetime has been unreasonably high for this model, approximately $0.01 (t/1140$ yr$)^{-1/2} \\msun$ yr$^{-1}$. In addition, a 20 km s$^{-1}$ shock, even at such density, would not adequately heat dust grains to temperatures observed. Dust heating is discussed in more detail in section~\\ref{grains}.  We are forced to consider inhomogeneous ejecta with a fast global shock to heat dust to observed temperatures, and slower shocks producing observed line emission. We propose the following picture: The shock swept through the low-density iron-nickel bubble interior early in the life of the SNR. It then encountered the dense, clumpy shell of the bubble, slowing down and further fragmenting the shell into dense clumps, which are still being overrun by slow shocks, currently 20 km s$^{-1}$. The global shock has now exited the iron-nickel bubble shell, and is propagating through the ejecta with relatively low ambient density. The speed of this shock is not well known, but 250 km s$^{-1}$ would be sufficient to heat the dust to the observed temperature of around 50 K (see below). Assuming pressure equilibrium between the dense clumps and the ambient ejecta, we derive a density contrast, given the difference in shock velocities, of $\\sim 150$. Support for this model can be inferred from {\\it HST} images of the nebula, as seen in Figure~\\ref{hstimage}, which shows [O III] filaments in the interior, not just in a shell. [O III] line profiles (M06) also do not match the shape that would be expected from a spherically symmetric expanding shell, i.e. a flat top. The slow shocks driven into the dense clumps are in rough pressure equilibrium with the fast shock driven into the less dense ejecta.  \\subsubsection{Progenitor Mass} \\label{progenitor}  We have compared the abundances of heavy elements listed in section~\\ref{lines} with the predicted abundances of \\cite{woosley95}, who consider abundance yields from core-collapse SNe ranging in mass from 11-40 $M_\\odot$ and metallicities between zero and solar. We consider models with both solar and 0.1 solar metallicity, as this range is most likely to reflect a massive star in the LMC. Our abundances listed are somewhat uncertain, and result from fits to optical and infrared line strengths. We considered the ratios of O to Ne, Mg, Si, and Fe. The data do not single out a particular model from \\cite{woosley95}, but ratios of heavy elements to oxygen do favor a low-to-medium-mass progenitor. High-mass progenitors ($\\gapprox$ 30 $M_\\odot$) are less favored, since they produce larger amounts of oxygen relative to other elements. This interpretation is consistent with that of \\cite{chevalier06}, who favored a type IIP explosion for this object based on observations of hydrogen in the spectrum. This is also consistent with the idea that type IIP SNe should result from the explosion of a single star of 8-25 $M_\\odot$ \\citep{woosley02}.  It is possible to quantify these results even further. If we assume a constant heavy-element mass flux through the radiative shocks for $10^{3}$ yr of 0.01 $M_\\odot$ yr$^{-1}$ with our abundances listed above, we get a total ejected mass in oxygen of $\\sim 3.5 M_\\odot$, though this number should only be considered accurate to a factor of a few, and is likely an upper limit. When compared with predictions from models, this value favors stars in the range of 20-25 $M_\\odot$. \\cite{maeder92} gives slightly different abundance yields for SNe, with lower overall oxygen abundances produced. In his model, high-mass stars ($\\gapprox 25 M_\\odot$) actually produce less oxygen than their lower-mass counterparts, due to mass-loss of outer layers and inability to synthesize O from He and C. However, these massive stars would be Wolf-Rayet stars, and can be ruled out based on the detection of hydrogen in optical spectra.   \\subsection{Dust}  One of the more obvious features of the continuum in 0540 as seen in Figure~\\ref{ll_sync} is the excess of emission above the extrapolated radio synchrotron spectrum at longer wavelengths. A similar excess has been observed in the Crab \\citep{temim06}, and has been attributed to warm dust. We have inferred the temperature and the amount of dust present, and have examined several possible mechanisms for grain heating.  In order to fit the long-wavelength excess above the continuum, it was necessary to remove contributions from emission lines and the underlying synchrotron continuum. The flux contributed by the lines is negligible, but their presence makes fitting of a model dust spectrum more difficult. We have thus clipped obvious emission lines and bad pixels out of the spectrum for this analysis.  \\subsubsection{Synchrotron Component}  In order to subtract the synchrotron component, it was necessary to produce a model synchrotron spectrum that includes the break in power-law indices from optical to radio. The synchrotron model used here is one of a class of simple outflow models in which various quantities are allowed to have power-law dependencies on radius: flow-tube width, flow velocity, gas density (where mass loading might allow a range of possibilities), and magnetic-field strength (Reynolds, in preparation).  Such models can produce synchrotron-loss steepening in spectral index both steeper and flatter than the homogeneous-source value 0.5 \\citep{reynolds06}.  Here the model, used for illustrative purposes, invokes a simple outflow geometry with conical flow tubes (width $w \\propto r$), mass increasing as radius (due presumably to mass loading), flow velocity decreasing as $r^{-2}$ roughly, and magnetic field as $r^{-1}$.  The initial magnetic field at the injection radius is $B_0 = 2.5 \\times 10^{-4}$ G.  This model predicts a decrease in size with frequency as $\\theta \\propto \\nu^{-0.34}$, which might be slow enough to be consistent with observations, especially as it might take place along the line of sight.  While this is not meant as a definitive model for 0540, it describes the data well as shown on Figure~\\ref{bbspectrum} and was used to estimate the synchrotron contribution.  The result of radiative losses on electrons above the break energy in a flat ($N(E) \\propto E^{-s}$ with $s < 2$) energy distribution is for such electrons to move to just below the break energy, where they can produce a perceptible ``bump''.  However, the ``bump'' is almost undetectable unless $s$ is very close to 0; for the 0540 value $s = 1.5$, there is almost no departure from the power-law below the break frequency.  The model in Figure~\\ref{bbspectrum} was calculated including the redistribution of electron energies, and it can be seen that the excess we observe below 24 $\\mu$m cannot be attributed to this cause.  \\subsubsection{Fitting the Dust Component}  This left us with a residual rising continuum that we then fit with a model dust spectrum. Since we presume that the dust present in 0540 would be newly formed ejecta dust, as seen in SN 1987A \\citep{ercolano07}, we have little {\\it a priori} knowledge about the grain-size distribution. However, since the wavelength of IR radiation is much larger than typical ISM grain sizes, we adopt a model with a single grain size, arbitrarily chosen to be $a=0.05$ $\\mu$m in radius. In any case, in the limit of $a$ $\\ll \\lambda$, the results are independent of the choice of grain radius. We also do not know the grain composition, as general results from the LMC should not apply to ejecta dust. We thus consider two models; a graphite dust model and the ``astronomical silicates'' model of \\citet{draine84}. We calculate the dust grain absorption cross section for both as a function of wavelength. We then fit a simple modified blackbody model (incorporating the grain absorption cross-section) to the data using a least-squares algorithm designed for this model. We obtain a dust mass of $\\sim 3 \\times 10^{-3}$ $\\msun$ at a temperature of $50 \\pm 8$ K for silicate dust, while the resulting fit to the temperature with graphite grains was slightly higher, $\\sim 65$ K, and the required dust mass was lower, $\\sim 1 \\times 10^{-3}$.  The errors on the dust temperature are estimates based on using different methods of removing lines and subtracting the background and the underlying synchrotron spectrum. The resultant dust spectrum is sensitive to these details. The dust mass should be considered uncertain, and is probably only accurate to within a factor of a few, as evidenced by the difference between derived masses for graphite and silicate grains. Our data do not allow us to distinguish between various dust compositions. It should also be noted that we are only sensitive to dust that has been warmed by the shock wave from the pulsar wind, and that there could be more dust that has not yet been shocked, and is still too cool to be detected. Thus, our mass estimate should be considered a lower limit.  \\subsubsection{Grain Heating Mechanisms} \\label{grains}  We now turn our attention to heating mechanisms for this dust. We first consider heating by the synchrotron radiation field from the PWN. Since the spectrum of the synchrotron radiation is known in the optical/UV portion of the spectrum and grain absorption cross-sections can be calculated as a function of wavelength, it is possible to estimate whether there is enough radiation to heat the dust to temperatures observed. We calculate the optical depth of the dust around the PWN, and integrate over all wavelengths from radio up 1 keV. Although the flux from the PWN is higher at longer wavelengths, most of the absorption occurs in the optical/UV portion of the spectrum, due to the steeply rising absorption cross-sections in this regime. We compare this number to the luminosity in dust derived from our dust model, $\\sim 5 \\times 10^{35}$ ergs s$^{-1}$. A simple calculation showed that the radiation available falls short by several orders of magnitude of what is necessary.  However, this method tells us nothing about the total amount of dust that could be present to absorb the synchrotron radiation. Thus, to further test this hypothesis, we calculated the temperature to which dust would be heated if it were exposed to such an ultraviolet radiation field. We find that dust would only be heated to $\\sim{20}$ K. If this were the source of the emission seen in IRS, it would predict a 70 $\\mu$m flux that is several orders of magnitude higher than the upper limit we have placed on emission there. Given that these order of magnitude estimates are drastically inconsistent with our observations, we consider heating by photons from the PWN to be ruled out.  We then considered the somewhat more exotic possibility of the observed excess arising from a protoplanetary disk around the pulsar, unassociated with the nebula. It has long been known that planets can form around pulsars \\citep{wolszczan92}, and the supposition has been that these planets arise from a protoplanetary disk around the pulsar, the source of which has been attributed to several mechnisms \\citep{bryden06}. Various surveys of known pulsars have been made in infrared and submillimeter wavelengths, but for the most part these surveys have only produced upper limits on the dust emission present.  However, \\citet{wang06} conducted a survey of neutron stars with IRAC and found a debris disk around the young isolated neutron star 4U 0142+61. The authors suggest that the age of the debris disk compared to the spin-down age of the pulsar favors a supernova fallback origin. The IRAC observations combined with K-band Keck-I observations suggest a multi-temperature thermal model with temperatures ranging from 700-1200 K, where the disk has inner and outer radii of 2.9 and 9.7 $\\rsun$, respectively. Using the same model the authors use \\citep{vrtilek90}, we calculate the necessary radius to reproduce observed fluxes for 0540 for a disk with temperature $\\sim{50}$ K to be on the order of $10^{4}$ $\\rsun$. A survey of disks around Anomalous X-ray Pulsars (AXPs) \\citep{durant05} found several candidates for fallback disks which consistently had IR(K-band)/X-ray flux ratios of order $10^{-4}$. Although we were not able to find any archival near-infrared observations of the PWN, we can make an estimate of this ratio by looking at the overall spectrum of the IRAC and optical points. An estimate of $5 \\times 10^{-2}$ is reasonable for this ratio in 0540, significantly different than that found in the AXPs.  Additionally, \\citet{wang07} observed 3 known AXPs with Spitzer, and found no mid-IR counterpart to any of them. Given these discrepancies between these cases and that of 0540, we do not believe that a protoplanetary disk around the pulsar is the origin of the far-IR excess.  What then is the cause? Collisional heating by hot gas heated by shocks driven into the ejecta can provide both a qualitative and quantitative explanation for the dust present. Grain heating rate, $\\cal{H}$, goes as  \\begin{equation} \\mathcal{H} \\propto n_ev_{e}T_{e} \\propto PT_{e}^{1/2}, \\end{equation} where $n_{e}$, $v_{e}$, and $T_{e}$ are electron density, velocity, and temperature, and $P$ is the pressure, $nT$. In the PWN, P is constant throughout the bubble, so that grain heating is more efficient in hotter gas. We find that the slow, radiative shocks are incapable of heating dust grains to temperatures much above $\\sim 25$ K. Faster shocks, and thus higher temperatures, are required to heat grains to observed temperatures.  To determine whether this is plausible, given the conditions in the object, it is necessary to make an estimate of the amount of gas that is still hot, i.e. capable of heating dust grains through collisions with electrons. The shock cooling time \\citep{mckee87} is  \\begin{equation} t_{cool} = 2.5 \\times 10^{10} v^{3}_{s7}/\\alpha \\rho_{0}, \\end{equation} where $v_{s7}$ is the shock speed in units of $10^{7}$ cm s$^{-1}$, $\\rho_{0}$ is the pre-shock density in amu cm$^{-3}$, and $\\alpha \\ge 1$ is a parameter for the enhancement of cooling due to higher metal content. We find that a shock with velocity $\\sim 250$ km s$^{-1}$ would effectively heat dust to 50 K, with a pre-shock density of $\\sim 8$ amu cm$^{-3}$, assuming the same pressure as in slow shocks. If the dust component is composed of graphite grains at $\\sim 65$ K, a slightly faster shock of 325 km s$^{-1}$ is required. Using the above equation, we find that the amount of hot gas is on the order of a few tenths of a solar mass. This yields dust-to-gas ratios of a few percent. Given the unknown dust content within the inner ejecta of a supernova, we believe this is a reasonable explanation.  As a check on the constraints of such a fast shock, we calculated the expected X-ray emission from such a shock, and found it to be below the upper limits of thermal X-ray emission observed from the PWN, except for very metal-rich ejecta.  \\subsection{Origin of O-rich Clumps} \\label{clumps}  \\cite{matzner99} considered a spherically-symmetric explosion of a 15 $M_\\odot$ RSG, and found that its He core and heavy element ejecta formed an approximately constant density, freely expanding ejecta core.  C05 rescaled their results to other values of ejecta mass $M_{ej}$ and kinetic energy $E_{51}$, arriving at the core density of  \\begin{equation} \\rho_ct^{3} = 2.4 \\times 10^{9} (M_{ej}/15 M_\\odot)^{5/2} E_{51}^{-3/2} {\\rm g\\ cm^{-3}\\ s^{3}}. \\end{equation} An additional compression is expected from the iron-nickel bubble effect. For the centrally-located Ni with mass $M_{Ni}$, the adjacent ejecta are expected to be swept up into a shell with velocity  \\begin{equation} V_1 = 975 (M_{Ni}/0.1 M_\\odot)^{1/5} (\\rho_c t^{3}/10^{9} {\\rm g\\ cm}^{-3}\\ {\\rm s}^{3})^{-1/5} {\\rm km\\ s^{-1}}. \\end{equation} The compression within the Fe-Ni bubble shell is at least by a factor of 7, expected in strong, radiation dominated shocks with $\\gamma = 4/3$. The shell density increases inward from a shock front to a contact discontinuity separating the shocked ejecta from the Fe-Ni bubble. In one dimensional hydrodynamical simulations, Wang (2005) finds an average shell compression by a factor of 20. The average shell density is then  \\begin{equation} \\rho_1 t^{3} = 4.8 \\times 10^{10} (M_{ej}/15 M_\\odot)^{5/2} E_{51}^{-3/2} {\\rm g\\ cm^{-3}\\ s^3}. \\end{equation} (Diffusion of radiation might reduce this compression by a modest factor of $\\le 2$ -- Wang 2005.) At the current remnant's age of 1140 yr, the shell density is  \\begin{equation} \\rho_1 = 1.0 \\times 10^{-21} (M_{ej}/15 M_\\odot)^{5/2} E_{51}^{-3/2} {\\rm g\\ cm^{-3}}. \\end{equation} Because the dense iron-nickel bubble shell has been accelerated by low-density gas within the bubble, the shell is subjected to the Rayleigh-Taylor instability, and we expect it to fragment into clumps. Within a factor of 2, their expected density is equal to the preshock density for the O-rich clumps in 0540. We conclude that these clumps are remnants of the iron-nickel bubble shell.  \\cite{matzner99} found a sharp density drop by a factor of 10 at the interface between the He core and the H envelope, with an approximately constant density through much of the H envelope. The envelope density $\\rho_{env}$ is then  \\begin{equation} \\rho_{env} t^{3} = 2.4 \\times 10^{8} (M_{ej}/15 M_\\odot)^{5/2} E_{51}^{-3/2} {\\rm g\\ cm^{-3}\\ s^3}, \\end{equation} 200 times less dense than the iron-nickel bubble shell. This density contrast is similar to the density contrast inferred between the O-rich clumps and the more tenuous inter-clump gas. It is likely that the PWN nebula expands now into the H envelope. Because the dense He core has been decelerated by the less dense H envelope during the SN explosion, the ensuing Rayleigh-Taylor instability led to a large-scale macroscopic mixing between them. As a result, we expect a two-phase medium ahead of the PWN shell, consisting of more tenuous H-rich gas and denser He-rich gas. It is possible that shocks driven into the He-rich gas became radiative; that could explain the presence of H and He recombination lines in optical spectra of 0540.  The dense iron-nickel bubble shell should contain not only O-rich ejecta, substantial amounts of He-rich gas are also expected. Slow (20 km s$^{-1}$) shocks driven into the dense He-rich gas may also become radiative; if so, they could produce strong lines of low ionization species. This could explain excess emission seen in optical and IR spectra for low ionization species (see discussion in \\S~4.2). More detailed shock models are necessary to determine whether or not our picture is consistent with observations.  The identification of dense O-rich clumps in 0540 with a compressed and fragmented shell swept up by the iron-nickel bubble has important implications for ejecta detection in SNRs. Dense O-rich clumps are expected to produce strong optical or X-ray emission, once shocked and heated by the reverse shock. The optical emission should be most prominent for remnants with a particularly dense ambient medium, either of circumstellar (e.g., Cas A) or interstellar (N132D) origin. The entire class of optically emitting O-rich remnants may owe its existence to the iron-nickel bubble effect. For ejecta expanding into less dense ambient medium, X-ray emission is expected instead since clumps will be reverse-shocked much later when their densities dropped significantly because of free expansion. The O-rich clumps such as seen in 0540, even when shocked to X-ray emitting temperatures 10,000 yr after the explosion, will have substantial ($\\sim 1$ cm$^{-3}$) electron densities and emission measures. Even old remnants should show O-rich ejecta in their interiors, in agreement with the accumulating evidence gathered by modern X-ray satellites. A good example is a 14,000 yr old SNR 0049 $-$73.6 in the SMC, where {\\it Chandra} imaging and spectroscopy revealed the presence of a clumpy O-rich ring in its interior \\citep{hendrick05}. Hendrick et al. interpreted this ring as the shell swept up by the iron-nickel bubble, based on mostly theoretical arguments. Observational evidence for the iron-nickel bubble effect in 0540 strengthens this interpretation for 0049$-$73.6, and possibly for many more mature SNRs with detected ejecta emission in their interiors.  Dust formation is most likely to occur where ejecta density is the highest. The dense O-rich clumps likely contain dust; this dust may survive the passage through the radiative shock. If it were mixed into the much hotter ambient medium, this surviving dust may be the source of the observed infrared emission."
    },
    "0807/0807.1770_arXiv.txt": {
        "abstract": "Non-gaussianity in the initial conditions of the universe is one of the most powerful mechanisms to discriminate among the competing theories of the early universe. Measurements using bispectrum of cosmic microwave background anisotropies are limited by the cosmic variance, i.e. available number of modes. Recent work has emphasized the possibility to probe non-gaussianity of local type using the scale dependence of large scale bias from highly biased tracers of large scale structure. However, this power spectrum method is also limited by cosmic variance, finite number of structures on the largest scales, and by the partial degeneracy with other cosmological parameters that can mimic the same effect. Here we propose an alternative method that solves both of these problems. It is based on the idea that on large scales halos are biased, but not stochastic, tracers of dark matter: by correlating a highly biased tracer of large scale structure against an unbiased tracer one eliminates the cosmic variance error, which can lead to a high signal to noise even from the structures comparable to the size of the survey. The error improvement on non-gaussianity parameter $f_{nl}$ relative to the power spectrum method scales as $(P\\bar{n}/2)^{1/2}$, where $P$ and $\\bar{n}$ is the power spectrum and the number density of the biased tracer, respectively. For luminous red galaxies at $z\\sim 0.5$ we get an error improvement of about a factor of three over the current constraints, assuming unbiased tracers are also available with a high number density. For an ideal survey out to $z \\sim 2$ the corresponding error reduction can be as large as a factor of seven, which should guarantee a detection of non-gaussianity from an all sky survey of this type. The improvements could be even larger if two high density tracers with a significantly different sensitivity to non-gaussianity can be identified and measured over a large volume. ",
        "introduction": "One of the major unsolved mysteries in cosmology is the creation of structure in the universe. There are many competing theories that differ in their predictions, some of which are accessible to cosmological observations today. Inflation \\cite{Starobinsky:1979ty,1981PhRvD..23..347G,1982PhLB..108..389L,1982PhRvL..48.1220A},  the oldest and most successful of these theories, predicts that correlations of initial structures are nearly scale invariant, adiabatic, and nearly gaussian \\cite{1981JETPL..33..532M,1982PhLB..115..295H,1982PhRvL..49.1110G,1982PhLB..117..175S,1983PhRvD..28..679B}. However, ever since the first detection of cosmic microwave background anisotropies in COBE it has been recognized that cosmic variance, a finite number of structures on the largest scales, provides a fundamental limitation to how well we can distinguish among competing models using two-point function correlations \\cite{1993MNRAS.263L..17S,1993ApJ...417L...9S}. This is because the primordial density field is a single realization of a random process and only its two-point function is specified by the theoretical predictions. For many discriminating tests such as the shape of the power spectrum on large scales, the cosmic variance limit is fundamental. As a consequence, many interesting questions open today may remain unanswered, such as the question of whether the low quadrupole in CMB is due to manifestation of some new physics on very large scales or just a chance fluctuation \\cite{2003ApJS..148..175S}.  Our desire to discriminate among the models motivates searching for other probes, specially those where the cosmic variance limit can be avoided. An important example where this limit can to some extent be avoided are primordial gravitational waves. This is due to the fact that there exists a polarization mode (so called B polarization) that only receives contributions from gravitational waves and not from the dominant scalar perturbations \\cite{1997ApJ...482....6S,1997PhRvL..78.2054S,1997PhRvL..78.2058K}. In this case the limit is determined by instrumental and other irreducible noise rather than by the cosmic variance. Inflation allows for a detectable level of gravitational waves on the largest scales and as a result primordial gravitational waves have been declared as the smoking gun for inflation \\cite{2001PhRvD..64l3522K}, based also on the fact that some competing models \\cite{2001PhRvD..64l3522K,2003PhRvL..91p1301K,2002Sci...296.1436S} do not predict detectable levels of gravity wave signal and its detection would rule them out. The fact that cosmic variance can be avoided and that competing theories differ in their predictions are the primary motivators behind the interest in CMB polarization as the next experimental frontier.  One would like to identify other probes that can discriminate among the competing theories. One of these is non-gaussianity and most of the models predict non-gaussianity of local type, $\\Phi=\\phi+f_{nl}\\phi^2$, where $\\Phi$ is the gravitational potential in the matter era and $\\phi$ is the corresponding primordial gaussian case. In the simplest models of inflation primordial non-gaussianity is predicted to be very small, $f_{nl} \\ll 1$. Alternatives to inflation, such as the new ekpyrotic/cyclic scenario naturally predict large non-gaussianity, $f_{nl} \\gg 1$ \\cite{2008PhRvL.100q1302B,2007JCAP...11..010C,Lehners:2007wc}, leading to non-gaussianity as the smoking gun for alternatives to the simplest models of inflation. However, all of the tests of non-gaussianity proposed so far suffer from cosmic variance limit. Until recently the most powerful method to place limits on $f_{nl}$ was based on bispectrum of cosmic microwave background (CMB), with the latest WMAP constraint $-9 < f_{nl}<111$ at 95\\% confidence level \\cite{2008arXiv0803.0547K}. The statistical power of this method is limited by the cosmic variance and with a better angular resolution one can sample more modes and thus improve the limits. For example, the limits should be improved significantly with the higher angular resolution Planck satellite, where 95\\% interval on $f_{nl}$ of 30 is expected \\cite{2008arXiv0803.4194C}, but this will be limited by astrophysical uncertainties such as the secondary anisotropies in CMB. Because the primordial CMB is damped on small scales it seems unlikely that the ultimate limits from CMB will be significantly improved over the current Planck predictions, but the promise of measuring high redshift 3-dimensional matter distribution with 21 cm transitions could improve the limits further, possibly to the level below unity \\cite{2006PhRvL..97z1301C}.  An alternative approach using clustering of biased tracers of structure on very large scales has recently been proposed \\cite{2008PhRvD..77l3514D}. It was shown that non-gaussianity leads to a very unique scale dependence of the large scale bias, one that increases strongly towards large scales, and whose amplitude scales with the bias of the tracer relative to the dark matter. One can therefore place the limits on $f_{nl}$ by comparing the scale dependence of the power spectrum of the biased tracer to the one expected in cosmological models under the assumption of a scale independent bias. Subsequent work explored further theoretical issues and prospects for the future \\cite{2008ApJ...677L..77M,2008arXiv0805.3580S,2008arXiv0806.1046A}. A first application of this method, which we will call the power spectrum method, has recently been presented using large scale clustering of quasar and luminous red galaxies (LRG) galaxy data from Sloan Digital Sky Survey (SDSS) \\cite{2008arXiv0805.3580S}. The result, $-29 < f_{nl}<69$ (95\\% c.l.), is already better than the latest CMB constraints from WMAP, suggesting this is a competitive method compared to the bispectrum from CMB and should be pursued further. Projections for the future suggest that a survey of $10^6$ LRGs out to $z \\sim 0.7$ over a quarter of the sky, such as the proposed SDSS-III, will improve these limits by another factor of two \\cite{2008PhRvD..77l3514D}. Further into the future, an all sky sample of highly biased tracers out to $ z\\sim 2$ with $10^8$ galaxies, such as those contemplated for the dark energy mission \\cite{2006astro.ph..9591A}, could reduce the width of the 95 \\% confidence interval to the level of 5-10 \\cite{2008arXiv0806.1061M,2008arXiv0806.1950C,2008arXiv0806.1046A}.  While these predictions for the future are impressive, it would be great if one could improve them further to reach the realm of a guaranteed detection, which is at the level of $f_{nl} \\sim 1$ \\cite{2004PhR...402..103B}. There are two main obstacles to this. First, this method suffers significantly from cosmic variance: this is because the signal is strongest on the largest scales where the cosmic variance, due to the finite number of realizations at a given scale within a given volume, makes the error on the power spectrum large. Relative error from each mode is of order unity, hence we can only detect the $f_{nl}$ signal if the relative change in power due to $f_{nl}$ is of order unity. Second limitation pointed out in \\cite{2008arXiv0805.3580S} is that for current data the effect of $f_{nl}$ is correlated with other cosmological parameters, such as matter density and primordial slope of the power spectrum, which also change the shape of the power spectrum on large scales. This is because over the limited range of scales accessible these parameters can partially mimic the effect of $f_{nl}$. It is particularly important for two-dimensional projections, such as the quasar and LRG samples based on photometric data used in \\cite{2008arXiv0805.3580S}. For a given multipole moment $l$ the k-space projection window is broad and the effect of $f_{nl}$ extends to fairly high $l$, but it does not have a very characteristic scale dependence that would allow one to separate it from the other parameters.  Above discussion suggests, firstly, that methods that can avoid the cosmic variance limit are particularly useful for future progress in discriminating between competing theories and, secondly,  that non-gaussianity is a key probe that can differentiate among them. In this paper we propose a method to probe non-gaussianity from large scale structure that circumvents the cosmic variance limit and also eliminates the problem of its degeneracy with other cosmological parameters. Instead of measuring the power spectrum of a tracer and comparing it to predictions based on cosmological models we propose to compare directly the density field of a biased tracer to the one of an unbiased tracer (or, more generally, to a tracer with a different bias, which may include less biased or even anti-biased tracers). The relative bias of the two tracers is scale independent in the absence of non-gaussianity, but picks up a scale dependence in its presence. The main advantage is that tracers are generally biased, but not stochastic, on very large scales. This means that the precision with which the relative bias between the two tracers can be determined is only limited by the noise, which is given by the Poisson sampling of the field, and not by the cosmic variance. Moreover, because we are directly comparing two density fields any scale dependence of relative bias can only be caused by non-gaussianity and not by other cosmological parameters, which cannot affect the amplitude ratio even if they can affect the power spectrum. Hence, there is no degeneracy between $f_{nl}$ and other cosmological parameters with this method. In the next section we quantify these statements, computing the error predictions for both methods and comparing them to each other.  We then apply it to some representative examples of present and future data sets, showing the improvements that can be expected. We finish by highlighting the requirements needed for the future surveys to maximize the power of this method. ",
        "conclusions": "Let us apply the above derived relations to some examples of interest. SDSS-III plans to do a spectroscopic survey of about $1.5\\times 10^6$ luminous red galaxies (LRGs) out to $z \\sim 0.7$ over a quarter of the sky. The expected number density is about $3 \\times 10^{-4}{\\rm (h/Mpc)^3}$ and average bias about $b=1.7$ \\cite{2007MNRAS.378..852P}, so $P_1(k=0.01{\\rm h/Mpc})=0.8 \\times 10^5{\\rm (Mpc/h)^3}$ and we may expect about a factor of 3 error improvement if the unbiased tracers are available and their shot noise can be neglected. This latter assumption is by no means easy to achieve, since lower bias galaxies live in lower mass halos, are typically fainter and SDSS-III will not target them spectroscopically. One way to pursue is using photometric samples, which is discussed further below. Another is to split the LRG sample into two luminosity bins, where we may expect for the brighter bin $\\bar{n}_1=10^{-4}{\\rm (h/Mpc)^3}$ and $b \\sim 2$, so $P_1(k=0.01{\\rm h/Mpc})=1.2 \\times 10^5{\\rm (Mpc/h)^3}$ at $z \\sim 0.5$, while for the fainter bin we may expect $n_2=2\\times 10^{-4}{\\rm (h/Mpc)^3}$ and $b=1.4$, in which case $d\\alpha/d f_{nl}$ is reduced by more than a factor of two and the noise is increased since $X_1=X_2=0.1$. In this case this method gives a comparable signal to noise to the power spectrum method. This may still be advantageous, since there is no degeneracy with other cosmological parameters and one can extend the analysis to small scales.  Looking further into the future, there are a number of planned spectroscopic surveys that will measure a high number of redshifts to higher redshifts \\cite{2008ExA...tmp...12C,2006astro.ph..6104P}. In this case one may expect larger improvements, specially if all of the most biased halos can be identified. For example, at $ z \\sim 1.8$ we may have $\\bar{n}_1 \\sim 3 \\times 10^{-3}{\\rm (h/Mpc)^3}$ halos with average bias $b=2$ and $P_1(k=0.01{\\rm h/Mpc}) \\sim 3 \\times 10^4{\\rm (Mpc/h)^3}$ and most of these should be detectable with the above mentioned surveys. Hence in this case the improvement can be a factor of 7 over the power spectrum method. Note that with the power spectrum method such an improvement would require a 50-fold increase in the volume of the survey. This is a particularly exciting prospect since the predicted 95\\% confidence level interval for a full sky survey to $z \\sim 2$ is of order $\\Delta f_{nl} \\sim 5-10$ \\cite{2008arXiv0806.1061M,2008arXiv0806.1950C,2008arXiv0806.1046A} and with the additional factor of 7 reduction of error we may be able achieve $\\Delta f_{nl} \\sim 1$ (95\\% c.l. interval), at which point we enter into regime of a guaranteed detection \\cite{Bartolo:2003bz}.  The benefits may be even greater for photometric surveys, where no spectroscopy is needed and hence can be done with a significantly smaller investment. On very large scales the advantages of a spectroscopic survey become less important, specially if the photometric redshift error is well below the largest scale of the survey and if, when several redshift slices are available, the cross-correlation information between photometric samples is included in analysis. Indeed, currently the strongest constraints come from the power spectrum method using LRG photometric sample and quasar photometric sample, both from SDSS \\cite{2008arXiv0805.3580S}. As mentioned above, the main limitation of photometric samples is that the projection along the line of sight causes the non-gaussianity effect to extend over a wide range of multipole moments, making the effect difficult to distinguish from other cosmological parameters using the power spectrum method \\cite{2008arXiv0805.3580S}.  This problem is eliminated using the method proposed here.  Another potential advantage is that it is easier to obtain a high density sample of unbiased tracers at the same redshift using photometric data. For the LRG sample that goes to $z \\sim 0.7$ the corresponding photometric low bias sample already exists from SDSS with well calibrated and reasonably accurate photometric redshifts and high number densities at the level of $\\bar{n}=10^{7}$ (number of galaxies per steradian) \\cite{2008MNRAS.386..781M}. For photometric LRGs with their number density of $\\bar{n}=5\\times 10^{5}$ and power $C_l=1.5\\times 10^{-4}$ at $l=6$ \\cite{2007MNRAS.378..852P} we find a possible improvement of a factor of 5 over the current LRG photometric constraints of $\\Delta f_{nl} \\sim 200$ for 95\\% confidence interval. A potential challenge of this approach is matching the redshift distribution of the two samples, since any mismatch gives rise to less than perfect correlations. This can be achieved by constructing a set of weights for each galaxy that takes into account its redshift error distribution. It is also possible to combine spectroscopic and photometric samples. For example, one can use a high bias spectroscopic sample, such as the above mentioned LRG spectroscopic survey from SDSS-III, and apply the radial weights to reproduce the redshift distribution of a photometric sample. Finally, we should also note that the ultimate unbiased tracer insensitive to $f_{nl}$ to correlate against can be the dark matter itself, as measured from weak lensing, although in this case the radial window is rather broad \\cite{2004MNRAS.350.1445P}. Future photometric surveys may allow for much more significant improvements. For example, PAU survey \\cite{2008arXiv0807.0535B} plans to measure galaxies with a number density of $\\bar{n}=10^{-3}{\\rm (h/Mpc)^3}$ out to $z \\sim 0.9$, so here both the biased and unbiased samples will be available. The expected accuracy of photometric redshifts, $dz \\sim 0.003(1+z)$, is sufficiently small that it can be ignored for the large scales of interest here. The expected error improvement will be about a factor of 3 over the power spectrum method and the predicted 95\\% c.l. interval of $\\Delta f_{nl} \\sim 5$ is approaching the levels of a guaranteed detection.  These predictions are based on the assumption that there is no stochasticity in the tracers, ie we assumed $r=1$ in the analysis. It is well known that halos and galaxies are not perfect tracers of the density field and on smaller scales stochasticity develops. On small scales and/or for high density samples this can become the limiting factor: as shown in equation \\ref{sn} this happens when $1-r^2$ exceeds $(\\bar{n}P)^{-1}$ and since we need the latter to be well below unity for the method to be superior over the power spectrum method this is a potential issue that needs further investigation. The large scale cross-correlation coefficient between different tracers has not been studied in much detail with numerical simulations. In \\cite{2004MNRAS.355..129S} the related cross-correlation between final dark matter and halos has been investigated, but the small box size used in the analysis makes the results noisy. Still, the results suggest that the cross-correlation coefficient is indeed small on very large scales, with $1-r^2<0.1$ for $k<0.05{\\rm h/Mpc}$ and likely to be even smaller on larger scales. An analysis with a much larger volume is in progress and preliminary results in configuration space suggest that $1-r^2 \\sim 0.1$ at 5 Mpc/h and $1-r^2 \\sim 0.01$ for $r>50$ Mpc/h \\cite{smithseljak}. This issue clearly requires further attention and more detailed studies using very large volume simulations with biased tracers will be needed to address it in detail to see if cross-correlation coefficient is sufficiently close to unity on very large scales to allow the several fold reduction of error promised by this method.  While we have focused on the biased tracers of large scale structure, it is also possible to look for the signature using the anti-biased tracers. For low mass halos the bias is asymptotically approaching $b \\sim 0.7$ \\cite{2004MNRAS.355..129S,2005MNRAS.363L..66G} and for this population the effect of $f_{nl}$ has the opposite sign from the biased population with an $f_{nl}$ dependence that is about a factor of 3-4 smaller than for $b=2$ population assuming $p=1$. In addition, such an anti-biased tracer has about a factor of 10 smaller amount of power. These two disadvantages can in principle be offset by the higher number density, which can be orders of magnitude larger for $b \\sim 0.7$ population relative to the $b \\sim 2$ population. We can also use these antibiased tracers instead of the unbiased tracers and comparing them to the biased tracers, which improves the sensitivity to $f_{nl}$ and reduces the shot noise relative to the unbiased tracer we assumed in our analysis.  More generally, recent studies suggest that bias can depend on quantities other than the halo mass \\cite{2005MNRAS.363L..66G,2007MNRAS.377L...5G}, and similarly one can also expect secondary parameters that may enhance or suppress the sensitivity to $f_{nl}$. Indeed, an extended Press-Schechter analysis suggests one such second parameter that suppresses the sensitivity to $f_{nl}$ of highly biased tracers is the recent merger activity \\cite{2008arXiv0805.3580S}. Observationally identifiable candidates that satisfy this criterion may for example be quasars. The challenge for the future is to find two tracers, one that exhibits a significant dependence on $f_{nl}$ and one that does not, both of which come with a sufficiently high number density to reduce the Poisson noise and both of which can be observationally identified within existing and future surveys. Given the potentially huge payoff of finding such tracers for the questions of interest to the fundamental theories of the universe this is a challenge that is worth exploring further both with observations as well as with numerical simulations and analytic methods.  I thank Nikhil Padmanabhan and An\\v ze Slosar for useful comments on the draft. U.S. is supported by the Packard Foundation and Swiss National Foundation under contract 200021-116696/1."
    },
    "0807/0807.4363_arXiv.txt": {
        "abstract": "A long-range fifth force coupled to dark matter can induce a coupling to ordinary matter if the dark matter interacts with Standard Model fields.  We consider constraints on such a scenario from both astrophysical observations and laboratory experiments. We also examine the case where the dark matter is a weakly interacting massive particle, and derive relations between the coupling to dark matter and the coupling to ordinary matter for different models.  Currently, this scenario is most tightly constrained by galactic dynamics, but improvements in E\\\"otv\\\"os experiments can probe unconstrained regions of parameter space. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0807/0807.0967_arXiv.txt": {
        "abstract": "{S.Co.P.E. is one of the four projects funded by the Italian Government in order to provide Southern Italy with a distributed computing infrastructure for fundamental science. Beside being aimed at building the infrastructure, S.Co.P.E. is also actively pursuing research in several areas among which astrophysics and observational cosmology. We shortly summarize the most significant results obtained in the first two years of the project and related to the development of middleware and Data Mining tools for the Virtual Observatory. ",
        "introduction": "S.Co.P.E. is a general purpose GRID infrastructure of the University Federico II in Naples funded through the Italian National Plan (PON) by the Italian Government to support both fundamental research and small/medium size companies. The infrastructure has been conceived as a metropolitan GRID, embedding different (and in some cases pre-existing) and heterogeneous computing centers each with its specific vocation: high energy physics, astrophysics, bioinformatics, chemistry and material sciences, electric engineering, social sciences. Its intrinsically multi-disciplinary nature renders the S.Co.P.E. an ideal test bed for innovative middleware solutions and for interoperable tools and applications finely tuned on the needs of a distributed computing environment. In what follows we shall shortly outline the main activities in the fields of astrophysics and observational cosmology and, in particular, we shall focus on: i) the ongoing efforts aimed at integrating the S.Co.P.E. GRID (hereafter SG) with the international Virtual Observatory (Sect.2), and ii) the implementation in the SG of the data mining (DM) VO-Neural package (Sect.3) which is developed in the framework of a collaboration with the Dept. of Astronomy at Caltech. In Sect. 4 we shortly outline a template scientific application and, finally, in Sect. 5, we outline some future developments. ",
        "conclusions": ""
    },
    "0807/0807.4914_arXiv.txt": {
        "abstract": "We discuss the relevance of the curvature radiation of pions in strongly magnetized pulsars or magnetars, and their implications for the production of TeV energy neutrinos detectable by cubic kilometer scale detectors, as well as high energy photons. ",
        "introduction": "\\label{sec:intro}  In pulsar and magnetar polar caps one expects the formation of vacuum or space-charge limited gaps (e.g. Harding and Lai, 2006), in which an electric field component parallel to the magnetic field accelerates particles to relativistic velocities.  The particles move away from the caps along the magnetic field lines, producing photons by curvature radiation, which at some height $h$ produce an $e^\\pm$ pair front that determines the upper extent of the gap (e.g. Ruderman and Sutherland, 1975, Harding and Muslimov, 2002). The caps with $\\Omega\\cdot B<0$ can accelerate positively charged protons or ions, whose interaction with X-rays from the polar cap will lead to very high energy (VHE, $E\\simg 10^{12}$ eV) neutrinos through photomeson interactions, $p\\gamma \\to \\pi^+ \\to \\mu^+\\nu_\\mu \\to e^+\\nu_e{\\bar \\nu}_\\mu$ (Zhang et al, 2003).  Here we explore the consequences for neutrino production of a different mechanism, the curvature radiation of pions by protons.  Relativistic protons interacting with a strong magnetic field can produce pions, a quantum treatment of this process being given by Zharkov (1964). Ginzburg and Zharkov (1965) used a semi-classical approach to emphasize the analogy between this pion radiation process and the usual synchrotron radiation of photons by protons.  This semi-classical method was further explored for $\\pi^0$ and $\\rho$ mesons by Tokuhisa and Kajino (1999) in the context of pulsars. The pion radiation mechanism is characterized by a parameter $\\chi=\\gamma^2 (B/B_Q)$ similar to that used for electron processes but with the proton mass replacing the electron mass, with $B_Q= m_p^2c^3/e\\hbar=1.5 \\times 10^{20}$ G. They showed that the synchrotron-like $\\pi^0$ radiation process becomes competitive with (proton) photon synchrotron radiation when the above parameter $\\chi \\simg 0.1$. The related synchrotron-like $\\pi^+$ radiation by protons in a magnetic field requires, however, a different semi-classical treatment, since the protons become neutrons in the process, and this was calculated by Herpay and Patk\\'os (2008).  In the astrophysical literature, particle acceleration in pulsars or magnetars is considered both near the polar caps (``inner\" or polar gap models), and further out near or beyond the light cylinder (``outer\" gap or wind models); e.g. see review by Harding and Lai (2006). Examples of the former are e.g. Blasi, Epstein and Olinto (2000), Arons (2003), etc.; while examples of the latter are, e.g. Sutherland (1979), Harding and Muslimov (2001, 2002, 2005), Thompson (2008), etc.. Here, we consider specifically the case of protons or ions accelerated in the inner regions, just above the polar caps of magnetars. In the inner acceleration region discussed here, the accelerated particles are expected to radiate away their transverse energy so efficiently that they are effectively constrained to the ground Landau level as they move along the curved field lines. The principal photon radiation loss is then curvature radiation (Ruderman and Sutherland, 1975; Harding and Lai, 2006). In this situation, the synchrotron-like pion radiation (which depends on the proton transverse energy) will be similarly suppressed. However, as for photon curvature radiation, the motion along the curved field lines produces an acceleration equivalent to a Larmor motion in a fictitious magnetic field, which will lead to curvature pion radiation (Berezinsky, Dolgov and Kachelriess, 1996). Unlike previous authors, in this paper we concentrate on $\\pi^\\pm$ production by the curvature-like mechanism of both protons and heavy nuclei, occurring in the immediate vicinity of the polar caps, and we explore its consequences for the production of VHE neutrinos, as well as their detectability by cubic kilometer Cherenkov detectors such as IceCube and KM3NeT. We compare the resulting neutrino fluxes to those from the photomeson process, and discuss the associated cosmic ray as well as GeV-TeV photon production.  In \\S \\ref{sec:model} we discuss an approximate acceleration and radiation loss model, including both photon and pion radiation losses. Simple time-scale arguments are presented to estimate the maximum particle energies possible and the ratio of photon to pion losses. In \\S \\ref{sec:num} we consider in more detail the equations of motion in self-consistent polar cap  magnetic field models, and calculate numerically the acceleration and energy loss rates in a general manner. These results are in good agreement with the approximate estimates of the previous section. In \\S 3 and \\S 4 we estimate the resulting neutrino luminosity from curvature pion decays and from photo-meson effects. In \\S \\ref{sec:disc} we discuss the possibility of detecting these effects in kilometer scale neutrino detectors and at GeV photon energies. ",
        "conclusions": "\\label{sec:disc}  Our discussion has assumed that the properties of the inner gaps near the polar caps of pulsars or magnetars allow protons to be accelerated to energies $\\simg 6\\times 10^8$ GeV. This is an open question, since the extent and properties of polar gaps (with ${\\vec \\Omega}\\cdot {\\vec B}\\leq 0$) including self-consistently pion radiation effects is unknown. Such studies, if undertaken, would have to take into account not only the effects of leptons acceleration from charged pion decays, but also high energy photons from neutral pion decay and their subsequent cascades.  Studies of space charge limited gaps with electron acceleration and pair production (e.g.  Harding and Muslimov, 2002; also Baring and Harding, 2002) suggest that Lorentz factors of order $10^9$ may be possible in fast-rotating, high field objects, in which case the effects discussed here may become important.  The curvature pion radiation mechanism is likely to be of interest for fast rotating magnetars which accelerate protons. This is because, from Fig.~\\ref{fig:tp}, we see that for magnetic field and period values $\\B15 \\P3^{-1} \\simg 1$ the pion curvature radiation dominates the photon curvature radiative cooling. For such magnetic field and period combinations, the acceleration of heavy ions results in the fragmentation of the ions before pion emission occurs, which is expected only from the individual protons resulting from the fragmentation after they have been further accelerated.  There is so far no observational evidence for millisecond  periods among the handful of known magnetars (e.g. Kaspi, 2007; Woods and Thompson, 2004), most of which are in our galaxy and have periods of seconds. From eq. (\\ref{eq:mupic}), we see that for $B\\simg 10^{15}$~G and millisecond periods, the muon event rate is so large that  they might have been detected by now, if the  magnetar were on-beam and at a distance $D\\siml$ Mpc.  The likeliest explanation is that such millisecond magnetars do not exist in our galaxy (or else they might be off-beam). At the same time, it is reassuring that  for periods $P\\simg 1$ s  the pion curvature efficiency predicted is $\\eta_\\pi=t_c/t_\\pi \\lll 1$, and the corresponding muon event rate in cubic kilometer detectors given by equation (\\ref{eq:mupic}) is negligible for known galactic magnetars.  An interesting possibility is that some fraction of core collapse supernovae lead, at least initially, to a millisecond magnetar. There is currently no observational evidence for this, but it is a plausible  hypothesis, among others from fast convective overturn dynamo arguments, e.g. Thompson and Duncan (1996).  In such objects, the envelope optical thickness to $\\nu N$ interactions is $\\tau_{\\nu  N} \\sim 10^{-6} (\\eps_\\nu /\\hbox{10 TeV}) (M_{env}/10\\msun) (v_{env}/0.1c)^{-2} (t/\\hbox{day})^{-2}$, and as long as the magnetic field and the rotation rate remain high, neutrinos produced by pion curvature and $p\\gamma$ interactions can escape without further reprocessing. Thus, core collapse supernovae which result initially in a millisecond magnetar with $B\\simg 10^{15}$ G are expected to undergo a brief period of intense pion and ultra-high energy neutrino production, which significantly exceeds any electromagnetic energy losses. If the fraction of core collapse supernovae leading to such objects were $0.05-0.1$, given the frequency of core collapse SNe in the LMC and in the local group, one could in principle expect some millisecond magnetars detectable in cubic kilometer detectors within timescales of years.  In this scenario one would also expect (\\S \\ref{sec:photomes}) a significant $\\siml 10$ MeV photon luminosity, which may be detectable by the GLAST GBM.  In summary, we have pointed out that pions produced by protons interacting with the curved strong magnetic fields of magnetars may be an important energy loss mechanism for the accelerated particles, leading to secondary leptons and photons which can affect the properties of the inner gaps.  Our model for the strong interaction processes involved here relies on the softness of the pion radiation process. Although the correctness of this assumption needs further, more detailed investigation, this level of approximation is justifiable here, given the larger uncertainties in the astrophysical model. Our results indicate that the effects analysed in this paper could lead to copious neutrino production in the TeV-nPeV range, which would provide interesting targets for cubic kilometer scale detectors.  {"
    },
    "0807/0807.4549_arXiv.txt": {
        "abstract": "NGC 6240 is a rare object in the local universe: an active merger remnant viewed at the point of merging where two active galactic nuclei are visible. We present IRAC data of this object, providing high sensitivity maps of the stellar and PAH distribution in this complicated system. We use photometry to analyze the variation in these distributions with radius and provide an SED in the four IRAC bands: 3.6\\,$\\mu$m, 4.5\\,$\\mu$m, 5.8\\,$\\mu$m and 8\\,$\\mu$m. We fit the radial profiles of the 3.6\\,$\\mu$m band to $r^{\\frac{1}{4}}$ and exponential profiles to evaluate the structure of the remnant. Finally, we compare the IRAC images with multi-wavelength data and examine how outflows in the X-ray, H$\\alpha$ and CO correlate with 8\\,$\\mu$m emission. The results support the general picture of NGC 6240 as a system experiencing a major merger and transitioning from a disk galaxy to a spheroid. The sensitivity of IRAC to low surface brightness mid-infrared features provides detailed information on the extended distributions of stars and dust in this rare system. ",
        "introduction": "Due to its proximity, NGC 6240 ($z=.0243$ \\citet{Solomon-et-al-1997}, $D=98$ Mpc for $H_{0}=75$ km s$^{-1}$ Mpc$^{-1}$) is one of the brightest objects in the IRAS All-Sky Survey and was one of the early identified LIRGs with a total infrared luminosity of $7.1\\times10^{11} L_{\\odot}$ \\citep{Wright-et-al-1984}.   Optical studies of NGC 6240 reveal a morphologically irregular galaxy with clear tidal extensions, two nuclei and large dust lanes, strongly suggesting a merging system \\citep{Zwicky-et-al-1961, Fried-Schulz-1983, Keel-1990, Gerssen-et-al-2004}. Near-IR data \\citep{Thronson-et-al-1990, Doyon-et-al-1994, Scoville-et-al-2000, Max-et-al-2005} and radio continuum \\citep{Carral-et-al-1990, Eales-et-al-1990, Colbert-et-al-1994, Beswick-et-al-2001} images confirm the double nuclei and show that they have considerable substructure. The observed distance between the the nuclei  depends on the wavelength at which they are observed, indicating large amounts of dust extinction in the center of the galaxy \\citep{Schulz-et-al-1993, Gerssen-et-al-2004}. H$\\alpha$ data shows strong evidence for nuclear winds, outflows and bubbles \\citep{Gerssen-et-al-2004}.  X-ray studies show a very active galaxy with two nuclei displaying AGN characteristics, indicating two black holes \\citep{Iwasawa-Comastri-1998, Vignati-et-al-1999, Komossa-et-al-2003}. Sub-millimeter array data reveal a molecular gas disk of mass $\\approx 10^{9}$ M$_{\\odot}$ in between the two nuclei with the possibility of outflows \\citep{Wang-Scoville-Sanders-1991, Tacconi-et-al-1999, Iono-et-al-2007}.  Merger remnants with double nuclei are rare in the nearby universe, suggesting that we are catching this object in a very short lived phase of its evolution. In addition, its infrared luminosity places it on border between the LIRG and a ULIRG classes. Though classified by their luminosity, the LIRG and ULIRG classes show differing characteristics besides their luminosity, including a higher occurrence of AGN and disturbed morphologies in ULIRGs \\citep{Veilleux-et-al-2002, Sanders-et-al-2004}. Given that NGC 6240 already shows many of these features, we may be observing a transition object between the LIRG and ULIRG classes.  Objects such as NGC 6240 are believed to be undergoing a merger between two comparable mass galaxies which is driving the remnant through evolution from a Seyfert type active galaxy to an AGN and finally a quasar \\citep{Norman-Scoville-1988, Sanders-et-al-1988, Hopkins-et-al-2006}. Models demonstrate a cycle in which a merger drives gas inflows that create starbursts and feed central super-massive black holes, triggering an active stage for the black hole. As the AGN ``turns on'' and contributes significant radiative feedback, it clears out the gas around the nuclei, and the AGN becomes visible as a quasar. However, this also depletes its fuel, leaving the AGN to die back to a quiescent phase. The cycle may start over again when the galaxy goes through another merger. \\citet{Hopkins-et-al-2008-redgals, Hopkins-et-al-2008-quasars} show this cycle is consistent with many observed properties of quasars and elliptical galaxies. This evolutionary sequence is short ($\\sim$ few Gyr) compared to the lifetime of a spiral or elliptical galaxy, so there are relatively few of them to study, and even fewer at the particular stage of merging where two AGN are observable. This makes them particularly interesting objects to study in detail.  We add Spitzer/IRAC \\citep{Fazio-et-al-2004} data of NGC 6240 to the existing observations of NGC 6240. Most infrared studies of NGC 6240 have focused on the nuclei \\citep{Klaas-et-al-2001, Lutz-et-al-2003, Armus-et-al-2006, Egami-et-al-2006}. We draw on IRAC's high sensitivity and multiple mid-infrared bands to investigate how stars and dust are distributed throughout the remnant. The IRAC band-passes single out emission from stars (3.6\\,$\\mu$m) and polycyclic aromatic hydrocarbon emission (PAH, a component of dust, 8\\,$\\mu$m), allowing us to dissect the structure of NGC 6240 by examining the distribution of stars and ISM separately \\citep{Pahre-et-al-2004}. All four IRAC bands are imaged on the same scale so that these components can be easily compared. We derive light profiles in each of the IRAC bands and calculate the SED. We also compare IRAC data to archive optical, X-ray and radio data. In \\S\\ref{sec:obs} we describe our observations and in \\S\\ref{sec:results} we describe the IRAC images and how they compare to data from other wavelengths. In \\S\\ref{sec:discuss} we present our photometry of the IRAC data and in \\S\\ref{sec:conc} we make some concluding remarks. ",
        "conclusions": "\\label{sec:conc}  NGC 6240 is an object experiencing a rarely observed stage of evolution, showcasing complex combinations of processes. We add Spitzer-IRAC imaging and photometry to the active study of this object. We show that stars and dust have very different distributions in the remnant, both regionally and radially from the central AGN. We examine how outflows in the X-ray, H$\\alpha$ and CO correlate with 8\\,$\\mu$m PAH emission. We show that though, when radially averaged, the stellar profile follows an   $r^{1/4}$ law surprisingly well, regions of the galaxy still show disk-like stellar profiles. All this supports the basic picture of a galaxy experiencing the effects of a major merger, which treats stars and gas/dust very differently, and can create a gas poor spheroid from a gas rich spiral galaxies. Though we can say little about the progenitors of this system, the sensitivity of IRAC data to low surface brightness emission adds detailed knowledge of the distribution of stars and gas in this remnant. This provides constraints for large samples of hydrodynamic simulations coupled with radiative transfer which, if they are able to match the characteristics of NGC 6240, may be able to determine its history and future."
    },
    "0807/0807.2159_arXiv.txt": {
        "abstract": "We study the evolution of homogeneous and isotropic, flat cosmological models within the general scalar-tensor theory of gravity with arbitrary coupling function and potential. After introducing the limit of general relativity we describe the details of the phase space geometry. Using the methods of dynamical systems for the decoupled equation of the Jordan frame scalar field we find the fixed points of flows in two cases: potential domination and matter domination. We present the conditions on the mathematical form of the coupling function and potential which determine the nature of the fixed points (attractor or other). There are two types of fixed points, both are characterized by cosmological evolution mimicking general relativity, but only one of the types is compatible with the Solar System PPN constraints. The phase space structure should also carry over to the Einstein frame as long as the transformation between the frames is regular which however is not the case for the latter (PPN compatible) fixed point. ",
        "introduction": "Scalar-tensor theories (STT) of gravitation \\cite{will, fm, faraoni} have emerged in different contexts of theoretical physics, e.g. in Kaluza-Klein type unifications, supergravity, and low energy approximations of string theories. The scalar-tensor action functional has been also used in topology as a mathematical tool for solving the geometrization conjecture of 3-manifolds via the Ricci flow \\cite{anderson}.  In cosmology STT has been invoked to model the accelerated expansion of inflation  \\cite{inflation} and dark energy \\cite{darkenergy} (for some recent papers see e.g. Ref. \\cite{darkenergy_recent}). However, observations in the Solar System tend to indicate that in the intermediate-range distances the present Universe around us is successfully described by the Einstein tensorial gravity alone \\cite{tests}. This means that only such models of scalar-tensor gravity are viable which in their late time cosmological evolution imply local observational consequences very close to those of Einstein's general relativity (GR) \\cite{tumty}.  The importance of the convergence of the STT solutions to those of GR at late times, guaranteeing agreement with the present post-Newtonian observations was recognized by Bekenstein and Meisels \\cite{bek} already long ago. Damour and Nordtvedt \\cite{dn0, dn} argued that for a wide class of homogeneous and isotropic scalar-tensor cosmological models there exists an attractor mechanism taking them to the limit of the corresponding general relativistic models. In subsequent papers the issue has been pursued by various approaches \\cite{asymptotics, bp, san, bs}.   The methods of dynamical systems provide natural tools to analyze the problem. Most previous studies which have considered STT cosmology as a dynamical system have focused upon examples with specific coupling functions, see Refs. \\cite{dynsys_einstein, mn, carloni_2008, meie3} for works based on the Einstein frame and Refs. \\cite{dynsys_jordan, gunzing_2001, abramo_2003, meie2, meie3} for the Jordan frame. In the Jordan frame the main properties of the general phase space geometry were outlined by Faraoni \\cite{faraoni1}, while Refs.  \\cite{burd_coley, faraoni3} consider general Jordan frame dynamics with particular attention to the de Sitter fixed point.  Although the equations of motion in the Einstein frame may have a simpler form, the cosmological observations are easier to interpret in the Jordan frame where matter is minimally coupled. In general there is a one-to-one correspondence between the phase spaces in the two conformal frames \\cite{dicke, weinberg, faraoni, flanagan, faraoni2}, provided that the transformations from one frame to another are regular. However, as it was pointed out in Ref. \\cite{meie3}, the differential equation relating the scalar fields in these two frames becomes singular at the limit of general relativity. Therefore utmost care must be exercised when addressing the question whether general relativity is an attractor, as the answer may depend on the frame chosen.  In this paper we take the Jordan frame and consider general scalar-tensor theories which contain two functional degrees of freedom, the coupling function $\\omega (\\Psi)$ and the scalar potential $V(\\Psi)$. We perform the dynamical systems analysis for the flat Friedmann-Lema\\^{\\i}tre-Robertson-Walker (FLRW) backgrounds with ideal barotropic fluid matter. Our strategy is to find the fixed points for the scalar field dynamics and compare these with the conditions of the limit of general relativity in the Solar System, as established by the parameterized post-Newtonian (PPN) formalism. Therefore, if the functional forms of $\\omega (\\Psi)$ and $V(\\Psi)$ are specified from some considerations (e.g. the compactification manifold), our results allow to determine the fixed points along with their type and thus immediately decide whether general relativity is an attractor, i.e. whether the model at hand is viable or not.   The plan of the paper is the following. In section 2 we write the field equations as a dynamical system and find the conditions which reduce these equations to the FLRW equations of general relativity (possibly with a cosmological constant), concluding that these conditions are marginally stricter than the standard PPN conditions. In section 3 we describe the general phase space and its physical subspace determined by the Friedmann constraint equation. In section 4 we consider the case of potential domination with (effectively) vanishing matter density, find the fixed points and determine their properties. For investigating the case of matter domination (with vanishing scalar potential) in section 5, we follow Refs. \\cite{dn} and \\cite{san} to introduce a new time parameter which allows us derive a decoupled ``master'' equation for the scalar field and again perform the fixed point analysis. In section 5 we briefly discuss comparison to some previous works and end with a speculation about the possible relevance of our results for the moduli stabilization problem in string theory. Section 6 gives a summary. ",
        "conclusions": "We have considered flat FLRW cosmological models in general scalar-tensor theories with arbitrary coupling function $\\omega(\\Psi)$ and scalar potential $V(\\Psi)$ in the Jordan frame. Using the methods of dynamical systems we have described the general geometry of the phase space and found the scalar field fixed points in two distinct asymptotic regimes: potential domination ($V\\not\\equiv 0, \\rho \\equiv 0$), and matter domination ($V \\equiv 0, \\rho \\not\\equiv 0 $). In nutshell there are two types of fixed points arising from different mechanisms: $\\Psi_{\\bullet}$ from a condition on the potential (equalling the local extremum of the Einstein frame potential) and $\\Psi_{\\star}$ from the singularity of the scalar field kinetic term. Approaching both types of fixed points the cosmological equations coincide with those of general relativity, yielding de Sitter expansion in the potential domination case and Friedmann evolution in the matter domination case. However, for the Solar System experiments in the PPN framework only the fixed points of $\\Psi_{\\star}$ type give predictions identical with those of general relativity. The nature of fixed points (attractor or otherwise) depends on the functional forms of $\\omega(\\Psi)$ and $V(\\Psi)$ according to corresponding eigenvalues given in Tables \\ref{V_fp_eigenvalues} and \\ref{matter_fp_eigenvalues}. Therefore, in Jordan frame analysis, general relativity is an attractor for a large class of scalar-tensor models, but not for all.  Provided the transformation relating the Jordan and Einstein frame is regular, there is an exact correspondence between the two frames and the Jordan frame phase space results should carry over to the Einstein frame. This is the case for the fixed point $\\Psi_{\\bullet}$. However, as the transformation of the scalar field fails to be regular in the limit of general relativity, the properties of the $\\Psi_{\\star}$ fixed point may be altered in the Einstein frame. To establish whether or how the correspondence holds in this case calls for a separate matching investigation in the Einstein frame.     \\bigskip {\\bf Acknowledgements}  We are grateful to the anonymous referee for clarifying questions. This work was supported by the Estonian Science Foundation Grant No. 7185 and by Estonian Ministry for Education and Science Support Grant No. SF0180013s07. MS also acknowledges the Estonian Science Foundation Post-doctoral research grant No. JD131."
    },
    "0807/0807.2645_arXiv.txt": {
        "abstract": "We present accurate Monte-Carlo calculations of Ly$\\alpha$ radiation pressure in a range of models which represent galaxies during various epochs of our Universe. We show that the radiation force that Ly$\\alpha$ photons exert on hydrogen gas in the neutral intergalactic medium (IGM), that surrounds minihalos that host the first stars, may exceed gravity by orders of magnitude and drive supersonic winds. Ly$\\alpha$ radiation pressure may also dominate over gravity in the neutral IGM that surrounds the HII regions produced by the first galaxies. However, the radiation force is likely too weak to result in supersonic outflows in this case. Furthermore, we show that Ly$\\alpha$ radiation pressure may drive outflows in the interstellar medium of star forming galaxies that reach hundreds of km s$^{-1}$. This mechanism could also operate at lower redshifts $z\\lsim 6$, and may have already been indirectly detected in the spectral line shape of observed Ly$\\alpha$ emission lines. ",
        "introduction": "\\label{sec:intro}  HII regions around massive stars convert a significant fraction of the total bolometric luminosity of young galaxies into Ly$\\alpha$ line emission \\citep{PP67,S03}. This Ly$\\alpha$ radiation can exert a large force on surrounding neutral gas, as the Ly$\\alpha$ transition has a cross-section that is $\\sim 7$ orders of magnitude larger than the Thomson cross-section, when averaged over a frequency band as wide as the resonance frequency itself \\citep[e.g.][]{Loeb01}. Not surprisingly, the impact of Ly$\\alpha$ radiation pressure on the formation of galaxies has been discussed extensively \\citep[e.g][]{Cox85,EF86,Bithell90,Haenelt95,Oh02,McKee07}, but the intricacies of Ly$\\alpha$ radiative transfer in 3D complicated an accurate numerical treatment of its dynamical effect on the gas. Nevertheless, an approximate estimate can be obtained from simple energy considerations as shown below.  Consider a self-gravitating gas cloud of total (baryons + dark matter) mass $M$ and radius $R$ that contains a central Ly$\\alpha$ source. The gravitational binding energy of the baryons inside the cloud, $E_B\\sim \\Omega_bGM^2/(\\Omega_m R)$, can be compared to the total energy in the Ly$\\alpha$ radiation field inside the cloud, $E_{\\alpha}=L_{\\alpha}\\times t_{\\rm trap}$. Here, $L_{\\alpha}$ is the Ly$\\alpha$ luminosity of the central source (in erg s$^{-1}$), and $t_{\\rm trap}$ is the typical trapping time of Ly$\\alpha$ photons in the cloud owing to scattering on hydrogen atoms. The Ly$\\alpha$ radiation pressure would unbind the baryonic gas from the cloud if $E_{\\alpha}> E_B$, i.e. $L_{\\alpha}>\\Omega_bGM^2/(\\Omega_mRt_{\\rm trap})$ \\citep[e.g.][]{Cox85,Bithell90,Oh02}. In this approach, $t_{\\rm trap}$ is one of the key parameters in setting the Ly$\\alpha$ radiation pressure. Calculations by \\citet{Adams75} imply that $t_{\\rm trap}\\sim 15 t_{\\rm light}$ for $3\\lsim \\log \\tau_0 \\lsim 5.5$, and $t_{\\rm trap}\\sim 15(\\tau_0/10^{5.5})^{1/3} t_{\\rm light}$ otherwise, for a static, uniform, infinite slab of material \\citep[also see Fig~1 of][]{Bo79}. Here $\\tau_0$ is the line center optical depth from the center to the edge of the slab, and $t_{\\rm light}$ is the light crossing time if the medium were transparent (i.e. $t_{\\rm light}=R/c$ in the case of the cloud described above). Note however, that the precise value of $t_{\\rm trap}$ depends on other factors including for example, the gas distribution (clumpiness and geometry), the velocity distribution of the gas, and the dust content of the cloud \\citep{Bo79}.  The Ly$\\alpha$ radiation pressure becomes comparable to gravity when \\begin{equation} L_{\\alpha,{\\rm 41}}=1.0\\Big{(}\\frac{M}{10^{8}M_{\\odot}}\\Big{)}^{4/3}\\Big{(}\\frac{16}{1+z}\\Big{)}^2\\Big{(}\\frac{15t_{\\rm light}}{t_{\\rm trap}}\\Big{)}, \\label{eq:lcrit} \\end{equation} where $L_{\\alpha}=L_{\\alpha,{\\rm 41}}\\times 10^{41}$ erg s$^{-1}$, and where we have substituted the virial radius of a galaxy mass $M$, $R_{\\rm vir}=0.97$ kpc $\\times(M/10^8M_{\\odot})^{1/3}$$(1+z/16)^{-1}$, for $R$ (Eq.~24 of Barkana \\& Loeb 2001). For comparison, a star forming galaxy can generate a Ly$\\alpha$ luminosity of $L_{\\alpha}=(10^{42}-10^{43})\\times({\\rm SFR}/M_{\\odot}\\hs{\\rm yr}^{-1})$ erg s$^{-1}$, where the precise conversion factor depends on the gas metallicity and the stellar initial mass function \\citep[e.g.][]{S03}. Therefore, a star formation rate of merely SFR$\\gsim 0.01$--$0.1M_{\\odot}$ yr$^{-1}$ is needed to generate a Ly$\\alpha$ luminosity that is capable of unbinding gas from a halo of mass $10^8-10^9M_{\\odot}$.  Halos of $\\la 10^9M_\\odot$ are very common at $z\\gsim 6$, and have a sufficiently large reservoir of baryons to sustain the above-mentioned star formation rates for a prolonged time. In this paper we provide a more detailed investigation of the magnitude of Ly$\\alpha$ radiation pressure in the environment of high-redshift star forming galaxies. In particular, we use a Ly$\\alpha$ Monte-Carlo radiative transfer code \\citep{mc} to compute Ly$\\alpha$ radiation pressure in a wider range of models. Our treatment of radiative transfer and our focus on the environment of high-redshift star forming galaxies, distinguish this paper from previous work. We will show that the radiation force exerted by Ly$\\alpha$ photons on neutral hydrogen gas can exceed the gravitational force that binds the gas to its host galaxy by orders of magnitude, and may drive supersonic outflows of neutral gas both in the intergalactic and the interstellar medium.  The outline of this paper is as follows. In \\S~\\ref{sec:code} we describe how Ly$\\alpha$ radiation pressure is computed in the Monte-Carlo radiative transfer code, and show the tests that are performed to test the accuracy of the code. In \\S~\\ref{sec:result}, we present our numerical results.  Finally, \\S~\\ref{sec:conc} summarizes the implications of our work and our main conclusions. The cosmological parameter values used throughout our discussion are $(\\Omega_m,\\Omega_{\\Lambda},\\Omega_b,h)=(0.27,0.73,0.042,0.70)$ \\citep{Komatsu08}. ",
        "conclusions": "\\label{sec:conc}  We have applied an existing Monte-Carlo Ly$\\alpha$ radiative transfer code \\citep[described and tested extensively in][]{mc} to the calculation of the pressure that is exerted by Ly$\\alpha$ photons on an optically thick medium. This code enabled us to perform (the first) direct, accurate calculations of Ly$\\alpha$ radiation pressure, which distinguishes this work from previous discussions on the importance of Ly$\\alpha$ radiation pressure in various astrophysical environments.  We have focused on a range of models which represent galaxies at different cosmological epochs. In \\S~\\ref{sec:firststar} we have shown that the Ly$\\alpha$ radiation pressure exerted on the neutral intergalactic medium (IGM) surrounding minihalos ($M_{\\rm tot}\\sim 10^6 M_{\\odot}$) in which the first stars form, can exceed gravity by 2--3 orders of magnitude (Fig~\\ref{fig:caseI}), and in principle accelerate the gas in the IGM to tens of km s$^{-1}$. Thus, Ly$\\alpha$ radiation pressure can launch supersonic winds in the IGM surrounding the first stars. Our analysis did not allow a detailed study of the consequences of these winds. A comprehensive study would require numerical simulations that capture the full IGM density field around minihalos in 3D and track the evolution of the shocks that may form in the IGM together with the Ly$\\alpha$ radiative transfer. In this paper, we have also shown that Ly$\\alpha$ radiation pressure is important in the neutral IGM that surrounds the HII regions produced by galaxies with a total halo mass of $M_{\\rm tot}=10^8M_{\\odot}$ (Fig~\\ref{fig:caseII}. These are the lowest mass, and hence the most abundant, halos in which gas can cool via atomic line excitation. Here, however, the absolute magnitude of the radiation force is too weak to drive the gas to supersonic velocities.  Finally, we have shown in \\S~\\ref{sec:wind} that the Ly$\\alpha$ radiation pressure exerted on neutral gas in the interstellar medium (ISM) of a galaxy can also have strong dynamical consequences. In particular, we have found that the Ly$\\alpha$ radiation force exerted on an expanding HI supershell can exceed gravity by orders of magnitude (Fig~\\ref{fig:caseIII}), for reasonable assumptions about the gravitational force. It is therefore possible that Ly$\\alpha$ radiation pressure plays an important role in determining the kinematics of HI supershells around starburst galaxies. We have demonstrated that the total Ly$\\alpha$ radiation force on a spherical HI supershell can be written as $F_{\\rm rad}=M_F L_{\\alpha}/c$, where the 'force-multiplication factor' $M_F$ relates to the average trapping time of Ly$\\alpha$ photons in the neutral medium. The factor $M_F$ can greatly exceed unity, as illustrated by Fig~\\ref{fig:fscat}. For comparison, the maximum possible radiation force due to continuum radiation\\footnote{ Continuum radiation may exert a force on steady-state outflows (with a constant mass ejection rate, $\\dot{M}$) around late-type stars that may significantly exceed $\\dot{M}\\Delta v\\gg L_{\\rm bol}/c$ \\citep[][]{salpeter74,Ivezic95}. This is not because of 'trapping' of continuum photons, but related to the propagation speed of the photons and the wind. We similarly expect the radiative force of trapped Ly$\\alpha$ photons in steady-state outflows to potentially exceed $\\dot{M}\\Delta v\\gg M_FL_{\\alpha}/c$ (and as argued in this paper, it is possible that $M_FL_{\\alpha}/c>L_{\\rm bol}/c$). Note though that these steady-state outflows are clearly different from those discussed in \\S~\\ref{sec:wind}, in which a well defined thin shell HI gas is physically separated from the central Ly$\\alpha$ source.}  is $L_{\\rm bol}/c$, in which $L_{\\rm bol}$ is the bolometric luminosity of the central galaxy.  For a typical star forming galaxy, $L_{\\alpha}\\sim 0.07L_{\\rm bol}$ \\citep[e.g.][]{PP67}, whereas for a galaxy that contains population III stars, $L_{\\alpha}\\sim 0.24L_{\\rm bol}$ \\citep{S03}. Hence, the Ly$\\alpha$ radiation pressure can dominate the maximum possible continuum radiation pressure if $M_F\\gsim 14$ (for a normal stellar population), a threshold which is easily exceeded at large column densities of relatively slow-moving HI shells (see Fig~\\ref{fig:fscat}).  The possibility that Ly$\\alpha$ radiation alone can result in a radiation force that exceeds $L_{\\rm bol}/c$ is important. \\citet{Murray05} have shown that the total momentum carried by radiation from a star-forming region can exceed the total momentum deposited by supernova explosions in it, and so galactic outflows may be driven predominantly by continuum radiation pressure. We have argued that Ly$\\alpha$ radiation pressure may in some cases be even more important than continuum radiation pressure, and thus provide the dominant source of pressure on neutral hydrogen in the ISM.  The important implication of our last result is that Ly$\\alpha$ radiation pressure may drive outflows of HI gas in the ISM. Observations of local starburst galaxies have shown that the presence of outflowing HI gas may be required to avoid complete destruction of the Ly$\\alpha$ radiation by dust and to allow its escape from the host galaxies \\citep{Kunth98,Hayes08,Ostlin08,Atek08}. At high redshifts, the Ly$\\alpha$ emission line of galaxies is often redshifted relative to other nebular recombination lines (such as H$\\alpha$) and metal absorption lines \\citep[e.g.][]{Pettini01,Shapley03}. Furthermore, the spectral shape of the Ly$\\alpha$ emission line is typically asymmetric, with emission extending well into the red wing of the line \\citep[e.g.][]{L95}. Both of these observations can be explained simultaneously if the observed Ly$\\alpha$ photons scatter off neutral hydrogen atoms in an outflowing 'supershell' of HI gas that surrounds the star forming regions \\citep{L95,T99,Ahn03,Ahn04,Verhamme06,Verhamme08}. The possibility that Ly$\\alpha$ radiation pressure may be important in determining the properties of expanding supershells is exciting, and is discussed in more detail in a companion paper (Dijkstra \\& Loeb 2008).   {\\bf Acknowledgments} This work is supported by in part by NASA grant NNX08AL43G, by FQXi, and by Harvard University funds. We thank Christian Tapken and an anonymous referee for helpful constructive comments."
    },
    "0807/0807.1316_arXiv.txt": {
        "abstract": "Due to their small radii, M-dwarfs are very promising targets to search for transiting super-Earths, with a planet of 2 Earth radii orbiting an M5 dwarf in the habitable zone giving rise to a 0.5\\% photometric signal, with a period of two weeks.  This can be detected from the ground using modest-aperture telescopes by targeting samples of nearby M-dwarfs.  Such planets would be very amenable to follow-up studies due to the brightness of the parent stars, and the favourable planet-star flux ratio.  MEarth is such a transit survey of $\\sim 2000$ nearby M-dwarfs.  Since the targets are distributed over the entire (Northern) sky, it is necessary to observe them individually, which will be done by using 8 independent 0.4m robotic telescopes, two of which have been in operation since December 2007 at the Fred Lawrence Whipple Observatory (FLWO) located on Mount Hopkins, Arizona.  We discuss the survey design and hardware, and report on the current status of the survey, and preliminary results obtained from the commissioning data. ",
        "introduction": "M-dwarfs are very favourable targets to search for transiting super-Earth exoplanets in the habitable zones of their parent stars (e.g. \\cite[Charbonneau \\& Deming 2007]{cd07}).  This is predominantly a result of the very low luminosities of M-dwarfs relative to solar type stars, which means the habitable zones reside at much smaller orbital distances.  For the present discussion, we assume an M5-dwarf star, with mass and radius $\\sim 0.25\\ \\msun$ and $0.25\\ \\rsun$ respectively, orbited by a planet of $7\\ \\mearth$ and $2\\ \\rearth$, representing a typical target of the MEarth survey.  Given these parameters, the M-dwarf has a luminosity $\\sim 1/200$ that of the Sun, and consequently, a planet receiving the same stellar insolation as the Earth would lie at $0.074\\ {\\rm AU}$, corresponding to an orbital period of $14.8\\ {\\rm days}$.  The geometric transit probability for such a planet is also raised by a factor of approximately three, from $0.5 \\%$ (for the Earth-Sun system) to $1.6\\ \\%$.  The small radii of M-dwarfs also give rise to substantially deeper transits ($0.5\\ \\%$ for the M5-dwarf, compared to $0.03\\ \\%$ for a solar-type host), and the combination of a smaller stellar mass and shorter orbital period increases the radial velocity semiamplitude, to $\\sim 10\\ \\mpers$ compared to $1.3\\ \\mpers$ for the Earth-Sun system.  The combination of these factors means that the detection of such planets around M-dwarfs by the transit method is feasible from the ground, using present-day observational techniques and detectors. MEarth is such a survey using modest commercially-available equipment, based on a design study published in \\cite{nc08}.  In this contribution, we summarise the results of the design study, present the novel mode of operation of the MEarth survey, and describe the implementation and current status of the hardware. ",
        "conclusions": ""
    },
    "0807/0807.4533_arXiv.txt": {
        "abstract": "We present theoretical modelling for the very rapid TeV variability of PKS 2155--304 observed recently by the H.E.S.S. experiment. To explain the light-curve, where at least five flaring events were well observed, we assume five independent components of a jet that are characterized by slightly different physical parameters. An additional, significantly larger component is used to explain the emission of the source at long time scales. This component dominates the emission in the X-ray range, whereas the other components are dominant in the TeV range. The model used for our simulation describes precisely the evolution of the particle energy spectrum inside each component and takes into account light travel time effects. We show that a relatively simple synchrotron self-Compton scenario may explain this very rapid variability. Moreover, we find that absorption of the TeV emission inside the components due to the pair creation process is negligible. ",
        "introduction": "Among the different models proposed to explain X-ray and gamma-ray emission of TeV blazars one distinguishes two classes: the leptonic and hadronic models (e.g. \\citealt{Krawczynski04}). This division depends on the assumption one makes about the particles that initially carry the energy that is to be converted into electromagnetic emission during the evolution of the source.  The aim of this paper is to explain the rapid variability in the simplest possible way in order to constrain the physical parameters of the source. Therefore, we have decided to use a relatively simple leptonic scenario.  A very basic leptonic model that can be used to explain the high energy emission of TeV blazars assumes a spherical source filled with a tangled magnetic field and relativistic electrons. This source is thought to travel with relativistic velocity at a distance of less than 1 pc from the central black hole and is usually assumed to be homogeneous.  The electrons inside the source produce synchrotron emission and also up-scatter part of this emission to TeV energies. This is the well known synchrotron self-Compton process (SSC) that is often applied to describe X-ray and gamma-ray observations of blazars (e.g. \\citealt{Bloom96}, \\citealt{Ghisellini96}, \\citealt{Inoue96}, \\citealt{Mastichiadis97}, \\citealt{Krawczynski00}).  The simplicity of this model allows to constrain some important physical parameters of the source (e.g. magnetic field strength, Doppler factor) directly from observations (e.g. \\citealt{Bednarek97}, \\citealt{Tavecchio98}, \\citealt{Katarzynski01}). On the other hand, some uncertainty remains in such estimations because the model is too simple; it does for example not take into account significant processes such as the particle evolution inside the source or light crossing time effects (hereafter LCTE). Therefore it is very important to analyse not only the emission observed at a given time but also the evolution of the emission, especially during the periods of activity. This may provide additional constraints for the models and result in a more precise estimation of the physical parameters.  Several leptonic models have been proposed to explain the time-dependent emission of TeV blazars (e.g. \\citealt{Dermer97}, \\citealt{Kirk98}, \\citealt{Coppi99}, \\citealt{Chiaberge99}, \\citealt{Kusunose00}, \\citealt{Kataoka00}, \\citealt{Sokolov04}). So far, no model is able to precisely describe SSC emission of an inhomogeneous or even a homogeneous source. The main problem appears in the description of the synchrotron radiation field inside the source that is inhomogeneous even in a spherically homogeneous slowly evolving source (\\citealt{Gould79}). Moreover, this radiation field depends on the particle energy distribution and the particle energy distribution depends in turn on the radiation field. This is certainly only true if the energy density of the radiation field is comparable to or larger than the energy density of the magnetic field, but this condition is fulfilled in TeV blazars. Finally, very rapid variability time scales are usually assumed to be comparable to the light crossing time at the source. This means that LCTE must be taken into account in calculations of the inverse-Compton scattering. In other words, the radiation field at a given time and a given position inside the source depends not only on the local physical conditions, but also contains photons from other parts of the source. These photons were already created at an earlier time and have propagated to the observed position.  We will call this process {\\it internal} LCTE. This effect is especially important for the calculation of the inverse-Compton emission and somewhat less for the synchrotron emission (although the radiation field has generally also an impact on the evolution of the electron energy spectrum). The internal LCTE was taken into account for the first time in the model proposed by \\citet{Sokolov04}. However, this model assumes dominance of the energy density of the magnetic field over that of the radiation field. In this case the impact of the radiation field on the electron energy spectrum is negligible.  A second LCTE, which we call {\\it external}, is equally important for both the synchrotron and the inverse-Compton emission. The external observer receives at a given time the emission produced by different parts of the source at different times in the comoving frame of the source. This is due to the different travel times that photons from different parts of the source need to reach the observer. If the emission level of the source does not change, the observer will always receive the same amount of radiation; if it changes on long time scales, differences in photon travel times can be neglected. However, if there is a change of the emission level with a duration shorter than the light crossing time of the source, the different travel times of the photons must be taken into account. One of the first attempts to account for this effect was made by \\citet{Chiaberge99}. It should be noted that in their model they neglected internal LCTE, which is a good enough approximation under certain conditions, as we will explain below in the description of the model.  Another important problem in time dependent SSC modelling is the description of the particle acceleration and evolution of the particle energy spectrum. It is widely believed that the particles are accelerated by shock waves created by colliding components of a jet. In such collisions, some fraction of the kinetic energy related to the bulk motion of the components is transformed into random kinetic energy of the particles. This energy is then radiated away. The acceleration and evolution of the spectrum are usually described by the kinetic equation, which is a partial differential equation (e.g. \\citealt{Kardashev62}). There are two general approaches to describe the acceleration in the equation.  In the first approach, the acceleration term in the equation describes the process inside a shock wave. This approach was used for example by \\citet{Kirk98}, where the evolution of the source was represented by two kinetic equations. The first equation described the acceleration inside the shock and the second the evolution of the particle spectrum in the downstream region of the shock. Particles were escaping into this region after the acceleration, creating a source for electromagnetic emission.  However, this is a rather complex approach and was applied only for the synchrotron emission.  The second approach is simpler.  There is no acceleration term, i.e. no description of the acceleration process. Instead, an ``injection term'' defines directly the result of the acceleration. One assumes a particle energy spectrum created by a hypothetical shock and describes it with a few free parameters. This approach was used for example in the model proposed by \\citet{Chiaberge99}, where different types of injection were investigated.  Since there is no perfect solution for modelling of variability in the framework of the leptonic SSC scenario, we have chosen here the relatively simple model proposed by \\citet{Chiaberge99}. The model, as we will show, provides very reasonable results while using only a minimum number of free parameters. ",
        "conclusions": "We have shown in this paper that the rapid TeV variability of PKS 2155--304 can be well explained using a standard SSC approach while taking into account the particle evolution and the external LCTE. In our approach the internal LCTE is of minor importance. The model we use requires in principle four free parameters to describe a single activity event. Therefore, we may compare the physical parameters derived from such modelling with the parameters derived form previous estimations of blazar emission (e.g. \\citealt{Tavecchio98}). This shows that our parameters are similar to other results (e.g., $B \\leq 0.1$ G and $\\delta = 20 \\to 30$); only the size ($R \\simeq 5 \\times 10^{14}$ cm) of the small sources is about one order of magnitude smaller than the values estimated so far."
    },
    "0807/0807.3125_arXiv.txt": {
        "abstract": "Current upper bounds of the neutron electric dipole moment constrain the physically observable quantum chromodynamic (QCD) vacuum angle $|\\bar\\theta| \\lesssim 10^{-11}$. Since QCD explains vast experimental data from the 100 MeV scale to the TeV scale, it is better to explain this smallness of $|\\bar\\theta|$ in the QCD framework, which is the strong \\Ca\\Pa~ problem. Now, there exist two plausible solutions to this problem, one of which  leads to the existence of the very light axion. The axion decay constant window, $10^9\\ {\\gev}\\lesssim F_a\\lesssim 10^{12}\\ \\gev$ for a ${\\cal O}(1)$ initial misalignment angle $\\theta_1$, has been obtained by astrophysical and cosmological data. For $F_a\\gtrsim 10^{12}$ GeV with $\\theta_1<{\\cal O}(1)$, axions may constitute a significant fraction of dark matter of the universe. The supersymmetrized axion solution of the strong \\Ca\\Pa~ problem introduces its superpartner the axino which might have affected the universe evolution significantly. Here, we review the very light axion (theory, supersymmetrization, and models) with the most recent particle, astrophysical and cosmological data, and present prospects for its discovery. ",
        "introduction": "Strong interaction phenomena has revealed that the discrete symmetries of charge conjugation \\Ca, parity \\Pa~ and time reversal \\Ti~ are separately good symmetries of nature. Therefore, quantum chromodynamics(QCD) based on the gauge group SU$(3)_c$ \\citep[Han, Nambu (1965), Bardeen, Fritszch, Gell-Mann (1972)]{Han65,Bardeen72} must respect any combinations of these discrete symmetries \\Ca, \\Pa~ and \\Ti~ to be accepted as the theory of strong interactions. Among these discrete symmetries, the \\Ca\\Pa~ symmetry is not necessarily respected in QCD due to the nonzero QCD vacuum angle $\\theta$, an issue known as the ``strong \\Ca\\Pa~ problem\". Since QCD is so successful phenomenologically, any possible solution to $\\lq\\lq$the strong  \\Ca\\Pa~ problem\" is expected to be realized in nature. Currently the most attractive solution leads to the existence of a very light axion \\citep[Kim (1979), Shifman, Vainstein, Zakharov (1980), Dine, Fischler, Srednicki (1981), Zhitnitskii(1981)]{Kim79,Shifman80,DFS81,Zhit80}. Searches for  QCD axions generated from the Sun \\citep[Andriamonje {\\it et. al.} (CERN Axion Search Telescope (CAST) Collaboration, 2007), Inoue {\\it et. al.} (Tokyo Axion Helioscope Collaboration, 2008)]{CAST07,Inoue08} and remnant axions from the early universe \\citep[Rosenberg (Axion Dark Matter Experiment(ADMX) Collaboration, 2004), Carosi (ADMX Collaboration, 2007)]{ADMX04,ADMX07} are presently ongoing.  The story of axions started with the QCD U(1) problem \\citep[Weinberg (1975)]{Wein75U1} which is now understood, having been solved by the 't Hooft determinental interaction \\citep['t Hooft (1976,1986)]{tHooft76,tHooftU1}. The determinental interaction is shown pictorially as the left diagram of Fig. \\ref{fig:U1tHooft} and the solution is shown as the shaded right diagram.  The strong interaction makes the quark bilinears condense with the vacuum expectation value (VEV) of order $v\\simeq 260$ MeV. The phase of this interaction $\\bar\\theta$ originates from the QCD vacuum angle which is known to be physical \\citep[Callan, Dashen, Gross (1976), Jackiw, Rebbi (1976)]{Callan76,Jackiw76} and contributes to the neutron electric dipole moment (NEDM) by an order $\\bar\\theta$ times the neutron size, which is absurdly large. Peccei and Quinn(PQ) observed that there exists a way to make $\\bar\\theta$ just a phase by introducing a symmetry, now called \\UPQ, and hence physical amplitudes do not depend on  $\\bar\\theta$ as in the massless quark case \\citep[Peccei, Quinn (1977)]{PQ77,PQ77b}. In the standard model (SM), this phase is a pseudoscalar Goldstone boson called the `axion' in the multitude of Higgs fields as noted in \\citep[Weinberg (1978), Wilczek (1978)]{Wein78,Wil78}. If the PQ idea is completed with Fig. \\ref{fig:U1tHooft}, this axion is exactly massless (but observable), and  $\\bar\\theta$ would behave  `unphysically' by the freedom of choosing an appropriate axion VEV, which was the original PQ idea. However, there exists subleading terms, proportional to one power of $m_q$, which close the quark lines with the current quark mass instead of a condensation. Then, an axion potential develops, and the axion becomes a pseudo-Goldstone boson. The axion solution of the strong \\Ca\\Pa~ problem is cosmological in that the axion VEV chooses $\\bar\\theta=0$ at the minimum of this axion potential. The currently allowed axion is very light and long lived.  \\begin{figure}[!] \\resizebox{0.9\\columnwidth}{!} {\\includegraphics{RMPfig01}} \\caption{The determinental interaction of light quarks. Chiral symmetry breaking introduces the anomalous $\\eta'$ mass term from the quark condensations.}\\label{fig:U1tHooft} \\end{figure}  The properties of the axion (denoted as $a$) are chiefly given by its decay constant $F_a$ which sets the scale of nonrenormalizable axion interactions through $a/F_a$. Initial axion searches placed $F_a$ far above the electroweak scale and additional stringent bounds on $F_a$ were obtained from the studies of stellar evolutions and cosmology \\citep[Kim (1987)]{Kimrev}. Axion astrophysics started by \\citep[Dicus, Kolb, Teplitz, Wagoner (1978, 1980)]{Dicus78,Dicus80}, using earlier ideas from \\citep[K. Sato, H. Sato (1975), K. Sato (1978)]{Sato75,Sato78}, now gives a stringent lower bound on the decay constant, $F_a\\ge 0.5\\times 10^9$ GeV, from the study of SN1987A \\citep[Turner (1990), Raffelt (1990)]{Turnerrev90,Raffeltrev90}. With this large decay constant, the axion flux from the Sun is a small fraction of the solar neutrino flux, but still may be detectable by the CAST experiment and by the Tokyo helioscope.  It is known that very light axions with $F_a$ in the $10^{12}$ GeV region (axion mass in the micro-eV range) can compose some part of cold dark matter(CDM) in the unverse \\citep[Preskill, Wise, Wilczek (1983), Abbott, Sikivie (1983), Dine, Fischler (1983)]{Preskill83,Abbott83,DineFish83}. The exact amount of the axion CDM depends on the initial axion misalignment angle $\\theta_1$ at the time of axion creation when the universe temperature was around the axion decay constant, $T\\sim F_a$. This observation puts the very light axion on the list of leading CDM candidate particles. If indeed these cosmic axions compose a significant fraction of CDM in the universe, they may be detectable by collecting axion-converted photons in the cavity type detectors \\citep[Sikivie (1983)]{SikDet} as tried by \\citep[DePanfilis {\\it et al.} (RBF Collaboration, 1987), Hagmann  {\\it et al.} (Univ. of Florida, 1990)]{DePanfilis87,Hagmann90} and now continuing at the ADMX experiment.  Cosmology with CDM was the leading candidate of the early universe in the 1980s \\citep[Blumenthal, Faber, Primack, Rees (1984), Kolb, Turner (1990), Weinberg (2008)]{Blum84,KolbTur90,Wein08}. Since then  this view has given way to the new cosmology with the discovery of dark energy(DE) in 1998 \\citep[Perlmutter {\\it et al.} (1998), Riess {\\it et. al.} (1998)]{Perlmutter98,Riess98}. The current view of the dominant components of the universe is $\\Omega_{\\rm CDM}\\simeq 0.23$ and $\\Omega_\\Lambda\\simeq 0.73$ with only a few percent consisting of baryons \\citep[Spergel {\\it et al.} (Wilkinson Microwave Anisotropy Probe(WMAP) Collaboration, 2007)]{WMAPsigma}. The most plausible dark matter (DM) candidates at present are the lightest supersymmetric(SUSY) particle(LSP), the axion, the axino, and the gravitino. Here, we will review mostly on the axion and its CDM-related possibility.  The need for DM was suggested as early as the 1930s \\citep[Zwicky (1933), Smith (1936)]{Zwicky33,Smith36}. Since then, evidences of nonluminous DM in the universe has been accumulating: examples of which include flat galactic rotation curves, Chandra satellite photos, and gravitational lensing effects. If the galactic bulge is the dominant mass in the galaxy, the velocity $v$ of a star located at $r$ from the center should rotate with  $v\\sim r^{-1/2}$. But the observed flat rotation curve \\citep[See, for example, McGaugh {\\it et al.} (2007),  and references therein]{flatrot07} violates this expectation and implies an extended mass in the halo as $\\rho(r)\\sim 1/r^2$. Also, the Chandra observation of X-ray and gravitational lensing images implies this matter profile around the bullet cluster \\citep[Clowe {\\it et al.} (Chandra Collaboration, 2006)]{Chandra06}. Circular gravitational lensing images \\citep[Jee {\\it et al.} (2007)]{gravlens07}  also support the existence of DM. The DM density around us the Solar system is usually taken as $\\rho_{\\rm DM}\\simeq (0.3-0.45)$ GeV/cm$^3$.  Current CDM candidates belong to either incoherent particles or coherent oscillations of spin-0 fields. From this disctinction, the bosonic collective motion such as the axion can be considered as CDM.  The popular incoherent CDM particles are the weakly interacting massive particles(WIMP) or decay products of WIMPs. A more frequently used independent distinction is thermal relics and nonthermal relics, but there is no strict correspondence relation between the incoherent--coherent particles and the thermal--nonthermal relics. WIMPs are massive particles with weak interaction cross sections, first discussed in terms of a heavy neutrino, corresponding to the RHS crossing point of Fig. \\ref{fig:LeeWein}(a) \\citep[Lee, Weinberg (1977)]{LeeWein77b}. The LHS crossing point corresponds to 10 eV neutrino \\citep[Marx, Szalay (1972), Cowsik, McClelland (1972)]{Marx72,Cowsik72}. WIMPs, such as the LSP, are thermal relics when their number density is determined by the freezeout temperature and are called nonthermal relics if their number density is determined by the other mechanism such as by the decay of heavier relics \\citep[Choi, Kim, Lee, Seto (2008)]{ChoiKY08prd}.  In Fig. \\ref{fig:LeeWein}(b), we sketch the axion energy density in terms of the axion mass. The shape is a flipped one from that of Fig. \\ref{fig:LeeWein}(a), because in the axion case the low mass and high mass regions contributes $\\Omega_a$ from different physics, one from the vacuum misalignment and the other from the hot thermal relics.  \\begin{figure}[!] \\resizebox{0.9\\columnwidth}{!} {\\includegraphics{RMPfig02}} \\caption{The Lee-Weinberg type plots for (a) the neutrino $\\Omega_\\nu h^2$  \\cite{KolbTur90} and (b) the axion $\\Omega_a h^2$, where $h$ is the present Huble constant in units of $\\rm 100\\, km\\, s^{-1}Mpc^{-1}$. The dash line of (a) is for $\\Omega_\\nu h^2=0.113$. In (b), it corresponds to the hadronic axion. The blue and red dashlines correspond to the CDM and hot DM limits, respectively. }\\label{fig:LeeWein} \\end{figure}  In addition to the heavy neutrino, SUSY with R-parity conservation allows the LSP to be just such a WIMP particle. The LSP interaction is $\\lq\\lq$weak\" since the interaction mediators (SUSY particles) are supposed to be in the 100 GeV range. For a WIMP to be a successful CDM candidate, usually the interaction cross section at the time of decoupling needs to be around \\citep[Kolb, Turner (1990), Spergel {\\it et al.} (WMAP Collaboration, 2007)]{KolbTur90,WMAPsigma} \\begin{align} &\\langle\\sigma_{\\rm int} v\\rangle|_{\\rm at\\ decoupling}\\approx 0.2\\times 10^{-26} ~{\\rm cm^3s}^{-1},\\nonumber\\\\ &\\quad\\quad {\\rm with}\\ \\Omega_m h^2\\simeq 0.113\\pm 0.009.\\label{WIMPcond} \\end{align} This is roughly the cross-section for the LSP from low energy SUSY, which is the reason why the DM community is so interested in the WIMP LSP. Some superweakly interacting particles such as gravitinos, axinos, and wimpzillas \\citep[Chung, Kolb, Riotto (1999)]{ChungWzilla} can be CDM candidates as well, but their cross sections do not fall in the range of Eq. (\\ref{WIMPcond}).  The CDM candidate particles are shown in the $\\sigma_{\\rm int}$ versus mass plane in Fig. \\ref{fig:sigmavsmass} with minor modification from that of \\citep[Roszkowski (2004)]{Roszfig04}. The incoherent fermions, such as the neutrino and left ends of the bars of the axino and gravitino correspond to the left crossing points of Fig. \\ref{fig:LeeWein}(a). The rest, except the axion, corresponds more or less to the right crossing points of Fig. \\ref{fig:LeeWein}(a) with the reheating after inflation considered if necessary.  Currently, there are experimental efforts to discover the LSP as predicted by SUSY models. Direct cosmological searches are also ongoing \\citep[Jungman, Kamionkowski, Griest (1996), Bertone, Hooper, Silk (2005), Bernabei {\\it et al.} (DAMA collaboration, 2003, 2008), Lee {\\it et al.} (KIMS Collaboration, 2007), Angle {\\it et al.} (XENON Collaboration, 2008), Behnke {\\it et al.} (COUPP Collaboration, 2008), Ahmed {\\it et al.} (CDMS Collaboration, 2008)]{Jungman96,Bertone05,Dama03,Dama08,KIMSCo07,Xenon10, COUPP08,CDMS08}. At the LHC, the probable LSP mass ranges will be looked for by the neutralino decay producing the LSP.   \\begin{figure}[!] \\resizebox{0.9\\columnwidth}{!} {\\includegraphics{RMPfig03}} \\caption{Some proposed particles in the interaction cross section versus the corresponding particle mass $m_i$ plane. The skeleton is taken from Ref. \\cite{Roszfig04}. The dashed curves represent schematic shapes of $\\Omega_i$ versus the corresponding particle mass $m_i$. The small red square-box corresponds to the hot DM hadronic axion. Two small outside squares (in the cyan and blue colors) in the axion region are marked to show the plausible (GUT and CDM) axions, respectively. The abundances of heavy axino, gravitino and wimpzilla depend on how inflation ends. }\\label{fig:sigmavsmass} \\end{figure}   It is known that density perturbations began growing much earlier than the time of recombination in order to become large enough to form galaxies in the young universe. For galaxy formation, therefore, DM is needed since proton density perturbations could not grow before the recombination time, but DM perturbations could. With DM, the equality point of radiation and matter energy densities can occur much earlier than the recombination time since DM is not prohibited in collapsing by Silk damping \\citep[Silk (1968)]{Silkdamp68}. If the WIMP mass and interaction cross section fall in the region allowed by Eq. (\\ref{WIMPcond}), it can be CDM. If the LSP is the only CDM component, then the LSP mass would give one number for the DM density, which may not be the case. Thus, even if the LSP is contributing to the CDM density, we may need the axion to account for the correct amount of CDM around us. This is possible in the anthropic scenario of very light axions because it is equally probable for the initial axion misalignment angle $\\theta_1$ to take any value between 0 and $\\pi$ \\citep[Tegmark, Aguirre, Rees, Wilczek (2006)]{Tegmark06}.  Here we review the axion, which is probably the most interesting Nambu-Goldstone boson \\citep[Nambu (1960), Nambu, Jona-Lasinio (1961), Goldstone (1961)]{Nambu60, Nambu61, Goldstone61}, as well as related issues. In Sect. \\ref{sec:StrongCPsols} we discuss the strong \\Ca\\Pa~ problem and its plausible solutions. In Sect. \\ref{sec:axions} we review the most attractive solution giving the very light axion and present the axion theory in terms of possible axion couplings defined with $c_1,c_2,$ and $c_3$ which will be used throughout this review. In Sect. \\ref{sec:axionsOuterspace} we present axion astrophysics and cosmology. Here we present a new number for the cosmic axion abundance in view of the recent more accurate data on light quark masses.  In Sect. \\ref{sec:AxionDet} we summarize the axion detection ideas and the ongoing axion detection experiments. In Sect. \\ref{sec:AxionModels} we summarize the proposed very light axion models, including superstring axions. Finally in Sect. \\ref{sec:Axinocosmology} we briefly discuss cosmology with the axino, the axion's superpartner.  If the axion was observed, it would mark one of the most profound elementary particle discoveries because it would confirm experimentally the instanton--based arguments of QCD. In addition, if it were shown to be consistent with a cosmologically significant amount of axions, the CDM idea of the bosonic collective-motion would also be confirmed experimentally. If SUSY is correct and the axion is the solution to the strong \\Ca\\Pa~ problem, it must affect the evolution of the universe as well. ",
        "conclusions": ""
    },
    "0807/0807.4828_arXiv.txt": {
        "abstract": "With more than 80 transits observed in the CoRoT light curve with a cadence of 32~s, CoRoT-Exo-2b provides an excellent case to search for the secondary eclipse of the planet, with an expected signal of less than 10$^{-4}$ in relative flux. The activity of the star causes a modulation on the flux that makes the detection of this signal challenging. We describe the technique used to seek for the secondary eclipse, that leads to a tentative 2.5$\\sigma$ detection of a 5.5$\\times$10$^{-5}$ eclipse. If the effect of the spots are not taken into account, the times of transit centers will also be affected. They could lead to an erroneous detection of periodic transit timing variations of $\\sim$20~s and with a 7.45~d period. By measuring the transit central times at different depths of the transit (transit bisectors), we show that there are no such periodic variations in the CoRoT-Exo-2b O-C residuals larger than $\\sim$10~s. ",
        "introduction": "Launched in December 2006, the CoRoT space mission is producing high quality light curves with outstanding duty cycles of better than 90\\% (see the Baglin et al. and Barge et al. papers in this volume). Out of the exoplanets detected up to date by this mission, CoRoT-Exo-2b (\\cite[Alonso et al., 2008]{alonso08}) is remarkable because of its relatively big size (1.47 R$_J$) and high mass (3.3 M$_J$). Moreover, with its 1.74~d period, CoRoT was able to record more than 80 transits of this planet, in a star that shows clear signs of activity. This activity is recognized because of 1) a $\\sim$2\\% modulation of the flux of the star with time, caused by the rotation of the star (with a period of $\\sim$4.5~d) and evolving star spots, and 2) the effect of star spots being occulted at almost every transit. Both effects are clearly seen in the Fig.~\\ref{fig1}. While carrying precious information about the activity of the star, these modulations of the light curve make the search for a secondary eclipse or the measurement of potential transit timing variations (TTV) a challenging task.  Secondary eclipses in the optical have been searched by different teams, providing upper limits to the albedos (or thermal emission in the infrared) of different exoplanets (e.g., \\cite[Charbonneau et al., 1999]{charb99} and  \\cite[Collier-Cameron et al. 1999]{collier99} using high S/N spectra, or \\cite[Rowe et al. 2008]{rowe08} using the MOST satellite). Both the results of these attempts and the theoretical models (e.g. \\cite[L\\'opez-Morales \\& Seager 2007]{merce07}, \\cite[Fortney et al. 2008]{fortney08}) predict a very low contribution of the planetary flux in the optical wavelengths; in the case of pM class planets as defined by \\cite[Fortney et al. 2008]{fortney08}, which are characterized by having an appreciable opacity due to Ti0 and VO gases, the planet's thermal emission in the optical is predicted to be more important than the reflected light from the star.  \\begin{figure}[t] \\begin{center} \\includegraphics[width=\\textwidth]{LC_combined.eps} \\caption{The CoRoT-Exo-2b light curve (top) and two selected regions showing the effects of the stellar activity, as a modulation of the light curve with the star rotational period (middle), and as alterations of the individual transits resulting from the occultation of stellar spots or active regions (bottom).} \\label{fig1} \\end{center} \\end{figure} ",
        "conclusions": "The activity of the host star of CoRoT-Exo-2b is limiting our capabilities to detect the secondary eclipse of this planet, and it can introduce spurious transit timing variation signals. We have made an analysis that shows a tentative 2.5$\\sigma$ detection of the secondary eclipse, and employed a new technique to properly measure the individual transit times when the star is active.  In the near future, the light curve of CoRoT-Exo-2b will still be further processed to correct for some residual noises, like the satellite jitter, that might help to improve the significance of the secondary eclipse detection. If the data were dominated by photometric photon noise, and not affected by the stellar activity, the expected S/N ratio for a secondary detection of the depth detected in this paper should be $\\sim$5~$\\sigma$. Additionally, the photometry of CoRoT-Exo-2b includes information on three different bandpasses (named red, green and blue). As we expect the secondary eclipse to be dominated by the thermal emission from the planet, this signal should be bigger in the red color than in the other two. Even if the correction for the satellite jitter is more difficult in the colored data than in the white light data analyzed in this paper, they might provide further constraints on the depth of the secondary eclipse. We plan to investigate also the behavior of several stars with similar brightness as CoRoT-Exo-2 that were observed with the same aperture shape, in order to search for potential uncorrected systematic noises.   \\begin{figure}[t] \\begin{center} \\includegraphics[width=9cm]{bisectors3.eps} \\caption{The transit bisectors of the 81 observed transits. A bigger dispersion of the centers of the transits is seen at the bottom of the transits, as these are the parts more affected by the stellar activity.} \\label{fig5} \\end{center} \\end{figure}"
    },
    "0807/0807.3255_arXiv.txt": {
        "abstract": "The recurrent nova RS Oph experienced an outburst in 2006, 21 years after its previous explosion in 1985, as expected. It was observed at almost all wavelengths, and important information about its properties is still being extracted. We present theoretical models of the explosion of this fascinating object, which indicate that the mass of the accreting white dwarf should be very close to the Chandrasekhar mass, to allow for such a short recurrence period. In addition, since models suggest that this nova ejects less mass than it accretes, it is an excellent candidate for a thermonuclear supernova explosion, in about $10^5-10^7$ years from now. We also analyze the emission of soft gamma-rays by RS Oph detected with the BAT instrument onboard Swift, and with the PCA onboard RXTE. We rule out that this emission has its origin in radioactive decays in the expanding nova envelope. ",
        "introduction": "\\label{intro}  Recurrent novae are defined as systems with more than one recorded nova outburst. There are few members of the recurrent novae subclass: just 10 (see \\citet{AM99} and review by \\citet{Miko08}), divided in two groups according to their orbital periods. RS Oph belongs to the long period group, with P$_{\\rm orb}=456$ days \\citep{Fekel00}, which are interpreted as interacting binaries  where a hot white dwarf accretes mass from a giant secondary; these recurrent novae are a subclass of symbiotic binaries. In contrast, classical novae occur in cataclysmic variables, a different type of binary system where a white dwarf accretes mass from a main sequence companion, overflowing its Roche lobe; typical periods of cataclysmic variables are about a few hours.  RS Oph has undergone various recorded outbursts (1898, 1933, 1958, 1967, 1985 and 2006, and two additional possible outbursts in 1907 and 1945). In the two last outbursts, a lot of observational data was gathered. The recent outburst in 2006 offered the opportunity to observe RS Oph at practically all wavelengths, radio, IR, optical, X-rays and soft gamma-rays. However, the very important UV range was missing, since there was not a satellite ``equivalent'' to IUE flying in 2006. It was a pity, since IUE made very important observations in the previous RS Oph eruption in 1985 \\citep{Shore96}, which could not be repeated in 2006. On the other hand, current space missions like Spitzer, RXTE, XMM-Newton, Chandra and Swift have provided a wealth of IR and X-rays data, of much better quality than what could be obtained at such energy ranges in the previous outburst. All in all, there is now an impressive data set; its interpretation  is a real challenge for theorists and it has provided already many exciting results.  It is worth noting that there is not yet a self consistent model of the RS Oph outburst itself, since it is not straightforward to find a set of initial conditions leading to a thermonuclear runaway (TNR) with a recurrence period as short as 21 years; classical novae recurrence periods are much longer ($\\ge 10^4-10^5$ years), well in agreement with the expected time to reach outburst conditions with typical accretion rates and initial white dwarf masses in cataclysmic variables. The only known way to reconcile a short interoutburst period -about some decades- with a TNR (i.e., explosive hydrogen burning in degenerate conditions on top of the white dwarf), is that a very massive white dwarf, close to the Chandrasekhar limit, accretes with a large mass accretion rate. This fact gives an additional interest to RS Oph and its relatives, since they are potential progenitors of thermonuclear -or type Ia- supernovae, in a relatively short time, if the white dwarf grows in mass in each nova eruption. Our evolutionary models indicate that ejected mass is smaller than accreted mass, so confirming that possibility.  In this paper we present our recent attempts to model the basic properties of the RS Oph recurrent nova outburst, with special emphasis on the ratio of ejected to accreted mass, the ejecta chemical composition and the potential gamma-ray emission. The role played by the initial conditions of the white dwarf is underlined. We also discuss the impossibility to explain the Swift/BAT hard X-soft gamma-rays detections as related with the continuum gamma-ray emission coming from electron-positron annihilation in the expanding ejecta. ",
        "conclusions": "\\label{conclu} Models of the RS Oph explosion should account for its main observed properties, e.g., very short recurrence period (21 years), small ejected mass, roughly solar chemical composition of the ejecta and maximum luminosities and ejecta velocities similar to those of classical novae. Hydrodynamic models of massive (larger than $\\sim$ 1.35 M$_\\odot$) accreting white dwarfs with large accretion rates - ranging from  $\\dot {\\rm M}= 10^{-8}$ to $2 \\times 10^{-7} {\\rm M_\\odot/year}$ - well reproduce these features. But there are still some aspects to be better understood, as the need of initially low white dwarf luminosities, as compared with some observations at quiescence. We want to point out a few important puzzles. First of all, very massive white dwarfs should be made of ONe (if they were born as massive white dwarfs); therefore, they pose a problem as scenarios for type Ia supernovae, because ONe white dwarfs rather collapse to a neutron star than explode, when reaching M$_{\\rm Chandra}$ (see \\citet{GCGB05} and references therein). One could envision that the white dwarf was initially born as a CO white dwarf, but then it should have been less massive than 1.1 M$_\\odot$ and be able to reach 1.35-1.38 M$_\\odot$ through successive recurrent nova eruptions. However, this seems unrealistic because models do not reproduce such behaviour. Second, the mass left after the recurrent nova explosion, i.e., M$_{\\rm acc}$- M$_{\\rm ejec}$ (see above) should be compatible with the duration of the super soft X-ray emission phase (around 60 days, see \\cite{HKL07}). According to our models \\citep{SH05}, the remnant envelope mass for a hydrogen mass fraction of 0.6 and mass of the white dwarf larger than 1.3 M$_\\odot$ should be around $10^{-7}$ M$_\\odot$, and surely smaller than $10^{-6}$ M$_\\odot$. Therefore, the mass of the white dwarf is more plausibly 1.38 than 1.35 M$_\\odot$ (see Table 1). Finally, the hard X/soft $\\gamma$-rays detected with Swift/BAT and RXTE in the early phases of the explosion could not come from electron-positron annihilation emission (with positrons coming from the decay of the $\\beta^+$-unstable nuclei $^{13}$N, $^{18}$F in the ejecta), both because of the tiny amount of radioactive nuclei, the too late appearence of such emission and its spectral shape. Deeper studies of both the thermonuclear runaway phase and the interaction between the ejecta and the dense stellar wind of the red giant are in progress, to better understand the wealth of observational data available for the fascinating RS Oph recurrent nova."
    },
    "0807/0807.1591_arXiv.txt": {
        "abstract": "We have analyzed sensitive high spatial resolution archival radio continuum data at 1.3, 2.0, 3.6 and 6.0 cm as well as the H$_2$O maser molecular line data obtained using the Very Large Array (VLA) in its hybrid AB configuration toward the high-mass star-forming region IRAS 17233-3606 (G351.78-0.54). We find nine compact radio sources associated with this region, six of them are new radio detections. We discuss the characteristics of these sources based mostly on their spectral indices and find that most of them appear to be optically thin or thick ultra- and hyper-compact HII regions ionized by B ZAMS stars. Furthermore, in a few cases  the radio emission may arise from optically thick dusty disks and/or cores, however more observations at different wavelengths are necessity to firmly confirm their true nature. In addition, we compared our centimeter maps with the mid-infrared images from the Spitzer Space Observatory GLIMPSE survey revealing a cluster of young protostars in the region together with multiple collimated outflows some of whom might be related with the compact centimeter objects. Finally, we find that one of these centimeter objects, VLA2d, is well centered in an apparent strong and compact north-south bipolar outflow traced by OH masers and we therefore suggest that this object maybe is energizing the latter. ",
        "introduction": "Attention to the high-mass star-forming region IRAS 17233-3606 (or G351.78-0.54) was first garnered by the detection of very strong maser emission from various species e.g. hydroxyl (OH) and water (H$_2$O) \\citep{CaswellHaynes1980, CaswellHaynes1983}. At the time of its detection (1980) it possibly contained the strongest known interstellar hydroxyl maser in the sky, with a peak intensity of 1000 Jy \\citep{CaswellHaynes1980}. Later on, intense class II methanol (CH$_3$OH) maser emission was detected \\citep{Menten1991}. From the IRAS fluxes and assuming a distance of 2 kpc it has been estimated that IRAS 17233-3606 has a bolometric luminosity of $\\sim$1.5 $\\times$ 10$^5$ L$_\\odot$ \\citep{HughesMacLeod1993}. However, more recent measurements with a better estimation of its kinematic near distance (of about 1 kpc) suggest a luminosity of 2.5 $\\times$ 10$^4$ L$_\\odot$ \\citep{MacLeodetal1998}. We here adopt a distance to IRAS 17233-3606 of 1 kpc, the near kinematic distance to the object reported by \\citet{MacLeodetal1998}. The reason that the far distance appears unlikely is because the source is located at an angular separation of more than 0.5$^\\circ$ from the Galactic plane.  \\citet{Argonetal2000}, \\citet{Argonetal2002} and \\citet{Fishetal2005} using VLA and VLBA radio observations of all the four hyperfine structure lines of hydroxyl (OH) around 1.6 GHz revealed that the maser spots clearly show a velocity gradient with a north-south orientation and in a range between -10.4 to 7 km s$^{-1}$. The masers are centered in the position R.A. = 17$^h$26$^m$42.7$^s $, decl.= -36$^\\circ$09$'$17.4$''$ (J2000.0). \\citep{Fishetal2005}. The large broad velocity gradient displayed by the OH masers, several times larger than in most sources \\citep{Argonetal2000} indicates that they could may be tracing a compact bipolar outflow.  \\citet{Norrisetal1993} found methanol maser spots coincident with the location of the OH masers that also appear to display a well-defined north-south velocity gradient. \\citet{ForsterCaswell1989} and \\citet{Forster1990} detected water maser spots about 3$''$ west of the hydroxyl and methanol maser location. The water masers cover a much broader velocity range, between $-38$ to +22 km s$^{-1}$ and seem to be tracing an ``expanding ring'' surrounding possibly a central massive protostar. \\citet{ForsterCaswell1989} proposed that the compact size ($\\sim$ 2$''$), wide velocity spread in the masers, and a lack of strong continuum emission indicate that this ``ring'' might be produced by a star in its early expansion phase. However, they also suggested that the geometry could  be due to one or multiple outflows.  \\begin{figure*}[ht] \\centering \\includegraphics[scale=0.32]{f1.eps}\\vspace{0.5cm} \\caption{\\scriptsize Spitzer infrared color images of the region IRAS 17233-3606. {\\bf Left:} Two image composite with red representing 8.0 $\\mu$m and green 4.5 $\\mu$m. The angular resolution for the 8.0 $\\mu$m and 4.5 $\\mu$m images are $\\sim$ 1$''$ and 2$''$, respectively. {\\bf Right:} 8.0 $\\mu$m image of the central part of the IRAS 17233-3606 region overlaid with the 6 cm continuum emission from the VLA. The contours are -3, 3, 4, 5, 6, 7, 8, 9, 10, 12, 14, 16, 18, 20, 40, 60, 80, 100, 150, and 200 times 0.1 mJy beam$^{-1}$, the rms noise of the image. The FWHM contour of the synthesized beam is shown in the bottom left corner of the image. The size of the synthesized beam is 1.39$''$ $\\times$ 0.88$''$ with a position angle (P.A.) of 47.3$^\\circ$ E of N. Please note the multiple powerful outflows and jets that emanate from this region traced by the 4.5 $\\mu$m band. This IRAC band (4.5 $\\mu$m) contains several H$_2$ transitions which trace shock-excited material \\citep{Smithetal2005}. The position of the yellow star in the right image marks the position of the radio binary star reported by  \\citet{HughesMacLeod1993} }\\vspace{1.0cm} \\label{fig1} \\end{figure*}  A few radio sources have been found in this region, \\citet{Haynesetal1979} and \\citet{CaswellHaynes1980} reported for the first time the possible presence of an HII region using observations at 2 cm. Subsequent observations at 2 cm by \\citet{Fixetal1982} confirmed this detection and found that this continuum source is extended and bright (310 mJy at 2 cm) and that it is located almost 10$''$ east of the OH/CH$_3$OH/H$_2$O maser zone. Later radio observations at 3.6 cm  with a better angular resolution also detected this extended continuum source and slightly resolved it in a compact HII region with a cometary morphology \\citep{Walshetal1998}.  \\citet{Fixetal1982} and \\citet{HughesMacLeod1993} report the detection of a very weak radio continuum source near to the location of the OH/CH$_3$OH/H$_2$O masers. The observations at 2 cm made by \\citet{HughesMacLeod1993} with a resolution of $0\\rlap.''3$ revealed that this object is actually a radio binary and that it is located very close to the maser positions.  \\citet{Fixetal1982} and \\citet{Walshetal1999} detected 2.0 $\\mu$m infrared sources toward IRAS 17233-3606. \\citet{Fixetal1982} reported an infrared object close to the location of the masers. Images by \\citet{Walshetal1999} find a source coincident with the OH/CH$_3$OH/H$_2$O masers and a source coincident with the extended HII region to the east. Finally, \\citet{DeBuizeretal2000} detected a compact and faint infrared object at 10 $\\mu$m located about 10$''$ east of the position of the masers. They also detected very diffuse and extended 18 $\\mu$m infrared emission about 3$''$ northeast of the OH/CH$_3$OH/H$_2$O maser location.  In this paper we present radio continuum observations made with the Very Large Array (VLA) toward the massive star forming region IRAS 17233-3606. We report the detection of new radio sources associated with this zone which are distributed among a cluster of strong infrared sources detected by the Spitzer Space Telescope. We discuss the nature of these radio objects based mainly on their radio spectral indices and counterparts at other wavelengths. In addition, we present imaging of H$_2$O maser emission associated with IRAS 17233-3606.  \\begin{deluxetable}{ccccc}[h] \\tablecolumns{5} \\tabletypesize{\\footnotesize} \\tablewidth{0pc} \\tablecaption{Summary of the Observations} \\label{table1} \\tablehead{ \\colhead{}& \\multicolumn{2}{c}{Angular Resolution$^b$} & \\colhead{Rms}& \\colhead{Largest Angular}\\\\ \\cline{2-3} \\colhead{} & \\colhead{Beam} & \\colhead{P.A.}& \\colhead{Noise}& \\colhead{Scale}\\\\ \\colhead{Band}& \\colhead{[arcsec]}& \\colhead{[degrees]}& \\colhead{[mJy Beam$^{-1}$]}& \\colhead{[arcsec]} } \\startdata 1.3cm & 0.27$\\times$0.25 & 7.1&0.45 & 4\\\\ 2.0cm & 0.51$\\times$0.42 & 9.7&0.28 & 7\\\\ 3.6cm & 1.48$\\times$0.43 & 40.3&0.06 & 15\\\\ 6.0cm & 1.39$\\times$0.88 & 43.7&0.03 & 20\\\\ Mas.$^a$ & 0.35$\\times$0.20 & 34.9&10 & 4 \\enddata \\tablecomments{ (a): The rms-noise is in a velocity channel.\\\\ (b): Mayor axis $\\times$ minor axis; position angle of mayor axis. } \\end{deluxetable}  \\begin{figure*}[!ht] \\centering \\includegraphics[scale=0.6]{f2.eps} \\caption{\\scriptsize VLA 1.3 cm continuum grey scale image overlaid with a blue-contour image at 2.0 cm of a central region of IRAS 17233-3606. The contours are -3, 3, 6, 9, 12, 15, 18, 20 and 25 times 0.30 mJy beam$^{-1}$, the rms noise of the image. The half-power contour of the synthesized beam is shown in the bottom left corner of the image. The size of the synthesized beam is 0.51$''$ $\\times$ 0.42$''$ with a P.A. of 9.7$^\\circ$ E of N. The grey scale bar on the top indicates the 1.3 cm continuum emission on mJy beam$^{-1}$. The crosses indicate the positions of the multiple radio and infrared sources presented in Figure 1. The 1.3 cm image was convolved to the same angular resolution as the 2.0 cm image, using the parameter \"UVTAPER\" from the task \"IMAGER\" of AIPS. }\\vspace{1.0cm} \\label{fig2} \\end{figure*} ",
        "conclusions": "In a region of about 1$'$ $\\times$ 1$'$, we detected a total of 5 compact  sources in one or more of the individual bands, as well as 16 water maser spots;  see Figures 1, 2, and 3.  However, two of the radio sources, VLA1 and VLA2, were resolved into 2 (VLA1a,b) and 4 (VLA2a,b,c,d) compact sources, respectively, see Figures 2 and 3. The observed parameters were determined from a linearized least-square fit to an elliptical Gaussian using the task IMFIT of AIPS.  Infrared two color band (8 and 4.5 $\\mu$m) images from the Spitzer Space Observatory's GLIMPSE survey \\citep{Benjaminetal2003} overlaid  with our 6 cm radio image of the IRAS 17233-3606 (both with similar angular reslotions) are presented in Figure 1. In this image a cluster of 10 bright  8 $\\mu$m infrared sources is detected. We give the positions and sizes of the sources in Table 2 obtained using the task IMFIT of AIPS. These infrared sources are possibly associated with the objects detected by \\citet{Fixetal1982}, \\citet{Walshetal1999} and \\citet{DeBuizeretal2000} at IR wavelengths. In a similar manner we give the observed parameters of the radio sources and the masers in Tables 3 and 4.  As it can be seen, some of these infrared sources appear to be related with the centimeter objects, however, the nature of these infrared objects will not be discussed in this article, they are presented here only for a reference.  In what follows, we discuss separately each of the 9 radio sources detected in this region. Our discussion of the spectral indices is based on the assumption that the flux densities did not change between the 1997 January 26 and 28 observations. Furthermore, due to the spectral index measurements presented here can be very sensitive to the largest angular size and the angular resolution of the observations, matching-beam observations of the sources at several frequencies are thus needed to discuss them in a very reliable way.  \\begin{figure*}[!ht] \\centering \\includegraphics[scale=0.4]{f3.eps} \\caption{\\scriptsize VLA 1.3 cm continuum image of the maser zone. The contours are -3, 3, 4, 5, 6, 7, 8, 9, 10, 11, 12, 13, 14 and 15 times 0.43 mJy beam$^{-1}$, the rms noise of the image. The half-power contour of the synthetized beam is shown in the bottom left corner of the image. The size of the beam is 0.27$''$ $\\times$ 0.25$''$ with a P.A. of 7.1$^\\circ$. The grey scale bar on the top indicates the 1.3 cm continuum emission on mJy beam$^{-1}$. The blue and red dots indicate the positions of the blue- and red-shifted OH maser spots, respectively  (Fish et al. 2005). Note that these masers appear to be tracing the innermost part of  a north-south compact outflow that shows a moderate velocity gradient, going from -10.4 to 7 km s$^{-1}$ (Fish et al. 2005) and that seems to emanate from the radio source VLA2d. The blue and red open squares (about 3$''$ west from the OH center maser) indicate the position of the water masers presented on this article (see Table 4) and that seem to be tracing an ``expanding ring'' (Foster 1990). The two black crosses mark the position of the objects VLA1a and b.  } \\label{fig3}\\vspace{1.5cm} \\end{figure*}  \\subsection{Classical Ultra-compact HII Regions}  \\subsubsection{VLA5}  As already mentioned in the introduction this source was discovered and reported as a faint HII region for the first time by \\citet{Haynesetal1979} and \\citet{CaswellHaynes1980}. Later observations with a better signal to noise and angular resolution by \\citet{Fixetal1982}, \\citet{Walshetal1998}, and \\citet{Argonetal2000} confirm its detection. However, this source is not related with the center of the OH, CH$_3$OH and H$_2$O masers ({\\it the maser zone}), but rather it is offset by about 10$''$ east , see Figures \\ref{fig2}, \\ref{fig3}, and the images from \\citet{Walshetal1998} and \\citet{Argonetal2000}.  Figure 2 shows the 1.3 and 2 cm continuum emission from the central portion of the IRAS 17233-3606 region, the HII region and the radio binary system reported by \\citet{HughesMacLeod1993} are shown in this image.  The HII region is resolved and shows a very clumpy cometary morphology. We give its deconvolved dimensions and flux densities for all bands in Table 3.  \\begin{figure*}[!ht] \\centering \\includegraphics[scale=0.35]{f4.eps} \\includegraphics[scale=0.35]{f5.eps}\\\\ \\includegraphics[scale=0.35]{f6.eps} \\includegraphics[scale=0.35]{f7.eps} \\caption{\\scriptsize Spectral energy distributions of the radio sources VLA2a-d from about 3 to 30 GHz (or 10 to 1 cm). The squares are detections. The squares with arrows are upper limits (4$\\sigma$). In cases where the error bars were smaller than the squares they are not presented.  The line is a least-squares power-law fit (of the form S$_\\nu$ $\\propto$ $\\nu^\\alpha$) to each spectrum. For VLA2a and b the flux densities at 14 and 23 GHz (2.1 and 1.3 cm) were obtained from maps convolved at a resolution of 1$''$. The spectrum for VLA2a can be also modeled by a two-component spectrum from a classical HII region plus thermal dust emission.} \\label{fig4} \\end{figure*}  The total 2 cm flux densities for this source reported by \\citet{Fixetal1982} and us (see Table 3) are 310 and 330 mJy, respectively. We consider the agreement very good. Furthermore, the 3.6 cm peak flux densities reported by \\citet{Walshetal1998}, and \\citet{Argonetal2000} also seem to be consistent with our measurements.  Taking the flux densities at 2 and 1.3 cm from Table 3, which both have a similar angular resolution (see Table 1), we estimated a very flat spectral index for this object, equal to 0.06, indicating thus that the emission might come from an optically thin UCHII region.  Following the derivation of \\citet{vanandmentenetal2005}, and assuming that the free-free emission at 2 cm is optically thin (which must be the case for the above calculation), that the ionized gas has a temperature of 10$^4$ K, that this source is located a distance of 1 kpc, and taken the values of the number of ionizing photons corresponding to the spectral type for a young star from \\citet{Panagiaetal1973}, we estimate that possibly the UCHII region is ionized by a B0.5 ZAMS star. This is again in very good agreement with the values and spectral types reported in \\citet{Walshetal1998} and \\citet{HughesMacLeod1993} .  From Figure 1 we can see that this HII region is coincident with the compact infrared sources, IRAC 8 and 9, suggesting that maybe they are the stars responsible for its ionization. However, both infrared sources seem to lie south and east of the UCHII region.  \\subsubsection{VLA3}  This source is first detected here and is imaged at 6 cm in Figure \\ref{fig1}. We note that VLA 3 at this wavelength is resolved and turns out to be an UCHII region with a shell-like morphology. We give its deconvolved dimensions and flux densities for all bands in Table 3.  In a similar way than VLA 5, taking the flux densities at 2 and 1.3 cm from Table 3 for VLA3 (both measurements with similar angular resolution), we estimated a spectral index for this object of $\\leq$ 1.1, consistent with the emission also coming from an optically thin HII region.  In addition, assuming again that the free-free emission at 2 cm is optically thin, that the ionized gas has a temperature of 10$^4$ K, and that this source is located  a distance of 1 kpc, we estimate that this HII region is possibly ionized by a B2 type ZAMS star.  From Figure 1, we also see that this UCHII region is coincident with one compact infrared source, IRAC 3.     \\subsection{On the Nature of VLA1 and VLA4}  These sources were only detected in our more sensitive maps, at 6 and 3.6 cm, see Figure \\ref{fig1} for an image at 6 cm. VLA4 is very compact and faint, while VLA1 is quite elongated in northeast-southwest orientation at 6 cm. Furthermore, VLA1 is resolved into two sources (VLA1a,b) in our 3.6 cm map with a better angular resolution. We give their flux densities and deconvolved sizes at this wavelength in Table 3.  From the values of the flux densities at 6 and 3.6 cm presented in Table 3, which also both have a similar angular resolution (see Table 1), we estimate a spectral index for VLA4 of $\\leq$ 1.0, and for VLA1a and b of $\\leq$ -0.8 suggesting that possibly the emission arises from optically thin HII regions or from ionized stellar winds (the latter case is more likely for VLA4). The radio emission from VLA1a and b also might be associated with synchrotron emission. This kind of emission is believed to trace gyrosynchrotron emission from active magnetospheres of young, low-mass stars. However, similar negative spectral indices also have found in sources associated with high- and low-mass stars and are believed to be produced in strong shocks, see \\citet{Garayetal1996}.  However, these flattened spectral indices could also be due to missing flux density in the higher angular resolution and higher frequency observations, see Table 1 where we include the largest scale structure to which the array is sensitive.  None of these radio sources have a counterpart at infrared wavelengths (8.0 and 4.5 $\\mu m$), see Figure 1. \\vspace{-0.5cm} \\begin{deluxetable*}{lcccccccc}[!h] \\tabletypesize{\\scriptsize} \\tablecolumns{9} \\tablewidth{0pc} \\tablecaption{Parameters of the Water Masers Spots} \\label{table3} \\centering \\tablehead{ \\colhead{Maser} & \\multicolumn{2}{c}{Position} & \\colhead{Central Vel.} & \\colhead{Deconvolved Size$^a$} & \\colhead{P.A.$^b$} & \\colhead{$\\int$ S$_\\nu$ dv}\\\\ \\cline{2-3} \\colhead{Feature}& \\colhead{$\\alpha_{2000}$}& \\colhead{$\\delta_{2000}$}& \\colhead{[Km s$^{-1}$]}& \\colhead{[arcsec]}& \\colhead{[deg.]}& \\colhead{[Jy beam$^{-1}$ km s$^{-1}$]} } \\startdata  1  & 17 26 42.155 & -36 09 11.76 & -1.0 & 0.14$\\pm$0.03 $\\times$ $\\leq$ 0.2        &  173$\\pm$80    &     2.1$\\pm$0.2 \\\\ 2  & 17 26 42.239 & -36 09 19.77 & -1.0 & 0.40$\\pm$0.01 $\\times$ $\\leq$ 0.2        &   30$\\pm$1     &    11.0$\\pm$0.2 \\\\ 3  & 17 26 42.280 & -36 09 17.57 &  0.5 & 0.49$\\pm$0.02 $\\times$ 0.16$\\pm$0.01     &  159.6$\\pm$0.2 &   181.2$\\pm$0.3 \\\\ 4  & 17 26 42.284 & -36 09 17.68 & -0.5 & 0.33$\\pm$0.01 $\\times$ $\\leq$ 0.3        &  140.8$\\pm$0.3 &   154.3$\\pm$0.3 \\\\ 5  & 17 26 42.326 & -36 09 16.46 & -1.5 & 0.06$\\pm$0.03 $\\times$ 0.01$\\pm$0.06     &  159$\\pm$2     &    89.8$\\pm$0.2 \\\\ 6  & 17 26 42.343 & -36 09 15.98 & 8/1.5/-19.5 & 0.28$\\pm$0.01 $\\times$ 0.02$\\pm$0.01 &   26$\\pm$2     &    13.6$\\pm$0.2 \\\\ 7  & 17 26 42.392 & -36 09 18.69 & -2.0 & $\\leq$ 0.5 $\\times$ $\\leq$ 0.5          &   --           &    23.0$\\pm$0.5 \\\\ 8  & 17 26 42.395 & -36 09 18.67 & 22.5/0.5/-10  & 0.330$\\pm$0.002 $\\times$ 0.052$\\pm$0.007&63.4$\\pm$0.5 &    64.3$\\pm$0.6 \\\\ 9  & 17 26 42.428 & -36 09 16.21 & -14.0 & 0.05$\\pm$0.01 $\\times$ $\\leq$ 0.02   &   35$\\pm$5     &    58.2$\\pm$0.2 \\\\ 10 & 17 26 42.438 & -36 09 18.57 & -2.5  & 0.37$\\pm$0.01 $\\times$ 0.05$\\pm$0.01 &   90.4$\\pm$0.3 &   138.3$\\pm$0.3 \\\\ 11 & 17 26 42.511 & -36 09 18.49 & -4.0  & 0.13$\\pm$0.01 $\\times$ 0.08$\\pm$0.01 &   27$\\pm$2     &    34.0$\\pm$0.2 \\\\ 12 & 17 26 42.539 & -36 09 17.28 & -3.5  & 0.13$\\pm$0.01 $\\times$ 0.02$\\pm$0.01 &  168.6$\\pm$0.5 &   119.0$\\pm$0.2 \\\\ 13 & 17 26 42.538 & -36 09 17.31 & -8.5/-13.5 & $\\leq$ 0.5 $\\times$ $\\leq$ 0.5  &   --           &    21.0$\\pm$0.5 \\\\ 14 & 17 26 42.798 & -36 09 20.35 & -11.5 & $\\leq$ 0.3 $\\times$ $\\leq$ 0.1           &   --           &     3.7$\\pm$0.5 \\\\ 15 & 17 26 42.902 & -36 09 09.06 & -5.5  & $\\leq$ 0.3 $\\times$ $\\leq$ 0.2           &   --           &     5.9$\\pm$0.5 \\\\ 16 & 17 26 42.955 & -36 09 09.15 & -4.5  & $\\leq$ 0.3 $\\times$ $\\leq$ 0.2           &   --           &     3.3$\\pm$0.5 \\\\ \\enddata \\tablecomments{Units of right ascension are hours, minutes, and seconds, and units of declination are degrees, arcminutes, and arcseconds.\\\\ (a): Mayor axis $\\times$ minor axis\\\\ (b): Position angle of mayor axis.} \\end{deluxetable*}   \\subsection{Radio Objects Associated with The Maser Zone}  In Figure 3 we show the 1.3 cm radio emission together with the position of the OH, CH$_3$OH and H$_2$O masers spots reported by \\citet{Fishetal2005}, \\citet{Forster1990} and \\citet{Caswelletal1997}. We found six compact radio sources in this region, VLA2a,b,c,d (VLA2a,b were detected for the first time by \\citet{HughesMacLeod1993} at 6 cm) and VLA1a,b. In Figure 2 we also show the radio sources detected by \\citet{HughesMacLeod1993}, VLA2a and VLA2b, they appear to form a binary system with a spatial separation of about 800 AU. However, the new observations show that possibly these sources are part of a multiple system. The sources VLA2c and d are only detected in our 1.3 cm images.  Figure 4 shows the SEDs for the centimeter compact sources VLA2a,b,c, and d. The flux densities values for the four wavelenghts were obtained from Table 3, however the values at 2 and 1.3 cm were obtained using a convolved map at a resolution of 1$''$, {\\it i.e.} with a angular resolution similar to those at 3.6 and 6 cm (see Table 1). We convolved the maps using the parameter \"UVTAPER\" from the task \"IMAGER\" of AIPS. This parameter set a Gaussian taper to weight down long baseline data points.  VLA2c and d have steeper spectral indices, $\\geq$ 2.2 consistent with either optically thick hyper-compact HII regions or with optically thick thermal dust emission from massive dusty cores and/or disks. If the emission is arising from thermal dust, the spatial sizes (about 300 AU) of these sources indicate that they are more likely to be compact circumstellar disks rather than dusty cores. VLA2b shows a negative spectral index $\\leq$ -0.15, possibly associated with synchrotron emission. The spectrum of the source VLA2a also can be modeled by a two-component spectrum where for the centimeter wavelengths is dominated for a classical HII region, while on the millimeter wavelengths is dominated by a component that rises rapidly with frequency. This component is likely to be associated with dust emission from a core or disk, see Figure 4.  From Figure 3 we note that the radio source VLA2d is well centered in an apparent compact bipolar north-south outflow traced by the OH maser spots. This again suggests that VLA2d maybe is a compact circumstellar disk rather than an hyper-compact HII region. We also note that there is a strong bow-shock about 40$''$ north of this compact outflow  and which is quite well aligned to this compact outflow traced by OH masers, see Figure 1. Possibly this bow-shock is part of an older ejection of this compact outflow. If this is the case, the outflow should have a dynamical age of  the order of 10$^4$ years, assuming a velocity of 10 km s$^{-1}$. Furthermore, \\citet{Leurinietal2008} found a strong north-south $^{12}$CO outflow using the Atacama Pathfinder EXperiment (APEX) radiotelescope and that is emanating from the maser zone, it is likely that this molecular compact outflow forms part of this bow-shock and the very compact outflow traced by the OH maser.  Finally, there is another highly collimated northwest-southeast jet that seems to be emanating from this region (see Figure 1), possibly also associated with these compact radio objects.  We do not find any 8 $\\mu$m infrared source associated with these four radio objects, indicating that maybe they are highly embedded in the molecular cloud. The extended 4.5 $\\mu$m source associated with this zone (see Figure 1) seems to be more likely associated with shocked gas possibly from the multiple outflows originating here   \\subsection{An Expanding Ring?}  As first suggested by \\citet{ForsterCaswell1989} and as can be observed in Figure 3 the water maser spots located to the east of the OH masers appear to be tracing a ``ring''  centered on the position R.A. = 17$^h$26$^m$42.4$^s$ decl.= -36$^\\circ$09$'$17.5$''$ (J2000.0) with a radius of approximately 1.5$''$. The maser spots associated with the ``ring'' are displaying broad red-shifted (-4 to 22.5 km s$^{-1}$) and blue-shifted velocities (-4.5 to -19.5 km s$^{-1}$), see Table 4. Moreover, systematically the blueshifted maser spots are located to the northeast, while the redshifted masers lie to the southwest. The velocities and size are in good agreement with those reported by \\citet{ForsterCaswell1989} and \\citet{Forster1990}. However, we think that such large radial velocity gradients ($\\Delta$V $\\sim$ 40 km s$^{-1}$) observed in the masers are too large to be explained in terms of a ring in expansion produced by an energetic stellar homogenous wind that compresses the ambient medium and drives a shock into it, and particularly since there is no a radio or infrared source. A possibility is that they are due to multiple outflows. \\citet{Torrellesetal2001} and \\citet{Uscangaetal2005} reported the presence of \"rings\" in expansion (on much smaller scales) traced by water masers in the high mass star forming regions Cepheus and W75N, respectively. The radial velocities of the masers are close to the ambient cloud velocity.  We therefore suggest that possibly the radio sources VLA2a and VLA1a that are well centered on the positions of the water masers might be exciting them. These radio sources thus perhaps are driving northeast-southwest compact outflows."
    },
    "0807/0807.1134_arXiv.txt": {
        "abstract": "We describe calculations for the formation of icy planets and debris disks at 30--150~AU around 1--3~\\msun\\ stars. Debris disk formation coincides with the formation of planetary systems. As protoplanets grow, they stir leftover planetesimals to large velocities. A cascade of collisions then grinds the leftovers to dust, forming an observable debris disk. Stellar lifetimes and the collisional cascade limit the growth of protoplanets. The maximum radius of icy planets, $r_{max} \\approx$ 1750 km, is remarkably independent of initial disk mass, stellar mass, and stellar age. These objects contain $\\lesssim$ 3\\%--4\\% of the initial mass in solid material. Collisional cascades produce debris disks with maximum luminosity $\\sim 2 \\times 10^{-3}$ times the stellar luminosity.  The peak 24~$\\mu$m excess varies from $\\sim$ 1\\% times the stellar photospheric flux for 1 \\msun\\ stars to $\\sim$ 50 times the stellar photospheric flux for 3 \\msun\\ stars. The peak 70--850 $\\mu$m excesses are $\\sim$ 30--100 times the stellar photospheric flux.  For all stars, the 24--160~$\\mu$m excesses rise at stellar ages of 5--20~Myr, peak at 10--50~Myr, and then decline. The decline is roughly a power law, $ f \\propto t^{-n}$ with $n \\approx$ 0.6--1.0. This predicted evolution agrees with published observations of A-type and solar-type stars. The observed far-IR color evolution of A-type stars also matches model predictions. ",
        "introduction": "During the past 25 years, observations from {\\it IRAS}, {\\it ISO}, and {\\it Spitzer} have revealed substantial mid-infrared (mid-IR) excesses associated with hundreds of normal main sequence stars \\citep[e.g.,][]{bac93,hab01,rie05,bry06,moor06,rhe07a}. Current samples include stars with spectral types A--M and ages $\\sim$ 5~Myr to $\\sim$ 10 Gyr \\citep[e.g.,][]{che05,kim05,rie05,bei06,su06,hil08}. Although binary stars and single stars in dense clusters and in the field are roughly equally likely to have IR excesses \\citep{stau05,su06,bry06,gor06,cur07a,sie07,tri07}, the frequency of excess emission declines from $\\sim$ 30\\%--40\\% for A-type stars \\citep{su06} to $\\sim$ 10\\%--20\\% for solar-type stars \\citep{gre03,tri08,mey08}. Thus, this phenomenon is common among main sequence stars and may depend on stellar mass.  High quality images demonstrate that dust orbiting the central star produces the excesses \\citep{smi84,bra04,sta04,kal05,mey07}. In $\\beta$ Pic and AU Mic, the dust is in a geometrically thin, edge-on disk with an outer radius of $a \\sim$ 200--1000~AU \\citep{smi84,tel88,gol93,kal04,liu04a,aug06}. In these disks, the small scale height of the dust, $H/a \\sim$ 0.1, is consistent with material in roughly circular orbits \\citep[e.g.,][and references therein]{bac93,kal04}.  Although broad tori of dust are visible in many other systems \\citep[e.g.][]{gre98,au99,hol03,kal06,su06,fit07}, narrow rings of dust produce the emission in $\\alpha$ PsA and HR 4796A \\citep{jay98,sch99,gre00,tel00,kal05b}. For systems with face-on rings and tori, the total emission constrains the scale height, $H/a \\sim$ 0.1. Thus, the dust in these systems is as highly flattened as the structures in $\\beta$ Pic and AU Mic.  Broadband spectral energy distributions constrain the luminosity, size, temperature, and total mass of the dust \\citep{bac93,lag00,den00,wol03}. Some stars have excesses from grains plausibly associated with the terrestrial zone \\citep[e.g.,][]{bei05,abs06,cur07b,rhe07b,mey08,lis07c}. Optically thin emission from cooler material with temperature $T \\sim$ 20--150~K is more typical \\citep[e.g.][]{su06,tri08,hil08}. For systems with submm observations, the measured fluxes suggest grains with sizes $\\sim$ 1~$\\mu$m--1~cm and total mass $\\sim$ 0.01~\\mearth\\ \\citep{liu04b,naj05,che06,wil06}. The grains have compositions similar to dust in the asteroid belt, comets, or the trans-Neptunian region of the Solar System \\citep{gru95,bro97,lis07a,lis07b}. Because the dust mass in these systems lies between the initial mass of solids in protostellar disks \\citep[$\\sim$ 100--1000~\\mearth;][]{nat2000,and05} and the dust mass in the Solar System \\citep[$\\lesssim 10^{-4}$ \\mearth;][]{hah02,lan02,nes06}, the dusty structures in these systems are often called `debris disks' \\citep{bac93,lag00}.  In addition to the dust properties, several other observations suggest plausible links between debris disks and the formation of planetary systems. Observations of A-type stars suggest a `rise and fall' of debris disk emission \\citep{cur08a}, with a clear increase in the typical 24~$\\mu$m excess at 5--10~Myr, a peak at 10--20~Myr, and a decline for $t \\gtrsim$ 20--30~Myr. The rise in debris disk emission roughly coincides with the disappearance of optically thick emission from protostellar disks \\citep[e.g.][]{hai01,sic06,her06,cur07a,cur07c}. The broad plateau occurs at a time when radiometric dating \\citep{yin02} and theory \\citep{cha01,kb06} suggest the Earth contained $\\sim$ 90\\% of its final mass. The decline of dusty debris around A-type stars is at roughly the same time as a gradual decrease in the cratering rate of objects in the Solar System \\citep{swi93,mel93,wad00,koe03}. These results suggest that the evolution of dust in debris disks parallels the evolution of larger solid objects in the Solar System.  Simple physical arguments also link debris disks with the formation of planetary systems. Because radiation removes 1--100~$\\mu$m grains on timescales shorter than the stellar age, some process replenishes the dust. To maintain the observed dust masses for long timescales, normal stars must have a large reservoir, $\\sim$ 10--100 \\mearth, of unseen objects that continuously collide at large velocities and fragment into smaller objects. Remnant material from planet formation satisfies both needs. The growth of 1000 km or larger planets in a disk of small grains naturally leaves behind an ensemble of `leftover' 1--10~km `planetesimals' on eccentric orbits \\citep[][2004b]{kb04a}. For a mass of $\\sim$ 10--100~\\mearth\\ in leftovers, high velocity collisions produce enough dust for most debris disks \\citep[e.g.,][]{bac93,hab01,kb04b}. If this interpretation is correct, debris disks provide conclusive evidence for the formation of Pluto-mass or larger planets around many, if not most, main sequence stars.  In addition to these considerations, numerical calculations suggest that an evolving swarm of 1--10~km planetesimals explains several observed trends in the properties of debris disks.  Starting with an ensemble of $\\lesssim$ 1 km-sized planetesimals, \\citet[][2004a, b, 2005]{kb02b} show that collisions and mergers form 500--1000~km-sized objects in 1--50 Myr. These protoplanets stir up leftover planetesimals along their orbits. Destructive collisions among the leftovers then produce a collisional cascade -- where collisions gradually grind large objects into smaller ones -- along with copious amounts of dust \\citep[see also][]{wil94,dur97,qui07}.  \\citet{dom03}, \\citet{wya07a}, and \\citet{loh08} show that collisional evolution in a belt of high velocity planetesimals naturally produces a dust luminosity that declines roughly inversely with time \\citep[see also][2004b, 2005]{kb02a}, explaining the observed time evolution -- $L_d \\propto t^{-n}$, with $n \\approx$ 0.5--1 -- suggested by recent observations of A-type stars \\citep[e.g.,][]{kal98,hab01,dec03,gre03,rie05,rhe07a}.  To account for the large observed range of IR excesses among stars of similar ages, \\citet{wya07a} propose belts with a range of initial masses and semimajor axes, as suggested from submm observations of protostellar disks \\citep[][2007b]{and05}.  Here, we continue to explore the evolution of dusty debris arising from planet formation in a disk of icy planetesimals. Our suite of calculations for disks at 30--150~AU around 1--3~\\msun\\ stars yields robust predictions for the maximum sizes of icy planets as a function of semimajor axis and stellar age. Results for the long-term evolution of IR excesses account for many fundamental aspects of the data.  These calculations are the first to explain the `rise and fall of debris disks' around A-type stars \\citep{cur08a} and the apparent peak in the 70--160~$\\mu$m excesses of G-type stars with ages of $\\sim$ 100~Myr \\citep{hil08}. Comparisons between our models and current observations suggest that the minimum stable grain size and the slope of the IR emissivity law are critical parameters.  The models suggest a set of further critical observations.  Spatially resolved images of debris disks around A-type and solar-type stars can improve our understanding of the minimum stable grain size. Larger samples of debris disks with high quality submm data from ALMA, Herschel, and SOFIA can place better constraints on the slope of the emissivity law. Together, these data can test our predictions for the time evolution of debris disk emission around 1--3 \\msun\\ stars and provide input for more complete calculations that include the formation and dynamical evolution of giant planets.  We outline our model in \\S2. We describe results for the formation of icy planets in \\S3 and the evolution of debris disks in \\S4.  After discussing several applications of our calculations in \\S5, we conclude with a brief summary in \\S6. ",
        "conclusions": "Our calculations provide a robust picture for the formation of planets and debris disks from a disk of icy planetesimals and set the context for the evolution of dusty debris in a dynamic system of planets. The results of this study provide a framework for interpreting existing observations of debris disks around 1--3 \\msun\\ stars and suggest new observational tests of this picture.  We describe a suite of numerical calculations of planets growing from ensembles of icy planetesimals at 30--150~AU in disks around 1--3~\\msun\\ stars. Using our hybrid multiannulus coagulation code, we solve for the evolution of sizes and orbits of objects with radii of $\\sim$ 1~m to $\\gtrsim$ 1000~km over the main sequence lifetime of the central star.  These results allow us to constrain the growth of planets as a function of disk mass, stellar mass, and semimajor axis.  Debris disk formation is coincident with the formation of a planetary system. All calculations of icy planet formation at 30--150~AU lead to a collisional cascade which produces copious amounts of dust on timescales of 5--30 Myr. This dust is observable throughout the lifetime of the central star. Because we consider a broad range of input parameters, we derive the time evolution of (i) dust produced in the collisional cascade and (ii) the IR and submm emission from this dust as a function of disk mass, stellar mass, and time.  We divide the rest of this section into (i) theoretical considerations, (ii) observable consequences, and (iii) observational tests. The theoretical considerations build on the highlights of icy planet formation in \\S3.4. Observable consequences of the calculations follow from the discussion in \\S4. The observational tests of the models are described in \\S5.  \\subsection{Theoretical Considerations}  \\begin{itemize}  \\item Icy planet formation at 30--150 AU is self-limiting. Starting with a swarm of $\\lesssim$ 1 km planetesimals, runaway growth produces a set of 100--500 km protoplanets. As the protoplanets grow, they stir up leftover planetesimals along their orbits. When the leftovers reach high $e$, collisions produce debris instead of mergers. Because protoplanets cannot accrete leftovers rapidly, a cascade of destructive collisions grinds the leftovers to dust. Poynting-Robertson drag and radiation pressure then remove the dust from the disk.  \\item The maximum sizes of icy planets at 30--150 AU are remarkably independent of initial disk mass, stellar mass, and stellar age. For disks with $x_m$ = 1/3 to 3 around 1--3 \\msun\\ stars with ages $t = 0.1 - 1 t_{ms}$, the typical planet has $r_{max} \\sim$ 1750~km and $m_{max} \\sim 0.005$~\\mearth. These objects contain $\\lesssim$ 3\\%--4\\% of the initial disk mass.  Although this result is also independent of the fragmentation parameters, the finite main sequence lifetimes of 1--3 \\msun\\ stars limits the formation of many large planets in the outer disk. Thus, the inner disk produces many more Pluto-mass planets than the outer disk (Tables \\ref{tab:rad1000.1} and \\ref{tab:rad1000.2}).  \\item For stars close to the main sequence turnoff, stellar lifetimes and the collisional cascade limit the mass in solid objects at 30--150~AU. In the inner disk, the collisional cascade removes most of the leftover planetesimals before the central star evolves off the main sequence. Thus, the typical mass in small objects is $\\sim$ 10\\% of the initial mass at 30--40~AU.  In the outer disk, smaller collision rates produce a slower cascade. Thus, the central star evolves off the main sequence with $\\sim$ 50\\% of the initial mass remaining in 1--10~km planetesimals at 125--150~AU.  \\item The collisional cascade produces copious amounts of dust. Dust begins to form during the transition from runaway to oligarchic growth ($t$ = 5--10~Myr), peaks when the first objects reach their maximum sizes ($t$ = 10--30~Myr), and then slowly declines $t \\gtrsim$ 30--50~Myr). The peak mass in 0.001--1~mm (0.001--1~m) particles is $\\sim$ 1--2 lunar masses ($\\sim$ 1~\\mearth). Disks with initial masses $x_m$ = 1/3 to 3 reach these peak masses when the age of the star is $\\sim$ 10\\% to 20\\% of its main sequence lifetime. Because the timescale to form dust is short ($\\sim$ 10--20~Myr), stars are surrounded by large disks of debris at 30--150~AU throughout their main sequence lifetimes.  \\item Radiative processes remove large amounts of mass from debris disks. Radiation pressure produces a radial wind of small particles containing $\\sim$ 60\\% to 70\\% of the mass removed from the disk. Poynting-Robertson drag pulls the rest of the lost mass into the inner disk. Because radiation pressure is more important than Poynting-Robertson drag when collision rates are large, we expect more wind (inner disk) emission earlier (later) in the evolution.  \\end{itemize}  \\subsection{Observable Consequences}  We derive clear observational consequences of the collisional cascade.  \\begin{itemize}  \\item The dusty debris from the collisional cascade is directly observable.  For disks around 1--3~\\msun\\ stars, the maximum fractional dust luminosity of $L_d / L_{\\star} \\sim 2 \\times 10^{-3}$ is comparable to the maximum dust luminosities of known debris disks \\citep{bac93,rie05,su06,rhe07a}. The dust temperature at the inner edge of a 30--150~AU disk scales with the temperature of the central star; thus, the predicted 24~$\\mu$m excess is very sensitive to the stellar mass. At 70~$\\mu$m, the predicted excesses scale roughly linearly with disk mass and stellar mass. The predicted 160--850~$\\mu$m excesses depend on the disk mass but are nearly independent of the stellar mass.  \\item For systems with little or no emission from terrestrial dust ([5.8]--[8] $\\lesssim$ 0.1), mid-IR color-color diagrams clearly distinguish debris disks around stars of different masses. In a [8]--[24] {\\it vs} [24]--[70] diagram, 2--3 \\msun\\ (1--2 \\msun) stars have red (blue) [8]--[24] and blue (red) [24]--[70] (Fig \\ref{fig:cc-1}). In both cases, the color scales with the initial disk mass (Fig \\ref{fig:cc-2}).  Optical colors and spectra generally provide good estimates for stellar mass; thus, these diagrams provide good tests of our model predictions.  \\end{itemize}  \\subsection{Observational Tests}  We compare our predictions with observations of A-type stars and solar-type stars.  \\begin{itemize}  \\item For A-type stars, our calculations are the first to explain the observed rise and fall of debris disk fluxes at 24~$\\mu$m \\citep[Fig. \\ref{fig:su1};][]{cur08a,cur08b}. In our picture, the rise in debris disk emission corresponds to the transition from runaway growth -- when mergers of small planetesimals produce larger protoplanets -- to oligarchic growth -- when the collisional cascade begins to grind leftover planetesimals into dust. When oligarchs in the inner disk are close to their maximum sizes of $\\sim$ 1750~km, the collisional cascade produces a maximum in debris disk emission. For a wide range of initial conditions, this maximum occurs at 10--20~Myr. As the collisional cascade moves out through the disk, smaller collision rates produce less dust which emits at lower temperatures. Thus, the 24~$\\mu$m excess falls with time. The predicted rate of decline, $t^{-n}$ with $n \\approx$ 0.6--0.8, is close to the observed rate \\citep[$n \\approx$ 0.5--1;][]{gre03,rie05,rhe07a}.  \\item At longer wavelengths, the maximum excess is larger and lasts longer than at 24~$\\mu$m. Predicted mid-IR colors also increase slowly with time. Although larger samples of A-type stars with 8~$\\mu$m photometry would provide a better test of our models, current data for the 70~$\\mu$m excess and the evolution of the [24]--[70] color agree with our predictions (Fig. \\ref{fig:su2}--\\ref{fig:su3}).  For 2--3 \\msun\\ stars near the main-sequence turnoff, our calculations also yield a clear maximum in the 850~$\\mu$m flux. Large surveys, such as the proposed JCMT Legacy Survey \\citep{mat07} and submm observations with ALMA, Herschel and SOFIA, can test this prediction.  \\item For solar-type stars, our models match observations of most sources. The predicted evolution of the 70--160~$\\mu$m excesses follows the observed rise at 10--100~Myr, the peak at $\\sim$ 30--100~Myr, and the decline at $\\gtrsim$ 300~Myr. Although $\\sim$  70\\% (55\\%) of observed debris disks have fluxes that lie within model predictions, our models underpredict fluxes for the brightest sources by a factor of 5--10. Fluxes for models with $r_2 \\approx$ 0.1~$\\mu$m and $q \\lesssim$ 0.7 provide better matches to these observations. To guide our choices for $r_2$ and $q$, we require spatially resolved images and submm fluxes for these objects.  \\item For 1--3 \\msun\\ stars with ages $\\sim$ 0.1--1 Gyr, current data suggest that solar-type stars have a larger range of far-IR excesses than A-type stars. In our models, faster debris disk evolution around A-type stars produces a smaller dispersion in far-IR excesses and colors for stars with ages of 100 Myr to 1 Gyr. Larger samples of debris disks can test this prediction in more detail.  \\item We also consider observations of Vega, the prototypical debris disk. If we adopt models with a blowout radius $r_2$ = 10~$\\mu$m, we can match observations with a standard collisional cascade within a broad torus at 80--200~AU. If the torus contains $\\sim$ 1--5 \\mearth\\ in large objects with $r \\gtrsim$ 1 cm, the cascade can generate (i) the observed ensemble of grains with $r \\sim$ 200--300~$\\mu$m within the torus and (ii) an outflowing wind of small grains with $r \\sim$ 1--50~$\\mu$m.  This conclusion differs from \\citet{su06}, who postulate a recent catastrophic collision between two large objects as the source of the dusty Vega wind. Although the complete destruction of two large icy objects can produce a massive outflowing wind, our results suggest that the dusty wind is short-lived and cannot be rapidly replenished by the observed population of larger objects.  We show that a steady-state collisional cascade can explain the {\\it Spitzer} data \\citep[see also][]{kb05}. If our interpretation is correct, sensitive observations at 1--10 mm should detect our proposed reservoir of larger objects.  \\end{itemize}  Matching other observations of debris disks requires more realism in our planet formation calculations. Adding binary companions and giant planets provides ways to modify the evolution of the collisional cascade and to impose structure on rings and tori \\citep[e.g.][]{wil02,mor05,qui06}. Extending the coagulation calculations to smaller sizes allows studies of the formation of winds and other large structures. Although these calculations have been prohibitively expensive in computing time, rapid advances in computing technology will make these additions possible in the next few years.  Based on the results described here and in \\citet{kb04b}, we conclude that debris disks are the inevitable outcome of icy planet formation in a disk of solid objects. The basic structures produced by this model -- broad tori and narrow rings of dust that propagate out through the disk \\citep{kb04b} -- are consistent with observations \\citep[e.g.][]{jay98,kal05,su06,fit07}. The model also explains the time evolution of mid-IR colors and fluxes for debris disks around A-type and solar-type stars."
    },
    "0807/0807.3307_arXiv.txt": {
        "abstract": "The Randall-Sundrum  model with a small curvature is considered in which five-dimensional Planck scale lies in the TeV region, and a spectrum of Kaluza-Klein gravitons reminds that in one flat extra dimension. The cross sections for interactions of ultra-high energy cosmic neutrinos with nucleons are calculated. It is shown that effects related with massive graviton excitations can be detected in deeply penetrating inclined air showers induced by these neutrinos. The expected number of air showers at the Auger Observatory is estimated as a function of two parameters of the model. ",
        "introduction": "One of the most important problem of the modern particle physics is a hierarchy problem, i.e. unnaturally large ratio of the gravity scale ($10^{19}$ GeV) to the electroweak scale ($10^2$ GeV). To solve this problem, theories with large extra dimensions have been proposed~\\cite{Arkani-Hamed:98}. However, they could only explain a huge value of the Planck mass by introducing another large scale, namely, a size of extra flat dimensions. Thus, the hierarchy problem was not really solved, but reformulated in terms of this new scale.  The model which does solve the problem most economically is the Randall-Sundrum (RS) model~\\cite{Randall:99} with a single extra dimension and warped background metric~\\cite{Kisselev:05}: \\begin{equation}\\label{metric} ds^2 = e^{2 \\kappa (\\pi r_c - |y|)} \\, \\eta_{\\mu \\nu} \\, dx^{\\mu} dx^{\\nu} + dy^2 \\;. \\end{equation} Here $y = r_c \\, \\theta$ ($-\\pi \\leq \\theta \\leq \\pi$), $r_c$ being the ``radius'' of the extra dimension, while $\\{x^{\\mu}\\}$, $\\mu = 0,1,2,3$, are the coordinates in four-dimensional space-time. The parameter $\\kappa$ defines the scalar curvature in five dimensions. Note that the points $(x^{\\mu}, y)$ and $(x^{\\mu}, -y)$ are identified, and the periodicity condition, $(x^{\\mu}, y) = (x_{\\mu}, y + 2 \\pi r_c)$, is imposed. The tensor $\\eta_{\\mu \\nu}$ is the Minkowski metric.  It is assumed that there are two 3-dimensional branes with equal and opposite tensions located at the point $y = 0$ (called the Plank brane) and point $y =  \\pi r_c$ (referred to as the TeV brane). All SM fields are confined to the TeV brane, while the gravity propagates in five dimensions. The following relation between the 4-dimensional (reduced) Planck mass, $\\bar{M}_{\\mathrm{Pl}}$, and (reduced) gravity scale in five dimensions, $\\bar{M}_5$, can be derived: \\begin{equation}\\label{hierarchy} \\bar{M}_{\\mathrm{Pl}}^2 = \\frac{\\bar{M}_5^3}{\\kappa} \\left(e^{2 \\pi \\kappa r_c} - 1 \\right) \\;. \\end{equation}  The masses of the Kaluza-Klein (KK) graviton excitations are proportional to the curvature parameter $\\kappa$: \\begin{equation}\\label{KK_masses} m_n = x_n \\, \\kappa , \\qquad n=1,2 \\ldots \\;, \\end{equation} where $x_n$ are zeros of the Bessel function $J_1(x)$. On the TeV brane, the zero graviton mode, $h^{(0)}_{\\mu \\nu}$, and massive graviton modes, $h^{(n)}_{\\mu \\nu}$, are coupled to the energy-momentum tensor of the matter, $T^{\\mu \\nu}$, as follows: \\begin{equation}\\label{Lagrangian} \\mathcal{L}_{\\mathrm{int}} = - \\frac{1}{\\bar{M}_{Pl}} \\, T^{\\mu \\nu} \\, h^{(0)}_{\\mu \\nu} - \\frac{1}{\\Lambda_{\\pi}} \\, T^{\\mu \\nu} \\, \\sum_{n=1}^{\\infty} h^{(n)}_{\\mu \\nu} \\;, \\end{equation} with \\begin{equation}\\label{Lambda} \\Lambda_{\\pi} = \\left( \\frac{\\bar{M}_5^{3}}{\\kappa} \\right)^{\\! 1/2} \\end{equation} being a physical scale on this brane.  Let us note that the metric \\eqref{metric} differs from that presented in the original paper~\\cite{Randall:99} in which both $\\bar{M}_5$ and $\\kappa$ have to be taken as large as the Planck mass ($\\bar{M}_5 \\sim \\kappa \\sim \\bar{M}_{\\mathrm{Pl}}$). Moreover, the size of the warped extra dimension should be extremely small ($r_c \\simeq 60 \\, l_{\\mathrm{Pl}}$).~\\cite{Kisselev:05} Thus, in order to explain the huge value of $\\bar{M}_{\\mathrm{Pl}}$ in such a scheme, one has to introduce new mass scales of the same order, namely, $\\bar{M}_5 $, $\\kappa$, and $r_c^{-1}$.  However, the hierarchy problem can be successfully solved in the RS scenario, but with the metric \\eqref{metric}. The equation \\eqref{hierarchy} allows us to consider the {\\em small curvature option} of the RS model~\\cite{Kisselev:05}-\\cite{Giudice:05}:% \\footnote{For numerical estimates, the region $0.5 \\mathrm{\\ GeV} \\leqslant \\kappa \\leqslant 1.5 \\mathrm{\\ GeV}$ will be used (see Fig.~\\ref{fig:event rate}).} \\begin{equation}\\label{small_curv} \\kappa \\ll \\bar{M}_5 \\sim 1 \\mathrm{\\ TeV} \\;. \\end{equation} In such a case, we get an almost continuous spectrum of low-mass graviton excitations with the small mass splitting $\\Delta m \\simeq \\pi \\kappa$. Note that in the standard scenario of the RS model~\\cite{Randall:99} one has a series of KK graviton resonances with the lightest one having a mass around 1 TeV.  The RS model with the large extra dimension has been checked by the DELPHI Collaboration~\\cite{LEP_limit}. The gravity effects were searched for by studying photon energy spectrum in the process $e^+e^- \\rightarrow \\gamma + E_{\\perp}\\hspace{-6mm} \\diagup \\hspace{2mm}$. The limit on $\\bar{M}_5$ obtained~\\cite{LEP_limit} is \\begin{equation}\\label{LEP_limit} \\bar{M}_5 > 0.92 \\mathrm{\\ TeV} \\;. \\end{equation} The search for large extra dimensions in the diphoton channel using data collected by the CDF and D\\O~Collaborations at $\\sqrt{s} = 1.96$ TeV are presented in Refs.~\\cite{Tevatron_ADD}. The measured $p_{\\perp}$-distributions are in a good agreement with the SM background, that allowed us to obtain the bound~\\cite{Kisselev:08}: \\begin{equation}\\label{Tevatron_limit} \\bar{M}_5 > 0.81 \\mathrm{\\ TeV}\\;. \\end{equation} The discovery limit of the LHC in the two-photon production (requiring a $5 \\sigma$ effect) has also been derived for two values of the integrated luminosity $\\mathcal{L}$~\\cite{Kisselev:08}: \\begin{equation}\\label{LHC_limit} \\bar{M}_5 = \\left\\{ \\begin{array}{rl} 6.3 \\mathrm{\\ TeV} \\;, & \\mathcal{L} = 100 \\mathrm{\\ fb}^{-1} \\\\ 5.1 \\mathrm{\\ TeV} \\;,  & \\mathcal{L} = 30 \\mathrm{\\ fb}^{-1} \\\\ \\end{array} \\right. \\end{equation}  Previously, the gravity effects in the RS model with the small curvature were already looked for in a number of different processes (see Refs.~\\cite{Giudice:05}-\\cite{Kisselev:08}). In the present paper we will estimate an upper bound on $\\bar{M}_5$ which can be reached by the Auger ground array in detecting quasi-horizontal air showers induced by ultra-high energy (UHE) cosmic neutrinos. ",
        "conclusions": ""
    },
    "0807/0807.0238_arXiv.txt": {
        "abstract": "We present multi-epoch, near-infrared imaging of the binary T5.0+T5.5~dwarf 2MASS~J15344984$-$2952274AB obtained with the Keck laser guide star adaptive optics system.  Our Keck data achieve sub-milliarcsecond relative astrometry and combined with an extensive (re-)analysis of archival \\HST\\ imaging, the total dataset spans $\\sim$50\\% the orbital period.  We use a Markov Chain Monte Carlo analysis to determine an orbital period of 15.1$_{-1.6}^{+3.1}$~yr and a semi-major axis of 2.3$_{-0.2}^{+0.3}$~AU.  We measure a total mass of $0.056 \\pm 0.003$~\\Msun\\ ($59\\pm3$~\\Mjup), where the largest uncertainty arises from the parallax.  This is the first field binary for which both components are confirmed to be substellar.  This is also the coolest and lowest mass binary with a dynamical mass determination to date. Using evolutionary models and accounting for the measurement covariances, we derive an age of \\hbox{0.78$\\pm$0.09~Gyr} and a mass ratio of $0.936_{-0.008}^{+0.012}$.  The relatively youthful age is consistent with the low tangential velocity of this system.  For the individual components, we find \\hbox{$\\Teff = 1028\\pm17$~K} and $978\\pm17$~K, $\\logg = 4.91\\pm0.04$ and $4.87\\pm0.04$ (cgs), and masses of $0.0287\\pm0.0016~\\Msun$ ($30.1\\pm1.7$~\\Mjup) and $0.0269\\pm0.0016~\\Msun$ ($28.2\\pm1.7$~\\Mjup).  These precise values generally agree with previous studies of T~dwarfs and affirm current theoretical models.  However, (1) the temperatures are about 100~K cooler than derived for similar field objects and suggest that the representative ages of field brown dwarfs may be overestimated. Similarly, (2) the H-R diagram positions are discrepant with current model predictions and taken at face value would overestimate the masses.  While this may arise from large errors in the luminosities and/or radii predicted by evolutionary models, the likely cause is a modest ($\\approx$100~K) overestimate in temperature of T~dwarfs determined from model atmospheres.  We elucidate future tests of theory as the sample of substellar dynamical masses increases.  In particular, we suggest that field brown dwarf binaries with dynamical masses (``mass benchmarks'') can serve as reference points for \\Teff\\ and \\logg\\ and thereby constrain ultracool atmosphere models, as good as or even better than single brown dwarfs with age estimates (``age benchmarks''). ",
        "introduction": "Over about the past decade, the parameter space of traditional stellar astrophysics has been greatly expanded with the discovery and characterization of brown dwarfs, objects that for most of their lifetimes are colder and less luminous than main-sequence stars. Despite ample progress in finding and characterizing brown dwarfs, very few direct measurements of their physical properties have been made so far.  In particular, dynamical masses for brown dwarfs are sorely needed to test the theoretical models over a wide range of parameter space.  In comparison to the $>$100 binary stars with direct mass determinations, dynamical masses have been measured for only a handful of objects clearly below the stellar/substellar boundary: \\begin{enumerate} \\item[1.] the secondary component of the M8.5+M9 visual binary Gl~569Bab, which itself is a companion to a young ($\\approx$100--300~Myr) field M2.5~dwarf \\citep{2000ApJ...529L..37M, 2001ApJ...560..390L, 2001ApJ...554L..67K, 2004astro.ph..7334O}; \\item[2.+3.] the very young ($\\sim$Myr) eclipsing M6.5+M6.5 binary brown dwarf 2MASS~J05352184$-$0546085 \\citep{2006Natur.440..311S}; \\item[4.] the secondary component of the GJ~802AB system, where the secondary has an estimated spectral type of L5 and the primary is a M5.5+M5.5 spectroscopic binary \\citep{2006ApJ...649..389P, 2006ApJ...650L.131L, gl802b-ireland}. \\end{enumerate} In addition, GJ~569Ba itself may be a unresolved binary brown dwarf \\citep{2006ApJ...644.1183S}, and the secondary component of the L0+L1.5 binary 2MASSW~J0746425+2000321AB, which appears to be an old ($\\gtrsim$1~Gyr) field system, has a mass near the stellar/substellar boundary \\citep{2001AJ....121..489R, 2004A&A...423..341B, 2006AJ....131..638G}.  About 100~ultracool visual binaries are known,\\footnote{We follow the popular convention that ``ultracool'' refers to objects of (integrated-light) spectral type M6 or later.} found with high angular resolution imaging surveys conducted by {\\sl Hubble Space Telescope} (\\HST; \\eg, \\citealp{2001AJ....121..489R, 2003AJ....126.1526B, 2003AJ....125.3302G, 2003ApJ...586..512B, 2006ApJS..166..585B}) and ground-based adaptive optics (AO) imaging (\\eg, \\citealp{2003ApJ...587..407C, 2005AJ....129.2849B, 2006astro.ph..5037L, 2008AJ....135..580R}; Liu \\etal, in prep). Only $\\approx$10\\% of these are binaries composed of the coldest class of brown dwarf, the T~dwarfs.  T~dwarfs are distinguished by their very red optical colors arising from pressure-broadened alkali resonance lines and very blue near-infrared (IR) colors from strong \\meth, \\htwoo, and collision-induced \\htwo\\ absorption \\citep[e.g.][]{1995Sci...270.1478O, geb01, kirk05}. These are the lowest luminosity and coolest objects directly detected outside of our solar system, with bolometric luminosities (\\Lbol) of $\\lesssim10^{-4.5}\\Lsun$ and effective temperatures (\\Teff) of $\\approx700-1300$~K \\citep[e.g.][]{2004AJ....127.2948V, gol04, 2006ApJ...639.1095B, 2007arXiv0705.2602L, 2007MNRAS.381.1400W, 2008A&A...482..961D}.  As such, analyzing their physical properties is an important pathway to understanding the properties of gas-giant extrasolar planets.  The subject of this paper is the T~dwarf 2MASS~J15344984$-$2952274AB, hereinafter \\twomassbin, which has an integrated-light infrared spectral type of T5 \\citep{burg01, 2005astro.ph.10090B}.\\footnote{\\citet{2005astro.ph.10090B} report integrated-light spectral types of both T5 and T5.5 for \\twomassbin\\ based on a spectrum obtained with the CTIO/OSIRIS instrument. Examination of their spectrum by us finds that T5 is the correct typing.  Also, a new near-IR spectrum obtained with the IRTF/SpeX spectrograph confirms an integrated-light spectral type of T5 (\\S~\\ref{sec:spex}).}  This object was first resolved as a 0.065\\arcsec\\ binary in August~2000 in \\HST/WFPC2 imaging \\citep{2003ApJ...586..512B}.  Among known visual ultracool binaries in the field, this system has the shortest estimated orbital period, only 4~years (see compilation in \\citealp{2006astro.ph..2122B}).  In combination with its very high quality parallax measurement of $73.6\\pm1.2$~mas \\citep{2003AJ....126..975T} and the fact that its two components are nearly equal-magnitude (indicating nearly equal masses), this system is a very appealing object for astrometric monitoring.  Laser guide star (LGS) AO provides a powerful tool for high angular resolution studies of brown dwarf binaries.  Through resonant scattering off the sodium layer at $\\sim$90~km altitude in the Earth's atmosphere, sodium LGS systems create an artificial star bright enough to serve as a wavefront reference for AO correction \\citep{1985A&A...152L..29F, 1987Natur.328..229T, 1994OSAJ...11..263H}. Thus, most of the sky can be made accessible to near diffraction-limited IR imaging from the largest existing ground-based telescopes.  We have previously used Keck LGS AO to discover that the nearby L~dwarf Kelu-1 is a binary system \\citep{2005astro.ph..8082L} and to identify the novel L+T binary SDSS~J1534+1615AB \\citep{2006astro.ph..5037L}. In regards to dynamical mass determinations, ground-based telescopes equipped with LGS AO can provide the necessary long-term platforms for synoptic monitoring of visual binaries, especially where the required amount of observing time at each epoch is relatively modest but many epochs are needed, in contrast to \\HST\\ where target acquistion can be slow and monitoring a populous sample over many epochs is quite telescope time-intensive.  We present here the results of multi-epoch imaging of \\twomassbin, observed as part of our ongoing high angular resolution study of ultracool binaries using LGS AO.  \\S~2 presents our Keck LGS AO observations and (re-)analysis of archival \\HST\\ imaging.  \\S~3 presents the resolved photometric properties of the binary and fitting of the orbit using a Markov Chain Monte Carlo method.  \\S~4 compares the resulting total mass against evolutionary models, and \\S~5 summarizes our findings.  Those readers interested solely in the results can focus on \\S~4 and~\\S~5. ",
        "conclusions": "We have determined the first dynamical mass for a binary T~dwarf, the T5.0+T5.5 system \\twomassbin, by combining six epochs of Keck LGS AO imaging from 2005--2008 with three epochs of \\HST\\ imaging obtained in 2000 and 2006.  Both datasets achieve milliarcsecond accuracy or better for the relative astrometry of the two components and are validated through extensive testing with images of simulated binaries. We employ a Markov Chain Monte Carlo analysis to determine the orbital parameters and their uncertainties.  The time baseline of our complete dataset covers about half of the total period.  We find that the orbital motion of the binary is viewed in an almost edge-on orientation and has a modest eccentricity. Our determination of a $15.1_{-1.6}^{+2.3}$~yr orbital period is significantly longer than the original 4-year estimate, as by chance the binary was at a very small projected separation when discovered in 2000.  The total mass of the system is $0.056\\pm0.003~\\Msun$ ($59\\pm4$~\\Mjup), including the uncertainty in the parallax.  This is the second brown dwarf binary directly confirmed, the first among the field population.  It is also the coolest and lowest mass binary with a dynamical mass determination to date.  With very accurate measurements of the total mass and the bolometric luminosity ($\\log(\\Lbol/\\Lsun) = -4.751\\pm0.011$), we use the Tucson and Lyon evolutionary tracks to determine the remaining physical properties for the system.  The two sets of models give largely consistent results, which highlights the difficult of distinguishing between them even with such precise observational data.  We average the model results to represent the final determinations.  We find a relatively youthful age for the system of 0.79$\\pm$0.09~Gyr (1$\\sigma$), consistent with its low tangential velocity relative to other field T~dwarfs. The remaining physical parameters of the individual components are then fully determined: radii of \\hbox{$0.0986\\pm0.0015~\\Rsun$} and \\hbox{$0.0993\\pm0.0017~\\Rsun$}, effective temperatures of $1028 \\pm 17$~K and $978 \\pm 17$~K, surface gravities of \\hbox{$\\logg = 4.91 \\pm 0.04$} and $4.87 \\pm 0.04$, and masses of $0.0287\\pm0.0016~\\Msun$ ($30.1\\pm1.7~\\Mjup$) and $0.0269\\pm0.0016~\\Msun$ ($28.2\\pm1.7~\\Mjup$) for components~A and~B, respectively.  We take care to account for the covariances inherent in the measurement uncertainties, by using a Monte Carlo approach to derive these physical quantities from the evolutionary models.  Our approach also assumes that the system is coeval and composed of only two components.  These precise determinations for 2MASS~J1534$-$2952A and~B are in general accord with the \\Teff\\ and \\logg\\ values found previously for field T~dwarfs based on model atmospheres and with the ages of T~dwarfs predicted by Monte Carlo simulations of the solar neighborhood.  However, upon closer scrutiny, there are two potential discrepancies with past studies.  Both suggest that the temperatures of field T~dwarfs may be overestimated by $\\approx$100~K, though we stress that the two discrepancies must arise from independent effects. (1)~The temperatures of 2MASS~J1534$-$2952A and~B appear to be cooler than field objects of comparable spectral type.  This resembles discrepancies previously noted by \\citet{2006astro.ph..7514M} and \\citet{2006astro.ph..9464L} for late-L/early-T dwarfs that are companions to young main-sequence stars.  They have hypothesized that the effect is due to the gravity sensitivity of the L/T transition. The fact that this discrepancy also occurs for \\twomassbin\\ suggests instead that the problem may arise from a factor of $\\approx6\\pm3$ overestimate in the adopted ages of field objects when determining their temperatures using evolutionary tracks.  Ages of $\\sim$0.3--1.0~Gyr are preferred based on this binary.  (2) The temperatures of 2MASS~J1534$-$2952A and~B are slightly cooler than inferred for other mid-T~dwarfs from model atmospheres.  Detailed analysis of the system's integrated-light and resolved spectra with model atmospheres is needed to directly assess the \\Teff\\ and \\logg\\ of the two components and to refine the comparison with the values derived from evolutionary models.  The positions of the two components on the H-R diagram are discrepant with theoretical evolutionary tracks corresponding to their individual masses.  In fact, taken at face value, using the H-R diagram positions to infer masses from the evolutionary tracks would lead to masses of $\\approx$0.05--0.06~\\Msun, about a factor of two larger than the actual measured masses.  While this discrepancy could stem from large systematic errors in the luminosities ($\\sim$50\\% errors) and/or radii ($\\sim$20\\% errors) predicted by evolutionary models, the likely cause is that temperatures from model atmospheres are too warm by $\\approx$100~K for mid-T dwarfs.  This highlights the need for continued improvements to the model atmospheres.  Future monitoring of \\twomassbin\\ will help to refine its orbit and its dynamical mass.  The orbital separation of the system is now rapidly decreasing and will not be readily resolvable again until around 2011.  At the same time, an improved parallax for the system will be required --- the uncertainty in the total mass from the orbit fitting is 3\\%, compared to the 5\\% that arises from the uncertainty in the parallax.  Radial velocity monitoring and/or absolute astrometry will directly determine the individual masses and test if the system is a higher order multiple.  However, given the very similar fluxes of the two components (implying nearly equal mass), individual mass measurements are unlikely to resolve the discordant H-R diagram position of the two components relative to evolutionary tracks.  This problem is likely driven by the systematic uncertainties in current model atmospheres for T~dwarfs.  The fundamental characteristic of the field population is that it spans a range of (largely unknown) ages.  However despite this uncertainty, field brown dwarf binaries can strongly test theoretical models, if analyzed appropriately.  These systems will be especially valuable in light of the current paucity of eclipsing field ultracool binaries and resolvable, short-period ultracool binaries in open clusters/groups.\\footnote{For reference, about 1~out of a~1000 stars are eclipsing binaries, whereas there are only $\\sim$600~L~and T~dwarfs known.}  Specifically, attempts to directly test different evolutionary tracks by placing ultracool objects on the H-R diagram (the ``H-R Diagram Test'') will be challenging, given the similarity between the tracks and the difficulty in independently determining \\Teff\\ with model atmospheres.  Instead, atmosphere models can be confronted against \\logg\\ and \\Teff\\ values for ultracool objects as derived from the evolutionary models, which can be exceptionally precise (the ``Benchmark Test'').  This approach has previously been applied to single brown dwarfs that are companions to stars of known age (``age benchmarks'').  We suggest that in an analogous fashion, field ultracool binaries with dynamical mass determinations (``mass benchmarks'') can test the model atmospheres.  In fact, given the plausible observational uncertainties, mass benchmarks are likely to provide stronger constraints (by a factor of $\\approx$5) on \\logg\\ and \\Teff\\ than age benchmarks, since dynamical masses can be determined far more accurately than ages for main-sequence field stars.  With the widespread advent of LGS AO on the largest ground-based telescopes, we can look forward to a rapid increase in dynamical mass determinations for low-mass field binaries and thus substantial advances in our understanding of the properties and evolution of substellar objects."
    },
    "0807/0807.5137_arXiv.txt": {
        "abstract": "{}{}{}{}{}  \\abstract {\\object{PG~1159-035}, a pre-white dwarf with $T_{\\mathrm{eff}}\\simeq 140\\,000$~K, is the prototype of the PG~1159 spectroscopic class and  the DOV pulsating class. Pulsating pre-white dwarf stars evolve rapidly: the effective surface temperature decreases rapidly, the envelope contracts and the inner structure experiences stratification due to gravitational settling. These changes in the star generate  variations in its oscillation periods. The measurement of temporal change in the oscillation periods, $\\dot P$, allows us to estimate directly rates of stellar evolutionary changes, such as the cooling rate and the  envelope contraction rate, providing a way to test and refine evolutionary models for pre-white dwarf pulsating stars. } {  Previously, only two pulsation modes of the highest amplitudes for PG~1159-035 have had their $\\dot P$ measured: the 516.0 s and the 539.3 s modes. We measured the $\\dot P$ of a larger number of pulsation modes, increasing the number of constraints for evolutionary studies of PG~1159-035. We attempted to use the secular variations in the periods of multiplets to calculate the variation in the rotational period, the envelope contraction rate, and the cooling rate of the star. } { The period variations were measured directly from the PG~1159-035 observational data and refined by the (O-C) method. } { We measured 27 pulsation mode period changes. The periods varied at rates of between 1 and 100 ms/yr, and several can be directly measured with a relative standard uncertainty below 10\\%. For the 516.0 s mode (the highest in amplitude) in particular, not only the value of $\\dot P$ can be measured directly with a relative standard uncertainty of 2\\%, but the second order period change, $\\ddot P$, can also be calculated reliably. By using the (O-C) method, we refined the $\\dot P$s and estimated the $\\ddot P$s for six other pulsation periods. As a first application, we calculated the change in the PG~1559-035 rotation period, ${\\dot P}_{\\rm rot}=(-2.13\\pm 0.05)\\times 10^{-6}\\,{\\rm ss}^{-1}$, the envelope contraction rate $\\dot R=(-2.2\\pm 0.5)\\times 10^{-13}\\,{\\rm R}_\\odot\\rm{s}^{-1}$, and the cooling rate $\\dot T=-1.42\\times 10^{-3} {\\rm Ks}^{-1}$. } {} ",
        "introduction": "%  \\object{PG~1159-035} is the prototype of two star classes: the PG1159 spectroscopic class and the DOV pulsating class, discovered by McGraw et al. (1979). The estimated temperature for PG~1159-035 is $140\\,000\\pm 5\\,000$ K (Werner et al. 1991; Dreizler et al. 1998; Jahn et al. 2007), which places this star in the pre-white dwarf sequence of the Hertzsprung-Russell diagram.  Analyzing the light curve obtained by the {\\it Whole Earth Telescope} (WET) (Nather et al. 1990) in the 1989 campaign, Winget et al. (1991) (henceforth W91) identified 122 pulsation modes  in PG~1159-035 with periods between 300 and 1000 seconds, and  spherical harmonic index $\\ell=1$ or $\\ell=2$. The excellent results allowed W91 to calculate the asteroseismological mass of the star ($M/M_\\odot \\simeq 0.59$)  and several other stellar parameters, such as the rotation period ($P_{\\rm rot} = 1.38\\pm 0.01$ days), the inclination of the rotation axis ($i \\simeq 60^o$), and a limit to the magnetic field strength ($B< 6000$ G). Costa et al. (2008), analyzing the PG~1159-035  combined data sets from  1983, 1985, 1989, 1993, and 2002, increased the number of identified pulsation modes to 198, the largest after the Sun, refining the determination of the rotation period ($P_{\\rm rot} = 1.3920\\pm 0.0008$ days) and other stellar parameters previously estimated by W91. The comparison of the Fourier transforms of light curves from different years indicates that the amplitudes of the pulsation modes change with time, and their intensities in the light curves can decrease to lower than the lower limit for reliable detection. Of the 198 detected modes, about 75\\% appear in only one of the FTs, about 25\\%   in two or more FTs, and only 14 of the 198 pulsation modes  (all of them with index $\\ell=1$) appear in the five FTs.   About $97\\%$ of all stars end their evolution as white dwarfs. The pre-white dwarf stage is one of the ``front doors'' to the white dwarf cooling sequence. There are two evolutionary channels that are known to occur prior to the white dwarf stage. The first channel involves the evolution of the star from the Horizontal Branch to white dwarf by means of  the Asymptotic Giant Branch (AGB), the Planetary Nebula phase, and PG1159 type stars. In the second channel, the star evolves directly from the Extended Horizontal Branch to the white dwarf stage, but without passing through  the Planetary Nebula phase.  In its evolution, a pre-white dwarf star can traverse the DOV (PG1159 or {\\it GW Vir}) instability strip, becoming a multiperiodic non-radial pulsating star, of periods between 100 and 1000 seconds. The exact values of the pulsation periods are defined by the mass, temperature, rotation, magnetic field, and structural characteristics of the star. While it passes through the DOV/PNNV instability strip, the star cools, its envelope contracts, and its interior experiences stratification due to the gravitational settling. As a consequence, the pulsation periods change with time. The hotter the star is, the quicker are the changes in its pulsation periods. For example, in \\object{G117-B15A}, a DAV white dwarf star with $T_{\\rm eff}\\simeq 12,000$ K, the highest amplitude period, 215 s, changes 1 second in 8 million  years, while in the hotter DOV star, PG~11559-035, with $T_{\\rm eff}\\simeq 140,000$ K, the 516 s period, the one of the highest amplitude, changes 1 second in only 350 years.   This work is a complementary study of the results obtained in Costa et al. (2008). We used the measured periods in the PG~1159-035 light curves of the 1983-2002 data sets to calculate the secular variation in  several pulsation modes of PG~1159-035, enlarging the number of constraints for future evolutionary studies of this star. As an immediate application, we show in Sect.~\\ref{sect:6} how the measured $\\dot P$s in a multiplet can be used to calculate the cooling rate of the star, the envelope contraction rate and the variation in the rotational period. Part of the procedure used here was proposed by Kawaler et al. (1985a), Kawaler (1986) and other authors, but was never applied to (pre-)white dwarf stars because this requires precise period and $\\dot P$ determinations. As highlighted in Sect.~\\ref{sect:6}, the calculations are based on many hypothesis about the interior of the star, but it is a first approach taking into account the  available observational data. PG 1159-035 is the first and only white dwarf in which we can apply these theories and test their limitations and reliability.    {\\scriptsize \\begin{table*} \\caption{Observed periods (in seconds) in each yearly data set.} \\label{table:1} {\\scriptsize \\begin{center} \\input{9118tb1} \\end{center} } \\end{table*} } ",
        "conclusions": "%  The main results of our present work are:  \\begin{enumerate}  \\item   The PG~1159-035 pulsation periods vary at rates of between 0.5 and 1.0 ms/years. For the WET data sets, the accuracy in the period determination is 10 ms for the least accurate cases (low amplitudes) and 1 ms for the most accurate ones, which enables the direct measurement of the pulsation periods after a few years of observation. After 19 years of observations, the $\\dot P$ of the 516.0 s mode (one of the modes with higher amplitude) was measured with a relative standard uncertainty ($\\sigma_{\\dot P}/\\,|{\\dot P}|$ of 2\\% and its $\\ddot P$ could be estimated with  relative standard uncertainty of 10\\%.  \\item   The $\\dot P$ of the 27 periods present in three or more yearly Fourier transforms were measured directly and the  values were refined using  (O-C) fitting. For seven of the pulsation modes, we used  third-order (O-C) fitting to calculate  $\\ddot P$. The measured $|\\ddot P|$ are between $1.4\\times 10^{-21}\\,{\\rm ss}^{-2}$ and $8.7\\times 10^{-19}\\,{\\rm ss}^{-2}$.  \\item   Using the $\\dot P$ of the $m=0$ and $m=+1$ components of the 517.1 s multiplet, we estimated the rotation period change to be $\\dot P_{\\rm rot}=(-2.13\\pm 0.05)\\times 10^{-6}\\,{\\rm ss}^{-1}$ or 67.2 s/yr.  \\item   From the $\\ddot P$ of the same multiplet, we estimated that the second order variation in the rotational period was $\\ddot P = (+9.0\\pm 0.3)\\times 10^{-14}\\,{\\rm ss}^{-2}$.  \\item   If we used the calculated values for $\\dot P_{\\rm rot}$, the PG~1159-035 contraction rate was $\\dot R \\simeq (-5\\pm 1)$ km/s for $R_\\star \\simeq (0.025\\pm 0.005)\\,R_\\odot$, assuming uniform rotation, and the cooling rate is $\\dot{T}_m \\simeq -45\\,000$ K/yr, assuming $T_m\\simeq T_{\\rm core} \\simeq 7.72\\times 10^7$ K.  \\end{enumerate}   Our results demonstrate that some periods decrease, while others increase, at different rates, even for the same multiplet components. Evolutionary models calculated by Kawaler \\& Bradley (1994) and  the La Plata group (C\\'orsico et al. 2008; Althaus et al. 2008) predict that the $\\dot P$ of ($\\ell=1$ and $m=0$) pulsation modes  must have different values and signs (some increase, while other decrease) due to the trapping of pulsation modes in different  layers of the star, but the calculated values do not fit the observed $\\dot P$s. The interior, evolution, and pulsation mechanisms in PG~1159 stars are not well known. New models have been developed by the La Plata group to fit the observed results and will be published in a future article.    A negative value for ${\\dot P}_{\\rm rot}$ is expected, because pre-white dwarf stars such as PG~1159-035 undergo rapid envelope contraction processes. With contraction, the stellar radius decreases. To conserve angular momentum, the angular rotation speed must increase. In other words, its rotation period decreases, and therefore ${\\dot P}_{\\rm rot}<0$. The shorter rotation periods observed in white dwarfs are of the order of few hours. This is the case, for example, for the DBV \\object{EC20058-5234} ($P_{\\rm rot} \\simeq 2$ h, Sullivan et al. 2008) and of the magnetic \\object{GD~356}  ($P_{\\rm rot} = 2.6$ h, Brinkworth et al. 2004). If the PG~1159-035 rotation period changed from its current value to $P_{\\rm rot} = 2$ h during the DOV phase ($\\sim 10^6$ yr), the variation rate would be ${\\dot P}_{\\rm rot} \\simeq -0.1$ s/yr. However, the calculations in Sect.~\\ref{sect:6.1} estimate that the rotation period is changing at a ${\\dot P}_{\\rm rot}\\simeq (-67\\pm 2)\\, \\rm{s/year}$ rate, far higher than expected and ${\\ddot P}_{\\rm rot}$ is too excessively high to be physically acceptable. A possible explanation is that the second order effects of the stellar rotation over the rotational splitting are non-negligible as assumed in Eq.~\\ref{eq:9}. The simplification of assuming a uniform rotation is also an error source in the ${\\dot P}_{\\rm rot}$ calculation, but the difference in the final result  must not be larger than a factor of 10 (Ostriker \\& Bodenheimer 1968). On the other hand, a high de/acceleration of the stellar rotation is expected if the star recently experienced a late thermal pulse (due to ejection of  matter from the star) and this may be the case for PG~1159-035. We note that if $|{\\dot P}_{\\rm rot}|$ is overestimated, both the contraction rate $|\\dot R|$ and the cooling rate $|\\dot T|$ must also be overestimated.  The accuracy in the determination of ${\\dot P}_{\\rm rot}$ by Eq.~\\ref{eq:14} depends strongly on the accuracies of $\\dot{P}_m$ and $\\dot{P}_0$, far more than on the period accuracies.   Although these first results are not fully understood, we have shown that several long-term campaigns appear to be reaching sufficient accuracy to interpret evolutionary changes. The last PG~1159-035 observational campaigns were carried out in 2002 by WET. New observations of PG~1159-035 in future years should allow the direct measurement of additional pulsation modes $\\dot P$ and $\\ddot P$ and the refining of derived values to date, improving the calculation of the PG~1159-035 evolutionary rates."
    },
    "0807/0807.2651_arXiv.txt": {
        "abstract": "{\\it Kepler} will monitor enough stars that it is likely to detect single transits of planets with periods longer than the mission lifetime. We show that by combining the {\\it Kepler} photometry of such transits with precise radial velocity (RV) observations taken over $\\sim$3 months, and assuming circular orbits, it is possible to estimate the periods of these transiting planets to better than 20\\% (for planets with radii greater than that of Neptune) and the masses to within a factor of 2 (for planet masses $m_{\\rm p} \\ge M_{\\rm Jup}$). We also explore the effects of eccentricity on our estimates of these uncertainties. ",
        "introduction": "Planets that transit their stars offer us the opportunity to study the physics of planetary atmospheres and interiors, which may help constrain theories of planet formation. From the photometric light curve, we can measure the planetary radius and also the orbital inclination, which when combined with radial velocity (RV) observations, allows us to measure the mass and density of the planet. All of the known transiting planets orbit so close to their parent stars that the stellar flux plays a major role in heating these planets. In contrast, the detection of transiting planets with longer periods ($P > 1\\, {\\rm yr}$) and consequently lower equilibrium temperatures would allow us to probe a completely different regime of stellar insolation, one more like that of Jupiter and Saturn whose energy budgets are dominated by their internal heat.  Because of its long mission lifetime ($L = 3.5\\, {\\rm yrs}$), continuous observations, and large number of target stars ($N \\simeq 10^5$), the {\\it Kepler} satellite (Borucki et al. 2004; Basri et al. 2005) has a unique opportunity to discover long-period transiting systems. Not only will {\\it Kepler} observe multiple transits of planets with periods up to the mission lifetime, but it is also likely to observe single transits of planets with periods longer than the mission lifetime. As the period of a system increases beyond $L/2$, the probability of observing more than one transit decreases until, for periods longer than $L$, only one transit will ever be observed.  For periods longer than $L$, the probability of seeing a single transit diminishes as $P^{-5/3}$.  Even with only a single transit observation from {\\it Kepler}, these long-period planets are invaluable.  Figure \\ref{f1} gives the total number of systems that {\\it Kepler} should observe that exhibit exactly one and exactly two transits during the mission lifetime of 3.5 years. The probability of observing a transit during the mission has been convolved with the observed fraction of planets as a function of semimajor axis for $a \\lesssim 3\\, {\\rm AU}$ and multiplied by the number of stars {\\it Kepler} will observe (100,000) to give the number of single transit systems and two transit systems that can be expected. Above $a \\gtrsim 3\\, {\\rm AU}$ the observed sample of planets as a function of semi-major axis is expected to be incomplete. We have extrapolated the number of expected single transit systems to larger semi-major axes by convolving the single transit probability distribution with the extrapolation of the fraction of planets as a function of semi-major axis by Cumming et al. (2008) that is constant in dN/dloga. The total expected numbers are 4.0 one transit events (5.7 with the extrapolation) and 5.6 two transit events.  \\begin{figure} \\begin{center} \\includegraphics[width=4in]{f1.eps} \\end{center} \\caption{The expected number of one (solid line) and two (dotted line) transit events. The probability of observing exactly one or exactly two transits during the mission lifetime is convolved with the known distribution of semi-major axes. The shaded area indicates where this distribution is expected to be incomplete, and the dashed line shows our extrapolation to larger semi-major axes based on the flat distribution fit by Cumming et al. (2008). \\label{f1}} \\end{figure}  \\begin{figure} \\begin{center} \\includegraphics[width=4in]{f2.eps} \\end{center} \\caption{The model lightcurve. $t_c$ is the time of transit center, $T$ is the duration of the full-width half-maximum depth of the transit, $\\tau$ is duration of ingress or egress, and $\\delta$ is the fractional depth of the transit. $t_1$--$t_4$ are the points of contact. $F_0 \\equiv 1$ is the out-of-transit flux, and $D$ is the total duration of the observations. \\label{f2}} \\end{figure} ",
        "conclusions": ""
    },
    "0807/0807.4331.txt": {
        "abstract": "In this article, we demonstrate that PKSJ 1037-2705 has a weak accretion flow luminosity, well below the Seyfert1/QSO dividing line, weak broad emission lines (BELs) and moderately powerful FRII extended radio emission. It is one of the few documented examples of a broad-line object in which the time averaged jet kinetic luminosity, $\\overline{Q}$, is larger than the total thermal luminosity (IR to X-ray) of the accretion flow, $L_{bol}$. The blazar nucleus dominates the optical and near ultraviolet emission and is a strong source of hard X-rays. The strong blazar emission indicates that the relativistic radio jet is presently active. The implication is that even weakly accreting AGN can create powerful jets. Kinetically dominated ($\\overline{Q}>L_{bol}$) broad-line objects provide important constraints on the relationship between the accretion flow and the jet production mechanism. ",
        "introduction": "It is unclear how the enormous stored energy in the radio lobes of powerful FRII radio sources is related to the thermal luminosity of the accreting gas that flows toward the central black hole. On the one hand, it seems reasonable to expect them to be unrelated, since the bulk of the lobe plasma was ejected from the central engine $\\sim 10^{7}$ years earlier than the optical flux from the accretion flow. Therefore, one expects to find many or most FRII sources with fossil lobes, no jets and a weak accretion flow. Yet surprisingly, low frequency surveys such as the 3C survey (which are biased towards sources with strong lobes, i.e. steep spectrum radio emission) contain many sources with FR II lobes in which high dynamic range observations reveal jets that can be traced back to the central quasar, within the resolution of the radio telescope ($\\sim$ 1kpc for nearby sources) and many others have strong radio cores and jets on VLBI (parsec) scales (eg, Cygnus A, 3C 175, 3C 334, 3C 215) \\citep{bri94,pun01}. On the other hand, there has been considerable research attempting to connect the accretion state with the jet kinetic luminosity, $Q$ - with the conflicting conclusions being the subject of much controversy \\citep{raw91,wil99,wan04,bor02}. A truly rigorous and unbiased analysis of the connection between the accretion state and $Q$ cannot exclude all of the optically weak FRII radio sources a priori, the narrow line radio galaxies (NLRGs). Many or all of the NLRGs are optically obscured by dusty molecular gas, so a proper treatment must probe inside this gas by looking at the Mid-IR emission. Alternatively, some NLRGs might have weak central AGN and are not significantly obscured. The Spitzer Telescope Mid-IR analysis of a large sample of nearby 3C FRII NLRGs in \\citet{ogl06} seemed to indicate that there was a roughly equal mix of FRII NLRGs that had no hidden quasar and those that had a hidden quasar. The implication is that there is a weak central AGN in many of the FRII objects, of which a large proportion have unobstructed lines of sight to the accretion disk. Conversely, Cleary et al 2007 analyzed an even larger sample of Spitzer observations of 3C FRII sources (NLRGS and QSOs) and concluded that orientation effects alone accounted for the differences in FRII NLRGs and FRII QSOs - i.e., for the most part NLRG had hidden QSOs of the same strength as in FRII QSOs. \\par Another closely related debate in the literature is the relationship of accretion luminosity, $L_{bol}$ (which is sometimes expressed in a related form in terms of the Eddington ratio, $R_{Edd}= L_{bol}/L_{Edd}$; where $L_{bol}$ is the bolometric thermal luminosity of the accretion flow and $L_{Edd}$ is the Eddington luminosity) to the jet kinetic luminosity, $Q$. For example, \\citet{bor02} concluded that a small $R_{Edd}$ is conducive to strong jets in quasars. The study in \\citet{wan04} was an attempt to correlate various properties of the accretion flow with the jet power. They concluded that $Q/L_{bol}$ and $L_{bol}/L_{Edd}$ were inversely correlated in blazars. The inverse correlation claimed between $Q/L_{bol}$ and $L_{bol}/L_{Edd}$ in \\citet{wan04}, although true, is a trivial consequence of the fact that $Q$ is very weakly correlated with $L_{bol}$ in quasars (the subpopulation of blazars in their sample that excluded BL-Lacs) and not unexpectedly, $L_{bol}/L_{Edd}$ and $L_{bol}$ are strongly correlated in quasars \\citep{tin05}. Thus, it was straightforward by means of a partial correlation analysis to demonstrate that inverse correlation between $Q/L_{bol}$ and $L_{bol}/L_{Edd}$ is a spurious correlation \\citep{tin05}. \\par Clearly, powerful extragalactic jets occur in high accretion thermal luminosity systems, the FRII quasars. In fact, some powerful FRII quasars have $L_{bol} > 10^{47} \\mathrm{ergs/sec}$ \\citep{tin05}. The fundamental question is whether low luminosity accretion flows can also power FRII jets. The answer to this question will constrain the dynamics of plausible central engines that can power FRII jets. In order to shed more light on these matters, we have actively searched for FRII radio sources with direct, unambiguous observational evidence of a low luminosity accretion flow below the Seyfert 1/QSO dividing line: a total thermal luminosity (IR to X-ray) of the accretion flow, $L_{bol}< 2\\times 10^{45}\\mathrm{ergs/s}$ (which is shown in section 3.2 to be equivalent to the conventional dividing line between Seyfert 1 and QSO broad-line objects, $M_{V}=-23$). To accomplish this, one requires strong evidence for a direct line of sight to the nucleus as in a nearly pole-on view, such as a blazar line of sight. Secondly, it is preferable to have two metrics of $L_{bol}$, since the optical/UV continuum is often highly contaminated by the high frequency synchrotron tail of the jet emission and is at best an upper bound to the optical/UV flux from the accretion flow. The line strength of one or more broad UV emission lines is a more direct way to measure the accretion flow luminosity in blazars. Previously, we reported on the most extreme blazar in this family, 3C 216 with $L_{bol}\\gtrsim 10^{44}\\mathrm{ergs/s}$ and a long term time averaged jet kinetic luminosity, $\\overline{Q}> 10^{46}\\mathrm{ergs/s}$, \\citep{pun07}. In fact it was shown that $\\overline{Q}> L_{Edd}$ in 3C 216. \\par This article describes the broadband properties of the core dominated radio source PKSJ 1037-2705. It was previously identified as a broadline AGN at z=0.567 from an optical spectrum with uncalibrated flux \\citep{gut06}. The radio observations presented in section 2 of this paper, indicate lobe emission with a 5 GHz luminosity typical of an FRII radio source. The implication is that the time averaged jet kinetic luminosity, $\\overline{Q}\\sim 5.0\\times 10^{44}\\mathrm{ergs/s}$ (In this paper we assume: $H_{0}$=70 km/s/Mpc, $\\Omega_{\\Lambda}=0.7$ and $\\Omega_{m}=0.3$). In section 3, new calibrated optical observations are presented that indicate a broad MgII emission line ($\\sim 4000$ km/s), and the line strength is consistent with a radio source that has a very low accretion flow thermal luminosity, $L_{bol}\\gtrsim 10^{44}\\mathrm{ergs/s}$ (well below the Seyfert 1/QSO dividing line). If it were not for the powerful jet seen in a pole-on orientation, PKSJ 1037-2705 would be a very ordinary Seyfert 1 galaxy: $L_{bol}\\gtrsim 10^{44} \\mathrm{ergs/s}$. ",
        "conclusions": "In this article, we established the following \\begin{enumerate} \\item PKSJ 1037-2705 is a blazar based on radio core dominance, a flat radio core spectrum, the large radio and X-ray variability, a steep optical spectrum and an inordinately large X-ray luminosity (relative to the UV and $L_{MgII}$). \\item PKSJ 1037-2705 has the extended radio flux typical of an FR II quasar. \\item PKSJ 1037-2705 is a broadline object that is likely to have an accretion flow luminosity an order of magnitude weaker than a quasar, i.e. it is a Seyfert 1 galaxy. \\item PKSJ 1037-2705 is likely to be kinetically dominated based on the spread in the estimated values of $\\overline{Q}$ and $L_{bol}$ that were determined in sections 2 and 3, respectively: $4.3/3.1 =1.4 < \\overline{Q}/L_{bol}< 9.3/0.82 = 11.3 $. \\end{enumerate} \\par The proof that Seyfert 1 AGN can host FRII jets is based on a few anecdotal examples. Note that these objects are distinct from broad line radio galaxies (BLRGs). The Mid-IR observations of \\citet{ogl06} indicate that FRII BLRGs have nuclei with quasar level luminosity which are often obscured by dusty gas. Even though Seyfert 1 FRII radio sources might not be exceptionally rare, demonstrating the existence of a single object explicitly is extremely challenging. For example, \\citet{ogl06} suggest that their weak Mid-IR subsample of NLRGs are hosted by low luminosity AGN, i.e., Seyfert 1 level thermal luminosity. However, \\citet{cle07} would argue that the absorbing columns to these NLRGs is so dense that even the IR can not escape and there is still a hidden quasar buried inside. \\par The existence of Seyfert 1 nuclei in FRII sources is important for understanding the connection between the accretion flow and the jet power, with two possible implications. Firstly, the kpc scale emission that was used to compute $\\overline{Q}$ was ejected from the AGN $\\sim 10^{6} - 10^{7}$ years earlier than $L_{bol}$ was created \\citep{wil99,blu00}. It is not clear if there should be a strong connection between the present value of $L_{bol}$ and the value of $L_{bol}$ $\\sim 10^{6} - 10^{7}$ years ago. On the one hand, strong variability has been detected in the thermal emission from Seyfert 1 nuclei with changes of a factor of 5 in luminosity in weeks or months \\citep{ant83,all85}. The observation of this type of variability has led researchers such as \\citet{all85,cut85,ant88} to conclude that the variability time scale for $L_{bol}$ can not be related to the thermal and viscous time scales of an accretion disk. On the other hand, there is evidence that some high redshift quasars have been active for $> 10^{7}$ years based on the $\\mathrm{He}^{+}$ Lyman $\\alpha$ \"proximity effect\" \\citep{jak03}. The ``proximity effect'' arises when a large bubble forms about the quasar in which helium is highly ionized by the quasar. One can use the size of the bubble to estimate the amount of time it would take for the photo-ionizing source to ionize the extended region. However, one must be cautious to extrapolate this result to much less luminous Seyfert 1 galaxies and it does not provide evidence that the photo-ionization source was steady within a factor of 10 during the entire $10^{7}$ years. With all this being said, it would be a huge stretch in reasoning to assume that in spite of the large short term variations that have been observed in some Seyfert 1 galaxies that $L_{bol}$ is constant to within a factor of 10 over $10^{7}$ years. Hence, there does not seem to be any compelling observation or underlying reason why the two quantities, $\\overline{Q}$ and $L_{bol}$ should be related, even if the accretion flow drives the jet. \\par Alternatively, the existence of Seyfert 1 nuclei in FR II radio sources could mean that the accretion rate is not that strongly coupled to the jet power. Unfortunately, as a consequence of the Doppler boosting there is no reliable contemporaneous estimate of $Q(t)$ that can be used to test the latter hypothesis (see the discussion in section 2.2). In the context of our previous work, we have now reported on 3 Seyfert 1 galaxies that have $\\overline{Q} \\geq L_{bol}$, PKSJ 1037-2705 ($L_{bol}\\gtrsim 10^{44}$ ergs/s), 3C 216 ($L_{bol}\\gtrsim 10^{44}$ ergs/s) and PKS 1622-253 ($L_{bol}\\approx 2\\times 10^{45}$ ergs/s) \\citep{pun07,rod05}. All three have convincing evidence of a contemporaneous powerful jet and a weak (Seyfert 1 level) $L_{bol}$. The evidence for an active jet in PKSJ 1037-2705 is a powerful peak in the synchrotron SED, $\\nu L_{\\nu}>10^{45}$ ergs/s, in the IR based on archival 2MASS data. Similarly, the two epochs of X-ray data indicates that it is plausible that the inverse Compton peak of the SED in the X-ray band, (0.3 keV - 10 keV in the rest frame) exceeds $10^{45}$ ergs/s during the high states (ROSAT epoch). 3C 216 has a strong peak in the synchrotron SED $>10^{46}$ ergs/s. PKS 1622-253 is one of the strongest EGRET gamma ray sources, the gamma-ray apparent luminosity has a time average value of $\\sim 10^{47}\\mathrm{ergs/sec}$ and flares at $\\sim 10^{48}\\mathrm{ergs/sec}$ \\citep{har99}. This circumstantial evidence tends to support the interpretation that even weak (Seyfert 1 level) accretion flows onto a black hole can produce a central engine for FRII jets. Thus, our study is supportive of the \\citet{ogl06} interpretation of the weak Mid-IR, FRII NLRG subsample, they contain low luminosity nuclei. \\par Finally, we comment on the conjecture in \\citet{bor02,mac03} that these three Seyfert I galaxies, possessing a large jet power (at the FRII level), are not unusual objects, based on the small value of $L_{bol}/L_{Edd}$ in Seyfert 1 galaxies compared to QSOs \\citep{sun89}. This conclusion is a natural consequence of the claim in \\citet{bor02,mac03} that a small value of $L_{bol}/L_{Edd}$ is conducive to jet formation. The implication is that Seyfert I galaxies (with their relatively low values of $L_{bol}/L_{Edd}$ compared to QSOs), possessing FR II radio power, should be a common state of AGN activity. However, this conjecture seems difficult to reconcile with the sparsity of known FR II level extended emission associated with broad line nuclei with Seyfert 1 level luminosity in optically selected samples (there is not a single example in the New General Catalog, Palomar-Green Survey or the Markarian catalog). Such a simple proposal is also at odds with the anecdotal case of PKS~0743$-$67 that was studied in detail in \\citet{tin05} for just this reason. PKS~0743$-$67 is an example of a quasar that has an ultra-luminous accretion flow, $L_{bol}>2\\times 10^{47}\\mathrm{ergs/s}$, and has a very high Eddington rate, $L_{bol}/L_{Edd}\\approx 1$. The jet kinetic luminosity was conservatively estimated at $\\overline{Q}= 4.1 \\times 10^{46}$ ergs/sec which is 2.5 times that of Cygnus A. However, this was a conservative lower bound and the radio data also supports $\\overline{Q}\\approx 10^{47}$ ergs/sec, i.e one of the most powerful radio sources in the known Universe. Furthermore, PKS~0743$-$67 is presently active as evidenced by the powerful ($> 1 Jy$) unresolved VLBI radio core. A low Eddington ratio is not likely to be determinant to FRII jet production."
    },
    "0807/0807.2467_arXiv.txt": {
        "abstract": "{In the context of exoplanet detection, the performance of coronagraphs is limited by wavefront errors.} {To remove efficiently the effects of these aberrations using a deformable mirror, the aberrations themselves must be measured in the science image to extremely high accuracy.} {The Self-Coherent Camera which is based on the principle of light incoherence between star and its environment can estimate these wavefront errors. This estimation is derived directly from the encoded speckles in the science image, avoiding differential errors due to beam separation and non common optics.} {Earth-like planet detection is modeled by numerical simulations with realistic assumptions for a space telescope.} {The Self-Coherent Camera is an attractive technique for future space telescopes. It is also one of the techniques under investigation for the E-ELT planet finder the so-called EPICS.} ",
        "introduction": "Very high contrast imaging is mandatory for the direct detection of exoplanets, which are typically a factor of between~$10^7$ and~$10^{10}$ fainter than their host and often located within a fraction of an arcsecond of their star. First of all, coronagraphs are required to suppress the overwhelming flux of the star but they are limited by wavefront errors in the upstream beam, which creates residual speckles in the science image. Adaptive optics must be used to correct for the effect of most of these aberrations. Some remain uncorrected generating quasi-static residual speckles\\,\\citep{Cavarroc06}. Interferential techniques take advantage of the incoherence between companion and stellar lights to measure these wavefront errors in the science image to high accuracy\\,\\citep{Codona04,Guyon04}. In this Letter, we describe such a technique called a Self-Coherent Camera\\,\\citep{Baudoz06}. Residual speckles in the science image, also called interferential image hereafter, are spatially encoded by fringes so that we can derive an estimation of wavefront errors to be corrected by a Deformable Mirror~(DM). Since the number of DM actuators is finite, this correction leaves residual speckles. Thus, after reaching the DM limit correction, we apply an image post-processing algorithm\\,\\citep{Galicher07} to achieve Earth-like planet imaging. Hereafter, we detail the SCC principle. Then, we describe the wavefront error estimator that we use. Finally, we present expected performances from space. ",
        "conclusions": "We have numerically demonstrated that the SCC associated with a $32 \\times 32$ DM  enables us to detect Earths from space in a few hours when using realistic assumptions~(zodiacal light, photon noise, read out noise, VLT pupil aberrations, and $20\\%$ bandwidth). SCC could be a good candidate to be implemented in the next generation of space telescopes. The technique involves two steps. We first use SCC to estimate wavefront errors and operate a DM that completes the correction in a few steps. To overcome the limitation linked to the DM uncorrectable high-order frequencies, we apply  to the final iteration image, the SCC post-processing algorithm. This post-processing has yet to be optimized.  SCC is one of the techniques under investigation for the E-ELT planet finder so-called EPICS. For this reason, we propose to consider the impact of different parameters, such as amplitude errors and turbulence residuals on the SCC performance. We will also test the compensation for amplitude errors proposed by~\\citet{Borde06}. A preliminary study, which assumes a more realistic coronagraph~(achromatic Four Quadrant Phase Mask), indicates that our algorithm converges but more slowly than with a perfect coronagraph. The quality of the reference beam should not be important for SCC because of the filtering by the pinhole and the reduction in the beam diameter~($D_R$), which induces a wide diffraction pattern in the focal plane. Experimental validations of the SCC technique are also planned soon.  We thank Pascal~Bord\\'{e} and Anthony~Boccaletti for useful discussions."
    },
    "0807/0807.4594_arXiv.txt": {
        "abstract": "{} { We explore the gamma-ray bursts' (GRBs') contributions in constraining the dark energy equation of state (EOS) at high ($1.8 < z < 7$) and at middle redshifts ($0.5 < z < 1.8$) and estimate how many GRBs are needed to get substantial constraints at high redshifts. } { We estimate the constraints with mock GRBs and mock type Ia supernovae (SNe Ia) for comparisons. When constraining the dark energy EOS in a certain redshift range, we allow the dark energy EOS parameter to vary only in that redshift bin and fix EOS parameters elsewhere to $-1$. } { We find that it is difficult to constrain the dark energy EOS beyond the redshifts of SNe Ia with GRBs unless some new luminosity relations for GRBs with smaller scatters are discovered. However, at middle redshifts, GRBs have comparable contributions with SNe Ia in constraining the dark energy EOS. } {}  \\keywords { cosmological parameters - Gamma rays: bursts } ",
        "introduction": "Since the discovery of the accelerating expansion of the universe~\\citep{Riess:1998cb, Perlmutter:1998np}, a lot of work has been done on constraining the behavior of dark energy. In contrast to simple parametric forms for dark energy equation of state (EOS), such as $w(z) = w_0+w'z$~\\citep{Cooray:1999da} and $w(z)=w_0+w_az/(1+z)$~\\citep{Chevallier:2000qy, Linder:2002et}, a nearly model-independent approach was introduced by~\\citet{Huterer:2004ch} to estimate the evolution of dark energy independently, and was adopted in analyses using type Ia supernova (SN Ia) data~\\citep{Riess:2006fw, Sullivan:2007pd, Sullivan:2007tx}. \\citet{Qi:2008zk} extended the approach to gamma-ray burst (GRB) luminosity data~\\citep{Schaefer:2006pa}, for which higher redshifts are accessible compared to SNe Ia. Though more stringent results have been obtained by this, the dark energy EOS at high redshifts, where we only have GRBs, is still totally unconstrained, except that it is most likely negative. This is primarily due to matter dominating dark energy at high redshifts and to dark energy becoming less important for determining the cosmic expansion; therefore, constraining it becomes more difficult. In this work we explore both how many GRBs are needed to get substantial constraints on dark energy EOS at high redshifts beyond those of SNe Ia and where the GRBs' contributions lie for a foreseeable number of GRBs. ",
        "conclusions": "We explored the GRBs' contributions in constraining the dark energy EOS at high redshifts ($1.8 < z < 7$) and at middle redshifts ($0.5 < z < 1.8$). When constraining the dark energy EOS in a certain redshift range, we allow the dark energy EOS parameter to vary only in that redshift bin and fix EOS parameters elsewhere to $-1$. We find that it is difficult to constrain the dark energy EOS parameters beyond the redshifts of SNe Ia with GRBs unless some new luminosity relations for GRBs with smaller scatters are discovered. However, at middle redshifts, GRBs have contributions comparable with SNe Ia in constraining the dark energy EOS."
    },
    "0807/0807.4077_arXiv.txt": {
        "abstract": "{Usually, a loose stellar distribution can be classified as an OB stellar group, an association, or a young open cluster. We make use of comparisons with the typical OB association Vul\\,OB1} {In the present paper we discuss the nature of Bochum\\,1, a typical example of an object affected by the above classification problem.} {Field-decontaminated 2MASS photometry is used to analyse Colour-Magnitude Diagrams (CMDs) and stellar radial density profiles (RDPs) of the structures present in the region of Bochum\\,1.} {The field-decontaminated CMD of Bochum\\,1 presents main sequence (MS) and pre-main sequence (PMS) stars. We report two new small angular-size, compact young clusters and one embedded cluster in the area of Bochum\\,1. Vul\\,OB1 harbours the young open cluster NGC\\,6823 and the very compact embedded cluster Cr\\,404. The Vul\\,OB1 association includes the H\\,II region Sh2-86, and its stellar content is younger ($\\approx3$\\,Myr) than that of Bochum\\,1 ($\\approx9$\\,Myr), which shows no gas emission. Bochum\\,1 harbours one of the newly found compact clusters as its core. The RDP of Bochum\\,1 is irregular and cannot be fitted by a King-like profile, which suggests important erosion or dispersion of stars from a primordial cluster. Similarly to Bochum\\,1, the decontaminated CMD of NGC\\,6823 presents conspicuous MS and PMS sequences. Taken separately, RDPs of MS and PMS stars follow a King-like profile. The core shows an important excess density of MS stars that mimics the profile of a post-core collapse cluster. At such young age, it can be explained by an excess of stars formed in the prominent core.} {The present study suggests that Bochum\\,1 is a star cluster fossil remain that might be dynamically evolving into an OB association. Bochum\\,1 can be a missing link connecting early star cluster dissolution with the formation of low-mass OB associations.} ",
        "introduction": "\\label{intro}  Associations are loose stellar systems that may contain as much as 2600 stars, as in Cyg\\,OB2 (\\citealt{Albacete02}; \\citealt{Kn00}). Although sharing a common origin and moving approximately in the same direction through the Galaxy, the member stars are gravitationally unbound. This definition encompasses a wide variety of objects, from the extended OB associations in spiral arms (\\citealt{Blaauw64}) to the post-T Tauri associations in the Solar neighbourhood (e.g. \\citealt{Torres00}). OB associations can be observed over a wide range of distances from the Sun, from the relatively nearby ($\\approx140$\\,pc) Scorpius-Centaurus Association (\\citealt{MA01}) to the sparse and large ($\\sim400$\\,pc) associations in the Large Magellanic Cloud and Andromeda (\\citealt{EfEl98}).  Most field stars appear to have been formed in stellar groups of different kinds, the OB associations in particular (Gomes et al. 1993; \\citealt{Massey95}). The rapid early gas removal is an efficient mechanism to drive cluster stars into the field and dissolve most of the very young star clusters in a time scale of $10-40$\\,Myr, depending on cluster mass and star-formation efficiency (e.g. \\citealt{GoBa06}). As a consequence of this infanticide, only about 5\\% (\\citealt{LL2003}) of the embedded clusters are able to dynamically evolve into bound open clusters (OCs).  The above aspects raise the fundamental issue of the detection of dispersed cluster debris, and the distinction (if possible) between genuine small associations from dispersing debris. The present scenario, where early-cluster disruption is favoured, may provide clues to the dispersion of large star-forming complexes into the field, and remaining young OC families (\\citealt{Piskunov06}; \\citealt{laF08}).  Bochum\\,1 was defined by \\citet{Moffat75} as a group of 8 OB stars from the catalogue of Southern Luminous Stars (LS) by \\citealt{StepSand71}. \\citet{Moffat75} carried out photoelectric observations with the ESO La Silla Bochum telescope, hence the object designation. They pointed out that the stars stand out from the background, and the Colour-Magnitude Diagram (CMD) indicated a common reddening and distance. They derived $\\ebv=0.55\\pm0.06$, $\\ds=4.06$\\,kpc and an O7 turn-off. Bochum\\,1 is located at $\\ell=193.43^\\circ$, $b=+3.40^\\circ$, and $\\rm\\alpha(J2000)=6^h25^m30^s$ and $\\rm\\delta(J2000)=19^\\circ46\\arcmin00\\arcsec$.  \\citet{Yadav03} observed Bochum\\,1 with CCD photometry and referred to it as an  OB association. They found that part of the bright stars in the area have a common proper motion, and a mass function that, within uncertainties, is comparable to that of \\citet{Salpeter55}. They derived $\\ebv=0.47\\pm0.10$, $\\ds=2.8$\\,kpc, and an age of 10\\,Myr. Bochum\\,1 is included as a star cluster in the WEBDA\\footnote{\\em http://www.univie.ac.at/webda - \\citet{Merm03}} database, which shows the values of \\citet{Yadav03}.  Recently, \\citet{FSRcat} presented a catalogue of star cluster candidates detected as stellar overdensities in the 2MASS\\footnote{The Two Micron All Sky Survey --- {\\em www.ipac.caltech.edu/2mass/releases/allsky/ }} catalogue. FSR\\,911 was considered to be the same object as Bochum\\,1, but the positions are not coincident (Fig.~\\ref{fig1}).  In this context, the following interesting questions arise: Is Bochum\\,1 a missing link between star clusters and stellar associations? A dispersing fossil remain of a star cluster where scattered early-type stars and a remnant core are observed? Or a proto association that will eventually show up only scattered stars?  In a recent series of studies we coupled the classical CMD method with the analysis of the stellar radial density profile (RDP) to obtain intrinsic cluster astrophysical parameters and, in uncertain cases, to establish the nature of the objects (e.g. \\citealt{BB08}; \\citealt{ProbFSR}; \\citealt{OldOCs}; \\citealt{BB07}; \\citealt{N4755}). Field decontamination was crucial in all these studies.  A previous discussion on the connection between massive young star cluster and a massive OB association was carried out by \\citet{Kn00} with Cyg\\,OB2. In the present study we apply these methods to Bochum\\,1 and other low-mass stellar systems in the area. We also analyse NGC\\,6823 in Vul\\,OB1, for comparison purposes among young stellar systems.  \\begin{table*} \\caption[]{General data on the star cluster/associations} \\label{tab1} \\renewcommand{\\tabcolsep}{1.05mm} \\renewcommand{\\arraystretch}{1.25} \\begin{tabular}{ccccccccc} \\hline\\hline $\\ell$&$b$&$\\alpha(2000)$&$\\delta(2000)$&\\multicolumn{2}{c}{Diameter}&\\rx&Designation&Comments\\\\ \\cline{5-6} ($^\\circ$)&($^\\circ$)&(hms)&($^\\circ\\,\\arcmin\\,\\arcsec$)&(\\arcmin)&(pc)&(\\arcmin)\\\\ (1)&(2)&(3)&(4)&(5)&(6)&(7)&(8)&(9)\\\\ \\hline \\multicolumn{9}{c}{The region of Bochum\\,1}\\\\ \\hline 192.44 & $+$3.41 & 6:25:30 & $+$19:46:00 & 26.0$^a$ & 32$^d$ &60&Bochum\\,1\\\\ 192.30 & $+$3.36 & 6:25:00 & $+$19:52:03 & 14.0$^b$ & 17$^e$ &30& FSR\\,911\\\\ 192.64 & $+$3.98 & 6:25:03 & $+$19:53:00 &  2.5$^c$ & 3.1$^e$ &30& Faint-star clump & Core of FSR\\,911?\\\\ 192.31 & $+$3.36 & 6:25:01 & $+$19:50:55 &  1.0$^c$ & 1.3$^d$ &30 &New\\,Cluster\\,1 & Core of FSR\\,911? Includes LS46 and LS47\\\\ 192.43 & $+$3.40 & 6:25:28 & $+$19:45:10 &  0.6$^c$ & 1.2$^e$ &&New\\,Cluster\\,2 & Core of Bochum\\,1? Includes LS51\\\\ 192.17 & $+$3.41 & 6:24:55 & $+$19:59:59 &  1.0$^c$ & 1.3$^d$ &30&New\\,Cluster\\,3 & Core of FSR\\,911? Includes IRAS\\,06219+2001\\\\ \\hline \\multicolumn{9}{c}{The region of NGC\\,6823}\\\\ \\hline 59.40 & $-$0.14 & 19:43:09 & $+$23:17:58 & 10.0 & 5.8$^d$ &30&NGC\\,6823,Cr\\,405,OCl-124\\\\ 59.14 & $-$0.11 & 19:42:28 & $+$23:05:13 &  0.8 & 0.5$^d$ &20&Cr\\,404,OCl-122 & Embedded cluster, includes IRAS\\,19403+2258\\\\ \\hline \\end{tabular} \\begin{list}{Table Notes.} \\item Cols.~1-4: Central coordinates. Angular diameters (Col.~5) from: (a) - the catalogue of \\citet{Dias02}; (b) - twice the tidal radius of \\citet{FSRcat}; (c) - estimated in the present study from 2MASS \\ks\\ images. Absolute apparent diameters (Col.~6) computed with distances from the Sun from (d) - this work (Table~\\ref{tab4}); (e) - assuming the same distance as that of Bochum\\,1. Col.~7: 2MASS extraction radius. \\end{list} \\end{table*}  This paper is structured as follows. In Sect.~\\ref{targets} we provide fundamental data on Bochum\\,1 and the other relevant stellar clusterings in its field, build the 2MASS CMDs and RDPs, and derive fundamental and structural parameters. In Sect.~\\ref{N6823} we apply the above methods to the template young OC NGC\\,6823. In Sect.~\\ref{MF} we build mass functions and compute the stellar content. In Sect.~\\ref{Disc} we discuss whether Bochum\\,1 is a genuine association or a dispersed young star cluster (or more than one). Concluding remarks are given in Sect.~\\ref{Conclu}. ",
        "conclusions": "\\label{Conclu}  Because of its sparse stellar distribution, previous studies have considered Bochum\\,1 as an OB stellar group, an open cluster, or an association. In the present paper we investigate its nature with field-star decontaminated 2MASS photometry and stellar radial density profiles. With an age of $\\approx3$\\,Myr, the Vul\\,OB1 association, which contains the young OC NGC\\,6823 and the very compact embedded cluster Cr\\,404, are used for comparisons with the young stellar content and stellar radial density profiles present in the Bochum\\,1 area.  Likewise NGC\\,6823, conspicuous sequences of MS and PMS stars are present in the decontaminated CMD of Bochum\\,1, consistent with the $\\approx9$\\,Myr of age. However, the RDP of Bochum\\,1 is irregular and does not follow a King-like profile, which suggests important erosion or dispersion of stars from a possible primordial cluster. The MS and PMS sequences of NGC\\,6823 produce King-like RDPs, but with an important excess of MS stars near the centre. Unlike Vul\\,Ob1, which includes the H\\,II region Sh2-86, no evident gas emission appears in Bochum\\,1.  We found two new small compact young clusters in the area of Bochum\\,1 (one of which may be its remnant core), as well as one embedded cluster. One of the small compact clusters is located $\\approx8\\arcmin$ to the NW of Bochum\\,1 (probably the remnant core of a large previous cluster). Together with a clump of faint stars somewhat to the north, both have been taken as the core of an extended cluster (FSR\\,911).  Available evidence shows that structurally, Bochum\\,1 is not a star cluster. A possible scenario points to a fossil remain of a star cluster, as suggested by the core (New\\,Cluster\\,2) and the radially decaying stellar density profile. In this context, Bochum\\,1 can be a missing link connecting young star cluster dissolution with the formation of low-mass OB associations. However, the processes that generate large OB associations such as Vul\\,OB1 are yet to be explored in more detail. Probably the difference between objects like Bochum\\,1, and the star cluster NGC\\,6823 arises from the mass stored in the objects. On the other hand, the erosion of a primordial cluster is not the only explanation for the irregular radial stellar density profile of Bochum\\,1. Indeed, star formation may occur along the border of a radially expanding density wave or ionisation front (e.g. \\citealt{Soria05}; \\citealt{Elme77}). In some cases, the neutral interstellar medium can be compressed by the expanding bubble above the stability criterion against gravitational collapse, which could trigger localised star formation, giving rise to a clumpy radial density profile.  The two compact New\\,Clusters\\,1 and 2, and the embedded New\\,Cluster\\,3, are new findings in the area of Bochum\\,1 and FSR\\,911. We show that Bochum\\,1 and FSR\\,911 are different objects. The optically identified small object New\\,Cluster\\,1 and the large infrared detected (Fig.~\\ref{fig1} and Tab.~\\ref{tab6}) may be the same object. However, probably because of the lack of decontamination, \\citet{FSRcat} overestimated dimensions of FSR\\,911. The whole ensemble may turn out to be a Rosetta Stone to decode early-dynamical evolution processes involving embedded clusters, young compact clusters, OCs and associations. However, to unravel all evolutionary connections we probably must wait for Gaia's\\footnote{http://www.rssd.esa.int/index.php?project=Gaia} deep proper motions."
    },
    "0807/0807.3657_arXiv.txt": {
        "abstract": "{Time-dependent cooling processes are of paramount importance in the evolution of astrophysical gaseous nebulae and, in particular, when radiative shocks are present. Given the recent improvements in resolution of the observational data, simulating these processes in a more realistic manner in magnetohydrodynamic (MHD) codes will provide a unique tool for model discrimination.} {The present work introduces a necessary set of tools that can be used to model radiative astrophysical flows in the optically-thin plasma limit. We aim to provide reliable and accurate predictions of emission line ratios and radiative cooling losses in astrophysical simulations of shocked flows. Moreover, we discuss numerical implementation aspects to ease future improvements and implementation in other MHD numerical codes.} {The most important source of radiative cooling for our plasma conditions comes from the collisionally-excited line radiation. We evolve a chemical network, including 29 ion species, to compute the ionization balance in non-equilibrium conditions. The numerical methods are implemented in the PLUTO code for astrophysical fluid dynamics and particular attention has been devoted to resolve accuracy and efficiency issues arising from cooling timescales considerably shorter than the dynamical ones. } {After a series of validations and tests, typical astrophysical setups are simulated in 1D and 2D, employing both the present cooling model and a simplified one. The influence of the cooling model on structure morphologies can become important, especially for emission line diagnostic purposes.} {The tests make us confident that the use of the presented detailed radiative cooling treatment will allow more accurate predictions in terms of emission line intensities and shock dynamics in various astrophysical setups.} ",
        "introduction": "Astrophysical gases emit thermal radiation while undergoing dynamical transformations. There are cases where the gas is so diluted that the typical timescales for cooling greatly exceed the dynamical ones but, in many instances, cooling and the related ionization/recombination processes for the emitting species become comparable to, or faster than, the dynamical evolution of the system and should be considered. Classical examples of intensively radiating gases are \\ion{H}{ii} regions, planetary nebulae, supernova remnants and star forming regions. Thus, when studying gas flows in such environments, particular care must be taken to treat the interplay between dynamics and radiation in the correct way. This is particularly true and crucial whenever radiative shocks are involved. When cooling is strong, the ionization fraction of the emitting species is far from that at equilibrium, but evolves so rapidly with time that it must be treated in a time-dependent fashion.  Under these conditions, the magnetohydrodynamic (MHD) equations, coupled with the equations describing the evolution of the emitting species and the radiative losses, must be solved by numerical means. The numerical problems posed by radiative cooling are particularly challenging whenever the advance time step of the integration is controlled by radiation/ionization rather than dynamics. This may happen in few grid points in the computational domain where the radiative losses are intense, e.g., right behind a shock front, where the cooling time becomes very small. Therefore, it is necessary to devise strategies able to deal with very different integration timescales: this is one of the points we address.  Time dependent ionization calculations were previously performed for gaseous nebulae by Marten \\& Szczerba (\\cite{MS97}) in hydrostatic conditions. Their approach is similar to our implementation of the ionization state treatment, however the ion species and implemented physical processes are in part different. Also, radiative cooling in optically-thin plasmas was previously investigated, and synthetic cooling functions were designed (e.g., Schmutzler \\& Tscharnutter \\cite{ST93}). Among the radiative numerical codes employed in astrophysics, one can quote the hydrocode YGUAZ\\`U (Raga et al. \\cite{RA00}), ASTROBEAR MHD code (Poludnenko et al. \\cite{PA05}, Berger \\& LeVeque \\cite{BL98}) and Virginia Hydrodynamics - 1 (VH-1, Sutherland et al. \\cite{SA03}, Blondin \\& Lufkin \\cite{BL93}). The MHD simulation code we use for our astrophysical applications -- PLUTO -- is a freely distributed application developed and maintained at the Turin University -- Turin Astronomical Observatory (Mignone et al. \\cite{MA07}). A previous numerical analysis about the evolution of radiative shocks in Young Stellar Object (YSO) jets (Massaglia et al. \\cite{MA05}) was carried out with PLUTO, using a simplified model for the radiative cooling losses, which evolved in time only the ionization fraction of hydrogen (c.f., Rossi et al. \\cite{RA97}). This model will be called from now on SNEq (Simplified Non-Equilibrium cooling).  We illustrate a more general treatment of atomic cooling and evolution of the ionization fraction of the emitting species, embedded in the PLUTO code as well, for use within MHD simulations of astrophysical interest. We will call this new cooling function MINEq (Multi-Ion Non-Equilibrium cooling). The main advantage of our approach is the full ionization state computation during the MHD simulation, which allows for better predictions of emission line intensities.  Section \\ref{general} contains a general overview of the adopted method and implementation of the treatment of radiative losses. Then, in Sect. \\ref{cfdesc}, a description of the physics of the cooling model can be found, followed in Sect. \\ref{testeq} by the validations and tests in equilibrium conditions. The numerical implementation and testing are discussed in Sect. \\ref{numerical}, while in Sect. \\ref{testint} we present some typical astrophysical applications. Technical details on ionization-recombination processes and numerical issues are presented in extended form in the Appendix. ",
        "conclusions": "After a series of tests and validation, the detailed treatment of radiative losses MINEq is now implemented in the MHD code PLUTO. The choice of the integration technique and the optimizations are, as far as we know, unique at the time of this writing and provide a high degree of accuracy in ion species abundances and radiative losses computation. Both theoretical and technical aspects of the current implementation, as well as testing process and applications were discussed in the previous sections.  A major advantage obtained by using MINEq, compared to previously employed SNEq, is that the line emission can be computed in conditions of non-equilibrium ionization for all species, more likely to be encountered in situations of rapid changes, as it is the case of shock waves.  As shown by the tests presented, the choice of the cooling model between MINEq and SNEq, has an increasing effect on the structure formation whenever temperatures exceed $25\\,000$ K, and has an important influence for the ionization fractions that reflects in the emission line computations. For a preliminary dynamical study, a tabulated cooling function that does not integrate any ion specie can be employed, but the relatively low computational cost of a cooling that evolves the hydrogen ionization in a time-dependent fashion makes the latter advisable in most cases. Whenever the purpose is the computation of emission line ratios, employing a detailed cooling like MINEq produces more reliable results as the ionization fractions are followed in non-equilibrium conditions, twhich are likely to be encountered in the real astrophysical objects.  An important feature of the cooling model implementation is that it is upgradeable, more ion species and other processes can be added (increasing also the temperature range of application). Also, it can be used as starting point for the integration of atomic chemistry and cooling processes to other MHD codes.  The newly-developed cooling function provides a powerful tool for investigating the stellar jets and gaseous nebulae. It is, in the current configuration, suited for the study of radiative shocks in stellar jets. High resolution, adaptive numerical simulations to predict line emission in YSO jets will be subject of forthcoming works.   \\appendix"
    },
    "0807/0807.0113_arXiv.txt": {
        "abstract": "We study the evolution of an accelerating hyperrelativistic shock under the presence of upstream inhomogeneities wrinkling the discontinuity surface. The investigation is conducted by means of numerical simulations using the PLUTO code for astrophysical fluid dynamics. The reliability and robustness of the code are demonstrated against well known results coming from the linear perturbation theory. We then follow the nonlinear evolution of two classes of perturbing upstream atmospheres and conclude that no lasting wrinkle can be preserved indefinitely by the flow. Finally we derive analytically a description of the geometrical effects of a turbulent upstream ambient on the discontinuity surface. ",
        "introduction": "{} There seems to be strong evidences that Gamma-Ray Bursts (GRBs -- see Piran (2005) for a review) involve flows of dense shells thrown by a dying compact star in the ambient medium at Lorentz factors $\\Gamma>10^2-10^3$. When the ejecta impact on the surrounding matter carried by a pre-existing stellar wind, a shock is formed and begins to propagate into a decreasing atmosphere, circumstance which leads, for sufficiently steep density profiles, to the shock acceleration. The length scale $k_0^{-1}$ on which the stellar atmosphere rarefies maybe reasonably much smaller than the distance from the center of the star, thus justifying the approximation of planar symmetry in studying the shock evolution.  The problem of a blast wave moving into a decreasing atmosphere has been analyzed both in its Newtonian~\\citep{gand,sakurai,raizer,grover,hardy} and relativistic~\\citep{bmk,best,perna,naka,pan,sari} regimes, and several self-similar solutions have been found for the flow in both power-law and exponentially shaped density profiles. Despite the importance of the issue, very few papers have been spent to study the stability of the system subject to wrinkling perturbations. In Newtonian regime~\\citet{chev1} and~\\citet{luo} studied the exponential atmosphere; power-law profiles have been considered in~\\citet{sari2}. \\citet{wang}, while taking into account relativistic effects, were not able to find any self-similarity in the perturbation analysis of a spherical blast wave propagating in a power-law atmosphere.  \\citet{palma} performed a linear stability analysis of a highly relativistic planar shock propagating in an exponential atmosphere and retrieved a self-similar solution for the first order problem. They obtained that, at least in the small perturbation limit, with respect to what happens in the Newtonian regime, the corrugation wavelength $k^{-1}$ can drop by a factor of $\\Gamma$ still giving rise to no sensible restoring effect in the flow, a behaviour reminiscent of the infinite wavelength case, even for small ratios $k_0/k$. This allows the instability of the downstream energy density to persist, thus delaying the saturation phase.  Of course, as the instability exits the linear regime, a numerical approach has to be adopted since the arguments that exclude the arising of stabilizing phenomena in the flow may become weaker and not so pertinent.  The plan of the paper is as follows. In \\S\\ref{par1} we review the analytical properties of the problem extensively discussed in~\\citet{perna} and~\\citet{palma}: these predictions provide useful benchmarks to which our numerical scheme can be compared. In \\S\\ref{par2} we describe the code used to integrate the relativistic hydrodynamics (RHD) equations; a specific subsection is reserved to explain how we overcome the relevant technical difficulties. \\S\\ref{par3} lists all the major tests through which we run the code before declaring it reliable for our purposes. In \\S\\ref{par4} we tackle the central subject of the paper, thus reporting several results of simulations dealing with nonlinear variations of the perturbations induced into the system in \\S\\ref{par3}. Lastly, in \\S\\ref{par5}, we will show that self-similarity is reached fast enough to allow for an analytical expression for the shock speed in a quite arbitrarily shaped atmosphere. By means of such a result we will develop a technique to calculate the shock position without making time expensive simulations and will apply it to describe the shape evolution of a planar shock impacting a turbulent upstream. Conclusions are drawn in the shape of an excursus in \\S\\ref{concl}. ",
        "conclusions": ""
    },
    "0807/0807.1786_arXiv.txt": {
        "abstract": "{} {The TeV BL Lac object Markarian 501 is a complex, core dominated radio source, with a one sided, twisting jet on parsec scales. In the present work, we attempt to extend our understanding of the source physics to regions of the radio jet which have not been accessed before. }  {We present new observations of Mrk 501 at 1.4 and 86 GHz. The 1.4 GHz data were obtained using the Very Large Array (VLA) and High Sensitivity Array (HSA) in November 2004, in full polarization, with a final r.m.s.\\ noise of $25\\, \\mu$Jy/beam in the HSA total intensity image; the 86 GHz observations were performed in October 2005 with the Global Millimeter VLBI Array (GMVA), providing an angular resolution as good as $110\\, \\mu\\mathrm{as} \\times \\, 40 \\, \\mu\\mathrm{as}$.}  {The sensitivity and resolution provided by the HSA make it possible to detect the jet up to $\\sim$700 milliarcseconds (corresponding to a projected linear size of $\\sim$500 pc) from its base, while the superior resolution of the 86 GHz GMVA observations probes the innermost regions of the jet down to $\\sim$ 200 Schwarzschild radii. The brightness temperature at the jet base is in excess of $6 \\times 10^{10}$ K.  We find evidence of limb brightening on physical scales from $\\la1$ pc to $\\sim 40$ pc. Polarization images and fits to the trend of jet width and brightness vs.\\ distance from the core reveal a magnetic field parallel to the jet axis.}  {} ",
        "introduction": "The study of extragalactic radio jets is an important area in astrophysics. In radio loud sources, jets contribute a large fraction of the total radiated power, and sustain the formation of energetic kiloparsec scale lobes.  While observational properties of jets are widely differentiated, they are present in high and low power sources, with some common features; on the parsec scale, they are relativistic in both types and they are also intrinsically identical in beamed and misaligned sources.  \\citet{gio01} have shown that the Lorentz factor in the parsec scale jet of low power FRI radio galaxies as well as of more powerful FR IIs are both in the range $\\Gamma = 3 - 10$. With these values, jets also appear intrinsically identical in beamed (BL Lac objects) and misaligned sources (FRI radio galaxies), if the former have jet axis oriented at an average viewing angle of $\\langle \\theta \\rangle = 18^\\circ \\pm 5^\\circ$ \\citep{gir04b}.  In the present paper we focus on the jet structure of the BL Lac source Markarian 501. This object is highly active and well-studied at all frequencies. Its activity and variability at high energy \\citep[as high as TeV regime,][]{qui96} seems to require high Doppler factors and consequently a small angle to the line of sight. In the radio band, centimeter VLBI observations have revealed a clear limb-brightened structure, beginning in the very inner jet, suggestive of a dual velocity structure \\citep[~hereinafter G04]{gir04a}. The complex limb-brightened structure makes component identification problematic and multi-epoch attempts to measure pattern speed conclude that it is not well defined \\citepalias{gir04a} or in any case at most subluminal \\citep{edw02}. These seem to be common features in TeV blazars \\citep{pin04,gir06}, and theoretical models have been proposed to reconcile them with the very high energy emission \\citep{ghi05,wan04}.  However, the results obtained so far still leave some major questions unanswered. For example, it is not at all clear whether the jet velocity structure is intrinsic or produced by the interaction with the surrounding medium. We want to understand why the properties of the radio jet on parsec scales are different from those needed to explain the $\\gamma-$ray emission in Mrk~501, as well as in other TeV blazars. Moreover, a change in regime must occur on much larger scales, since the large scale structure of the source is known to be symmetric rather than one-sided \\citep[e.g.,][]{ulv83}. We want to investigate if this transition is smooth and what is the configuration of the magnetic field in the outer jet, which needs sensitive images in total intensity and polarization.  In order to search for an answer to such questions, we need to go beyond the capability of the instruments available for ordinary centimeter wavelength VLBI, which provides information only for the region between $\\sim 1$ and $\\sim 100$ milliarcseconds. Smaller and larger scale regions remain inaccessible because of inadequate resolution and sensitivity, respectively.  Improvements in the technical and organizational issues are now offering to astronomers VLBI arrays of unprecedented resolution and sensitivity, such as the High Sensitivity Array (HSA\\footnote{{\\tt http://www.nrao.edu/hsa/}}), and the Global mm-VLBI Array (GMVA\\footnote{{\\tt http://www.mpifr-bonn.mpg.de/div/vlbi/globalmm/}}).  Thanks to its proximity and brightness, Mrk 501 is an ideal laboratory for experiments using these advanced VLBI techniques: it is at $z = 0.034$ (1 mas = 0.67 pc, using $H_0 = 70$ km s$^{-1}$ Mpc$^{-1}$); the total flux density at 5\\,GHz is $S_5 = 1.4$\\,Jy; the Schwarzschild radius for its central black hole is estimated around $1 R_S = 10^{-4}$\\,pc ($1.4 \\times 10^{-4}$\\,mas), if we adopt $M_\\mathrm{BH} = 10^9\\,M_\\odot$ \\citep{rie03}.  Using these new facilities, we can therefore access regions never studied previously: the jet base with the GMVA, and the faint, resolved jet region at $>100$\\,mas with HSA.   \\begin{figure} \\resizebox{\\hsize}{!}{\\includegraphics{9784f1.eps}} \\caption{Visibility amplitude vs.\\ $(u,v)-$radius for the GMVA observations.\\label{f.uv.gmva}} \\end{figure}  In \\S \\ref{sec:observations} we describe the instruments used for our new observations, along with the data reduction methods required by high frequency. Results are presented in \\S 3 and discussed in \\S \\ref{sec:discussion}. We present our conclusions in \\S \\ref{sec:conclusions}. ",
        "conclusions": "\\label{sec:conclusions}  We have successfully explored new regions in the remarkable jet of Mrk~501. Thanks to the great sensitivity of the HSA, we reveal that the VLBI jet is one-sided (and therefore in the relativistic regime) out to at least 500 parsecs from the core. The polarization vectors are clearly aligned with the jet spine, suggesting that the magnetic field is orthogonal to the jet main axis. This is also in agreement with the results of the adiabatic fit to the jet brightness and width as a function of distance from the core. Limb brightening -- already detected on intermediate scales by VSOP observations \\citepalias{gir04a} -- is now visible on HSA transverse profiles at $\\sim 60$ mas from the core, and is likely present even on the sub-parsec scales imaged by the GMVA.  Despite its presumed weakness, Mrk~501 has in fact been clearly detected by the GMVA on sub-milliarcsecond scales. This result is encouraging given the performance of the existing mm-VLBI array and suggests that not only the brightest AGN can be studied on the smallest scales. Present and future upgrades of the array (e.g., the installation of new receivers at Plateau de Bure, inclusion of new or existing telescopes such as the 40\\,m at Yebes, Spain) are expected to make the instrument even more sensitive and reliable. The brightness temperature of the most compact component is about $6 \\times 10^{10}$\\,K. This region has a linear size of $0.020 \\times 0.014$\\,pc or, in terms of gravitational radii, $200 \\times 140 R_S$. Significant emission is also revealed in the short baselines between the most sensitive telescopes, and awaits proper imaging with increased fidelity."
    },
    "0807/0807.3783_arXiv.txt": {
        "abstract": "On 2005 January 15, the active region AR10720 produced an X1.2 solar flare that induced high levels of seismicity into the photospheric layers. The seismic source was detected using helioseismic holography and analysed in detail in Paper~I. Egression power maps at 6 mHz with a 2 mHz bandwidth revealed a compact acoustic source strongly correlated with the footpoints of the  coronal loop that hosted the flare. We present a magneto-siesmic study of this active region in order to understand, for the first time, the magnetic topological structure of a coronal field that hosts an acoustically active solar flare. The accompanying analysis attempts to answer questions such as: Can the magnetic field act as a barrier and prevent seismic waves from spreading away from the focus of the sunquake? And, what is the most {\\em efficient} magnetic structure that would facilitate the development of a strong seismic source in the photosphere? ",
        "introduction": "Our understanding of the acoustics of solar flares has been greatly improved in recent years through a combination of observational and computational techniques. It was \\citet{w1972} who first suggested that solar flares could release acoustic noise into the solar interior. \\citet{kz1995} simulated this phenomenon for the first time, and soon after, \\citet{kz1998} discovered the first seismic event, in the form of ripples, propagating away from the flare of 1996 July 9.  With the advancement of local helioseismic techniques such as helioseismic holography \\citep{lb2000}, we have now detected numerous seismic sources of varying size and intensity produced by M- and X-class flares \\citep{dbl1999,dl2005,donea2006,moradi07,betal2007,Martinez-Oliveros2007}. Extended work on this field has also been continued  by \\citet{kz2006, zharkova2007, Martinez-Oliveros2008}.  During the impulsive phase of a flare, the coronal magnetic energy is transferred down into the photosphere and further into the solar interior. This energy is then refracted back to the solar surface within approximately 50 Mm of the source and within an hour of the beginning of the flare.  The surface manifestation of this phenomenon is the appearance of ``ripples'' on the solar surface, which we identify as sunquakes. It is interesting to note that the majority of flares do not generate sunquakes. Most large flares are seismically inactive, which suggests that the strong magnetic fields of the hosting active regions may substantially alter the behaviour of helioseismic signals emerging from below. To date, the magnetism of solar seismic regions has not been studied in depth. \\citet{bdl1987} and \\citet{b1995} observed that sunspots partially absorb wave energy and shift the phase of the oscillations. A long line of theoretical developments then followed \\citep{cb1993,cbz1994,bc1997,cc2003,cc2005} which has shown that near-surface conversion to slow magnetoacoustic waves is predominantly responsible for the absorption.  Furthermore, \\citet{schunker2005}  confirmed through observations that magnetic forces should be of particular significance for acoustic signatures in penumbral regions, where the magnetic field is significantly inclined from vertical. \\citet{sh2005} also found that a sizable proportion of magnetic field variations occur in the penumbral regions of flaring sunspots. Remarkably, the majority of seismic sources induced by flares are also located either inside, or within close proximity to the penumbra. These observations possibly suggest a new mechanism which may be driving seismic waves at the photospheric level.  Indeed, \\citet{hfw2007} have recently introduced the idea of  the coupling of flare energy into a seismic wave, namely the ``McClymont magnetic jerk'', produced during the impulsive phase of acoustically active flares. They estimated the mechanical work that would be done on the photosphere by a sudden coronal restructuring. Their energy estimates are similar to those based on our helioseismic observations.  During January 11\\,--\\,20 2005, AR10720 produced 5 X-class solar flares, including an X7.1 on January 20 which produced an intense solar proton storm. However, the Michelson Doppler Imager (MDI) onboard the Solar and Heliospheric Observatory (SOHO) instrument provided helioseismic observations only for the X1.2 flare of January 15. This flare was situated at N14E08 on the solar surface. The detection of the powerful seismic transient of 2005 January 15 was initially reported by \\citet{betal2006} and \\citet{Moradi2006a,Moradi2006b}. The properties of the seismic waves generated by the event were later analysed by \\citet{kz2006}.  \\citet{moradi07} (Paper I) extensively analysed the sunquake of 2005 January 15 and compared the acoustic signatures with other supporting observations.  They also compared certain qualities exhibited by the flare with all other known acoustically active flares. The coincidence between strong compact acoustic source and nearby signatures of hard X-ray emission is remarkable.  This and the spatial coincidence of the acoustic emission with the sudden white-light signature, suggests that the sudden heating of the low photosphere results in seismic waves at the solar surface.  \\citet{moradi07} further suggested that a detailed examination of the heated magnetic photosphere is needed to complete our studies.  In this paper we will analyse the magneto-seismic activity of AR10720. Specifically, we shall investigate the role of the photospheric and coronal magnetic fields in generating the seismic waves based on the results of \\citet{hfw2007}.  We will also use vector magnetograms of AR10720 to analyse the evolution and dynamics of the photospheric magnetic field.  The structure of this paper is as follows. Section 2 outlines the observational data used for our analyses. Section 3 presents the location of the seismic source in AR10720 with references to \\citet{moradi07} for details. Section 4 outlines the line-of-sight magnetic transients in AR10720 associated with the X1.2 solar flare.  Section 5 shows the coronal magnetic field reconstruction models of AR10720.  In the final section, we conclude with a discussion of the magnetism of the seismic region based on what we have learnt from our analyses. ",
        "conclusions": "We have shown that the seismic area in AR10720 is located just beneath a complex coronal loop structure with highly twisted lines, which means that  the photospheric impact was significant in the region where the twisting allowed a maximum storage of energy. This twisting  caused the interaction between neighbouring low-lying loops, triggering the flare. This can be seen in the SOHO-EIT images (Figure~\\ref{potential}), where the flaring of the loops is visible along the field lines.  A main question we want to ask now is: how was this sunquake produced? That is, what was the mechanism for transporting the flare energy efficiently, from the reconnection site downward into the chromosphere?  Let us describe a number of possible mechanisms that could trigger a seismic source at the photospheric level during a flare and discuss the likehood that these mechanisms can explain the sunquake of the 2005 January 15 flare.  The first mechanism \\citep{kz1998} propose that seismic emission into the solar interior in sunquakes is the continuation of a chromospheric shock and condensation resulting from explosive ablation of the chromosphere and propagating downward through the photosphere into the underlying solar interior. Chromospheric shocks are well known under such circumstances, based on red-shifted H$\\alpha$ emission at the flare site at the onset of the flare.  The simulations were worked out at length by \\citet{fcm1985a,fcm1985b,fcm1985c} and others since.  The hypothesis that the photospheric emission is a direct continuation of such shocks was considered by \\citet{dl2005} and \\citet{kz2006}. For the flare studied in this paper, the hydrodynamic impact of the photosphere was clearly significant since, amazingly, this X1.2 type flare triggered a very powerful seismic source and visible seismic waves (see Paper I and \\citet{kz2006} for details). The spatial coincidence between the HXR emission and the seismic source leads us to connect the two processes, and conclude that the high-energy electrons played an important role.  However, we have to look at the statistics of acoustically active events \\citep{betal2008} and acknowledge that most solar flares do not produce sunquakes. This leads us to believe that, for the majority of flares, strong radiative damping depletes the chromospheric transient before its arrival at the low photosphere.  Therefore, we need to look for alternative mechanisms to explain the excitation of seismic sources.   In Paper~I we proposed a second mechanism. We stated that the coincidence between the locations of sudden white-light emission and seismic emission in all acoustically active flares (including the 2005 January 15) suggests that a substantial component of the seismic emission seen is a result of sudden heating of the low photosphere associated with the observed excess of visible continuum emission (radiative back-warming). The origin of white-light emission would have to be entirely in the chromosphere, where energetic electrons dissipate their energy \\citep{metcalf1990,zk1991,zk1993}, mainly by ionizing previously neutral chromospheric hydrogen approximately to the depth of the temperature minimum. It appears that the low photosphere itself would be significantly heated as well. This is primarily the result of Balmer and Paschen continuum edge recombination radiation from the overlying ionised chromospheric medium, approximately half of which we assume radiates downward and into the underlying photosphere.  \\citet{dl2005}, \\citet{donea2006}, and \\citet{moradi07} have  analysed this process in detail. \\citet{cd2006} also affirm that the white-light flare signatures highlight the importance of radiative back-warming in transporting the energy to the low photosphere when direct heating by beam electrons is impossible.  A third possible mechanism proposed by \\citet{zharkova2007}, states that high-energy protons, can directly deposit energy in the photosphere, inducing a seismic source . However, for the flare of 2005 January 15 there is no indication of high-energy protons that could directly supply the energy on which the acoustic emission depends. Likewise, energetic electrons consistent with HXR signatures, seem to be unable to penetrate into the low photosphere in anywhere near sufficient numbers required to account for the direct heating needed by the seismic sources \\citep{mcs1990}.    A fourth mechanism \\citep{hfw2007} suggested that the ``McClymont magnetic jerk'' can account for the seismic activity of some flares.  Here, we want to apply the relations of \\citet{hfw2007} for the seismic area of AR10720 in order to determine whether the ``McClymont magnetic jerk'' can account for the seismic activity of the 2005 January 15 flare.  For a line-of-sight MDI magnetic field change of 60~G, as measured in the region where the main kernel of the acoustic seismic appeared (area $\\sim 6$ Mm $ \\times \\, 9 $~Mm), the total Lorentz force, for $B_z \\sim 400$~G, is $2 \\times 10^{21}$~dyne ($\\delta f_z \\sim 2.4 \\times 10^3$~dyne/cm$^2$ $\\cdot$ $1.2 \\times 10^{18}$~cm$^2$). In Paper I we observed that the photospheric impact produced a depression of about 10 km. Using this, the maximum work done by the Lorentz force on the photosphere is estimated at $\\sim 2 \\times 10^{27}$~ergs, which is twice the energy needed by the entire seismic source to oscillate at a frequency centered at 6 mHz within a 2 mHz band.  From Paper I we extract that the seismic kernel accounted for $\\sim 45$\\% of the total egression power (estimated at $\\sim 1 \\times 10^{27}$~ergs) integrated over the region encompassing the entire flare signature (kernel plus the lenticular diffuse component).  Of course, the inferred number is just an upper limit based on many uncertainties of the local physics. We conclude that the ``McClymont magnetic jerk'' may explain the formation of the acoustic kernel, but does not explain the diffuse lenticular element of seismic activity surrounding the main kernel, which is distributed along the neutral line up to $\\sim$15~Mm east and $\\sim$30~Mm west of the kernel \\citep{moradi07}. The fact that the erratic motion of the HXR sources is observed only above the acoustic kernelled area sustains this assumption.   We  note that if integrating over the full area of the seismic source (including the diffuse lenticular acoustic emission surrounding the main kernel) the change in magnetic field is very small (about 5~G), which is understandable, because the full area of the seismic source spans negative and positive magnetic polarities.  The area is also permeated by field lines from loops of different orientation, making the local magnetic geometry much more complicated. We also emphasize that the seismic area of the solar flare of 2005 January 15 has magnetic loops of a very large inclination angle, positioning the reconnection site close to photospheric levels.  According to \\citet{h2000} and \\citet{hfw2007}, one expects that the field in the photosphere should become ``more horizontal'', as a result of the coronal magnetic field contraction that follows the decrease in the coronal magnetic energy.  Limited by the existing observations, we cannot say whether the overall field structure of AR10720 had tilted even more during the flaring.  Clearly, we cannot definitively affirm that for this complicated structure of AR10720 the Lorentz force is the main triggering mechanism for this quake (we remind the reader that the X1.2 flare of 2005 January 15 generated the most powerful solar seismic source detected so far). We believe that in reality, a combination of all the above mechanisms may be required to describe the entire phenomenon.  For simpler magnetic field configurations, where seismic sources have been also identified \\citep{dl2005, donea2006} as localised acoustic kernels at the location of moving hard X-ray footpoints, we expect the ``McClymont jerk'' mechanism to work efficiently in parallel with the chromospheric shocks driven by sudden, thick-target heating of the upper and middle chromosphere \\citep{kz1998,dl2005} and the ``back-warming'' mechanism.   \\begin{figure} \\begin{center} \\hspace*{-1cm}\\includegraphics[width=1.35\\columnwidth, angle=0]{figs/eit1600.eps} \\caption{TRACE image at 1600~\\AA taken at 00:12:35~UT. The observable feature in the image resembles the structure observed in the extrapolation.} \\label{imaeit} \\end{center} \\end{figure}"
    },
    "0807/0807.4676_arXiv.txt": {
        "abstract": "{ The South Pole Acoustic Test Setup (SPATS) has been built to evaluate the acoustic characteristics of the South Pole ice in the 10 to 100 kHz frequency range so that the feasibility and specific design of an acoustic neutrino detection array at South Pole can be evaluated. SPATS consists of three vertical strings that were deployed in the upper 400 meters of the South Pole ice cap in January 2007, using the upper part of IceCube holes. The strings form a triangular array with the longest baseline 421 meters. Each of them has 7 stages with one transmitter and one sensor module. Both are equipped with piezoelectric ceramic elements in order to produce or detect sound. Analog signals are brought to the surface on electric cables where they are digitized by a PC-based data acquisition system. The data from all three strings are collected on a master-PC in a central facility, from which they are sent to the northern hemisphere via a satellite link or locally stored on tape. A technical overview of the SPATS detector and its performance is presented.  }   \\begin{document} ",
        "introduction": " ",
        "conclusions": "The South Pole Acoustic Test Setup has been succesfully deployed in the 06/07 austral summer after extensive testing of all components. All 21 deployed acoustic stages were calibrated in water. The data are of good quality.  A fourth acoustic string is being prepared for deployment in the 07/08 austral summer. 3 SPATS-like acoustic stages will be deployed to a maximum depth of 500 m. Other acoustic devices are currently under construction; the goal is to deploy these second generation devices at the other available breakouts of the fourth string. Continued R\\&D efforts work towards a possible installation of a 100 km$^{2}$ scale acoustic array in the South Pole ice as part of a hybrid EeV neutrino detector."
    },
    "0807/0807.1862_arXiv.txt": {
        "abstract": "We present an empirical assessment of the use of broadband optical colours as age indicators for unresolved extragalactic clusters and investigate stochastic sampling effects on integrated colours.   We use the integrated properties of Galactic open clusters as models for unresolved extragalactic clusters.  The population synthesis code {\\it Starburst99} \\citep{lei99} and four optical colours were used to estimate how well we can recover the ages of 62 well-studied Galactic open clusters with published ages.   We provide a method for estimating the ages of unresolved clusters and for reliably determining the uncertainties in the age estimates.  Our results support earlier conclusions based on comparisons to synthetic clusters, namely the (U$-$B) colour is critical to the estimation of the ages of star forming regions.    We compare the observed optical colours with those obtained from {\\it Starburst99} using the published ages and get good agreement.  The scatter in the (B$-$V)$_{observed}-$(B$-$V)$_{model}$ is larger for lower luminosity clusters, perhaps due to stochastic effects. ",
        "introduction": "Thousands of luminous young star clusters have been discovered in external galaxies by the Hubble Space Telescope (HST) (e.g., \\citealp{hol92,hol96,whi93,whi95,meu95,joh99,ben02,kee03,han03,whi03,wei04}, and references therein).  Population synthesis analyses show that these clusters are sometimes very young, with ages of a few to a few tens of Myr.  Such age information can be extremely valuable, because it provides clues to cluster formation processes as well as cluster destruction mechanisms. A radial gradient in the ages of clusters within a galaxy can tell us about gas inflow driven by a bar or an interaction.  For interacting galaxies, comparison of cluster ages with dynamical models can provide information about star forming mechanisms. For example, gas compression along the Arp 107 tidal arm may have caused a gradient of ages along the arm (e.g., \\citealp{smi05a}), while in M51, a burst of cluster formation occurred at around the time of the latest passage of the companion (\\citealp{bas05,lee05}).  Cluster synthesis studies have led to the suggestion that there is a high `infant mortality' in star clusters, with many dissolving within 10 Myr (e.g., \\citealp{bas05,fal05}). Star clusters are themselves sometimes clustered into complexes with characteristic radii of $\\sim$1 kpc (e.g., \\citealp{zha01,lar04}), and the more massive clusters tend to be located  near the centre of these complexes, suggesting cluster merging.  Accurate age dating of young clusters associated with `ultraluminous X-ray sources' (ULXs) can help distinguish between stellar-mass and intermediate-mass (100$-$1000 M$_{\\sun}$) black holes for the origin of the X-ray emission (e.g., \\citealp{smi05b}).   Stellar population synthesis models have evolved dramatically in the last several  decades, since the pioneering works where data on globular clusters and individual giant stars were used to fill gaps in the evolutionary tracks to make early colour models (e.g.\\ \\citealp{tin68,tin72,str78}).  Recently, several groups have introduced evolutionary synthesis codes, e.g., \\citet{bru93} (B\\&C), \\citet{fri94} (GALEV), \\citet{fio97} (PEGASE), and \\citet{lei99} (Starburst99).  These new models vary in terms of their input physics, stellar spectral libraries and extinction laws.  In general, however, there is good agreement among these models (e.g., \\citealp{cha96,vaz05}).   Unfortunately, however, at the present time, it is not clear how accurate age estimates based on these models are for unresolved extragalactic clusters.  In most cases spectra of the targets are not available, so the amount of extinction and metal abundance are unknown.  In these cases, one typically compares some combination of UV, optical, and/or near-IR broadband colours to model cluster spectral energy distributions (SED) to simultaneously estimate age and extinction assuming some metallicity (e.g. \\citealp{pas03,han03,han07,smi08}).  There are several parameters that can affect the colour of a cluster other than age.  Reddening plays a very important role, as does the chemical composition.  However, in some age ranges and colours, reddening and chemical composition are degenerate with age. Furthermore, for low mass clusters ($\\la10^{5}$ M$_{\\odot}$), the observed integrated colours  can be affected by stochastic sampling of the initial mass function (IMF) (see for example, \\citealp{cer06,cer04,cer03}  and references therein).  For example, the random addition of a small number of high mass stars will affect the integrated colour of a cluster.  This suggests that the observed integrated colours of low mass  clusters may not be good indicators of age when compared to the integrated colours of a stellar  population model with a fully sampled IMF.   It is not clear how well a particular set of colours can predict the age of a cluster.  Are some colour sets better suited than others?  What uncertainties can be expected because of the choice of colours in the comparison?  What uncertainties can be expected because of the assumptions made in generating the model SEDs?   Several authors have investigated the use of broadband colours as age indicators by comparison of model colours to the colours of synthetic clusters (e.g. \\citealp{gil02,and04,deg05}).   We use an alternative method by comparing model colours to the integrated colours of resolved Galactic open clusters (OCs).  We compare the published integrated colours of well-studied OCs to a set of population synthesis models.  These clusters have published ages previously determined by the turn-off of the zero-age main sequence on the H-R diagram.  This work parallels that of \\citet{pes08}, who do a similar analysis of Magellanic Cloud clusters using both optical and near-IR data.   The present paper is organized as follows.  In \\S2 we describe our sample of OCs.  We describe the population synthesis models used in this study in \\S3 and the data analysis in \\S4.   Predicting the ages of the OCs, predicting the amount of extinction and the effects of metallicity, stochastic sampling effects and cluster dissolution are discussed in \\S5.  Finally, we summarize in \\S6. ",
        "conclusions": "We present an empirical assessment of the use of broadband colours as age indicators for unresolved extragalactic clusters and investigate stochastic sampling effects on integrated colours.  The population synthesis code {\\it Starburst99} \\citep{lei99} and four optical colours were used to estimate how well we can recover the ages of 62 well-studied Galactic open clusters with published ages.  We conclude the following:  1)  Galactic open clusters can be used for testing the integrated properties from population synthesis and serve as reasonable benchmarks for future assessments of age-dating methods.  2)  The (U$-$B) colour is critical.  Only colour combinations that included (U$-$B) resulted in good age and extinction predictions, consistent with previous results.  3)  Only with (U$-$B) included were the predicted age uncertainties reasonably constrained.  4)  Changes of a factor of $\\sim$2 in assumed metal abundance do not result in significantly different predictions of cluster age.  This indicates that the uncertainties in predicting the age of a cluster resulting from the age-reddening degeneracy dominate over the other sources of degeneracy in these optical bands, over our age range and metallicity range.  Another possibility is that uncertainties in the measured colours (and hence the predicted ages) dwarf the metallicity-age degeneracy.  5)  A $\\chi^2$ minimization and a $\\Delta\\chi^2$ defined to give 68\\% confidence levels provide reliable age estimates, and more importantly, reliable age uncertainties.  The difference in the photometric ages and the ages derived from the HR-diagrams of our selected cluster sample with both the (U$-$B) and (B$-$V) colours are smaller than 0.5 dex for 79\\% of the clusters, with a maximum difference of 1.5 dex.  6)  It is likely that the uncertainties in the measured colours and/or ages are large enough (0.14 mag and $\\sim20\\%$ respectively) that the effects of stochastic sampling are washed out.  A more detailed analysis of the stellar content of these resolved OCs will be required to further investigate the extent to which stochastic sampling of the IMF affects the integrated broadband colours.  7)  If stochastic sampling effects are observationally characterized in  this sample of OCs, they will provide a good bench mark for testing future population synthesis models."
    },
    "0807/0807.4730_arXiv.txt": {
        "abstract": "We model the transport of cosmic ray nuclei in the Galaxy by means of a new numerical code. Differently from previous numerical models we account for a generic spatial distribution of the diffusion coefficient. We found that in the case of radially uniform diffusion, the main secondary/primary ratios (B/C, N/O and sub-Fe/Fe) and the  modulated antiproton spectrum match consistently the available observations. Convection and re-acceleration do not seem to be required in the energy range we consider: $1 \\le E \\le 10^3~\\GeV/{\\rm nucleon}$. We generalize these results accounting for radial dependence of the diffusion coefficient, which is assumed to trace that of the cosmic ray sources. While this does not affect the prediction of secondary/primary ratios, the simulated longitude profile of the diffuse $\\gamma$-ray emission is significantly different from the uniform case and may agree with EGRET measurements without invoking {\\it ad hoc} assumptions on the galactic gas density distribution. ",
        "introduction": "Understanding the origin and propagation of Cosmic Rays (CR) in the Galaxy is an intriguing puzzle which requires the combination of many different observations over a wide range of energies. While simplified propagation models, most noticeably the {\\it leaky-box} and the modified slab models, succeeded interpreting the main features of the CR nuclear composition and energy spectra for $1 \\simleq E \\simleq 100~\\GeV/{\\rm n}$\\footnote{This unit of measure corresponds to the energy $E$ competing on average to each nucleon within a nucleus with $A$ nucleons and energy $A\\times E$.}, more realistic diffusion models are called for to provide a comprehensive description of multi-channel observations (including heavy nuclei, electrons,  $\\gamma$-rays and antimatter particles) accounting for the growing amount of available astrophysical data  (see \\cite{Strong:2007nh}  for a comprehensive review).  Two main approaches have been developed so far: analytical (or semi-analytical) diffusion models (see e.g.~\\cite{Berezinsky:book} and ref.s therein), which solve the CR transport equation by assuming simplified distributions for the sources and the  interstellar gas, and fully numerical diffusion models. Well known recent realizations of those two approaches are respectively the {\\it two-zone model} \\cite{Maurin:01,Maurin:02} and the GALPROP code \\cite{Strong:98,Strong:04,GALPROPweb}. In the case of GALPROP, the adoption of  realistic gas and radiation field distributions allows to model also the spectrum and angular distribution of the $\\gamma$-ray secondary emission.  Although these models represent a significant step forward with respect to previous analyses, they still perform a number of simplifications with respect to a more realistic physical scenario. Most noticeably, they assume diffusion to be statistically isotropic and homogenous, i.e.~they adopt the same diffusion coefficient value (or at most two values in the two-zone model) all over the propagation volume.  However, such assumptions may not always be justified, as diffusion coefficients generally depend on the regular magnetic field orientation and on the ratio between the regular and chaotic magnetic field energy densities. Although these quantities are poorly known, several observations and theoretical arguments suggest that they are far from being spatially homogeneous in the Galaxy (see sec.~\\ref{sec:diffusion} for more details). This may have relevant consequences for the CR spatial distribution in the Galaxy, for the angular distribution of the secondary $\\gamma$-ray and neutrino emissions \\cite{Evoli:2007iy} and to interpret the CR anisotropy.  In order to be able to test some of those effects, as well as to verify previous results which have been derived in the literature under more conventional conditions, we developed a new numerical code, DRAGON (Diffusion of cosmic RAys in Galaxy modelizatiON).   % DRAGON is especially designed to  account for a spatially in-homogeneous  and an-isotropic diffusion coefficient. In its present version it allow to model CR nuclei transport at energies $E_{\\rm min}\\simgeq 1~\\GeV$/n as well as the secondary $\\gamma$-ray and neutrino emission produced by their interaction with the interstellar medium (ISM). In this work we disregard CR convection and re-acceleration (i.e.~we work in a plain diffusion (PD) regime) and show that most relevant measurements can be reproduced under these conditions. Above $E_{\\rm min}$ we expect that no other physical input than source spectra, diffusion and fragmentation processes can determine secondary/primary ratios, hence a comparison of our prediction with experimental data should allow to fix the slope delta of the diffusion coefficient for some assumed slope of the CR injection spectra (see \\cite{Castellina:2005ub} for a detailed discussion about this issue).  In our analysis we will mainly refer to measurements of the secondary/primary flux ratios of several nuclear species (the most relevant are B/C, N/O and sub-Fe/Fe) and the antiproton and ${\\bar p}/p$ spectra, performed by several satellite and balloon experiments.  In order to test our code, we firstly study the conventional case of a uniform diffusion coefficient. Afterwards, we will analyze the previously unconsidered case in which $D$ grows exponentially with the distance from the Galactic Plane (GP) and  traces the radial distribution of supernova remnants (SNR).  This paper is structured as follows: in sec.~\\ref{sec:diffusion} we give several theoretical and observational motivations for considering the effects of in-homogeneous diffusion on the CR distribution. Moreover, we explain why it is worth studying perpendicular, rather than isotropic diffusion. In sec.~\\ref{sec:code} we describe our numerical model and our main assumptions on several variables entering the CR transport equations. In particular, we describe our assumed CR source and target gas distributions, and the nuclear cross section models we exploit. Then, we describe our main results on the physics of CRs, by studying secondary/primary ratios and individual species spectra in sec.~\\ref{sec:test}. We will show that the CR spatial distribution can be significantly affected by in-homogeneous diffusion and discuss how this can have relevant consequence for the solution of the CR gradient problem, as we already argued in \\cite{Evoli:2007iy}. Although to simulate a detailed map of the $\\gamma$-ray diffuse emission, suitable to comparison with the one the FERMI observatory will soon provide \\cite{GLAST}, is beyond the aims of this work, in sec.~\\ref{sec:gamma} we will model the longitudinal profile of this emission along the GP showing that under reasonable condition it matches EGRET measurements. Finally, in sec.~\\ref{sec:conclusion} we draw our conclusions. ",
        "conclusions": "\\label{sec:conclusion}  We modeled CR transport in the Galaxy assuming a plain diffusion model. We study two main cases: in the first one the diffusion coefficient $D$ is assumed to be uniform along the Galactic Plane, while in the second case we consider, for the fist time, a $D(r)$ which traces the radial profile of SNRs (which we assume to be the CR sources).  The fact that we can consistently reproduce the observed antiproton spectra and the main secondary/primary nuclear ratios for  $E \\simgeq 1~\\GeV$/n makes us quite confident of the validity of our approach. By using only the B/C, C/O and N/O data we found that the preferred range ($1\\sigma$) of  values of the slope of diffusion coefficient is 0.43-0.65. The best fit value is  $\\delta \\simeq 0.57$. This is in agreement with findings of other authors. A Kolmogorov spectrum is disfavored and re-acceleration seems to be unnecessary to interpret data above $1-2~\\GeV$/n. Forthcoming experiments like CREAM \\cite{CREAM} and TRACER \\cite{TRACER}  for what concerns nuclei and PAMELA \\cite{PAMELA} and AMS \\cite{AMS} for antiprotons may soon allow to strengthen this conclusion by improving both statistics and quality of data.  While in both cases (uniform and radially dependent $D$) we obtain substantially the same successful predictions for what concerns nuclei and antiprotons reaching the Earth, the corresponding CR primary spatial distributions in the Galaxy can be considerably different. This may have a number of interesting effects, including a possible role in the solution of the problem which  plain diffusion models face predicting a too high CR anisotropy above 100 TeV.  We focused here on the effects on the expected secondary $\\gamma$-ray diffuse emission. We showed that the longitude distribution of that emission can be significantly affected by in-homogeneous diffusion. In \\cite{Evoli:2007iy} we already noticed that the effect goes in the right direction to provide a viable solution of the CR gradient problem.  Here we provide further theoretical arguments in favour of a radially dependent diffusion coefficient and succeed reproducing EGRET observations for $4 < E < 10~\\GeV$ and $|b| < 1^\\circ$ for a choice of the relevant diffusion parameters which allow to match the B/C and the antiproton spectrum. The extension of our predictions to larger latitudes would require to implement in DRAGON electron propagation (and losses) and more detailed gas and radiation distributions which we plan to do in a forthcoming paper."
    },
    "0807/0807.1115_arXiv.txt": {
        "abstract": "\\Spitzer\\ MIPS images in the Bo\\\"otes field of the NOAO Deep Wide-Field Survey have revealed a class of extremely dust obscured galaxy (DOG) at $z\\sim2$.  The DOGs are defined by very red optical to mid-IR (observed-frame) colors, $R - [24\\; \\mu m] > 14$ mag, i.e. $f_{\\nu} (24\\; \\mu m) / f_{\\nu} (R) > 1000$.  They are Ultra-Luminous Infrared Galaxies with $L_{8-1000 \\mu m}>10^{12}-10^{14} \\;L_{\\odot}$, but typically have very faint optical (rest-frame UV) fluxes.  We imaged three DOGs with the Keck Laser Guide Star Adaptive Optics (LGSAO) system, obtaining $\\sim0.06\\arcsec$ resolution in the $K'$-band.  One system was dominated by a point source, while the other two were clearly resolved.  Of the resolved sources, one can be modeled as a exponential disk system.  The other is consistent with a de Vaucouleurs profile typical of  elliptical galaxies.   The non-parametric measures of their concentration and asymmetry, show the DOGs to be both compact and smooth.  The AO images rule out double nuclei with separations of greater than $0.1\\arcsec$ ($< 1$ kpc at $z=2$), making it unlikely that ongoing major mergers (mass ratios of 1/3 and greater) are triggering the high IR luminosities.  By contrast, high resolution images of $z\\sim2$ SCUBA sources tend to show multiple components and a higher degree of asymmetry.  We compare near-IR morphologies of the DOGs with a set of $z=1$ luminous infrared galaxies (LIRGs; $L_{IR} \\sim 10^{11} \\; L_{\\odot}$) imaged with Keck LGSAO by the Center for Adaptive Optics Treasury Survey.  The DOGs in our sample have significantly smaller effective radii, $\\sim1/4$ the size of the $z=1$ LIRGs, and tend towards higher concentrations.  The small sizes and high concentrations may help explain the globally obscured rest-frame blue-to-UV emission of the DOGs. ",
        "introduction": "The \\Spitzer\\ Space Telescope \\citep{Werner04} has revealed large numbers of luminous and ultra-luminous infrared galaxies (LIRGs with $L_{IR} \\sim 10^{11-12} \\; L_{\\odot}$, and ULIRGs with $L_{IR} > 10^{12} \\; L_{\\odot}$) in the distant universe \\citep[e.g.,][]{Bell05, LeFloch05}.  Dust heating by star formation and active galactic nuclei (AGN) within the LIRGs and ULIRGs result in large mid-infrared fluxes detectable to high redshift by the \\Spitzer\\ Multiband Imaging Photometer \\citep[MIPS;][]{Rieke04}.  Because much of the energy in these systems has been reprocessed by dust, optical (rest-frame UV) studies have underestimated their bolometric luminosities \\citep{Bell05}.  At the extreme are sources that are highly obscured at observed optical wavelengths.  Recent MIPS \\tf\\ images in the Bo\\\"otes field of the NOAO Deep Wide-Field Survey \\citep[NDWFS;][]{JannuziDey99}  have revealed a class of extremely dust obscured galaxy \\citep[DOG;][]{Houck05, Weedman06,Brand07}.  These systems were selected to have \\tf\\ flux densities in excess of 0.3 mJy, and optical to IR colors redder than $R-[24] >14$ mag, roughly a magnitude redder than Arp 220 at any redshift \\citep{Dey08}.    Spectroscopic redshift confirmations for 86 DOGs were obtained with a combination of optical spectroscopy \\citep[Keck DEIMOS and LRIS;][]{Desai07},  near-IR spectroscopy \\citep[Keck NIRSPEC;][]{Brand07}, and mid-IR spectroscopy \\citep[\\Spitzer\\ IRS;][Higdon et al., in preparation]{Houck05, Weedman06}.  The redshift distribution of these DOGs is Gaussian, with a mean of $z=1.99$, a sigma of $\\sigma_z =0.5$, and a redshift range of $z=0.8 - 3.2$ \\citep{Dey08}.   The rest-frame UV to mid-IR spectral energy distributions (SEDs) of DOGs range from power laws, rising into the mid-IR, to those that show a ``bump'' at rest wavelength $1.6 \\; \\mu m$ (observed in \\Spitzer\\ IRAC 3 - 8 $\\mu m$ bands; Desai et al. in preparation, Dey et al. 2008).  A possible explanation for these two classes is that the power law sources are AGN dominated, whereas the SEDs of the ``bump'' sources are primarily powered by star formation.  For instance,  spectroscopic $H\\alpha$ observations of 10 power law dominated, dust obscured galaxies in Bo\\\"otes showed that all 10 harbor an AGN \\citep{Brand07}.  In reality, most DOGs are likely to contain some combination of both vigorous star formation and AGN activity \\citep{Fiore08}.  While spectra and SEDs may reveal the primary power sources for the DOGs, they do not necessarily reveal the triggering mechanism.  Are the DOGs undergoing violent mergers that drive both central star formation and AGN?  Alternatively, are the DOGs isolated massive galaxies seen at the time of their first collapse?  Are there other processes contributing such as bar formation or minor mergers?  Deep, spatially resolved imaging of the DOGs should help distinguish between some of these scenarios.  For instance, recent merging activity may be revealed by multiple nuclei or tidal tails \\citep[e.g.,][]{Melbourneetal05}.  Smooth compact systems may be more readily explained by a primordial collapse model \\citep[e.g.,][]{Zirm07}. Imaging may also be useful for identifying strong central point sources suggesting AGN activity.   In order to observe the morphology of the underlying stellar population, rest-frame optical to near-IR imaging is preferable to imaging at shorter wavelengths.  It is less affected by ongoing star formation and dust obscuration compared with rest-frame UV imaging.  Rest-frame UV, however, is valuable for identifying sites of recent star formation and unobscured AGN, which may or may not track the underlying stellar distribution.    We have undertaken two efforts to obtain resolved images of the DOGs.  This paper reports on high spatial resolution ($0.06\\arcsec$) Keck laser guide star (LGS) adaptive optics (AO) $K'$-band imaging (rest-frame optical to near-IR) of three DOGs.  While this sample is small, it is the first and only set of high resolution $K$-band images available for the DOGs.  Bussman et al. (in preparation) will report on \\emph{Hubble Space Telescope}\\ (\\HST) imaging of an additional 30 DOGs.  The Bussman et al. data include ACS  and WFPC2 $I$-band imaging which probe rest-frame blue to UV, and which will be valuable for detecting the sites of ongoing star formation.  The Bussman et al. data also include NICMOS $H$-band imaging (rest-frame optical), which will be more sensitive to the underlying stellar population.  In comparison to NICMOS, the Keck AO data presented in this paper probe longer wavelengths, and have roughly a factor of 3 higher spatial resolution, providing an unprecedented view of the stellar distributions in the cores of the DOGs.     %  Section 2 describes the sample selection, the observations, and the data reduction.  Section 3 describes the morphologies of the DOGs and provides measurements of their structural parameters.  A discussion and comparison of the DOGs with two additional samples of high redshift dusty galaxies is given in Section 4.   Throughout, we report Vega magnitudes and assume a $\\Lambda$ cold-dark-matter cosmology: a flat universe with $H_0= 70$ km/s/Mpc, $\\Omega_m=0.3$, and $\\Omega_{\\Lambda}=0.7$. ",
        "conclusions": "To place the three DOGs in our sample into context, we compare their morphological parameters with two additional samples of intermediate to high redshift, dusty galaxies.  The first is a sample of of $z=2$ SCUBA sources \\citep[sub-mm sources,][]{Pope05} observed in the Great Observatories Origins Deep Survey  north field \\citep[GOODS-N;][]{Giavalisco04}.  Although SCUBA sources have a different selection criteria, they are also very dusty.  The sub-mm emission is thought to arise from cool dust in very strongly star-bursting galaxies. The second is a sample of $z\\sim1$ LIRGs from GOODS-S observed in the $K'$-band with Keck AO by the Center for Adaptive Optics Treasury Survey \\citep[CATS;][]{Melbourne08}.  These two samples where chosen because they contain dusty galaxies, and because they span the redshift range of the DOGs.  Figure \\ref{fig:cas} is a plot of asymmetry vs. concentration for the DOGs in our sample (stars).  Also plotted are the GOODS-N SCUBA sources (triangles) and the GOODS-S LIRGs (circles).  The C and A measures for the two AO PSF stars are shown as well (squares). For the SCUBA sources, we only plot C and A measurements of galaxies that were reported to have high signal-to-noise ratio \\citep{Pope05}.  The A and C measurements for SCUBA sources were made in the reddest of the GOODS \\HST\\ ACS bands, the $z$-band.  These images are  significantly bluer than our $K'$-band images of the DOGs.  However, \\citet{Pope05} suggest that there is little change in the A and C measurements through to the near-IR $H$-band (based on a small subset for which NICMOS $H$-band is available).  They propose that the high asymmetry in the SCUBA sample indicates significant recent merging.  Also plotted on Figure \\ref{fig:cas} are the A and C measurements of a set of $z\\sim1$ LIRGs from CATS \\citep{Melbourne08}.  For the LIRGs, both $K'$- band (filled circles), and $z$-band (open symbols) asymmetries are shown.  The $K'$-band asymmetries of the LIRGs  are low in comparison with the SCUBA sources.  This was anticipated because \\citet{Melbourne08} demonstrated that the morphologies of these LIRGs are primarily disk-like and unlikely to have had a recent major merger.  Interestingly, if we measure the asymmetries of the LIRGs in bluer bands, for instance the $z$-band (rest-frame optical to UV),  we find that they scatter towards higher asymmetries, similar to the SCUBA sources.   The LIRGs are more asymmetric in the bluer bands because of widespread star formation and dust obscuration, rather than merging. This suggests that interpretation of high asymmetry is difficult, especially when looking at rest-frame blue to UV images.  The two resolved DOGs in our sample have, on average, lower asymmetry than the SCUBA sources.  Their asymmetry measures are similar to the $K'$-band asymmetries of the LIRGs. This may be evidence against recent merging within the DOGs.    In order to quantify this we generate model merger systems and analyze them with GALFIT.  The models are based on the image of DOG 1.  We add two copies of the DOG 1 image together after shifting spatially and scaling to different flux ratios.  Figure \\ref{fig:model_sep20} shows model mergers generated with an $0.2\\arcsec$ separation. This separation is well below the resolution of seeing limited images and at or below the resolution of \\HST\\ NICMOS in the near-IR. As with Figure \\ref{fig:dogs}, the left column shows the merger images.  The middle column shows the best fitting single Sersic model for each merger, provided by GALFIT.  The right hand column gives the residual difference between the two.  For comparison, the top row shows the image of DOG 1 without any merger added, and a single Sersic model subtracted off.  Successive rows show mergers with decreasing flux ratios, of 1/1, 1/3, 1/5, and 1/10.  Visually, at these separations, 1/5 ratio mergers are easily detected in the residual images, but 1/10 ratio mergers are not obvious.  When we cut the separation in half to $0.1\\arcsec$, 1/5 ratio mergers become more difficult  to identify, but 1/3 ratio mergers are still relatively easy to detect.  Thus the AO images rule out 1/5 ratio mergers at $0.2\\arcsec$ separations, and 1/3 ratio mergers at $0.1\\arcsec$ separations.  If the DOGs in our sample are at $z=2$, $0.1\\arcsec$ translates to a physical separation of 0.8 kpc.  This is a strong constraint against ongoing major mergers as the trigger for the DOGs' IR energy.  Note, we can not rule out late stage mergers in which the nuclei have already merged completely.   While the LIRGs and the resolved DOGs in our sample have similar $K'$-band asymmetries, the DOGs have, on average, higher concentrations.  The DOGs also tend to be more highly concentrated than the SCUBA sources.  We will discuss possible implications for the high concentration below, but first we compare the physical sizes of the DOGs and the $z=1$ LIRGs.  In addition to measuring non-parametric morphologies, we have measured 2D galaxy profiles for the DOGs.  As we showed in the previous section, the three DOGs in our sample show a  range of morphologies from disk-like (DOG 1), to elliptical-like (DOG 2), to point-like (DOG 3).  Figure \\ref{fig:sersic} shows the best fit single Sersic parameters for the DOGs in our sample (stars).  We make the same measurements on the AO data for the $z\\sim1$ LIRGs (circles).  Apparent sizes for the two resolved DOGs are a factor of 4 - 5 smaller than for LIRGs with similar Sersic indices or $K$-band magnitudes.  We do not know the redshifts for the three DOGs in our sample, therefore we cannot directly compare physical radii. We do, however, know that DOGs are typically at $z\\sim2$ and span a range from $z\\sim1$ to $3$.  If the DOGs in our sample happen to be at the same redshift as the LIRGs ($z\\sim1$),  then their physical radii are smaller than the LIRGs by about factor of four.  Above $z=1$, apparent size remains relatively constant with redshift for a fixed physical size.  This means that even if the DOGs are at $z=3$ they still have sizes about a factor of four smaller than the LIRGs of $z=1$.  In order to test whether the small sizes observed for the DOGs are the result of an observational bias against measuring objects with extended radii, we embedded model galaxies into the DOG frames.  These galaxies were designed to have the same $K$-band fluxes as the DOGs, but with radii similar to the LIRGs.  While these model objects were lower surface brightness than the actual DOGs, we were able to accurately measure their sizes and luminosities with GALFIT to within better than 10\\%.   This argues that the DOGs really do have small effective radii. It is a further indication that the DOGs are not the same type of objects as the $z=1$ LIRGs.  So what are the DOGs?  The DOGs have very red rest-frame UV to IR luminosities, a magnitude redder than Arp 220.  This suggests that, not only do they contain significant amounts of dust, but the dust may be ubiquitous, blocking the majority of UV emitting sources.  While our sample is currently very small, it may turn out that in order to be selected as a DOG, a galaxy must have a high concentration and small physical size.  Otherwise the likelihood of observing some UV light, unobscured by dust, increases.  This hypothesis will be interesting to explore with larger samples.  Another key piece of information for explaining the nature of the DOGs may be their SEDs.  At least two of the DOGs show power-law SEDs suggestive of AGN activity.  Oddly enough those two DOGs do not show a significant central point source in the AO image.  Both of these sources, however, have more highly concentrated light profiles than the LIRGs.  This may mean that they contain a dust enshrouded AGN, increasing their concentration, even though the AGN is not obvious in the images.  Based on a stacking analysis of very deep x-ray data in the Chandra Deep Field South \\citet{Fiore08} show that DOGs in GOODS-S are likely to contain dust obscured AGN.  DOG 3 does show point-like morphology suggesting that it too may contain a strong central AGN.  Surprisingly, DOG 3 shows the least power-law like SED.  The ``bump'' in the SED of DOG 3 may be indicating that there is also strong star formation in this galaxy, and while the morphology is point like, the  C and A measures for DOG 3 suggest that it is not a pure point source galaxy.  Clearly we need to increase the sample size in order to better understand the relationship between morphology and SED type in the DOG population.  In addition to morphology, spatially resolved spectroscopy and kinematics may help place further constraints on AGN activity or mergers in the DOG sample. It may be possible to obtain such data with AO-fed integral field unit (IFU) spectrographs on large telescopes \\citep[e.g.,][]{Wright07}.  Spatially resolved, rest-frame optical spectroscopy targeting $H\\alpha$ and [N II]  can reveal the presence and location of AGN \\citep[e.g.,][]{Brand07}.  In addition, IFU spectroscopy can provide a measure of the kinematics within a galaxy.  While one of our systems has a disk-like light profile, we do not know if it is a normal rotating disk.  Spatially resolved spectroscopy might reveal that this disk-like system is actually kinematically disturbed in some way.  We have demonstrated the power of high spatial resolution AO imaging for measuring the morphologies of dust obscured galaxies beyond $z=1$.  We plan to continue imaging the DOGs with Keck AO, and hope to add kinematic information with AO-fed integral field spectrographs.  These future observations may untangle the nature of the DOGs and help explain the formation of today's massive galaxies."
    },
    "0807/0807.2446_arXiv.txt": {
        "abstract": "We present the BMW-{\\it Chandra}  source catalogue drawn from essentially all {\\it Chandra}   ACIS-I pointed observations with an exposure time in excess of 10\\,ks public as of March 2003 (136 observations). Using the wavelet detection algorithm developed by Lazzati et al.\\ (1999) and Campana et al.\\ (1999), which can characterise both point-like and extended sources, we identified 21325 sources. Among them, 16758 are serendipitous, i.e.\\ not associated with the targets of the pointings, and do not require a non-automated analysis. This makes our catalogue the largest compilation of {\\it Chandra} sources to date. The 0.5--10\\,keV absorption corrected fluxes of these sources range from $\\sim 3\\times 10^{-16}$ to $9\\times10^{-12}$ erg cm$^{-2}$ s$^{-1}$ with a median of $7\\times 10^{-15}$ erg cm$^{-2}$ s$^{-1}$. The catalogue consists of count rates and relative errors in three energy bands (total, 0.5--7\\,keV; soft, 0.5--2\\,keV; and hard, 2--7\\,keV), and source positions relative to the highest signal-to-noise detection among the three bands. The wavelet algorithm also provides an estimate of the extension of the source. We include information drawn from the headers of the original files, as well, and extracted source counts in four additional energy bands, SB1 (0.5--1\\,keV), SB2 (1--2\\,keV), HB1 (2--4\\,keV), and HB2 (4--7\\,keV). We computed the sky coverage for the full catalogue and for a subset at high Galactic latitude ($\\mid b \\mid \\, > 20^{\\circ}$). The complete catalogue provides a sky coverage in the soft band (0.5--2\\,keV, S/N =3) of $\\sim 8$ deg$^2$ at a limiting flux of $\\sim 10^{-13}$  erg cm$^{-2}$ s$^{-1}$, and $\\sim 2$ deg$^2$ at a limiting flux of $\\sim 10^{-15}$ erg cm$^{-2}$ s$^{-1}$. Furthermore, we present the results of the cross-match with existing catalogues at different wavelengths (FIRST, IRAS, 2MASS, GSC2, and ChaMP). The total numbers of matches with the FIRST, IRASPSC, 2MASS, and GSC2 catalogues obtained after a closest-distance selection are 13, 87, 6700, and 4485, respectively. ",
        "introduction": "The use of a wavelet transform (WT) as an X-ray detection algorithm was pioneered by Rosati et al.\\ (1995; 1998) for the detection of extended sources in the Roentgen Satellite ({\\it ROSAT\\/}) Position Sensitive Proportional Counter (PSPC) fields and subsequently adopted by many groups (\\citealt{Grebenevea95,Damianiea97b,Pislarea97,Vikhlininea98,Lazzatiea98,Freemanea02}). Differently from other WT-based algorithms, the Brera Multi-scale Wavelet (BMW, \\citealt{Lazzatiea99,Campanaea99}) algorithm, which was developed to analyse {\\it ROSAT} High Resolution Imager (HRI) data, automatically characterises each source through a multi-source $\\chi^2$ fitting with respect to a Gaussian model in the wavelet space. For these reasons it has also proven to perform well in crowded fields and in conditions of very low background (\\citealt{Lazzatiea99}).  The BMW was used to produce the BMW-HRI catalogue (\\citealt{Lazzatiea99,Campanaea99,Panzeraea03}), a catalogue of $\\sim 29\\,000$ sources detected with a probability of $\\ga 4.2\\sigma$ and with a sky coverage of 732 deg$^2$ down to a limiting flux of $10^{-12}$ erg cm$^{-2}$ s$^{-1}$ and 10 deg$^2$ down to a limiting flux of $10^{-14}$ erg cm$^{-2}$ s$^{-1}$.  The BMW-HRI  is being  currently used for  a number of  scientific projects. Among  preliminary  results,  an  analysis of  X-ray  detected  sources without obvious counterparts at other wavelengths, (a.k.a. blank field  sources, \\citealt{Chieregatoea05}) has  been  carried out with  the   aim  of  identifying   unusual  objects,  as  well   as  a serendipitous  X-ray  survey  of  Lyon-Meudon  Extragalactic  Database (LEDA) normal galaxies (\\citealt{Tajerea05}).    The BMW algorithm was modified to support the analysis of {\\it Chandra} Advanced CCD Imaging Spectrometer (ACIS) images (\\citealt{Morettiea02}), and has led to interesting results on the nature of the cosmic X-ray background (CXB, \\citealt{Campanaea01}; \\citealt{Morettiea03}). Given the reliability and versatility of the BMW algorithm, we decided to apply it to a large sample of {\\it Chandra} ACIS-I images, to take full advantage of the superb spatial resolution [$\\sim 0\\farcs5$ point-spread function (PSF) on-axis] of {\\it Chandra} while being able to automatically analyse crowded fields and/or with very low background. We thus produced the Brera Multi-scale Wavelet {\\it Chandra} source catalogue (BMW-C).  In this  paper we present a pre-release of the BMW-C, which is based on a subset of the whole  {\\it Chandra} ACIS observations dataset, roughly  corresponding to  the first  three years  of  operations.  Like the BMW-HRI, our catalogue provides source positions, count rates, extensions and relative errors. In addition, for all bright (100 counts in the 0.5--7\\,keV band) sources in the catalogue we extracted light curves which will be exploited in a search for periodic and non-periodic variability, as well as spectra to have an immediate spectral classification.  In Sect.~\\ref{bmwc:datasample} we describe the selection criteria of the {\\it Chandra} fields and the resulting data sample. In Sect.~\\ref{bmwc:dataproc} we describe the data processing, which includes data screening and correction, event list selections and energy band selections. In Sect.~\\ref{bmwc:detect} we describe the wavelet detection algorithm  and its application to the {\\it Chandra} fields as well as the resulting source catalogue, the BMW-{\\it Chandra} catalogue. In Sect.~\\ref{bmwc:character} we report the properties of the source sample, including their spatial extent fluxes, and the definition of our serendipitous sub-sample. In Sect.~\\ref{bmwc:skycov_how} we calculate the sky coverage of our survey. In Sect.~\\ref{bmwc:champcat} we compare our catalogue with the {\\it Chandra} Multiwavelength Project (ChaMP, \\citealt[][and references therein]{Champ3}) data. In Sect.~\\ref{bmwc:xcorr} we describe preliminary results of the cross-matches between the BMW-{\\it Chandra} and other catalogues at different wavelengths. Finally, in Sect.~\\ref{bmwc:summary} we summarise our work and highlight the plan for future exploitation of the catalogue. ",
        "conclusions": "} %  We presented the BMW-C source catalogue drawn from 136 {\\it Chandra} ACIS-I pointed observations with an exposure time in excess of 10\\,ks public as of March 2003. The full catalogue comprises 21325 sources, 16834 of which do not require a more in-depth, non-automated analysis (i.e.\\ not associated with bright and/or extended sources), including the pointed ones. Among them, 16758 are serendipitous. This makes our catalogue the largest compilation of {\\it Chandra} sources to date. The 0.5--10\\,keV absorption corrected fluxes of these sources range from $\\sim 3\\times 10^{-16}$ to $9\\times10^{-12}$ erg cm$^{-2}$ s$^{-1}$ with a median of $7\\times 10^{-15}$ erg cm$^{-2}$ s$^{-1}$.  The catalogue includes count rates and relative errors in three energy bands (total, 0.5--7\\,keV; soft, 0.5--2\\,keV; and hard band, 2--7\\,keV; the source positions are relative to the highest signal-to-noise detection among the three bands), as well as source counts extracted in four additional energy bands, SB1 (0.5--1\\,keV), SB2 (1.0--2.0\\,keV), HB1 (2.0--4.0\\,keV) and HB2 (4.0--7.0\\,keV), and information drawn from the headers of the original files. The WT algorithm also provides an estimate of the extension of the source which we refined with a $\\sigma$-clipping method.   We computed the sky coverage for the full catalogue and for a subset at high Galactic latitude ($\\mid b \\mid \\, > 20^{\\circ}$). The complete catalogue provides a sky coverage in the soft band (0.5--2\\,keV, S/N =3) of $\\sim 8$ deg$^2$ at a limiting flux of $\\sim 10^{-13}$  erg cm$^{-2}$ s$^{-1}$, and $\\sim 2$ deg$^2$ at a limiting flux of $\\sim 10^{-15}$ erg cm$^{-2}$ s$^{-1}$.  We also presented the results of the cross-match with existing catalogues at different wavelengths (FIRST, IRAS, 2MASS, GSC2 and ChaMP).  Among the scientific applications of the catalogue are: 1) search for periodic and non-periodic variability using light curves extracted for bright sources; 2) the optical/IR follow-up of a list of galaxy cluster candidates drawn from our sub-sample of $\\sim 300$ extended sources; 3) analysis of blank fields, i.e.\\ X-ray detected sources without counterparts at other wavelengths; 4) optical/IR follow-up of peculiar sources, such as isolated neutron stars candidates (ultra-soft sources) and  heavily absorbed sources (ultra-hard sources, not observed in the soft bands).  The current version of the BMW-C source catalogue, (as well as additional information and data) is available at the Brera Observatory site: {\\tt http://www.brera.inaf.it/BMC/bmc\\_home.html}."
    },
    "0807/0807.5099_arXiv.txt": {
        "abstract": "{} {\\iras\\ is a young massive star forming region with an intense outflow activity. Outflows are frequently studied in the near-infrared (NIR) since the \\Hm\\ emission in this wavelength range often traces the shocked molecular gas. However, the mechanisms behind the \\Hm\\ emission detected in \\iras\\ have not been clarified yet. We present here spatially resolved NIR spectroscopy which allows us to verify whether the \\Hm\\ emission originates from thermal emission in shock fronts or from fluorescence excitation by non-ionizing UV photons. Moreover, NIR spectroscopy also offers the possibility of studying the characteristics of the putative driving source(s) of the \\Hm\\ emission by the detection of photospheric and circumstellar spectral features, and of the environmental conditions (e.g. extinction).} {We obtained long-slit, intermediate-resolution, NIR spectra of \\iras\\ using LIRIS, the NIR imager/spectrographer mounted on the 4.2\\,m William Herschel Telescope. As a complement, we also obtained $J$, $H$ and $K_s$ images with the Las Campanas 2.5\\,m Du Pont Telescope, and archival mid-infrared (MIR) \\spitzer\\ images at 3.6, 4.5, 5.8 and 8.0\\,\\um.} {We confirm the shocked nature of the \\Hm\\ emission, with an excitation temperature of about 2000~K, based on the analysis of relevant \\Hm\\ line ratios, ortho-to-para ratios and excitation diagrams. We have also identified objects with very different properties and evolutionary stages in \\iras. The most massive source at millimeter wavelengths, mm1, with a mass of a few tens of solar masses, has a bright NIR (and MIR) counterpart.  This suggests that emission --probably coming through a cavity created by one of the outflows present in the region, or from the outflow cavity itself-- is leaking at these wavelengths. The second most massive millimeter source, mm2, is only detected at $\\lambda \\gtrsim 6$\\,\\um, suggesting that it could be a high-mass protostar still in its main accretion phase. The NIR spectra of some neighboring sources show CO first-overtone bandhead emission which is associated with neutral material located in the inner regions of the circumstellar environment of YSOs. } {} ",
        "introduction": "\\iras\\ is a young massive star forming region which has been extensively studied in a series of papers during the last few years. It is considered as a prime example of the complexity of high-mass star formation due to the intense outflow activity ongoing there. The 1.2 mm dust continuum map of \\iras\\ obtained with the IRAM 30\\,m telescope at low spatial resolution reveals two massive adjacent star forming clumps roughly aligned in a north-south direction and separated by $\\sim$53\\arcsec\\ \\citep[][]{beuther02c}. The southern clump, with a mass of about 840~$M_\\odot$, is more massive than the northern clump, with a mass of 190~$M_\\odot$. Observations in CO(2--1) with 11\\arcsec\\ angular resolution show that each clump is associated with a bipolar molecular outflow in the east-west direction \\citep{beuther02d}. The southern clump is resolved into at least four cores in the 2.6\\,mm continuum using the Plateau de Bure Interferometer, whereas the northern clump is resolved into two cores \\citep{beuther03}. These authors disentangle at least seven (possibly even nine) bipolar outflows in this region with accretion rates of the order of $10^{-4} M_\\odot$\\,yr$^{-1}$, sufficiently high to overcome the radiation pressure and form massive stars via continuous disk-accretion processes. Higher spatial resolution observations at 1.3\\,mm, also with the Plateau de Bure Interferometer, show that these two clumps split up into a even larger number of sub-structures, with at least 12 cores in each one \\citep{beuther04}. These cores have masses ranging between 1.7 and 25\\,M$_\\odot$, and visual extinctions of up to 1000 magnitudes.  At the very center of the southern clump and coinciding with the most prominent of the mm sources, \\cite{beuther02b} detected H$_2$O and \\ion{Class}{ii} CH$_3$OH masers, which are considered to be closely associated with the earliest stages of massive star formation. Only at this position is there evidence for a weak 1~mJy unresoved (at angular resolution of 0\\farcs7) cm continuum source \\citep{sridharan02}.  \\cite{beuther02} reported the detection of hard X-ray emission from a number of point sources in the vicinity of \\iras. They conclude that this X-ray emission is due to intermediate-mass pre-main-sequence Herbig Ae/Be stars or their precursors.  So far, only the kinematic distance to this object is known and there exists a distance ambiguity, with a near distance of 2.1~kpc and a far distance of 6.4~kpc \\citep{sridharan02}. However, authors like \\cite{beuther02} consider this source to be located at its near kinematic distance because the derived outflow parameters are unreasonably high if the far kinematic distance is assumed \\citep[see also][]{beuther02d}. Other arguments that strengthen the selection of the near kinematic distance are the following. (1) The CS(2--1) profile for \\iras\\ \\citep{bronfman96} is fairly narrow (with a FWHM of 3.1~km/s) as compared with other sources that also have a complex outflow structure, for instance IRAS\\,18507$+$0121, with a linewidth of 6.8~km/s at a distance of 3.5~kpc \\citep{shepherd07}. At a distance of 2.1~kpc, the 40\\arcsec\\ beam of these observations is probably detecting only the mm clump and not the lobes observed in CO(2--1). At the far distance of 6.4~kpc, however, the red and blue lobes of the outflow would have been detected as wings in the CS(2--1) profile. (2) \\cite{bronfman96} observed four IRAS point sources with far-infrared colors typical of ultracompact \\HII\\ regions, including \\iras, which have very similar (LSR) velocities. These sources are within a galactic extent of about 1.7\\degr, or about 60~pc at a distance of 2.1~kpc. This is the typical size of a giant molecular cloud (GMC), and it is likely that these four sources belong to the same GMC. At the far distance of 6.4~kpc, on the contrary, there would have to be a series of GMCs with the same velocity, or a single GMC with an extent of 180~pc, which is unlikely. (3)  \\cite{watson03} observed \\iras\\ in the H110$\\alpha$ and H$_2$CO lines in order to solve the distance ambiguity. No emission was detected in H110$\\alpha$, but they detected H$_2$CO in absorption with a (LSR) velocity of 22.6~km/s and a linewidth of 4.7~km/s. No absorption is detected at larger velocities.  We thus assume a distance of 2.1~kpc. At this distance, the infrared luminosity of \\iras\\ is $10^4$~L$_\\odot$.  Outflows are frequently studied in the near-infrared (NIR), since the \\Hm\\ emission in this range often traces the hot \\mbox{($<2000$~K)} shocked molecular gas. In both low-luminosity and high-luminosity outflows, the \\Hm\\ emission delineates single or multiple bow-shaped features associated with the leading edge of the CO emission and, in some cases, also traces collimated jets and cavity walls \\cite[e.g. see the review by][]{richer00}. \\cite{beuther03} mapped the \\Hm\\ emission around \\iras\\ with the 3.5\\,m telescope on Calar Alto and used this information (combined with high-spatial resolution CO observations) in their interpretation of the outflows occurring in this region. However, the excitation mechanism behind this \\Hm\\ emission has not yet been determined. With this aim in mind, we obtained long-slit, intermediate-resolution, NIR spectra of \\iras\\ using LIRIS, the NIR imager/spectrographer mounted on the 4.2\\,m William Herschel Telescope (WHT). These observations help us verify whether the \\Hm\\ emission is produced by thermal emission in shock fronts or by fluorescence excitation by non-ionizing UV photons. These mechanisms can be distinguished since they preferentially populate different levels producing different \\Hm\\ line ratios. NIR spectroscopy also offers the possibility of studying the characteristics of the putative driving source(s) of the \\Hm\\ emission by the detection of photospheric and circumstellar spectral features. Moreover, it allows us to investigate environmental conditions (e.g. extinction) towards \\iras.  As a complement, we also obtained $J$, $H$ and $K_s$ images with the Las Campanas 2.5\\,m Du Pont Telescope. Combined with archival \\spitzer\\ data at 3.6, 4.5, 5.8 and 8.0~\\um, obtained by the GLIMPSE Legacy Science Program \\citep{benjamin03}, these images are helpful to detect optically obscured, deeply embedded young stars and protostellar candidates, excess emission from their circumstellar matter and, at earlier stages, from their infalling envelopes.  This paper is structured as follows. Section~\\ref{sect:data} describes the observations and data reduction, and Sect.~\\ref{sect:results} presents the photometric and spectroscopic results. Finally, Sect.~\\ref{sect:conclusions} presents the discussion and main conclusions. ",
        "conclusions": ""
    },
    "0807/0807.5050_arXiv.txt": {
        "abstract": "Metal-free stars are assumed to evolve at constant mass because of the very low stellar winds. This leads to large CO-core mass at the end of the evolution, so primordial stars with an initial mass between 25 and 85 $M_{\\odot}$ are expected to end as direct black holes, the explosion energy being too weak to remove the full envelope.  We show that when rotation enters into play, some mass is lost because the stars are prone to reach the critical velocity during the main sequence evolution. Contrarily to what happens in the case of very low- but non zero-metallicity stars, the enrichment of the envelope by rotational mixing is very small and the total mass lost remains modest. The compactness of the primordial stars lead to a very inefficient transport of the angular momentum inside the star, so the profile of $\\Omega(r)$ is close to $\\Omega\\,r^2=$ const. As the core contracts, the rotation rate increases, and the star ends its life with a fast spinning core. Such a configuration has been shown to modify substantially the dynamics of the explosion. Where one expected a weak explosion or none at all, rotation might boost the explosion energy and drive a robust supernova. This will have important consequences in the way primordial stars enriched the early Universe. ",
        "introduction": "Population III (Pop III) stars occupy a key position in the history of the Universe, as a link between the pure H-He Universe at the beginning of times, and the metal-rich one we observe nowadays around us. Their evolution is important, because during their life they form the first heavy elements, but their death is also important, because this is the moment at which the newly synthesised elements are released in the surrounding medium, starting the chemical enrichment of the Universe.  While the first studies on the star formation at $Z=0$ concluded that only very massive objects could form \\citep{bcl02,abn02}, later findings show that the radiative feedback of the very first stars could lead to a second generation of metal-free stars with lower initial mass \\citep{omyosh03,oshabwn05,greifbr06}. Could these stars contribute to the chemical enrichment of the Universe? The consensual picture is that most of them are directly swallowed by a black hole. ",
        "conclusions": "The lack of metals affects strongly the stellar evolution at $Z=0$. One of the main result is to reduce dramatically the radiative mass loss. This means that rotating stars will keep all their angular momentum until the end of their evolution. At the pre-supernova stage, they will present a fast-rotating iron core, so we expect the collapse to be highly aspherical. This may boost the explosion energy and allow a strong supernova explosion where only a weak one or none was expected.  Note that the inclusion of magnetic fields could alter these results by providing the core-envelope coupling which is lacking in Pop III stars. However, in that case, a strong mass loss would be expected \\citep{kauai08}, so the evolution of the star would anyway be strongly affected with regard to the consensual picture."
    },
    "0807/0807.2670_arXiv.txt": {
        "abstract": "{The recent discovery of a group of trans-neptunian objects (TNOs) in a narrow region of the orbital parameter space and with surfaces composed of almost pure water ice, being 2003 EL$_{61}$  its largest member, promises new and interesting results about the formation and evolution of the TNb and the outer Solar System.} {The aim of this paper is to obtain information of the surface properties of two members of this group ((24835) 1995 SM$_{55}$, (120178) 2003 OP$_{32}$) and three potential members (2003 UZ$_{117}$, (120347) 2004 SB$_{60}$ and 2005 UQ$_{513}$) and to use that in order to confirm or reject their association.} {We obtained visible spectra of five TNOs using the 3.58m Telescopio Nazionale Galileo at the ``Roque de los Muchachos Observatory'' (La Palma, Spain)} {The spectra of the five TNOs are featureless within the uncertainties and with colors from slightly blue to red ($-2 < S' < 18$\\%$/0.1\\mu$m). No signatures of any absorption are found.} {We confirm the association of 1995 SM$_{55}$ and 2003 OP$_{32}$ with the group of 2003 EL$_{61}$ as their spectra are almost identical to that of 2003 EL$_{61}$. Only one of the three candidates, 2003 UZ$_{117}$,  can be considered as a possible member of  the EL$_{61}$-group, as its visible spectrum is compatible with a spectrum of a surface composed of almost pure water ice and no complex organics. The other two, 2004 SB$_{60}$ and 2005 UQ$_{513}$ are red and must be considered as interlopers.} ",
        "introduction": "TNO 2003 EL$_{61}$  is the largest member of a group of TNOs (hereafter EL$_{61}$-group) with orbits in a narrow region of the orbital parameter space (41.6 $<$a$<$ 43.6 AU, 25.8 $<$i$<$ 28.2 deg., 0.10 $<$e$<$ 0.19) (Pinilla-Alonso et al., \\cite{Pinilla-AlonsoRR43}) and with surfaces composed of almost pure water ice what means: (a) The visible is featureless within the S/N. There is no clear evidence of any absorption reported for other TNOs; (b) The visible is neutral ; (c) It presents two deep absorption bands centered at 1.5 and 2.0 $\\mu$m, indicative of water ice and compatible with the presence of the crystalline phase. Brown et al. (\\cite{Brownfamily}) claim that this group is a family of fragments product of a giant collision happened in the trans-neptunian belt (TNb).  The first identified members of the group are: (136108) 2003 EL$_{61}$, (55636) 2002 TX$_{300}$, (145453) 2005 RR$_{43}$, (120178) 2003 OP$_{32}$, (19308) 1996 TO$_{66}$, and (24835) 1995 SM$_{55}$. The identification of the members of this group is possible thanks to the study of their spectra as all of them show the same characteristics. Near infrared spectra have been published for all these objects and they show very deep water ice absorption bands (Pinilla-Alonso et al. \\cite{Pinilla-AlonsoRR43} and references therein). Visible spectra have been published only for 2003 EL$_{61}$, 2002 TX$_{300}$, and 2005 RR$_{43}$ (Tegler et al. \\cite{TegEL61}, Licandro et al. \\cite{LicTX}, Pinilla-Alonso et al., \\cite{Pinilla-AlonsoRR43}) and are all similar, featureless with an almost neutral slope.  Spectroscopy is the best way to study the surface of these objects but for those that are too faint, photometry is a usefull tool. Visible colors of objects 2003 OP$_{32}$, 1995 SM$_{55}$ and 1996 TO$_{66}$ are indicative of a neutral spectral slope in this spectral region (MBOSS database http://www.sc.eso.org/~ohainaut/MBOSS/). Another object, (86047) 1999 OY$_{3}$ has been identified as a member of the group as its visible and near-infrared colors are compatible with the characteristics of a typical spectrum of a member of the group (% Ragozzine et al. \\cite{Ragozzine} and Hainaut \\& Delsanti \\cite{MBOSS}).\\\\  Fitting scattering models to the visible and near-infrared spectra of 2005 RR$_{43}$ and 2003 EL$_{61}$ Pinilla-Alonso et al. (\\cite{Pinilla-AlonsoRR43} and \\cite{Pinilla-AlonsoEL61}) show that their surfaces are composed of a mixture of amorphous and crystalline water ice and discard the presence of a significant amount of other components. In particular they discard the presence of complex organics typically red in color produced by the irradiation of hydrocarbons and/or alcohols. Furthermore, they conclude that 2003 EL$_{61}$ (and probably the other members of the group) has a significant smaller fraction of carbon chains on its surface than the other TNOs. This carbon-depleted population of TNOs is located in an unstable region of the TNb crossed by resonances and could be the source of some of the carbon-depleted comets already noticed by A'Hearn et al. (\\cite{AHearn95}). This possibility makes this group particularly intriguing.  Pinilla-Alonso et al. (\\cite{Pinilla-AlonsoEL61}) also show that, if the surface of 2003 EL$_{61}$ is the product of a large collision, the relative percentage of amorphous and crystalline ice obtained from the models, implies that this collision should have happened more than 10$^{8}$ years ago, in agreement with dynamical predictions (Ragozzine \\& Brown \\cite{Ragozzine}) and laboratory experiments (Zheng et al. \\cite{ZenghWater}).  Assuming that this population is the product of a giant collision, Ragozzine \\& Brown (\\cite{Ragozzine}) applied techniques developed for the study of the asteroid families to the study of the objects in the vicinity of 2003 EL$_{61}$. They compute the minimum ejection velocity of potential fragments required to change their orbital parameters to their actual values ($\\Delta$v$_{min}$). In a second approximation they also consider diffusion in resonances ($\\delta$v$_{min}$, see  Ragozzine \\& Brown \\cite{Ragozzine} for details). They find that all the objects in the group can be explained by a velocity dispersion of 150 m s$^{-1}$ from a single collision location and considering diffusion in eccentricity in resonances. From these dynamical considerations, they even identify a group of TNOs that are potential candidates to be part of the group.  The study of the surface properties of these objects can be used to identify other possible members of the EL$_{61}$-group and to identify other TNOs that, having similar dynamical properties, are not members, this is important to our understanding of the origin of EL$_{61}$-group.  In this paper, we present visible spectra of three of the candidates in Ragozzine \\& Brown (\\cite{Ragozzine}), 2003 UZ$_{117}$, (120347) 2004 SB$_{60}$ and 2005 UQ$_{513}$, and of two of the already known members of the EL$_{61}$ group without published visible spectroscopy, (24835) 1995 SM$_{55}$, (120178) 2003 OP$_{32}$. ",
        "conclusions": "We present visible spectra of 5 TNOs in the neighborhood of 2003 EL$_{61}$ obtained with the 3.58m Telescopio Nazionale Galileo (La Palma, Spain). Two of them are members of the EL$_{61}$-group (1995 SM$_{55}$, 2003 OP$_{32}$) and do not have previous published visible spectroscopy. The other three are potential group members (2003 UZ$_{117}$,2004 SB$_{60}$ and 2005 UQ$_{513}$) according to Ragozzine \\& Brown (\\cite{Ragozzine}).  The spectra of the five TNOs are featureless within the uncertainties, and with colors from slightly blue to red ($-2 < S' < 18$\\%$/0.1\\mu$m). No signatures of any absorption band are found.  Only one of the three candidates, 2003 UZ$_{117}$, can be considered as a possible member of the EL$_{61}$-group considering that its visible spectrum is compatible with a spectrum of a surface composed of almost pure water ice and no significant amount of complex organics, but near-infrared photometry or spectroscopy is needed to confirm the presence of water ice. The other two, 2004 SB$_{60}$. and 2005 UQ$_{513}$  are red ($S'$= 13 and 18\\%$/0.1\\mu$m respectively) thus they cannot be members of the EL$_{61}$-group. The surface of these objects probably contains a significative fraction of complex organics."
    },
    "0807/0807.4260_arXiv.txt": {
        "abstract": "{ Several studies have proposed that the preferred model of the comoving spatial 3-hypersurface of the Universe may be a Poincar\\'e dodecahedral space (PDS) rather than a simply connected, infinite, flat space. } { Here, we aim to improve the surface of last scattering (SLS) optimal cross-correlation method and apply this to observational data and simulations. } { For a given ``generalised'' PDS orientation, we analytically derive the formulae required to exclude points on the sky that cannot be members of close SLS-SLS cross-pairs. These enable more efficient pair selection without sacrificing the uniformity of the underlying selection process. For a sufficiently small matched circle size $\\alpha$ and a fixed number of randomly placed points selected for a cross-correlation estimate, the calculation time is decreased and the number of pairs per separation bin is increased. Using this faster method, and including the smallest separation bin when testing correlations, (i) we recalculate Monte Carlo Markov Chains (MCMC) on the five-year Wilkinson Microwave Anisotropy Probe (WMAP) data; and (ii) we seek PDS solutions in a small number of Gaussian random fluctuation (GRF) simulations in order to further explore the statistical significance of the PDS hypothesis. } { For $5\\ddeg < \\alpha < 60\\ddeg$, a calculation speed-up of 3--10 is obtained. (i) The best estimates of the PDS parameters for the five-year WMAP data are similar to those for the three-year data. (ii) Comparison of the optimal solutions found by the MCMC chains in the observational map to those found in the simulated maps yields a slightly stronger rejection of the simply connected model using $\\alpha$ rather than the twist angle $\\phi$. The best estimate of $\\alpha$ implies that, {{\\em given a large-scale auto-correlation as weak as that observed,}} the PDS-like cross-correlation signal in the WMAP data is expected with a probability of less than about 10\\%. The expected distribution of $\\phi$ from the GRF simulations is not uniform on $[-\\pi,\\pi]$. } {Using this faster algorithm, we find that the previous PDS parameter estimates are stable to the update to five-year WMAP data. Moreover, for an infinite, flat, cosmic concordance model with Gaussian random fluctuations, the chance of finding {{\\em both}} (a) a large-scale auto-correlation as weak as observed and (b) a PDS-like signal similar to what is observed is less than about 0.015\\% to 1.25\\%.  } ",
        "introduction": "\\label{s-intro}  During the past half-decade, attention has focussed on the Poincar\\'e Dodecahedral Space (PDS) as a potentially better model of comoving space than the infinite flat model \\nocite{LumNat03,RLCMB04,Aurich2005a,Aurich2005b,Gundermann2005,KeyCSS06,NJ07,Caillerie07,LR08,RBSG08}({Luminet} {et~al.} 2003; {Roukema} {et~al.} 2004; {Aurich} {et~al.} 2005a, 2005b; {Gundermann} 2005; {Key} {et~al.} 2007; {Niarchou} \\& {Jaffe} 2007; {Caillerie} {et~al.} 2007; {Lew} \\& {Roukema} 2008; {Roukema} {et~al.} 2008). If the hypothesis that the comoving spatial section of the Universe is a PDS is correct, then it should be possible to estimate the astronomical coordinates of the fundamental domain, as has been tried by \\nocite{RBSG08}{Roukema} {et~al.} (2008) (hereafter, RBSG08). Moreover, successive improvements in data and analysis methods should yield successively more precise and more accurate estimates of these astronomical coordinates.   Here, we aim to improve on the method of optimising the cross-correlation $\\ximc$ \\postrefereechanges{[Eq.~(1), RBSG08]} of cosmic microwave background (CMB) temperature fluctuations between copies of the surface of last scattering (SLS) presented in \\nocite{RBSG08}{Roukema} {et~al.} (2008) and also applied by \\nocite{Aurich08a}{Aurich} (2008). We investigate whether or not this can lead to improved parameter estimates. We also test a small number of simulations to see whether the infinite flat model can reproduce the observational signal.  In RBSG08, the method was introduced by pointing out that algebraically, it is an extension of using the identified circles principle first published by \\nocite{Corn96,Corn98b}{Cornish} {et~al.} (1996, 1998). Here, we make use of the fact that the cross-correlation method is not only an algebraic extension of the identified circles principle, but its relation to the identified circles principle can also be interpreted in terms of identified {\\em annuli}. This leads to algebraic/trigonometric relations that enable a faster calculation of cross-correlations, without sacrificing the uniformity of the underlying random selection of points on the sky. Moreover, use of these relations also increases the numbers of pairs per separation bin, leading to less relative Poisson noise in individual cross-correlation estimates. The annulus interpretation and the associated relations are presented in \\SSS\\ref{s-method-anglimits}.  In \\SSS\\ref{s-method-fiveyr}, use of this faster method to recalculate Monte Carlo Markov Chains (MCMC) for the five-year release of the WMAP data is described. Since the new method increases the numbers of pairs per separation bin, the smallest separation bin available when estimating correlations should not be as noisy as without this new method, so we include this smallest separation bin.  Use of this faster method also makes it practical to perform a consistent analysis of both the WMAP map and a small number of Gaussian random fluctuation (GRF) WMAP simulated maps, in order to estimate the expectation of a PDS-like signal under the assumption of a simply connected model. In order to model this, we need to consider the following.  One of the key motivations from WMAP data for studying multiply connected models has been the lack of structure on angular scales larger than about 60{\\ddeg}, as has been noted by several authors \\nocite{WMAPSpergel,LumNat03,Aurich2005b}(e.g.  {Spergel} {et~al.} 2003; {Luminet} {et~al.} 2003; {Aurich} {et~al.} 2005b). As was noted in RBSG08, in particular in Eq.~(10), if the PDS model is incorrect, then cross-correlations of temperature fluctuations at pairs of points that are implied (by the PDS hypothesis) to be spatially close should, on average, be much smaller than the auto-correlations at the same small spatial separation scale. This is because if the PDS model is wrong, then these pairs are in reality widely separated rather than closely separated, so their correlations should, on average, be small.  However, this is only a statistical statement: even for the case that the PDS model is physically wrong, it is possible that some orientations of a PDS model, by chance, happen to give cross-correlations that violate Eq.~(10) of RBSG08. How probable is it for chance large-scale correlations in particular directions to mimic a PDS-like signal? This clearly depends on the amplitude of the large length-scale auto-correlation. If this amplitude is low or high, then the chance of finding PDS-like signals using the optimal cross-correlation method should be low or high respectively. Since the range of matched circle sizes $\\alpha$ studied in RBSG08 included circles with $ \\alpha \\le 60\\ddeg$, most of the cross-correlations are from pairs of points separated observationally by angles $\\gtapprox 60\\ddeg \\sim 1\\; \\mbox{rad}$.  The latter is the minimum scale above which \\nocite{WMAPSpergel}{Spergel} {et~al.} (2003) quantified the lack of temperature-temperature auto-correlations at large scales in the WMAP data, with their parameter $S$, \\postrefereechanges{defined \\begin{equation} S = \\int^{1/2}_{-1} [C(\\theta)]^2 \\mathrm{d} \\cos \\theta, \\label{e-Sxi-spergel} \\end{equation} where $C(\\theta)$ is the two-point auto-correlation function at an angular separation of $\\theta$ on the SLS} [Sect.~7, Eq.~(9), \\nocite{WMAPSpergel}{Spergel} {et~al.} 2003]. \\postrefereechanges{Hereafter}, we call this $\\Sxi$, in order to distinguish it from the $S$ parameter used for matched circle analyses. The two parameters have very different meanings.  The chance of observing $\\Sxi$ to be as small as that observed was estimated by \\nocite{WMAPSpergel}{Spergel} {et~al.} (2003) as 0.15\\% for an infinite, flat, cosmic concordance model with a fixed spectral index of density perturbations. \\nocite{EfstNoProb03b}{Efstathiou} (2004), using what should be a more accurate method if we assume a simply connected model, estimated the chances to be much higher, from 3\\% to 12.5\\% depending on which sky map is analysed and which (if any) galactic cut mask is used (Table~5, \\nocite{EfstNoProb03b}{Efstathiou} 2004).  We can quantify the relation between large length-scale auto-correlations and the chance of a PDS signal in the case that the PDS is wrong by writing that {\\em the lower the observed value of $\\Sxi$, the less likely it is that a non-PDS model will give a PDS-like signal.} Hence, since we expect that the two properties, a low $\\Sxi$ value and a low chance of a PDS-like cross-correlation signal occurring in a flat, infinite, cosmic concordance model, are related, it would not be useful to estimate their probabilities independently. Instead, what is of interest to investigate is the chance that a PDS-like cross-correlation signal occurs, {\\em given that $\\Sxi$ is as low as that which is observed.} Conservatively, we can calculate an upper limit to this conditional probability if we use simulated sky maps with $\\Sxi$ values a little higher than the observed estimate of $\\Sxi$.  Moreover, here we use the kp2 galactic cut mask, for which \\nocite{EfstNoProb03b}{Efstathiou} (2004) finds the highest probability for the observed $\\Sxi$ value to occur in an infinite, flat, cosmic concordance model, so we should obtain a conservative upper limit to the probability of both (a) an $\\Sxi$ value as low as that observed and (b) a PDS-like cross-correlation signal similar to that observed occurring. This is described in more detail in \\SSS\\ref{s-method-GRF}.    Results are presented in \\SSS\\ref{s-results}, discussed in \\SSS\\ref{s-disc} and conclusions are given in \\SSS\\ref{s-conclu}. For general background on spherical, multiply connected spaces, see \\nocite{Weeks2001}{Weeks} (2001), \\nocite{GausSph01}{Gausmann} {et~al.} (2001), \\nocite{LehSph02}{Lehoucq} {et~al.} (2002) and \\nocite{RiazSph03}{Riazuelo} {et~al.} (2004). See the references cited above for background on cosmic topology. Comoving coordinates are used when discussing distances (i.e. ``proper distances'' at the present epoch, \\nocite{Wein72}{Weinberg} 1972, equivalent to ``conformal time'' if $c=1$) and the Hubble constant is written $H_0 \\equiv 100 h$\\kms/Mpc. ",
        "conclusions": "\\label{s-conclu}  By use of some spherical trigonometry, it is possible to speed up the Markov Chain Monte Carlo cross-correlation method of testing a cosmic topology hypothesis described in RBSG08 by a factor of about 3--10. This is shown in Fig.~\\ref{f-alpha_p} and Eq.~(\\ref{e-alpha-pm}). This could, in principle, make it practical to make calculations at higher resolution than before. However, the physical interpretation of the calculations would be ambiguous because of several systematic effects listed above.  Moreover, for low matched circle sizes and low maximum pair separations, use of Eq.~(\\ref{e-alpha-pm}) can increase the numbers of pairs per separation bin, thereby decreasing the noise. A further improvement in the method and the relation between this method as a test of ``matched annuli'' and tests of matched circles are discussed above in \\SSS\\ref{s-disc-speed}.  Applying the faster method to the ILC and TOH versions of the WMAP five-year data, we find little significant change in the best estimate parameters for a Poincar\\'e dodecahedral space model of the Universe, compared to those given in RBSG08. Depending on the minimum pseudo-probability level used in the MCMC chains that is used for estimating the twist $\\phi$, the optimal value of $\\phi$ is a few degrees greater and a few degrees less than $+36.0\\ddeg$ in the ILC and TOH maps respectively (Table~\\ref{t-alpha-phi-5yr}).  We also applied the faster method to a small number of simulated skies. The WMAP observations confirmed the generic cosmic topology prediction \\nocite{Star93,Stevens93}({Starobinsky} 1993; {Stevens} {et~al.} 1993) of a cutoff in structure statistics such as the temperature-temperature fluctuation auto-correlation on large angular or spatial scales. Here, we estimated the weakness of the large scale auto-correlations using a statistic $\\Sxi$ [Eq.~(\\ref{e-defn-Sxi})] similar to that of \\nocite{WMAPSpergel}{Spergel} {et~al.} (2003). The low observed value of $\\Sxi$ implies that cross-correlations on these scales should usually be weak, so that MCMC chains as used here should have difficulty finding a region of parameter space with high optimal cross-correlations. Hence, to conservatively test the infinite, flat, concordance model hypothesis, on the assumption that the real observations have a low value of $\\Sxi$ due to being a single realisation of a random process, we analysed simulated skies which use the \\nocite{Hinshaw06}{Hinshaw} {et~al.} (2007) $C_l$ estimates and randomised phases of the spherical harmonics. The observational map was calculated from the three least contaminated frequency WMAP bands, Q, V and W, and the Gaussian random fluctuation simulations were created using an equivalent analysis pipeline \\nocite{LR08}({Lew} \\& {Roukema} 2008).  For simulated and observed skies with the kp2 Galactic cut, these simulations using the \\nocite{Hinshaw06}{Hinshaw} {et~al.} (2007) $C_l$ estimates generally yield overestimates of $\\Sxi$ for the cut sky.  Since we prefer to analyse the cut sky in order to minimise galactic contamination effects, it was necessary to select those simulations with the lowest values of $\\Sxi$, so that the simulations were statistically as similar as possible to the observations outside of the kp2 cut.  The results of running MCMC chains on the 20 simulations with the lowest $\\Sxi$ estimates from among 50 simulations were that only two simulations gave optimal cross-correlation solutions which did not escape to the lower limit in circle size $\\alpha$, where small number statistics favour fortuitous cross-correlations. This indicates a conditional probability of finding a non-noise solution of about 10\\% [Eq.~(\\ref{e-prob-PDS-given-Sxi_alpha})].  The distribution of optimal twists $\\phi$ from the simulations showed that despite the low value of $\\Sxi$, the expected distribution of $\\phi$ in the simulations was not uniform. Instead, it is consistent with a Gaussian distribution centred \\postrefereechanges{near} $\\phi=0$, of width $\\sim 33$--$38\\ddeg$. Possible reasons for the nature of this distribution are discussed above, in Sections~\\ref{s-why-phi-symmetric} and \\ref{s-why-phi-distbn}.   Assuming this numerical, Gaussian fit to the expected (for an infinite, flat model) distribution of $\\phi$, and using the estimate of $\\phi$ obtained from the observational map (Table~\\ref{t-alpha-phi-INC3}) that is most discrepant from $\\pi/5$, we found that $\\phi$ could be expected to be as close to $\\pm\\pi/5$ as the observational value is close to $+\\pi/5$ with a probability of about 22\\% [Eq.~(\\ref{e-prob-PDS-given-Sxi_phi})].  Both of these probabilities are conditional on $\\Sxi$ being low, which itself is unusual for an infinite, flat model \\nocite{WMAPSpergel,EfstNoProb03b}(e.g.,  {Spergel} {et~al.} 2003; {Efstathiou} 2004). Hence, for an infinite, flat, cosmic concordance model with Gaussian random fluctuations, we find that comparing an observational map with simulational maps gives an estimate of the chance of finding both (a) a large scale autocorrelation as weak as that observed, and (b) a PDS-like, optimal cross-correlation signal similar to that observed to be about 0.015\\% or 1.25\\%  for the \\nocite{WMAPSpergel}{Spergel} {et~al.} (2003) or \\nocite{EfstNoProb03b}{Efstathiou} (2004) estimates of the probability of low $\\Sxi$ respectively [Eq.~(\\ref{e-prob-final})]."
    },
    "0807/0807.1738.txt": {
        "abstract": "We consider resonant absorption in a spectral line in the outflowing plasma within several tens of Schwarzschild radii from a compact object. We take into account both Doppler and gravitational shifting effects and re-formulate the theory of P-Cygni profiles in these new circumstances. It is found that a spectral line may have multiple absorption and emission components depending on how far the region of interaction is from the compact object and what is the distribution of velocity and opacity. Profiles of spectral lines produced near a neutron star or a black hole can be strongly distorted by Doppler blue-, or red-shifting, and gravitational red-shifting. These profiles may have both red- and blue-shifted absorption troughs. The result should be contrasted with classical P-Cygni profiles which consist of red-shifted emission and blue-shifted absorption features.  We suggest this property of line profiles to have complicated narrow absorption and emission components in the presence of strong gravity may help to study spectroscopically the innermost parts of an outflow. ",
        "introduction": "It is widely accepted that radiation produced in the vicinity of compact objects can bear an imprint of the strong gravitational field, of either a stellar mass black hole (BH) or a neutron star (NS) in Galactic X-ray binary systems, or of a super-massive BH in case of active galactic nuclei (AGN). Thus the problem of direct detection of line emission from regions close to compact objects is of great importance. One of the most important piece of evidence for the presence of an accretion disk, and also the most extensively theoretically investigated subject in this field, is a broadened, red-shifted spectral feature attributed to a fluorescent ${\\rm K}\\alpha$ emission line observed in the X-ray spectra of some Seyfert 1 galaxies. It is believed to be formed due to the reprocessing of X-Rays emitted by corona in the inner parts of a relatively cold accretion disk (see e.g. ~\\cite{Fabian2000}). If this is the case, then the line forming region translates to a broad region in velocity and frequency  space making a line profile severely skewed by relativistic effects of the accreting plasma. This strong distortion of the spectral line makes detailed diagnostics of physical conditions in the plasma using this line difficult.  The grating spectrographs on the X-ray telescopes {\\it Chandra} and {\\it XMM-Newton} provide unprecedented spectral resolution up to $\\sim10\\, {\\rm keV}$. Thus, by means of {\\it Chandra} High Energy Grating observations, a narrow \"core\" of the Fe ${\\rm K}\\alpha$ line was detected in many type I AGNs (\\cite{Yaqoob2003}). In order to explain this spectral feature in these sources it is necessary to invoke both reflection from an accretion disk and also emission from distant regions (such as an obscuring torus).  Observations of several quasars and Seyfert galaxies reveal narrow X-ray absorption lines indicating matter outflowing, falling onto, and orbiting near the BH. These features are interpreted as Fe absorption lines, blue-, and red-shifted relative to the rest-frame frequency.  In some cases, observed spectral features are interpreted as being gravitationally red-shifted absorption lines. Tentative evidence has been found for an absorption feature which presumably originates from resonant scattering of iron within an accretion flow in the strong gravitational field of the BH in the Seyfert 1 galaxy NGC 3516 (~\\cite{Nandra99}). In the paper by \\citet{Yaqoob} it was suggested that the detection of the narrow red-shifted absorption line superimposed on the red wing of a broad Fe ${\\rm K\\alpha}$ line in the quasar E1821+643 is due to resonance absorption by  Fe {XXV} or Fe XXVI, gravitationally red-shifted from region within 10-20 $r_g$ from the BH. ~\\citet{Reeves} interpreted an absorption features from the quasar PG 1211+143 as being produced by ionized iron and either being Doppler red-shifted due to the plasma falling onto the black hole, or due to gravitational red-shifting, or both. ~\\citet{Matt}  also reported on the possible detection of a transient absorption line from the ionized Fe from the quasar Q0056-363. If the observed red-shift is only gravitational then it has been argued  the absorbing gas lies within $\\sim 5\\,r_g$ from the BH. In the paper by \\cite{Dadina05} the red- and blue-shifted absorption lines are found in the X-ray spectrum of the Seyfert 1 galaxy Mrk 509. The data suggest these features are formed in a flow moving with velocity, $v\\sim 0.15-0.2\\,c$, or/and being gravitationally red-shifted.  Much smaller astrophysical objects which may potentially demonstrate gravitationally red-shifted lines are neutron stars. It is widely believed that X-ray bursts are produced as neutron stars are accreting plasma from their massive companions in close binary systems. It was also suggested that measurement of the gravitationally red-shifted absorption lines during an X-ray burst can yield the mass-to-radius relation for the neutron star and thus provides the required constraint to the equation of state of the NS interior. \\citet{Cottam} reported on the observation of several absorption features in the X-ray burst spectra of the low-mass X-ray binary EXO0748-676. Prominent absorption features were attributed to the red-shifted lines of Fe {XXVI}, Fe {XXV}, and O {VIII}. It was speculated that both gravitational red-shift and possibly a slow outflow are responsible for the observed features.   From the above examples one may conclude that there exists the possibility to observe gravitationally red-shifted narrow spectral features from diverse classes of objects.  With an exception of a purely photospheric line, it is natural to assume that plasma which produces line emission is moving with respect to the source of radiation (photosphere or accretion disk).  The line emitting plasma may be in the form of a diffuse wind and may or may not have a clumpy structure. In both cases one may hope to acquire important information from analyzing the profiles of such lines, and to learn about the strength of the gravitational field,  distance, velocity field etc.  Many normal, luminous stars possess strong mass loss. In many cases it is possible to derive information about the wind by analyzing the P-Cygni profiles formed in these winds. These peculiar profiles have long been known. The first paper which contained an explanation of P-Cygni profiles on the basis of the Doppler effect was that of ~\\citet{Beals}, who gave a basic explanation of the observed line profiles from novae and Wolf-Rayet stars. Since then, these profiles have been studied both observationally and theoretically in numerous papers. For example, \\citet{Morton}, extensively investigated the P-Cygni ultraviolet resonant line profiles from winds of early supergiants.  A theoretical breakthrough was made in the paper by \\citet{Sob}, in which he recognized that if there exists a velocity gradient along the line of sight, the radiation transfer problem becomes purely local. That is, distant regions in the flow can no longer exchange information using photons with frequencies within the width of the line. The radiation which is seen at a certain frequency within a line profile by a distant observer comes from a surface of equal line-of-sight velocity. (see Sect. 3.1). Following this, significant attention has been paid to computing and analyzing P-Cygni profiles and also to investigating their dependence on fundamental parameters of the flow (see e.g. the atlas of P-Cygni profiles by \\cite{CastorLamers}) .   The main goal in this paper is to calculate the line profile produced by a plasma which is moving in the strong gravitational field of a compact object. Thus, we want to obtain an answer to the question of how much the gravitational red-shifting can change the observed line profile in comparison with the standard case of P-Cygni line. Expecting that the gravitational red-shifting may introduce some characteristic distortion to the P-Cygni profile, we wish to investigate what kind of the distortion is made and whether it can be used to deduce information about the wind and compact object system itself.  The plan of the paper is as follows: in Sect.2 basic assumptions are described about physical conditions and geometry of the line forming region. Here the objectives and limitations of the approach are defined; taking into account gravitational red-shifting,  in Sect.3 we first calculate a spectral line optical depth. This is done in the Sobolev approximation which we adopt throughout the paper. In this section equal frequency surfaces (EFS) are re-defined to include the gravitational red-shifting effect and the shapes of such surfaces are calculated for several typical velocity profiles. The results obtained in this section are extensively used in the  calculation of the radiation field within a spectral line in Sect.4 or when we numerically calculate line profiles in Sect.5. At the end of the paper we provide a discussion and summarize the results. ",
        "conclusions": "Our goal in this paper has been to study shapes of spectral lines from plasma which is rapidly moving in the vicinity of a neutron star or a black hole. The latter case can account for both stellar (Galactic Black Hole Candidates) and supper-massive BHs (AGN). In the well studied case of winds from normal hot stars, a rapidly moving wind interacts with the continuum radiation of a star and produces a P-Cygni profile. From the theory of stellar winds it is known that in the case of an arbitrary spherically-symmetrical distribution of plasma which is moving with a gradually increasing velocity, an absorption line that is blue-shifted with respect to the emission line is observed. This is widely interpreted as a fingerprint of moving plasma in a variety of astrophysical situations.  In this paper, shapes of line profiles for several velocity and opacity distributions were calculated taking into account gravitational red-shifting.  Strong gravitational red-shifting helps the radiation to escape efficiently and to interact with matter only locally. This was already established to be dynamically important. The papers by \\citet{Dora03} and \\citet{DN} addressed a problem of plasma acceleration driven by radiation pressure in spectral lines provided that the flow is launched in the vicinity of a compact object. It was shown that it is important to include gravitational red-shifting in the calculations of the radiation force. The radiation pressure depends on $d\\phi/dr$ and $dv/dr$, while in the case of O-type star wind it depends only on $dv/dr$. Proximity to compact object, i.e. strong ionizing radiation, makes it difficult for heavy ions to survive against being completely stripped of electrons, so the efficiency of this additional mechanism in real situation depends on the simultaneous solution for the ionization balance and  accounting on other possible mechanisms (such as clumping) to prevent over-ionization of the flow.  We are also motivated by the current accumulating evidence for the gravitationally red-shifted narrow absorption features in many AGN spectra as well as by observations of gravitationally red-shifted absorption lines in the X-ray burst spectra of neutron stars. Potentially interesting example of the latter is presented by the observations of the gravitationally red-shifted absorption lines in the X-ray burst spectra of EXO0748-676 (\\citet{Cottam}). If the interpretation is correct, these features are produced close to the compact object, where extreme conditions are coupled with high amplitude fluctuations of the radiation field and shortest dynamical time scale. As a result, these lines may be highly variable or/and transient. The non-detection of gravitationally red-shifted lines in the observations of the other good looking candidate GS, 1826-24 (\\cite{Cottam2}) may be an example.  Given the diverse nature of these objects, it is desirable to consider a model with a minimum number of free parameters and in so doing we assume our outflow to be spherically-symmetric and exposed to the continuum radiation of a spherical core.  Realistic models should consider a departure of the outflow and the radiation field from spherical symmetry. Bending of the photon's trajectories in the strong gravitational field of the compact object, may play some role. In our calculations we did not take into account because we were concerned with regions of the flow located approximately at radii $> 10 r_g$ .  In our calculations the idea of the equal frequencies surfaces (EFS), plays a major role. The strong gravitational field changes the shape and locus of such a surface. Their topology is complex; the gravitational field strongly distorts the P-Cygni profile. Some branches of the EFSs in the {\\it red-shifted} part of the spectra are found to be in front of the core, meaning the possibility for the absorption component to be observed as {\\it red-shifted}, with respect to the emission.  From numerical calculations, which are second order accurate in $v/c$, it is established that a superposition of Doppler and gravitational shifting of frequency can distort the P-Cygni profile in such a way that blue- and red-shifted absorption features are observed simultaneously. Often the red-shifted absorption line is superimposed on the emission wing. This effect is, of course, strongest close to the compact object. However, the emission part is also stronger there, and it may smear the absorption trough. Necessarily, the second absorption arises if the velocity profile is not very steep and at the same time the line forming region is situated within several tens of $r_g$. However the former requirement is not crucial, if the the line opacity peaks close to the compact object. Profiles with more than one emission and two absorption features are possible. However, in a model of pure line scattering, those additional (i.e supplementary to W-shape profile) features would be considerably weaker.    Different modifications of  a ''stellar'' type velocity laws were adopted. Such velocity profiles describe stationary spherically-symmetric hydrodynamical flows.  The well known example of such is the $v\\sim\\sqrt{1-1/x}$ profile. Note that different modifications of this law were obtained by different authors using  both theoretical and observational arguments (see e.g. \\cite{CastorLamers}). Thus we adopted a generalized form of this law, $v\\sim (1-1/x^{\\alpha_1})^{\\alpha_2}$. In case of non-magnetic accretion disk winds, the azimuthal component of the velocity quickly becomes much smaller than the radial one, and such a  velocity profile describes the real distribution of the line-of-sight velocity reasonably well.  During X-ray bursts, the situation can be more complicated. During the burst, the temperature at the stellar photosphere may quickly approach $T\\sim 10^9K$ and the electron scattering occurs in the Klein-Nishina regime. The reduced cross section allows for higher radiation flux (locally below Eddington) to diffuse out to larger radii where temperature is smaller and the radiation flux is locally supper Eddington. A quasi-stationary wind may result. Calculations of such winds must self-consistently account on the radiation transfer and relativistic corrections and demonstrate stellar type velocity profiles (\\cite{QuinnPacz, Nobili94}.)   To approximate the explosively accelerated plasma of X-ray bursters we adopt a linear velocity law. Making use of a Hubble law, ${\\bf v}={\\bf r}/t$ drastically simplifies the radiation transfer calculations in supernova shells (\\cite{Karp77}).  Two parameterized opacity laws and a linear, and generalized form of the velocity law, $v\\sim (1-1/x^{\\alpha_1})^{\\alpha_2}$ were considered.   Our results for all considered velocity and opacity distributions show that in particular circumstances, i.e. proper velocity and opacity laws, strong gravity, the observed line profile consists of  two absorption troughs separated by the emission component. The red-shifted absorption line can be weak and can be superimposed on the emission wing. Some of the strongest W-shape profiles were found for the linear velocity law. Surprisingly enough, no W-shape profiles are found for a $v\\sim\\sqrt{1-1/x}$ law. In this case the red-shifted absorption almost exactly compensates for the blue-ward emission peak, producing a sharp edge in the emission (Fig.~\\ref{prof_stel1}).  However, in many other cases of different $\\alpha_1$,$\\alpha_2$, W-shape profiles are also found (Fig.~\\ref{prof_stel2} - \\ref{prof_stel3}). If the opacity is distributed as $\\tau\\sim (1-v/V^\\infty)^k$, ($k>0$), the red-shifted absorption line in most cases is superimposed on the red-shifted emission. However, if the opacity peaks at smaller velocities, $\\tau\\sim 1/v^k$, distinct W-shape profiles are found again. Thus, the results suggest that the most favorable conditions realize when the opacity peaks at large gravitational red-shifts and in such a case the velocity profile is probably less important. The further the distribution of the opacity is from that the more important is the velocity law and the relative importance of the Doppler effect.  If some observed spectral feature is interpreted as being formed in the flow in a proximity of a strongly gravitating compact object, then our results suggest that the assumption of the distribution of the opacity (optical depth law) is the most critical one. Such distribution should be calculated simultaneously with the ionization balance calculations in the background of the pre-calculated hydrodynamical model.  The separation between red- and blue-shifted absorption features is a function only of the dynamics and relative importance of gravity in the line forming region.  The most robust prediction of our model is the possibility of the double absorption trough as a result of a concurrence between gravitational red-shifting and Doppler effects. Our calculations suggest that the particular shape and intensity of the emission component which separates these absorption lines depend sensitively on assumed parameters.  Gravitationally red-shifted absorption lines form in places close to the compact object. The emission component is necessarily formed in plasma occupying much larger volume. Thus, these features are formed in places which are possibly strongly separated in space. Thus, from the perspective of future observations, it would be interesting to look for correlated variability of different components of the profile between each other and with the continuum."
    },
    "0807/0807.3992_arXiv.txt": {
        "abstract": "A two-components model for Broad Line Region (BLR) of Active Galactic Nuclei (AGN) has been suggested for many years but not widely accepted (e.g., Hu et al. 2008; Sulentic et al. 2000; Brotherton et al. 1996; Mason et al. 1996). This model indicates that the broad line can be described with superposition of two Gaussian components (Very Broad Gaussian Component (VBGC) and InterMediate Gaussian Component (IMGC)) which are from two physically distinct regions; i.e., Very Broad Line Region (VBLR) and InterMediate Line Region (IMLR). We select a SDSS sample to further confirm this model and give detailed analysis to the geometry, density and evolution of these two regions. Micro-lensing result of BLR in J1131-1231 and some unexplained phenomena in Reverberation Mapping (RM) experiment provide supportive evidence for this model. Our results indicate that the radius obtained from the emission line RM normally corresponds to the radius of the VBLR, and the existence of the IMGC may affect the measurement of the black hole masses in AGNs. The deviation of NLS1s from the M-sigma relation and the Type II AGN fraction as a function of luminosity can be explained in this model in a coherent way. The evolution of the two emission regions may be related to the evolutionary stages of the broad line regions of AGNs from NLS1s to BLS1s. Based on the results presented here, a unified picture of hierarchical evolution of black hole, dust torus and galaxy is proposed. ",
        "introduction": "Type I AGNs are often classified into two subclasses according to the Full Width at Half Maximum (FWHM) of their broad H$\\beta$ lines. Those AGNs with FWHM greater than approximately 2000 km$\\cdot$s$^{-1}$ are called Broad Line Seyfert 1 (BLS1), and those with FWHM less than approximately 2000 km$\\cdot$s$^{-1}$ are called Narrow Line Seyfert 1 (NLS1). There is a tight correlation between the black hole mass and the stellar velocity dispersion of the bulge of normal galaxies, the so called M-sigma relation (e.g. Magorrian et al. 1998; Gebhardt,  et al. 2000;  Tremaine et al. 2002; Marconi \\& Hunt 2003; Ferrarese \\& Ford 2005). BLS1s are found also to follow this relation well (Greene \\& Ho 2006). However, NLS1s seem to deviate from this relation; they seem to have much smaller masses or much higher stellar velocity dispersions (e.g., Wang \\& Lu 2001; Bian \\& Zhao 2004; Zhou et al. 2006; Komossa \\& Xu 2007). {\\it We call this under-massive black hole problem of NLS1s.} Since there are a lot of uncertainties on measuring the black hole mass and sigma (Decarli et al. 2008; Komossa \\& Xu 2007; Komossa 2008; Komossa et al. 2008), it is difficult to determine whether the deviation reflects an intrinsic mass or/and sigma difference between the NLS1s and BLS1s, or it is only caused by the measurements. NLS1s also show softer X-ray spectra, strong Fe II emission, rapid continuum variation, and high accretion rate; the last one may partly be caused by the incorrect measurement of the black hole mass (e.g., Grupe \\& Mathur 2004; Komossa 2008). However, Williams et al. (2004) found some optically defined NLS1s are different from those defined by X-ray selected NLS1s, suggesting that strong, ultrasoft X-ray emission is not a universal characteristic of NLS1s and challenging the current paradigm that NLS1s accrete near the Eddington limit.  In the simple unified model for AGN, Type I and Type II objects differ only in terms of the angle between the observer's line of sight and the normal axis of a dusty torus (Urry \\& Padovani 1995). SEDs of QSOs also imply that the location of the inner wall of torus is determined by the sublimation of dust by the central radiation (e.g. Kobayashi et al. 1993; Elvis et al. 1994). It is widely accepted that $R_{\\rm{ torus}}\\propto L^{0.5}$ (e.g., Elitzur 2006; Peterson 2007; Nenkova 2008a), consistent with that predicted by the toy model of Krolik \\& Begelman (1988). Reverberation mapping (RM) based on infrared emission also supports this relation (Suganuma et al. 2006).  Many observations show that the fraction of type I AGNs increases with luminosity (Ueda et al. 2003; Hasinger 2004, 2005; Steffen et al. 2003, 2004; Barger et al. 2005; Wang et al. 2005), i.e., brighter AGNs are mostly type I AGNs. This is also equivalent observationally to the situation that it is less probable to find very luminous type II AGNs (mostly Seyfert II galaxies) in a random sample of AGNs. {\\it We call this under-populated luminous type II problem of Seyfert galaxies.} However, the standard receding torus model with constant height first suggested by Lawrence (1991) failed to describe this luminosity fraction, and thus a modified receding torus whose height slightly increases with luminosity with the relation $h\\propto L^{0.23}$ is needed (Simpson 2005). Based on the luminosity fraction for Type I AGN, Wang et al. (2005) obtained a relation between covering factor of the torus and luminosity, $\\log{C}=-0.17\\log{L_{x}}-0.36$. The variation of covering factor with luminosity may be totally caused by the receding of the inner torus (Wang et al. 2005).  Clearly both the broad line regions and the dust tori of AGNs show strong dependence of luminosity, yet the physical connections between these two key ingredients of AGN's unification scheme are not yet fully understood. It is also not clear if the luminosity dependence of dust torus and broad line region is simultaneously related to both problems of under-massive black holes in NLS1 galaxies and under-populated luminous type II Seyfert galaxies. Progress towards solving these problems may shed new lights to the understanding the hierarchical evolution of AGN and its host galaxy, as well as the feeding, fueling and growth of supermassive black holes. In this paper we attempt to address both problems in a coherent way by studying a sample of SDSS AGNs with strong broad emission lines.  The geometry and kinematics of broad line region in AGN have been studied for about three decades but far from been fully understood. There are mainly two kinds of views to interpret the structure of BLR based on the profiles of balmer lines. One interprets the profile as a Gaussian/Lorentz profile and BLR as a extended region has different projected velocity distributions (e.g., Goncalves et al. 1999; Collin et al. 2006). The other model is that the profile is a superposition of double Gaussian profiles (Intermediate Component+Very Broad Component in Hu et al. 2008; Broad Component+Very Broad Component in Sulentic et al, 2000 and Marziani et al. 2009. These two components will be re-named as InterMediate Gaussian Component (IMGC)+Very Broad Gaussian Component(VBGC) in this paper.), which come from two physically distinct emission regions due to their substantial differences in the line widths (e.g., Hu et al. 2008; Sulentic et al. 2000; Brotherton et al. 1996; Mason et al. 1996); we call these two regions Very Broad Line Region (VBLR) and InterMediate Line Region (IMLR). For the two-components model, the results from these papers are not consistent with each other and the physical interpretations to these two emission regions are also different. Sulentic et al (2000) suggested that VBGC is likely to arise in an optically thin region close to the central source which is slightly red-shifted, whereas Hu et al. (2008) concluded that IMGC is systematically red-shifted and may come from the inflow. In addition, the study of partially obscured quasars suggests that there is an inner narrow line region covered by dust which may be consistent with our IMLR (Zhang et al. 2009). The existence of these two emission regions need to be confirmed further and their dynamical and physical properties also need to be further studied. For this purpose, we select a SDSS sample with strong H$\\alpha$, H$\\beta$ and H$\\gamma$ lines, and decompose these balmer lines based on the two-components model. Our goal is to provide further evidence for this model, carry out more detailed analysis about the dynamics and evolution of these two emission regions, and ultimately to understand the above two problems (The under-massive black hole problem of NLS1s and The under-populated luminous type II problem of AGNs) in a coherent way.  This paper is organized in the following way. Detailed decomposition and FWHM measurement are described in Section 2. In Section 3 we present H$\\alpha$, H$\\beta$, and H$\\gamma$ decomposition and statistical analysis. In section 4 we confirm the two-components model based on the line decomposition and give detailed analysis of these two emission regions. Section 5 presents other supporting evidence for this model. Section 6 shows that this model provides a coherent interpretation to the two problems. Section 7 is conclusion and discussion. The continuum luminosity and some other parameters of AGNs are directly taken from Table 2 in La Mura et al. (2007). FWHM is defined as FWHM of the whole line, FWHMi means FWHM of the IMGC and FWHMb represents FWHM of the VBGC when they are not clear in the context. The unit is km s$^{-1}$ for the ``width\" of all lines and components, as well as the inferred velocity, throughout this manuscript. ",
        "conclusions": "We conclude that the decomposition of broad H$\\beta$, H$\\alpha$, and H$\\gamma$ line of the AGNs confirms the two component model of BLR which has been suggested by several previous studies (e.g., Brotherton 1996; Sulentic et al. 2000; Hu et al. 2008.). We have made detailed analysis about the two emission regions (VBLR and IMLR) based on the Balmer line decomposition, and find other supportive evidence for this model. Our main conclusions are:  \\begin{enumerate}  \\item The two Gaussian components exhibit evolutions with increasing FWHM (we note in passing that because the two components show quite different dependence with FWHM and luminosity as to be shown in the following, we rule out the possibility that the dependence is caused by systematic biases in the decomposition process). The evolution is much stronger for the H$\\beta$ and H$\\gamma$ lines. Our results offer strong evidence of the evolution of the broad line region (consists of a IMLR and a VBLR) evolution from NLS1s to BLS1s. We obtain the luminosity dependence for the radius of IMLR, $R_{\\rm IMLR}\\propto L^{0.37}$, if the luminosity dependence for the VBLR is taken as $R_{\\rm VBLR}\\propto L^{0.52}$ (Kaspi et al. 2005). The two emission regions have a trend to merge into one region with luminosity increasing. Balmer decrement and the Baldwin effect in IMLR indicate that it has a flattened geometry, higher density and contains more dust, compared to the VBLR. The receding velocity of IMLR is consistent with dust torus. Therefore, we suggest the IMLR is the hot inner skin of torus. A cartoon of the evolution of BLR emerge from these analysis.  \\item There are other evidence in support of this two component model. The study of micro-lensing provides possible evidence for the existence of IMLR. The double-peaked CCCD from the RM data of Mrk 79 also provides possible evidence. Simulations suggest that the narrower RMS spectra of broad emission lines from many AGNs may be consistent with our model, although more work needs to be done to establish this as the case. The existence of a weak IMGC in the broad emission lines of many sources with RM measurements may cause systematic biases for the measurements of the radius of the VBLR, thus biasing the Radius-Luminosity relation for the VBLR. It will be helpful to decompose each broad emission line into two Gaussian components as we have done here, and then do cross-correlation analysis between the continuum and each of the two components, in order to measure the Radius-Luminosity relations for the two components independently.  \\item In our model, only the VBGC should be used to estimate the black hole mass, and the radius measured by reverberation mapping based on emission lines normally represents the radius of this region. After correction for black hole masses, NLS1s still have smaller black hole masses (compared to BLS1s) but normal accretion rate in units of Eddington rate. Therefore, the black hole mass increases from NLS1s to BLS1s by following the M-sigma relation established for normal galaxies.  We obtain the luminosity dependence for the height of IMLR as $h\\propto L^{0.25}$. It can well explain the luminosity function of AGN, if the decreasing fraction of type II AGNs for higher luminosity is due to completely the decreasing covering angle of the dusty torus to the central irradiation source (Simpson. 2005).  Therefore both the problem of under-massive black hole in NLS1s and the problem of under-populated luminous type II Seyfert galaxies can be understood properly if our model is true, still more concrete evidences for this model are needed.  \\end{enumerate}  Hu et al. (2008) found evidence that IMLR is related to inflow towards the VBLR. Combining with this conclusion, we suggest that the inflow from the inner boundary region of AGN's dusty torus may provide the supply to the accretion disk surrounding the central black hole. The strong and positive luminosity dependence of the geometry of IMLR suggests that the dust sublimation by the central accreting black hole's radiation dominates the structure and evolution of IMLR. If IMLR is related to inflow (Hu et al. 2008), we may be able to further suggest that the inflow is caused by the dust sublimation, i.e., a consequence of the feed-back of the black hole's accretion and radiation. Because IMLR is also ionized similarly to VBLR, the IMLR induced viscosity allows efficient angular momentum transfer to drive gas inflow from the dust torus and consequently fuel the accretion flow onto the central black hole. In this scenario, the accretion flow is self-regulated by the radiation from the accretion disk through irradiation to the dust torus. Therefore the growth of the supermassive black hole is at the expense of consuming the material in the dust torus during the AGN phase; this is consistent with the observation that for very low luminosity AGNs, the luminosity decreases with decreasing absorption column, i.e., the AGNs in their last stages are running out accretion material supplied by the torus (Zhang \\& Soria et al. 2009). Of course not all material in the accretion flow falls into the black hole horizon to increase the black hole'a mass, since accretion winds and outflows are common in AGNs.  After the material in the dust torus is completely consumed, the AGN phase will be turned off and the galaxy becomes a normal and inactive galaxy. Indeed many AGNs in low luminosity (because of low accretion rate and low radiation efficiency) show very little, or even no signs of torus and/or broad line region; there is also no evidence for dust torus or broad line region in the centers of normal and inactive galaxies, including the Milky Way. Therefore the AGN's dust torus is the missing link or bridge between the coeval growth of a black hole and its host galaxy. This would require that during the merging of two galaxies, a dust torus is first formed, perhaps due to the residual orbital angular momentum of the two galaxies. The dust torus then fuels the accretion and growth of the supermassive black hole through the self-regulation of irradiation to the dust torus by the accretion disk. The initial trigger to this self-regulation process may be Bondi accretion of gas with negligible angular momentum, or the low level AGN activities of the two black holes in the two parent galaxies before the merger. This evolutionary scenario is illustrated in Fig.\\ref{fig:evolution}.  The above scenario is generally consistent with that proposed by Wang \\& Zhang (2007),  but also with some important difference. We suggest that the torus is formed by the merging of two galaxies and disappears after each AGN cycle; the different appearance (mainly geometry) of torus in different types of AGNs are mostly due to the self-regulation of the accretion and irradiation of the AGN. Therefore in our scenario, torus evolution is fast (only lasting for one episode of AGN activity) and synchronized with that of the broad line region, whose evolution is also dominated by the luminosity of the AGN. {\\it Our torus evolution is mostly hierarchical.} For example, although NLS1s also have generally smaller black hole masses, we do not find NLS1s deviate from the M-sigma relation more than the BLS1s after black hole mass correction made here. Therefore in our scenario NLS1s are produced by mergers of smaller galaxies compared to BLS1s; NLS1 may or may not show up as BLS1s in the future, depending upon if more galaxy mergers grow them up in the future. In the scenario of Wang \\& Zhang (2005), NLS1s are in their early growth stage and will grow to become BLS1 during this particular AGN cycle. {\\it Therefore their torus evolution is mostly secular.} In practice, both hierarchical and secular evolutions should be needed for the black hole, torus and host galaxy. It is natural that hierarchical evolution dominates at high redshifts where merger rate is very high, and secular evolution dominates at low redshifts."
    },
    "0807/0807.1735_arXiv.txt": {
        "abstract": "Using a sample of 43 bright ($V<16.1$, distance $<13$ kpc) RR Lyrae stars (RRLS) from the QUEST survey with spectroscopic radial velocities and metallicities, we find that several separate halo substructures contribute to the Virgo overdensity (VOD).  While there is little evidence for halo substructure in the spatial distribution of these stars, their distribution in radial velocity reveals two moving groups.  These results are reinforced when the sample is combined with a sample of blue horizontal branch stars that were identified in the SDSS, and the combined sample provides evidence for one additional moving group.  These groups correspond to peaks in the radial velocity distribution of a sample of F type main-sequence stars that was recently observed in the same directon by SEGUE, although in one case the RRLS and F star groups may not lie at the same distance. One of the new substructures has a very narrow range in metallicity, which is more consistent with it being the debris from a destroyed globular cluster than from a dwarf galaxy.  A small concentration of stars have radial velocities that are similar to the Virgo Stellar Stream (VSS) that was identified previously in a fainter sample of RRLS.  Our results suggest that this feature extends to distances as short as $\\sim 12$ kpc from its previous detection at $\\sim 19$ kpc.  None of the new groups and only one star in the sample have velocities that are consistent with membership in the leading tidal stream from the Sagittarius Dwarf Spheroidal Galaxy, which some authors have suggested is the origin of the VOD. ",
        "introduction": "}   Over the past several decades, a large number of investigations of the stars and the globular clusters in the Galactic halo have concluded that the halo formed chaotically and over a long period of time through the accretion and disruption of dwarf galaxies (e.g., Searle \\& Zinn 1978, Norris 1986; Carney et al. 1990; Majewski 1992; Mackey \\& Gilmore 2004; see also Freeman \\& Bland-Hawthorn 2002 for a review). \\nocite{sz78,norris86,carney90,majewski92,mackey04,freeman02} There is evidence, however, that the inner and outer parts of the halo have sufficiently different properties that at least two mechanisms are required to explain the halo \\citep[e.g.][]{hartwick87,zinn93,majewski93,chiba00,kinman07,miceli07,lee07,carollo07}. Accretion appears to have dominated the formation of the outer halo, but the flattening and prograde rotation of the inner halo and the old ages of the globular clusters in the inner halo suggest that its formation involved the dissipative merging of gas clouds at an early epoch \\citep[see][and references therein]{carollo07}.  While this dual halo picture is supported by substantial evidence, it continues to be important to see what fraction of the halo can be explained by, for example, an accretion-only scenario.  The most direct evidence for accretion comes from the detection of halo substructures consisting of co-moving streams of stars.  The best documented of these are the streams originating in the Sagittarius (Sgr) dwarf spheroidal (dSph) galaxy, which cross much of the sky \\citep[][among others]{yanny00,ivezic00,ibata01a,vivas01,vivas05,newberg02,majewski03}. Several other substructures have been recently detected in the halo of the Milky Way \\citep[][among others]{newberg02,ibata03,majewski03,yanny03,rocha04,vivas06,belokurov06} and in the halo of M31 \\citep{ibata01b,ferguson02,morrison03,zucker04,kalirai06,fardal07}. These streams provide strong evidence in favor of the hierarchical model of the formation of galaxies, in which large galaxies such as the Milky Way formed through multiple mergers with smaller building blocks.  Simulations of the hierarchical picture predict that the Milky Way has experienced a large number of minor mergers \\citep[see for example][]{bullock05} and show that the ones that occurred during the past few Gyrs. may be recognizable today as stellar streams. While these simulations appear to be consistent with the current observations of the outer halo \\citep{bell07}, the search for more streams and better documentation of the known ones remains important.  We report here an investigation of the halo substructure in the direction of Virgo.  The first sign of an overdensity of halo stars in this direction came from the QUEST RR Lyrae star (RRLS) survey \\citep{vivas01}, which revealed a small overdensity near 13 hr right ascension ($\\alpha$), which was the west boundary of the survey as it existed then.  As the survey was extended, it became clear that this feature is considerably larger and is centered near $\\alpha \\sim 12.4$ hr \\citep{vivas02,vivas03}. In the jargon of the QUEST survey it became known as the {\\it \"12.4 hr clump\"} (declination ($\\delta$) was not specified because of the narrow range of the survey around $\\delta = -1\\degr$).  The RRLS in the clump cover a range of distances from the Sun ($D$) with the greatest concentration at $D \\sim 19$ kpc \\citep{vivas06}.  Using photometry from the SDSS, Newberg et al. (2002) independently discovered an overdensity in F type main-sequence stars in approximately the same direction and at a similar distance, which they called S297+63-20.5.  They suggested that it may be related to the small group of RRLS identified earlier by \\citet{vivas01}, and several authors have subsequently suggested that the 12.4 hr clump of RRLS and S297+63-20.5 are the same substructure or at least related. Also using the SDSS photometric catalogue, Juric et al (2005, 2008) \\nocite{juric05old,juric05} later identified a very large overdensity of main sequence stars toward Virgo (the {\\it Virgo Overdensity} (VOD) ).  On the sky, the VOD encompasses both the 12.4 hr clump and S297+63-20.5. But, according to \\citet{juric05}, the VOD is spread along the line of sight from $\\sim 6$ to $\\geq 20$ kpc, which is a larger range than either the 12.4h clump or S297+63-20.5. The spectroscopy of RRLS in the 12.4 hr clump by \\citet{duffau06} provided the first evidence of a moving group in the Virgo direction, which they called the {\\it Virgo Stellar Stream} (VSS).  The recent spectroscopy of F type main-sequence stars in S297+63-20.5 by \\citet{newberg07} revealed a rich structure in velocity space, but with a major peak that nearly coincides with the velocity of the VSS. They tentatively conclude that S297+63-20.5, the VSS, and the VOD are likely associated, if not the same. On the other hand, several authors \\citep{majewski03,martinez04,martinez07} have argued that the Virgo substructures are parts of Sgr streams, while others \\citep{newberg02,newberg07,juric05,duffau06} have rejected this hypothesis and suggested instead that they are due to a separate merger event involving a dwarf galaxy.  We are carrying out a program to observe spectroscopically a large number of QUEST RRLS in order to measure their radial velocities and metallicities.  RRLS are arguably the halo tracer that can be most precisely located in space, and with the addition of radial velocities of even modest precision ($\\pm 20$ \\kms), they become powerful probes of phase space.  This wonderful property of RRLS comes with a price, for they are expensive observationally.  Complete light curves are necessary to obtain both mean magnitudes (hence, precise distances) and ephemerides. The later are necessary to isolate the motion of the star in the sky from the pulsational velocity of the star by means of the fitting of a radial velocity curve.  In this paper, we present spectroscopic observations of bright RRLS ($V_0 < 16.1$, equivalent to D $\\lesssim 12.5$ kpc) in the part of the QUEST survey that overlaps with the VOD.  After examining these data for halo substructure, we discuss the relationships between our results and the previous detections of substructures in Virgo and also consider the possibility that they are related to streams from the Sgr dSph galaxy.  The RRLS surveys by \\citet{keller07} and by \\citet{wilhelm07} provide important data on parts of the VOD that are not covered by the QUEST survey. ",
        "conclusions": "}  Kinematical information is crucial to detect substructures in regions where no obvious spatial overdensities are present.  Taken together, our RRLS sample, the BHB star sample of \\citet{sirko04a}, and the F star sample of \\citet{newberg07} paint a consistent picture that several separate halo substructures exist at $D \\lesssim 13$ kpc in the part of the VOD that is covered by the QUEST survey. The RRLS sample of \\citet{duffau06} and the same BHB and F star samples show that some of these substructures extend to larger $D$. None of these substructures seem to be related with the Sgr leading tail.  The analysis of the SDSS photometric data by \\citet{newberg02,newberg07} make a compelling case that the region of highest stellar density in S297+63-20.5 lies at $D \\sim 18$ kpc.  This is identical to within the errors of the region of highest density in the 12.4 hr clump of QUEST RRLS, which is where \\citet{duffau06} discovered the VSS in velocity space.  The near coincidence in $\\langle V_{gsr} \\rangle$ between the VSS and the largest peak in $V_{gsr}$ that \\citet{newberg07} observed in S297+63-20.5 provides strong evidence that they are the same.  As noted by \\citet{newberg07}, the small offset ($\\sim 30$ \\kms) could be due to measurement error and/or by the small sizes of the samples of spectroscopically observed stars.  It is also possible that it is due to a gradient in $V_{gsr}$ with $D$ (see above).  \\citet{newberg07} found the same velocity peak in their two spectroscopic fields, which indicates that the S297+63-20.5/ VSS has considerable angular extent on the sky.  Other evidence for this is provided by the recent RRLS surveys by \\citet{keller07} and \\citet{wilhelm07}. \\citet{keller07} found 2 clumps of RRLS south of the VOD region (see Figure~\\ref{fig-zone}) at coordinates $(l,b) = (301\\degr,56\\degr)$ and $(324\\degr,49\\degr)$ and distances 16 and 19 kpc respectively. Velocity information is not yet available for these clumps.  \\citet{wilhelm07} identified an excess of candidate RRLS having $V_{gsr} \\sim 100$ {\\kms } in the North Galactic Cap ($b>70\\degr$). The results presented here indicate that the S297+63-20.5/VSS feature may also extend along the line of sight to distances as close as $\\sim 11.8$ kpc.  We conclude that the VOD contains at least one relatively large structure, S297+63-20.5 /VSS, which is intersperced with smaller halo substructures that are unrelated to S297+63-20.5/VSS or to each other. While these smaller structures might be peculiar to the VOD region, it is more likely that they are common in the outer halo.  They are better documented in the VOD direction simply because the presence of S297+63-20.5/ VSS has drawn attention to this region.  In the case of our sample of RRLS, the spatial distribution of the stars provided only very marginal evidence of any substructure \\citep{vivas06}.  The picture changed dramatically when we used spatial and velocity data together.  There is already strong evidence from the spatial distribution of stars in the SDSS that much, if not all, of the outer halo was accreted from the destroyed of satellite galaxies \\citep{bell07}.  Our investigation of the VOD region suggests that the addition of velocity data will make the case even stronger."
    },
    "0807/0807.3730_arXiv.txt": {
        "abstract": "The driving mechanism of the AGB superwind has become controversial in recent years. The efficacy of dust-driven mass loss has been queried. Spitzer observation of AGB stars in Local Group Galaxies show the surprising result that at low metallicity,  AGB mass loss occurs at low luminosity, possibly lower than in the Galaxy, but only for carbon-rich stars.  Oxygen-rich stars in the Galaxy and in lower metallicity galaxies have similar mass-loss rates only at high luminosities.  To explain this dichotomy, we propose that the superwind has a dual trigger. The superwind starts either when sufficient excess carbon builds up for efficient formation of carbonaceous dust (which we propose occurs when $X_{\\rm CO} = \\rm (C - O)/ O_\\odot = 0.1$), or when the luminosity reaches a value sufficient for a silicate-dust-driven wind (proposed at $L = 10^4 Z^{-4/3}\\,\\rm L_\\odot)$. We show that this dual trigger fits the current observational constraints: the luminosity at which the superwind begins, and the predominance of carbon superwind star at low metallicity.  We use stellar evolution models to check the consistency of our explanations and present detailed predictions of the luminosities at which the superwind is triggered for different metallicities and initial stellar masses. ",
        "introduction": "The late phases of stellar evolution are crucial  to galactic evolution. Stars in this phase eject much of their gas into space, enriched by the products of nuclear burning and newly formed dust. Of all the stars which have died in the Universe so far, over 90\\% are low and intermediate-mass stars ($\\rm M\\sim 1$--$8\\,\\rm M_\\odot$: LIMS).  In consequence, LIMS are important contributors to the enrichment of the Interstellar Medium (ISM) by  sources internal to the galaxy.  LIMS eject between 20 and 80 per cent of their mass during the so-called {\\it superwind} phase which terminates the high luminosity asymptotic giant branch (AGB). This near-catastrophic but non-explosive wind removes effectively the entire hydrogen-rich envelope. During the superwind, dust forms in the ejecta, and the star tends to become self-obscured.  After this final burst of activity, the star remains as a hot, compact white dwarf. The expanding circumstellar ejecta become transparent, and briefly form an ionised planetary nebula, before dispersing into the interstellar medium.  At the low effective temperature of AGB stars, the ejecta are largely molecular.  The first stable and abundant molecule to form is CO: this locks up the least abundant of the C and O atoms. Other molecules and dust form from the remaining atoms. Dredge-up of primary carbon (produced by triple-$\\alpha$ burning) causes the C/O ratio to increase over time.  The chemistry changes dramatically when the C/O ratio passes unity, and this affects both the gas-phase species (Millar et al. 2000, Willacy et al. 1997) and the dust composition (Treffers et al. 1974).  In carbon stars, the C/O ratio (by number) exceeds unity, and carbon-rich molecules (e.g. C$_2$H$_2$, HCN) and carbonaceous dust result, whilst in oxygen stars (C/O$<1$), metal oxides and silicate dust form. This dichotomy is not complete: in the inner shell of AGB stars shocks can lead to the formation of some molecules of the 'wrong' type (Decin et al. 2008).  The mass-loss mechanism of AGB stars has become controversial.  A two-step process is likely involved. First, shocks due to pulsations from the star extend the atmosphere (Jones et al., 1981) leading to dust formation.  Second, radiation pressure on the dust drives an outflow; the gas is carried along through friction with the dust particles. Pulsations can drive mass-loss rates up to about $10^{-6}M_{\\odot}$yr$^{-1}$; the much higher rates observed are explained by dust-driven winds (Bowen \\&\\ Wilson 1991).  This simple distinction between weak pulsation winds and dusty superwinds (Winters et al. 2000) has recently been disputed. Woitke (2006) find in their models that the dust opacity of silicates is insufficient to drive a wind, while for carbon dust the winds become very chaotic (Woitke \\& Niccolini 2005).  For low-metallicity stars, dust-driven winds may fail due to the lack of dust (Zijlstra 2004): theoretical work by Bowen \\& Willson (1991) predicts that for metallicities below [Fe/H]$=-1$ the wind becomes pulsation-driven.  Mattson et al. (2008) argue that pulsation can drive a strong wind at low metallicity, set by the pulsation energy of the star. Recent theoretical models produce similar mass-loss rates for AGB stars with SMC, LMC and solar metallicities (Wachter et al. 2008).  The effect of dust should be most noticeable when comparing stars at different metallicity. This has become possible with the {\\it Spitzer Space Telescope} (Werner et al. 2004), which provided the sensitivity necessary to obtain mid-infrared spectra and derive mass-loss rates for the whole range of AGB masses and luminosities in the Magellanic Clouds and a few other Local Group Galaxies. These galaxies trace in general metal-poor stars with a range of metallicity and ages, and the known distances means that absolute mass-loss rates and luminosities can be established. The expectation for reduced mass loss efficiency at low Z has not been confirmed.  Instead, the {\\it Spitzer} surveys have shown that high mass loss rates with dusty winds can be reached even at very low $Z$, as seen in the Fornax dSph galaxy at [Fe/H]$\\sim -1$ (Lagadec et al. 2008).  One unexpected result of the Spitzer surveys is the almost complete dominance of carbon stars among the mass-losing stars in these low metallicity galaxies.  The estimated mass-loss rates for these carbon stars are very similar to the ones observed in Galactic AGB stars, while the few oxygen-rich AGB stars observed have mass-loss rates one or two orders of magnitude lower than their Galactic counterparts (Sloan et al. 2008).  All this seems to indicate that the abundance of carbon and oxygen is an important parameter for the mass-loss from AGB stars at low metallicity.  We will discuss in the present letter at what point during the stellar evolution the superwind begins, and how it is triggered. ",
        "conclusions": "In this letter, we made use of a Spitzer spectroscopic survey of mass-losing AGB star in Local Group low metallicity galaxies to study the superwind from the stars. We have shown that the luminosity at which the superwind is observed is decreasing with lower metallicity for carbon-rich stars. Oxygen-rich stars in the Galaxy and in lower metallicity galaxies have similar mass-loss rates for similar luminosities.  To explain this dichotomy, we propose that the main agent to trigger the superwind in low metallicity carbon stars is the abundance of carbon. When the C/O ratio becomes larger than unity, carbon-rich dust forms and the radiation pressure on this dust becomes more efficient.  We used stellar evolution models to check the consistency of our explanations and made some prediction of the luminosities at which the superwind should be triggered for different metallicities and stars initial masses."
    },
    "0807/0807.3512_arXiv.txt": {
        "abstract": "{} {We report here on the first detection at X-ray wavelengths of the Supernova Remnant (SNR) G337.8$-$0.1, carried out with the XMM-{\\it Newton} Observatory.} {Using the X-ray observations, we studied the X-ray morphology of the remnant at different energy ranges, analysed the spectral properties and investigated a possible variable behavior.} {The SNR shows a diffuse filled-center structure in the X-ray region with an absence of a compact source in its center. We find a high column density of $N_{\\rm H} >$ 6.9$\\times$10$^{22}$ cm$^{-2}$, which supports a relatively distant location ($d \\geq$ 7 kpc). The X-ray spectrum exhibits emission lines, indicating that the X-ray emission has a thin thermal plasma origin, and is well represented by a non-equilibrium ionization (NEI) plasma model. The X-ray characteristics and well-known radio parameters show that G337.8$-$0.1 belongs to the emerging class of mixed-morphology (MM) SNRs.} {} ",
        "introduction": "G337.8$-$0.1 (also known as Kes 41) was discovered with the Molonglo telescope at 408 MHz (Shaver \\& Goss 1970). This object, located in the southern galactic plane, was reported in the Molonglo Observatory Synthesis Telescope (MOST) Catalog (Whiteoak \\& Green 1996) as a bright small-diameter SNR and distorted shell. The remnant is 6'$\\times$9' in size, distinctly elongated in the northeast-southwest direction. Based on HI observations, Caswell et al. (1975) placed the SNR beyond the tangent point at 7.9 kpc. Koralesky et al. (1998) detected maser emission in the object at $-45$ km s$^{-1}$, implying a far kinematic distance of 12.4 kpc. This fact is characteristic of a shock interaction with dense molecular gas (Green et al. 1997).  The MOST flux density of S$_{\\rm 843 MHz}$=18 Jy, combined with S$_{\\rm 408 MHz}$=26 Jy (Shaver \\& Goss 1970), gives a spectral index of $\\alpha \\sim -0.51$ for the remnant.  At X-ray frequencies the object was never detected by early X-ray missions. This non-detection is likely a result of the absorption of the typically soft (i.e., $<$ 2 keV) X-ray emission by the high column density of gas and dust in the Galactic plane. As a part of an effort to identify  and study the X-ray emission of the SNR G337.8$-$0.1 we have used the greatly enhanced X-ray sensitivity provided by the XMM-{\\it Newton} telescope. An instrument with these characteristics offers a unique opportunity to study heavily obscured SNRs.  In this {\\it Letter} we report the first detection at X-ray wavelengths of SNR G337.8$-$0.1. The object is characterized by an apparent centrally brightened X-ray morphology and a spectrum which suggests that the bright central emission is thermal in nature. The structure of the paper is as follows: in Sec.~\\ref{observations} we describe the XMM-{\\it Newton} observations and data reduction. X-ray analysis and results are presented in Sec.~\\ref{analysis}. In Sec.~\\ref{discusion} we discuss the implications of our results and summarize the main conclusions. ",
        "conclusions": "\\label{discusion}  The SNR G337.8$-$0.1 is characterized by an apparent centrally filled X-ray morphology not well correlated with the radio shell, and thermally dominated X-ray emission. Thus, Kes 41 could be classified as a thermal composite (or mixed-morphology) remnant. Moreover, at radio frequencies, the emission is non-thermal (Whiteoak \\& Green 1996) and can be interpreted as synchrotron radiation from accelerated high-energy electrons at the SNR shock. These characteristics place the remnant in the emerging class of SNR that includes objects such as W44 (Jones et al. 1993), 3C 391 (Rho \\& Petre 1996); G272.2$-$3.2 (Harrus et al. 2001) and G290.1$-$0.8 (Slane et al. 2002). A list containing SNRs with these characteristics was proposed by Rho \\& Petre (1998).  Although the evolutionary states that led to the observed properties in these SNRs is not well understood, it is widely believed that these peculiar characteristics are linked to inhomogeneities in the ISM. Several possible scenarios have been introduced in the past to explain thermal X-ray radiation inside radio shells of SNRs, i) cloudlet evaporation in the SNR interior (White \\& Long 1991), ii) thermal conduction smoothing out the temperature gradient across the SNR and enhancing the central density (Cox et al. 1999), iii) a radiatively cooled rim with a hot interior (Harrus et al. 1997), and iv) possible collisions with molecular clouds (Safi-Harb et al. 2005).  G337.8$-$0.1 lies in a complex region of the sky, possibly adjacent to a very massive giant molecular cloud (Dame et al. 1986) and close to the associated maser emission detected at $-45$ km s$^{-1}$ by Koralesky et al. (1998). If the SNR is indeed propagating in a region with a large number of clouds, the most suitable scenario for describing the SNR evolution is that developed by White \\& Long (1991). In such remnants, the clouds evaporate by saturated conduction (Cowie \\& McKee, 1977) and the X-ray morphologies may be very different from ordinary shell-like SNRs.  Based on the X-ray properties of G337.8$-$0.1, specifically the filled-center morphology and detected lines in the spectrum, a thermal interpretation of the X-ray emission is the most plausible to describe the SNR evolution. Using the X-ray image of G337.8$-$0.1 we estimated the volume $V$ of the X-ray emitting plasma, assuming that the plasma fills an ellipsoid similar to the region from which the X-ray spectra was extracted (with a diameter of 5'$\\times$3'.2). For a distance range from 7 to 12 kpc, we then obtained a volume of (2.8-14.8)$\\times$10$^{58}$ cm$^{3}$. Based on the emission measure (EM) determined by the spectral fitting, we can estimate the electron density of the plasma, n$_{e}$, by n$_{e}$=$\\sqrt{EM/V}$, which varies between 0.20 cm$^{-3}$ and 0.26 cm$^{-3}$. In this case, the number density of the nucleons was simply assumed to be the same as that of electrons. The range of age $t$ was then determined from the ionization timescale, $\\tau$, by $t$=$\\tau$/n$_{e}$. Therefore, the elapsed time after the plasma was heated is within (12.000-16.000) yr. The total mass of the plasma $M_{total}$ was estimated by $M_{total}$=n$_{e}$V\\,m$_{\\rm H}$ = (5\\,-\\,32) M$_{\\odot}$, where m$_{\\rm H}$ is the mass of a hydrogen atom. These results are consistent with the expected values for a middle-age SNR.  One of the most detailed studies of middle-age and center-filled thermal emission in a SNR interacting with a dense molecular cloud is the W44 model (Cox et al. 1999; Shelton et al, 1999; Shelton et al. 2004). The proximity to molecular clouds of this SNR suggests that the environmental interstellar matter, where the SNR is propagating, is relatively dense. Using Cox's model we have roughly estimated some dynamical characteristics (i.e., the shock radius $R_{\\rm s}$ and shell velocity $v_{\\rm shell}$) for G337.8$-$0.1. Assuming an energy for the explosion of $\\sim$ 10$^{51}$ ergs, we found that $R_{\\rm s}$= (8-10) pc and $v_{\\rm shell}$= (180-200) km s$^{-1}$. These values are consistent with the angular size and range of distance for the SNR under consideration. Thus, G337.8$-$0.1 could be, a SNR with characteristics similar to W44.  Finally, we have compared the X-ray luminosity of G337.8$-$0.1 with other proto-typical mixed-morphology SNRs. The unabsorbed X-ray luminosities in the 0.5-10 keV band of G337.8$-$0.1 for a range of distance between 7 and 12 kpc is L$_{\\rm X} \\sim$ 4.5$\\times$10$^{34}$ erg/s and 1.3$\\times$10$^{35}$ erg/s, respectively. The first value is similar to the X-ray luminosity of W28 (L$_{\\rm X}$$\\sim$6$\\times$10$^{34}$ erg s$^{-1}$; Rho \\& Borkowski, 2002), G272.2-3.2 (L$_{\\rm X} \\sim$ 5.2$\\times$10$^{34}$ erg s$^{-1}$; Harrus et al. 2001), and Kes 32 (L$_{\\rm  X}$ $\\sim$ 3.2$\\times$10$^{34}$ erg s$^{-1}$; Vink 2008) for the same energy range. This suggests that a distance around 7 kpc for G337.8$-$0.1 is most probable.  In summary, in this work we reported the first detection at X-ray wavelengths of the SNR G337.8$-$0.1. The radio and X-ray properties show that the object displays the three basics attributes found in mixed-morphology SNRs, i.e. centrally filled X-ray morphology with lines in its spectra, a shell-like appearance at radio frequencies, and the absence of a prominent, central, compact source in radio and X-ray energies. If the SNR is indeed propagating in a cloudy medium, the interaction of the shock-front with the adjacent molecular material could be responsible for the gamma-ray emission detected by the EGRET telescope.  The detailed properties of the SNR are still unknown because of the low photon statistics in the XMM-{\\it Newton} data. Future X-ray observations with a higher spatial resolution and good photon statistics like Chandra or Suzaku will reveal the relation between the X-ray intensity peak and the radio emission."
    },
    "0807/0807.3724_arXiv.txt": {
        "abstract": "We present HST/NICMOS Paschen $\\alpha$ (Pa$\\alpha$) images and low and high resolution IRS spectra of photoevaporating disk-tail systems originally detected at 24$\\mu$m near O stars. We find no Pa$\\alpha$ emission in any of the systems. The resulting upper limits correspond to about 2$-$3 $\\times$ 10$^{-6}$ M$_{\\odot}$ of mass in hydrogen in the tails suggesting that the gas is severely depleted. The IRAC data and the low resolution 5$-$12 $\\mu$m IRS spectra provide evidence for an inner disk while high resolution long wavelength (14$-$30 $\\mu$m) IRS spectra confirm the presence of a gas free ``tail'' that consists  of $\\sim$ 0.01 to $\\sim$ 1 $\\mu$m dust grains originating in the outer parts of the circumstellar disks. Overall our observations support theoretical predictions in which photoevaporation removes the gas relatively quickly ($\\leq 10^5$ yrs) from the outer region of a protoplanetary disk but leaves an inner more robust and possibly gas-rich disk component of radius 5-10 AU. With the gas gone, larger solid bodies in the outer disk can experience a high rate of collisions and produce elevated amounts of dust. This dust is being stripped from the system by the photon pressure of the O star to form a gas-free dusty tail. ",
        "introduction": "The timescale at which a protoplanetary disk loses its primordial gas, which dominates the total disk mass, is a crucial factor regarding the outcome of the planet forming process in the vicinity of high mass stars. The photoevaporation of protoplanetary disks by external UV radiation in an environment dominated by high mass O-stars is a particularly potent mechanism for removing disk material. The timescale of photoevaporation is governed by the amount of UV radiation from the nearby O-stars, which significantly erodes the protoplanetary disk in less than a million years. Studies of disk erosion agree that the outer part of the disk evaporates quickly; however according to, e.g., \\citet{Holl04} and \\citet{Adam04} the erosion of the disk slows down significantly when the size reaches 10 AU. Other authors \\citep[e.g.][]{John98,Mats03b} have found that the photoevaporation process can shrink the radius of the disk below 1 AU within $10^6$ years and it can even destroy the whole disk when it is coupled with internal viscous evolution. Therefore, the presence of high mass stars might be lethal for forming planets since the EUV and FUV radiation of a nearby O-star can disrupt the protoplanetary disk on a short time scale (a couple of times $10^5$ yr) and stop the planet formation process. On the other hand the radiation might trigger the formation of rocky planets by accelerating the dust settling and grain growth in the midplane of the disk \\citep{Thro05}.   \\citet{Balo06} reported the discovery of three comet like structures near different O-type stars (with the comet-like tails pointing away from the central high mass star) in 24 $\\mu$m {\\it Spitzer}/MIPS images. Adopting a radiation-pressure driven outflow disk model \\citep{Su05} they concluded that these objects are photoevaporating protoplanetary disks where dust entrained in the photoevaporative flow is responsible for the mid-infrared tails detected at 24 $\\mu$m. The tails extend up to 0.2 pc from the evaporation working surface, so they are convenient for detailed study.  To investigate this phenomenon further and to get a better understanding of the nature of these objects, we have searched these three systems for the gas expected to be a significant component of the flow. We used Pa$\\alpha$ as an ideal tracer of the photoevaporation process, since it is the strongest emission line in the near infrared and is not strongly affected by extinction. We also report mid-IR spectra to provide a better understanding of the physical processes in the disk and to help determine the properties of the dust. ",
        "conclusions": "We present HST/NICMOS Pa$\\alpha$ images and IRS spectra of cometary structures detected in {\\it Spitzer}/MIPS 24 $\\mu$m images. We estimate an upper limit to the amount of gas in the comets' disk and tail and find that the gas-to-dust mass ratio is much lower than the value observed in the ISM: the tails are essentially gas free. Using this new observation we are able to constrain the flow velocity and thus the mass loss rate. The new mass loss rates allow us to estimate the timescale on which the phenomenon occurs ($10^5 - 10^6$ yr). The short timescale favors photoevaporation models predicting quick removal of gas from the outer parts of the disk.  The ``comet tails'' are produced from the outer regions of the disks, where larger grains collide at an elevated rate generating second generation dust. These small grains are then ejected due to photon pressure from the nearby O-star. The SED of the sources shows excess emission between 3 and 8 $\\mu$m, in agreement with the IRS low resolution spectra. This emission indicates that there is an inner disk that survives the photoevaporation process for longer than $10^5$ years, as predicted by the photoevaporation models."
    },
    "0807/0807.1721_arXiv.txt": {
        "abstract": "After launch, the Advanced CCD Imaging Spectrometer (ACIS), a focal plane instrument on the {\\it Chandra X-ray Observatory}, suffered radiation damage from exposure to soft protons during passages through the Earth's radiation belts. An effect of the damage was to increase the charge transfer inefficiency (CTI) of the front illuminated CCDs. As part of the initial damage assessment, the focal plane was warmed from the operating temperature of --100$^\\circ$~C to +30$^\\circ$~C which unexpectedly further increased the CTI. We report results of ACIS CCD irradiation experiments in the lab aimed at better understanding this reverse annealing process. Six CCDs were irradiated cold by protons ranging in energy from 100~keV to 400~keV, and then subjected to simulated bakeouts in one of three annealing cycles. We present results of these lab experiments, compare them to our previous experiences on the ground and in flight, and derive limits on the annealing time constants. ",
        "introduction": "\\label{sec:intro}  The Advanced CCD Imaging Spectrometer (ACIS), one of two focal plane science instruments on the {\\it Chandra X-ray Observatory}, utilizes frame-transfer charge-coupled devices (CCDs) of two types, front- and back-illuminated (FI and BI).  Soon after launch it was discovered that the FI CCDs had suffered radiation damage from exposure to soft protons scattered off the Observatory's grazing-incidence optics during passages through the Earth's radiation belts\\cite{gyp00,odell}.  The BI CCDs are less sensitive to the low energy particles which damaged the FI CCDs because the particles cannot reach the transfer channel.  In an attempt to anneal the radiation damage, the focal plane was warmed from the operating temperature of --100$^\\circ$~C to +30$^\\circ$~C.  Unexpectedly, the warmup had the reverse effect and caused a substantial increase in the charge transfer inefficiency.  To better understand this anomalous annealing, a laboratory experiment was developed.  Our initial experiment in low temperature irradiation and room-temperature annealing took place in September 2002 and was reported in Ref.~\\citenum{bakepaper1} (Paper I).  A single CCD was irradiated with 120~keV protons while at a temperature of --100$^\\circ$~C, then warmed to +30$^\\circ$~C for 8~hours.  The entire irradiation/annealing cycle was designed to duplicate the experience of the flight instrument.  In addition to the expected substantial increase in CTI due to the initial irradiation, we found the warmup itself produced an additional factor of $\\approx 2.5$ increase in CTI.  The room temperature bakeout also produced dramatic changes in the de-trapping times of the electron traps that cause CTI.  While the initial experiment qualitatively duplicated the ACIS on-orbit experience, the relative magnitude of the reverse annealing was larger in the ground data by a factor of four.  We speculated that the room-temperature annealing changed the trap population, converting traps with relatively benign re-emission time constants to more harmful ones.  In particular, the ACIS CCDs contain carbon impurities.  Irradiation of the CCDs creates vacancies and silicon interstitials which are highly mobile at ACIS operating temperatures.  The silicon interstitials can exchange places with the carbon forming a carbon interstitial ($C_i$) which is not mobile at ACIS operating temperatures and does not affect CTI.  At room temperatures, however, the carbon interstitials are free to form defects that can negatively affect CTI, such as $C_i$-$P_s$.  In this paper we describe a follow-up experiment, undertaken in May 2005, with multiple CCDs, irradiation/annealing profiles, and proton energies to better explore the physics behind the anomalous annealing observed in Paper I.  Our goals were to understand why our ground test results differed from the on-orbit bakeout experience, to test our model of the ``reverse annealing'' phenomenon and to gain greater confidence in predicting CCD changes resulting from any future ACIS bakeout.  By irradiating multiple CCDs we can determine the importance of chip-to-chip variations.  The energy dependence of annealing can be tested using multiple proton energies and different annealing profiles can be used to study the timescales of the phenomenon.  Section~\\ref{sec:exper} describes the experimental setup and charge transfer inefficiency data analysis.  Section~\\ref{sec:results} examines the properties of the anomalous annealing, while section~\\ref{sec:disc} compares these results to our previous experiment. ",
        "conclusions": "We report the results of laboratory experiments of cold irradiation and room-temperature annealing on ACIS CCDs.  Six CCDs were irradiated cold, at --100$^\\circ$~C, by low-energy protons with energies from 100 to 400~keV.  The CCDs were then annealed at +30$^\\circ$~C in one of three annealing cycles.  We used the fractional CTI increase due to annealing, $R$, to study the dependence on proton energy and intrinsic properties.  The measured value of $R$, averaged over all CCD quadrants was 0.17 with a standard deviation of 0.09.  This variation is larger than, but roughly consistent with the variance in concentration of carbon impurities in the ACIS CCDs.  We found a mild dependence on proton energy, increasing from 100 to 180~keV, then decreasing to 400~keV.  The long duration annealing tests showed that the processes which increase CTI appear to be complete after eight hours.  The results were generally consistent with ACIS flight experience, but highly discrepant with our original laboratory experiment.  We believe this is due to differences in CCD temperature.  To better understand the characteristics of the traps that cause the annealing-induced CTI increase, we plan to continue study of these data.  The results of the multi-temperature isochronal annealing, with an apparent CTI {\\it decrease} at the lower temperatures (0$^\\circ$~C to +20$^\\circ$~C) before the final increase at +30$^\\circ$~C, are puzzling and we do not currently have an explanation for this behavior.  It is likely that study of the charge in the pixels trailing the X-ray events may prove illuminating, as it was in Paper I, indicating changes in the re-emission timescales of the trap population at each annealing stage.  The trailing charge will also help further explore the differences and similarities between the two ground experiments and the flight case."
    },
    "0807/0807.1983_arXiv.txt": {
        "abstract": "In the present study time evolution of quasi-spherical polytropic accretion flow with toroidal magnetic field is investigated. The study especially focused the astrophysically important case in which the adiabatic exponent $\\gamma=5/3$. In this scenario, it was assumed that the angular momentum transport is due to viscous turbulence and used $\\alpha$-prescription for kinematic coefficient of viscosity. The equations of accretion flow are solved in a simplified one-dimensional model that neglects the latitudinal dependence of the flow. In order to solve the integrated equations which govern the dynamical behavior of the accretion flow, self-similar solution was used. The solution provides some insight into the dynamics of quasi-spherical accretion flow and avoids many of the strictures of the steady self-similar solution. The effect of the toroidal magnetic field is considered with additional variable $\\beta[=p_{mag}/p_{gas}]$, where $p_{mag}$ and $p_{gas}$ are the magnetic and gas pressure, respectively. The solution indicates a transonic point in the accretion flow, that this point approaches to central object by adding strength of the magnetic field. Also, by adding strength of the magnetic field, the radial-thickness of the disk decreases and the disk compresses. It was analytically indicated that the radial velocity is only a function of Alfv'en velocity. The model implies that the flow has differential rotation and is sub-Keplerian at all radii. ",
        "introduction": "Accretion is the main source of energy in many astrophysical objects including different types of binary stars, binary X-ray sources, quasars, and Active Galactic Nuclei (AGN). Though the first development of accretion theory started a long time ago (Bondi \\& Hoyle 1944; Bondi 1952), intensive development of the theory began after the discovery of the first X-ray sources (Giacconi et al. 1962) and quasars (Schmidt 1963). Since the removal of the angular momentum process operates on slower timescales as compared to free-fall time, the infalling gas with sufficiently high angular momentum can form a disklike structure around a central gravitating body that can be thin or thick depending upon their geometrical shapes. The models of thin accretion disks are perhaps better developed and seem to have good observational basis (Shakura \\& Sunyaev 1973). However, for the thick accretion disks, no fully developed model exists, and there remain many theoretical uncertainties about their structure and stability (Banerjee et al. 1995; Ghanbari \\& Abbassi 2004; Ghanbari et al. 2007).  It is thought that accretion disks, whether in star-forming regions, in X-ray binaries, in cataclysmic variables, or in the centers of active galactic nuclei, are likely to be threaded by magnetic fields. Consequently, the role of magnetic fields has been analyzed in detail by a number of investigators (Blandford \\& Znajek 1977; Lubow et al. 1994; Banerjee et al. 1995; Shadmehri 2004).A mechanism for angular momentum transport is another key ingredient in theory of accretion processes and still many theoretical uncertainties remain about its nature. As originally pointed out in Lynden-Bell (1969) and Shakura \\& Sunyaev (1973), a magnetic field can also contribute to the angular momentum transport. A robust mechanism of the excitation of magnetohydrodynamical (MHD) turbulence was shown to operate in accretion disks due to the Magneto-Rotational Instability (MRI) (Balbus \\& Hawley 1998; Machida et al. 1999; Begelman \\& Pringle 2007).  The toroidal magnetic fields have been observed in the outer regions of YSO discs ((Aitken et al. 1993; Wright et al. 1993; Greaves et al. 1997) and in the Galactic center (Novak et al. 2003; Chuss et al. 2003). Accretion disks containing toroidal magnetic field have been studied by several authors (Fukue \\& Okada 1990; Geroyannis \\& Sidiras 1992,1993,1995; Banerjee et al. 1995; Terquem \\& Papaloizou 1996; Machida et al. 1999; Liffman \\& Bardou 1999; Rempel 2006; Begelman \\& Pringle 2007; Akizuki and Fukue 2006). Fukue and Okada (1990) examined the oscillations of a gaseous disk which were penetrated by toroidal magnetic fields. Geroyannis and Sidiras (1992,1993) described differentially rotating polytropic models distorted by toroidal magnetic fields. Also, Geroyannis and Sidiras (1995) considered dissipative effects by viscous friction of differentially rotating visco-polytropic models that were further distorted by a toroidal magnetic field. Banerjee et al. (1995) presented a toroidal magnetic field that was generated by interaction rotating plasma and dipolar magnetic field of central object; they showed that toroidal magnetic field has an important effect in structure of the disk. Terquem and Papaloizou (1996) studied the a linear stability of a differentially rotating disk containing a purely toroidal magnetic field. They presented disks containing a purely toroidal magnetic field are always found to be unstable. Machida et al. (1999) considered three-dimensional global magnetohydrodynamical simulation of a torus treated by toroidal magnetic fields. Akizuki and Fukue (2006, hereafter AF) examined the effect of toroidal magnetic field on a viscous gaseous disk around a central object under an advection dominated stage. Assuming steady and axisymmetric flow and using steady self-similar method, they presented the nature of the disk was significantly different from that of the weakly magnetized case.\\\\ In this study, we want to explore how the dynamic of a rotating and accreting viscous gas depends on its toroidal magnetic field. By solving MHD equations for accreting gases that are self-similar in time, we will answer this question. We assume that turbulent viscosity is due to angular momentum transport of the fluid and there is efficient radiation cooling in the flow. This paper is organized as follows. In section 2, the general problem of constructing a model for quasi-spherical magnetized polytropic accretion flow is defined. In section 3, self-similar method for solving the integrated equations which govern the dynamical behavior of the accreting gas is utilized. The summary of the model is presented in section 4. ",
        "conclusions": ""
    },
    "0807/0807.0382_arXiv.txt": {
        "abstract": "\\noindent We argue that the curvaton decay takes place most naturally by way of a broad parametric resonance. The mechanism is analogous to resonant inflaton decay but does not require any tuning of the curvaton coupling strength to other scalar fields. For low scale inflation and a correspondingly low mass scale for the curvaton, we speculate on observable consequences including the possibility of stochastic gravitational waves ",
        "introduction": "The curvaton scenario \\cite{curvaton, Lyth, Moroi,earlier} provides an alternative mechanism for producing the primordial density perturbations. The curvaton is a light field which is subdominant during inflation but whose energy density will eventually become an important (if not actually the dominant) energy component in the universe. Because of its lightness, during inflation the curvaton acquires fluctuations which at the time of curvaton decay are transformed into curvature perturbations.  So far, the decay of the curvaton has mostly been treated as a perturbative process, see e.g.  \\cite{curvatondecay,LUW}, where the curvaton condensate decays into relativistic particles, determined by a rate $\\Gamma$ and occurring roughly when $H\\sim \\Gamma$. This happens at the stage when the curvaton field is oscillating about the minimun of its potential, whence it behaves like pressureless dust so that its relative energy density grows with respect to radiation produced by the inflaton decay. The curvature pertubation induced by the curvaton decay is mostly adiabatic but can give rise to a large non-gaussianity \\cite{LUW,curvatonng} if the curvaton energy density is too small\\footnote{Strictly speaking, this holds only for a quadratic curvaton potential, see \\cite{sami}.}; this implies a lower limit on the fraction of the energy density in the curvaton condensate at the decay time. Another generic feature of curvaton models with no mixed inflaton and curvaton perturbations is the small amplitude tensor perturbations \\cite{Lyth, LUW}; this is a consequence of the low inflationary scale required to keep the inflaton perturbations negligible. \\footnote{ See however \\cite{vaihkonenetal} where it is demonstrated  that the amplitude of CMB-scale tensor perturbations in the curvaton  model can be directly proportional to the amount of the produced non-gaussianity due to second order effects.}  However, since the curvaton is a coherently oscillating condensate, it is possible to excite other bose fields in a regime of parametric resonance, in a manner similar to inflationary preheating \\cite{preheating,linde}. As a consequence, the curvaton decay process could have important cosmological ramifications for e.g. non-gaussianities, which are known to be enhanced \\cite{enhance, arttu} in some preheating scenarios, as well as for stochastic gravitational waves, as will be discussed below. Preheating is a violent process in which the occupation numbers of fields coupled to the curvaton grow exponentially. In effect, dense lumps of decay products are created very rapidly and randomly in the real space, and as these highly inhomogeneous concentrates of energy density collide, there arise large-amplitude, high-frequency gravitational waves, as has been discussed recently \\cite{dufeau, gw}. These do not leave a signal on the CMB but could possibly be observed in current or future gravitational wave detectors such as LIGO \\cite{ligo}, LISA \\cite{lisa} or BBO \\cite{bbo}.  In the present paper we study the possibility of preheating in generic curvaton models and explore the main consequences of such a scenario (preheating has been discussed in the context of MSSM flat direction curvatons in \\cite{postma}). Indeed, the curvaton has to couple to other fields if it is to decay; if it couples to other scalars, then there would be two options. Either the additional scalars would be light and shifted from their minima during inflation, in which case they would simply be additional curvatons, or they could be massive and driven to their minima. In the first case one still would have to find a way for the curvaton(s) to decay. However, in the latter case there can arise a parametric resonance during curvaton oscillations with the ensuing rapid decay of the curvaton.  We find that the preheating process in the curvaton scenario is qualitatively similar to the preheating after inflation although the resonance dynamics can be somewhat altered. This is because the curvaton, unlike the inflaton, could remain subdominant during the entire preheating stage. Moreover, in the scenario studied here the backreaction of the excited bose fields always becomes significant at the end of preheating whereas in the standard inflationary case the preheating may terminate already at an earlier stage. Since the gaussian part of curvaton perturbations on superhorizon scales is not affected by preheating and the final decay of the curvaton will eventually take place in a perturbative reheating process, the possible preheating stage does not alter the standard predictions of the curvaton model \\cite{LUW}. However, interestingly enough we find that preheating may generate stochastic gravitational waves with an amplitude much larger than the inflationary tensor perturbations in the curvaton model, and we discuss the feasibility of observing such signatures of curvaton preheating. The study of the non-gaussianities is left for a future work. ",
        "conclusions": "Up to now, discussions of curvaton dynamics have mostly focused on perturbative decay of the curvaton condensate. Here we point out that if the curvaton is about to decay directly into bosonic degrees of freedom, it is highly likely that there will be a non-perturbative preheating stage before the perturbative decay and reheating. The curvaton preheating is qualitatively similar to the conventional inflationary preheating although minor differences can arise due to the fact that the curvaton does not need to dominate the universe during the preheating stage. We have also found that the end of the curvaton preheating is always dominated by the effects of backreaction of the amplified light bose fields. This is different from the inflationary case, where the importance of the backreaction depends on the parameters. Although we have not discussed the details of backreaction, we expect to find no significant differences as compared to the inflationary case here either. This is because the possible differences between the curvaton preheating and the inflationary preheating stem from differences in the expansion of the universe which are unimportant for the final stages of preheating. However, to complete the analysis, one should study numerically the thermalization of the universe after the end of preheating. Since the curvature perturbations generated in the curvaton scenario are sensitive to the evolution history before the curvaton decay, these stages might potentially affect even the gaussian perturbations and therefore yield interesting constraints on the curvaton model itself.  Numerics is also needed to resolve the issue of stochastic gravitational waves. Based on simple estimates we have postulated that models with a sufficiently low inflationary scale could produce gravitational waves with directly detectable frequencies and amplitudes. For example, supersymmetric models \\cite{flat} with the curvaton mass around the TeV scale could give rise to observable gravitational waves provided that the inflationary scale is not several orders of magnitude higher. However, our results are still preliminary and a full numerical simulation, subject to a future research, is required before any firm conclusions can be drawn.  Likewise, the possible enhancement of the non-gaussianities in curvaton preheating is best addressed by numerical simulations. Since curvaton preheating is very similar to the inflaton preheating, we may expect enhancements similar as has been found numerically by Chambers and Rajantie \\cite{arttu}. It therefore seems that curvaton decay by a parametric resonance, in addition to being a natural stage in the curvaton scenario, could also have interesting observational consequences.  \\vskip30pt  \\centerline{\\textbf{Acknowledgements}} \\vskip10pt We wish to thank J.~F.~Dufaux for critical comments as well as A.~Rajantie and A.~Tranberg for illuminating discussions. This work was supported by the EU 6th Framework Marie Curie Research and Training network \"UniverseNet\" (MRTN-CT-2006-035863) and partly by Academy of Finland grant 114419. S.N. is supported by the GRASPANP Graduate School."
    },
    "0807/0807.2949_arXiv.txt": {
        "abstract": "{} { The purpose of this study is to find transient X-ray sources in M~31, and to investigate and classify their nature. } { Three X-ray transients were observed with \\textit{Swift}. For each of the three X-ray transients we use the \\textit{Swift} X-ray and optical data together with observations from \\textit{XMM-Newton} and \\textit{Chandra} to investigate the lightcurves and the spectra of the outburst, and thereby to identify the source types. } { The outburst of XMMU J004215.8+411924 lasted for about one month. The source had a hard power-law spectrum with a photon index of 1.6. It was previously identified as a Be/X-ray binary based on the optical identification with a star. However, we show that with improved source coordinates it is clear that the optical source is not the counterpart to the X-ray source. The source SWIFT J004217.3+411532 had a bright outburst, after which it slowly decayed over half a year. The spectrum was soft, corresponding to a thermal accretion disk with innermost temperature of $\\sim250-600$ eV. The source was not seen in the optical, and the soft spectrum indicates that the source is most likely a black hole low mass X-ray binary. M31N 2006-11a is a nova that was previously observed in the optical. We detected it both in X-rays and UV with \\textit{Swift} $\\sim$half a year after the optical maximum, after which it decayed below the \\textit{Swift} detection threshold within a month. The spectrum of the X-ray transient can be modelled by a black-body with a temperature of 50 eV. We use catalogues of X-ray transients in M~31 to estimate their rate, and we find a lower limit of 9 yr$^{-1}$. } {} ",
        "introduction": "Being the nearest large spiral galaxy, M~31 is an important target for the study of X-ray point sources in galaxies. In a large number of studies, X-ray point sources in M~31 have been investigated and catalogued, using \\textit{Einstein} \\citep{Trinchieri}, \\textit{ROSAT} \\citep{Primini,Supper}, \\textit{XMM-Newton} \\citep[e.g.][]{Trudolyubov2,Pietsch} and \\textit{Chandra} \\citep[e.g.][]{Kong,Kaaret,Williams1,Williams2,Voss}.\\\\ An interesting subset of these sources are the transients. These are sources with quiescent luminosities below the detection limit ($\\sim 10^{35}$ erg s$^{-1}$) that display bright outbursts typically lasting from weeks to a few months. From Galactic observations it is well-known that there are three types of objects that can show such outbursts. Optical studies of soft X-ray transients have shown that they contain an accreting compact object more massive than 3 $M_{\\odot}$, revealing the presence of black hole accretors \\citep{McClintock}. In contrast the harder sources typically show pulsations, indicating that the accreting object is a neutron star \\citep[e.g.][]{Tanaka,Campana}. The third class is the thermal emission from classical novae in optical decline, showing up as super-soft X-ray transients \\citep{MacDonald,Starrfield,Pietsch2007}. Monitoring programs have shown that 6--12 transients appear per year in the bulge of M~31 \\citep{Trudolyubov}, and within the timescale of the X-ray observations, only a small fraction of them have been observed to recur.\\\\ Here we report on the detection and analysis of three transient X-ray sources in the bulge of M~31, using data from \\textit{Swift}, \\textit{Chandra} and \\textit{XMM-Newton}.  \\begin{table*} \\begin{center} \\caption{The list of the observations used for the analysis of the three X-ray transients, TR1 (XMMU J004215.8+411924), TR2 (SWIFT J004217.4+411532) and TR3 (M31N 2006-11a). In the columns of the three sources, a \"+\" signifies that the source was detected in this observation, whereas a \"-\" signifies that it was not detected.} \\label{tab:obs} \\begin{tabular}{llccccccc} \\hline\\hline Observation &Date & Observatory & Instrument & Obs-ID & TR1 & TR2 & TR3\\\\ \\hline 1 &2006 Jun 05 & Chandra & HRC-I & 7283       & - & - & -\\\\ 2 &2006 Jul 02 & XMM     & EPIC  & 0405320501 & - & - & -\\\\ 3 &2006 Jul 31 & Chandra & ACIS-I& 7139       & + & - & -\\\\ 4 &2006 Aug 09 & XMM     & EPIC  & 0405320601 & + & - & -\\\\ 5 &2006 Sep 01 & Swift   & XRT   & 00030802001& + & + & -\\\\ 6 &2006 Sep 11 & Swift   & XRT   & 00030804001& - & + & -\\\\ 7 &2006 Sep 24 & Chandra & ACIS-I& 7140     &-&+& \\\\ 8 &2006 Sep 30 & Chandra & HRC-I & 7284\t   & - & + & -\\\\ 9 &2006 Nov 13 & Chandra & HRC-I & 7285       & - & + & -\\\\ 10& 2006 Dec 4 & Chandra & ACIS-I & 7064 &  & + & \\\\ 11&2006 Dec 31 & XMM     & EPIC  & 0405320701 & - & + & -\\\\ 12&2007 Jan 16 & XMM     & EPIC  & 0405320801 & - & + & -\\\\ 13&2007 Mar 11 & Chandra & HRC-I & 7286       & - & - & -\\\\ 14&2007 Jun 01 & Swift   & XRT   & 00035336001& - & - & +\\\\ 15&2007 Jul 12 & Swift   & XRT   & 00030968001& - & - & -\\\\ 16&2007 Jul 13 & Swift   & XRT   & 00030968002& - & - & -\\\\ 17&2007 Jul 18 & Swift   & XRT   & 00030968004& - & - & -\\\\ \\hline \\end{tabular} \\end{center} \\end{table*} ",
        "conclusions": "For the first time, \\textit{Swift} was used to observe X-ray transients in M~31. We analysed three sources, two of which were discovered from the \\textit{Swift} observations, thereby showing the possibilities of using \\textit{Swift} for such studies. We furthermore made use of all available \\textit{Chandra} and \\textit{XMM-Newton} observations to analyse the developments in luminosity and spectrum for each of the sources. We investigated the rate of transients in M~31 and we found a minimum rate of 9 yr$^{-1}$, and that $\\sim$half of the sources have a hard spectrum."
    },
    "0807/0807.0457_arXiv.txt": {
        "abstract": "A model of inflation is proposed in which compact extra dimensions allow a graceful exit without recourse to flat potentials or super-Planckian field values. Though bubbles of true vacuum are too sparse to uniformly reheat the Universe by colliding with each other, a compact dimension enables a single bubble  to uniformly reheat by colliding with itself. This mechanism, which generates an approximately scale invariant perturbation spectrum, requires that inflation be driven by a bulk field, that vacuum decay be slow, and that the extra dimension be at least a hundred times larger than the false vacuum Hubble length. ",
        "introduction": "Many theories posit compact extra dimensions for reasons unrelated to cosmology \\cite{ued,add,rs} --- this paper will show that in doing so you may get a free bonus: a graceful exit from inflation without recourse to flat potentials or super-Planckian field values, and with a mechanism for generating approximately scale invariant perturbations.  The Universe is very flat. The spatial curvature radius, $l_{curv}  (\\textrm{today})$, is observed to be no less than ten times the current Hubble length, $H_0^{-1}$ \\cite{wmap}. Using the entropy density, $S$, we can construct the dimensionless (and, for adiabatic evolution, conserved) quantity \\begin{equation} l_{curv} \\,  S^{1/3} > e^{70}. \\label{entropy} \\end{equation} This number is huge. Inflation explains why.  \\begin{figure}[h] \\centering \\includegraphics[width=.4\\textwidth]{GE3.pdf} \\centering \\caption{The inflaton potential both for old inflation and for boom and bust inflation.} \\label{fig1} \\end{figure}  In old inflation \\cite{guth}, the Universe gets temporally stuck in a false vacuum during an early first order phase transition. While in this false vacuum the Universe (adiabatically) inflates, safely diluting primordial entropy and monopole density, and solving the horizon problem. Inflation ends, locally, when the false vacuum decays by nucleating bubbles of true vacuum \\cite{coleman} and the Universe (nonadiabatically) reheats, locally, when many such bubbles collide with one another. But old inflation does not work: if inflation is to last sufficiently long to do its job then decay must be slow, so bubbles are sparse, collisions are rare, and percolation fails \\cite{guth,gw}.  New inflation \\cite{slowroll} does work. It instead recovers the observable Universe from a single bubble --- the interior of a bubble has the geometry of an open Friedmann-Robertson-Walker (FRW) Universe. New inflation adds an exceptionally flat segment on the true vacuum side of the potential barrier so that inflation continues after nucleation while the field rolls slowly down. This post-nucleation period of inflation continues until the field reheats at the bottom of the potential, by which time the curvature radius has grown sufficiently to satisfy Eq.~(\\ref{entropy}).  In this paper we introduce a new model of inflation, which we call boom and bust inflation. Like old inflation it makes do with an unexceptional inflaton potential. Like old inflation it reheats via bubble collision. But like new inflation it recovers the visible Universe from a single bubble. To do so it borrows structure that has been independently invoked: compact extra dimensions\\footnote{For other inflationary scenarios involving bubbles in extra dimensions, see \\cite{TandB}.}. ",
        "conclusions": ""
    },
    "0807/0807.2452_arXiv.txt": {
        "abstract": "We present measurements of the luminosity and mass functions of low-mass stars constructed from a catalog of matched Sloan Digital Sky Survey (SDSS) and 2 Micron All Sky Survey (2MASS) detections. This photometric catalog contains more than 25,000 matched SDSS and 2MASS point sources spanning $\\sim$30 square degrees on the sky. We have obtained follow-up spectroscopy, complete to J=16, of more than 500 low mass dwarf candidates within a 1 square degree sub-sample, and thousands of additional dwarf candidates in the remaining 29 square degrees. This spectroscopic sample verifies that the photometric sample is complete, uncontaminated, and unbiased at the 99\\% level globally, and at the 95\\% level in each color range.  We use this sample to derive the luminosity and mass functions of low-mass stars over nearly a decade in mass (0.7 M$_{\\odot} >$ M$_* >$ 0.1 M$_{\\odot}$).  The luminosity function of the Galactic disk is statistically consistent with that measured from volume complete samples in the solar neighborhood.  We find that the {\\it logarithmically binned} mass function is best fit with an M$_c$=0.29 log-normal distribution, with a 90\\% confidence interval of M$_c$=0.20--0.50.  These 90\\% confidence intervals correspond to {\\it linearly binned} mass functions peaking between 0.27 M$_{\\odot}$ and 0.12 M$_{\\odot}$, where the best fit MF turns over at 0.17 M$_{\\odot}$.  A power law fit to the entire mass range sampled here, however, returns a best fit of $\\alpha$=1.1 (where the Salpeter slope is $\\alpha$ = 2.35); a broken power law returns $\\alpha$=2.04 at masses greater than log M = -0.5 (M=0.32 M$_{\\odot}$), and $\\alpha$=0.2 at lower masses.  These results agree well with most previous investigations, though differences in the analytic formalisms adopted to describe those mass functions, as well as the range over which the data are fit, can give the false impression of disagreement.  Given the richness of modern-day astronomical datasets, we are entering the regime whereby stronger conclusions can be drawn by comparing the actual datapoints measured in different mass functions, rather than the results of analytic analyses that impose structure on the data {\\it a priori}.  Having validated this method to generate a low-mass luminosity function from matched SDSS/2MASS datasets, future studies will extend this technique to the entirety of the SDSS footprint. ",
        "introduction": "The mass function (MF) is a fundamental property of stellar systems, describing the number of stars as a function of stellar mass.  The MF is thus a statistical measure of the end result of the star formation process.  Succinctly characterizing a stellar population, the MF also informs our understanding of the structure and dynamical evolution of stellar clusters, the Milky Way and other galaxies.   The first measurement of the MF was conducted by \\citet{Salpeter1955}; for higher mass stars, this seminal result remains essentially unchanged to the present day. Salpeter found that the MF can be fit as a power law, formally expressed as:  \\begin{equation} \\label{fundamental mass function} \\Psi(M) = \\frac{dN}{dM} \\propto M^{-\\alpha}~~stars~~pc^{-3}~~M_{\\odot}^{-1} \\end{equation}  \\noindent where $\\alpha =$ 2.35 is known as the `Salpeter slope'.  Power law MFs which find a larger fraction of low-mass stars have larger values of alpha and are expressed as being `steeper', while mass functions that find a lower fraction of low-mass stars have smaller values of $\\alpha$ and are described as `flatter'.  Other functional forms have been suggested to characterize the MF. In particular, a number of investigators \\citep{Miller1979,Chabrier2005a} present log normal MFs, which can be expressed as  \\begin{equation} \\label{eq:bestfitlognormalmf} \\xi(M) = \\frac{dN}{d~{\\rm log}~M} = A~{\\rm exp} \\bigg( \\frac{- ( {\\rm log}~M - {\\rm log}~M_c)^2}{2 \\sigma^2} \\bigg) \\end{equation}  Whichever functional form they adopt, most studies agree that the stellar MF appears to reach a peak at a few tenths of a solar mass.  Power law characterizations of the MF require a flattening from a Salpeter slope for M $>$ 1 M$_{\\odot}$ to $\\alpha \\sim$0.5--1.7 below 1~$M_{\\odot}$, while \\citet{Chabrier2005a} finds a log-normal MF with a characteristic mass of M$_c$=0.2 is required to reproduce the stellar MF for a volume complete sample in the solar neighborhood.  We refer the reader to comprehensive reviews by \\citet{Scalo1986}, \\citet{Reid2000}, \\citet{Kroupa2002}, \\citet{Chabrier2003}, and \\citet{Corbelli2005} for a more detailed understanding of the vast literature devoted to measuring the stellar MF.  Several processes have been advanced to explain the shape of the MF; examples include gravitational fragmentation \\citep{Klessen1998}, competitive accretion \\citep{Larson1992}, truncation of mass accretion due to radiative or dynamical feedback \\citep{Silk1995}, star-star interactions \\citep{Reipurth2001} and the primordial distribution of clump masses within molecular clouds \\citep{Padoan2002}. The efficiency of each mechanism is influenced by physical variables, such as the metallicity and magnetic field strength of the parent molecular cloud, the local stellar density, and the intensity of the surrounding radiation field.  These effects may ultimately result in observable MF variations as a function of environment, but at present the observed variations can be largely attributed to uncertainties arising from finite sample sizes and systematic differences in methodology \\citep{Chabrier2005a}, though intriguing signs of MF variations with environment may be emerging \\citep[e.g., Taurus; ][]{Luhman2004}.  Previous studies of the Galactic disk MF have been limited, however, to samples of a few thousand low-mass stars (with low-mass stars defined for the purposes of this paper as M$_{*} <$ 0.7 M$_{\\odot}$), due to the inability to simultaneously obtain a deep and wide photometric sample. \\citet{Hawley2002} demonstrated that Sloan Digital Sky Survey (SDSS) and Two Micron All Sky Survey (2MASS) photometry of low-mass stars shows monotonic behavior across a wide range of colors. The accurate, multi-color catalogs produced by SDSS and 2MASS thereby allow the identification and characterization of millions of low-mass stars in the local Galactic neighborhood, enlarging previous photometric samples of field low-mass dwarfs by several orders of magnitude.  The statistical power of such a sample makes an SDSS/2MASS catalog of low-mass dwarfs a promising avenue for improving our measurement of the low-mass stellar luminosity and mass functions.  To test and calibrate this technique before applying it to the entirety of the SDSS and 2MASS databases, we have performed a combined photometric and spectroscopic study of nearly 30,000 stars detected in $\\sim$30 square degrees of overlap between the SDSS and 2MASS footprints.  We describe in \\S \\ref{merge} the construction of this sample by matching the SDSS, 2MASS, and Guide Star Catalog (GSC) photometric databases. In \\S \\ref{specsec:overall} we present the spectroscopic follow-up of more than a third of the sample.  We verify the completeness limit of the photometric sample in \\S 4, where we also analyze the spectroscopic sample to understand the level of contamination and bias within the photometric catalog.  We use this photometric catalog to measure the luminosity function of low-mass stars in the Galactic neighborhood in \\S 5, derive a measurement of the low-mass stellar mass function in \\S 6, and summarize our conclusions in \\S 7.  An appendix to this work presents native or transformed SDSS photometry for stars with measured trigonometric parallaxes, useful for constraining the empirical SDSS/2MASS color-magnitude relation applied in \\S 5.  We note that a full and complete description of our analysis necessarily results in a rather lengthy paper; we therefore suggest that many readers may wish to skip directly to \\S 5 \\& 6 to concentrate upon the core scientific findings of this work, and then return to read \\S 2, 3 \\& 4 if a deeper understanding of the assembly of the photometric sample and the effects of incompleteness, contamination and bias is desired. ",
        "conclusions": "\\begin{enumerate}  \\item{We have identified a sample of more than 13,000 candidate low-mass stars detected within 30 square degrees of 2MASS, SDSS and/or GSC imaging. Empirical tests verify the sample is more than 99\\% complete to $J=16.2$, and more than 90\\% complete at even the reddest colors.}  \\item{Analysis of the spectroscopic sample reveals that exotic contaminants, such as QSOs, extragalactic stellar populations, CVs, and white dwarf/M dwarf pairs, comprise less than 0.3\\% of our matched sample.}  \\item{Photometric $i-J$ color provides a reliable proxy for stellar spectral type and T$_{eff}$; we measure dispersions of 1 subclass for a given 0.1 magnitude slice in $i-J$, or of 0.1 magnitude in $i-J$ for objects with a given MK spectral type.}  \\item{In a small percentage of cases (20/9624; 0.02\\%), a star's observed spectral type differs significantly from that expected on the basis of its photometric $i-J$ color. This usually occurs when the 2MASS and SDSS detections of a close visual binary are improperly matched, and we have identified an $i-z$ vs. $i-J$ color cut which can be used to eliminate such spurious detections from the sample.}  \\item{Analysis of the kinematic and spectroscopic properties of our sample reveal 88 candidate subdwarfs, representing a subdwarf fraction of 0.68\\% for our broader sample.  As the bulk of the candidate subdwarfs (69/88) were identified kinematically, and proper motion catalogs are biased towards objects with large space motions such as subdwarfs, this is an upper limit to the true subdwarf fraction of our photometric sample.}  \\item{TRILEGAL simulations of the Galactic stellar population predict giant stars will contaminate our sample at the 4\\% level globally, but with giants exceeding 10\\% of stars with $i-J \\sim$ 1.15.  Candidate giant stars identified in our sample using spectral indices verify the colors of giants predicted by the TRILEGAL simulation.  Restricting our analysis to sources with $J >$ 12 and $i-J >$ 1.25, however, reduce contamination of our sample by giant stars to a managable level (less than 2\\% in the bluest, most contaminated color bin) while sacrificing only a handful of bona fide low-mass dwarf stars.}  \\item{After correcting for unresolved binaries and Malmquist effects, we find the luminosity function of the Galactic disk is consistent with that measured from volume complete samples of the solar neighborhood. The dominant systematic uncertainty in our analysis is due to the remaining uncertainty in the SDSS/2MASS CMR, which translates into a 70\\% uncertainty at the peak of the stellar LF.  Systematic effects due to the binary correction prescription and the adopted Galactic model are comparatively small, at the 10-20\\% level.}  \\item{The logarithmically binned mass function measured from our sample peaks at 0.3 M$_{\\odot}$ and agrees within 1$\\sigma$ for all but one data point of the mass function measured by \\citet{Reid2002} from their volume complete sample.  A log normal characterization with A=0.1 \\citep[in units of stars (Log M$_{\\odot}$)$^{-1} pc^{-3}$ for comparison with the values reported by ][]{Chabrier2003}, M$_c$=0.29, and $\\sigma$=0.32 provides an adequate description of the data, with 90\\% confidence intervals for each parameter of A=0.08--0.12 , M$_c$=0.20--0.50, and $\\sigma$=0.22--0.37. These 90\\% confidence intervals correspond to {\\it linearly binned} mass functions peaking between 0.27 M$_{\\odot}$ and 0.12 M$_{\\odot}$, where the best fit MF turns over at 0.17 M$_{\\odot}$.  A power law fit to the entire mass range sampled here, however, returns a best fit of $\\alpha$=1.1 (where the Salpeter slope is $\\alpha$ = 2.35); a broken power law returns $\\alpha$=2.04 at masses greater than log M = -0.5 (M=0.32 M$_{\\odot}$), and $\\alpha$=0.2 at lower masses.}  \\item{We emphasize that comparisions of analytic MF fits, rather than the underlying data, can give the false impression of MF variations even when the underlying data agree well. We are entering the regime where we create more noise than signal by comparing analytic descriptions of the MF, which impose structure on the data {\\it a priori}, rather than comparing the actual MF datapoints themselves.}  \\end{enumerate}"
    },
    "0807/0807.1101_arXiv.txt": {
        "abstract": "We study the level of primordial non-Gaussianity in slow-roll two-field inflation. Using an analytic formula for the nonlinear parameter $\\fnl$ in the case of a sum or product separable potential, we find that it is possible to generate significant non-Gaussianity even during slow-roll inflation with Gaussian perturbations at Hubble exit. In this paper we give the general conditions to obtain large non-Gaussianity and calculate the level of fine-tuning required to obtain this. We present explicit models in which the non-Gaussianity at the end of inflation can exceed the current observational bound of $|\\fnl|\\lesssim100$. ",
        "introduction": "Standard slow-roll single-field inflation generates a quasi-scale invariant spectrum of Gaussian, adiabatic perturbations. While this is in agreement with current data, string theory and SUSY generically contain many scalar fields and it is important to test how these extra fields may change the predictions of the simplest models. In particular, the simplest models predict a level of non-Gaussianity which is too small to observe in the foreseeable future \\cite{Maldacena:2002vr,non-gaussian1,Seery:2005gb}, so any detection would be of great significance.  In this paper we focus on the possibility to obtain a large level of non-Gaussianity during slow-roll inflation and derive general conditions for when it may be large. We present simple, explicit two-field models which can saturate the observational bound. Although there are  many models of inflation which generate a large non-Gaussianity~\\cite{Alabidi:2006hg,Bartolo:2001cw,Enqvist:2004ey,Jokinen:2005by,Chen,cline, Misra:2008tx}, few of them do so during inflation except by breaking slow-roll, e.g.~with a kink in the potential \\cite{Chen:2006xjb}  or having a non-standard kinetic term, e.g.~DBI inflation \\cite{DBI}. Other widely considered methods to generate a large non-Gaussianity include the curvaton scenario \\cite{curvaton}, modulated reheating \\cite{modulatedreheating} or an inhomogeneous end to inflation \\cite{Bernardeau:2002jf,endofinflation,Alabidi:endinf,Sasaki:2008uc,Dutta:2008if}. All these scenarios require an additional light scalar field that doesn't affect the dynamics during inflation but that becomes important either at the end of inflation or later.  Several authors have considered the possibility of generating non-Gaussianity during inflation for a separable potential. In particular Vernizzi and Wands calculated a general formula for the non linearity parameter $f_{NL}$ which parameterises the bispectrum in the case of a sum separable potential \\cite{Vernizzi:2006ve}, this was later extended to the case of a product separable potential including non-canonical kinetic terms~\\cite{Choi:2007su} and a sum separable potential for an arbitrary number of fields \\cite{Battefeld:2006sz}. This has been further extended to the trispectrum (4-point function) \\cite{Seery:2006js}. For other approaches, see for example \\cite{Rigopoulos:2005us,Yokoyama:2007uu}.  Rather less work, however, has been done on analysing the formulas which have been calculated and are present in the literature. The results are rather long and depend on many parameters.  In general they appear to give a slow-roll suppressed non-Gaussianity, subject to several simplifying assumptions that are additional to the slow-roll assumption and that are not always valid. Here we carefully consider several explicit models of two-field inflation and scan the parameter space more generally than has been done before. We show that it is possible to generate an extremely large non-Gaussianity during slow-roll inflation even when the field perturbations at Hubble exit are Gaussian. However this is only possible for specific values of the model parameters and initial conditions. We consider under which conditions a general model can generate a large non-Gaussianity and show that this always requires some fine tuning of initial conditions.  Although previous papers have shown that it is possible to generate a narrow ``spike\" of large non-Gaussianity while the inflaton trajectory turns a corner, the non-Gaussianity decays again very quickly after the corner \\cite{Vernizzi:2006ve,Rigopoulos:2005us}. This spike of non-Gaussianity is associated with a temporary ``jump\" in the slow roll parameters \\cite{Yokoyama:2007uu}. In the models we consider the non-Gaussianity is large over many $e$-foldings of inflation and all of the slow roll parameters remain smaller than unity.   The paper is organised as follows: First we define relevant quantities and introduce some notation (Section \\ref{sec:theory}). Then, in the following two sections we derive the general conditions to generate a large non-Gaussianity during inflation, in Section~\\ref{Section_product} for a product separable potential and in Section~\\ref{Section_sum} for a sum separable potential. In Section~\\ref{Section_examples} we give specific examples for product and sum potentials which can generate large non-Gaussianity. This includes a two-field model of hybrid inflation in Section~\\ref{sec:hybrid}. In Section~\\ref{Section_noncanonical} we extend the previous results to a generalised action with non-canonical kinetic terms. We conclude in Section~\\ref{Section_conclusion}. ",
        "conclusions": "\\label{Section_conclusion} We have made an in depth investigation into the level of non-Gaussianity during two-field slow-roll inflation. We have shown that it is possible to generate a large level of non-Gaussianity during inflation without violating slow roll and when the inflaton field perturbations are Gaussian at Hubble exit.  The general conditions for generating a large non-Gaussianity show that the inflaton potential must have a specific shape so that the angle of the background trajectory can grow by about two orders of magnitude or more during inflation. In the case of an inflation potential made of a product of two quadratic potentials, this condition is not possible so that we conclude this model cannot generate a large non-Gaussianity during slow roll inflation. When the angle of the background trajectory grows sufficiently (in relative terms), we still need one of the fields to dominate throughout inflation, yet the remaining field can not remain full insignificant. For typical values of the slow roll parameters, we require that initially the subdominant field, say $\\vp$, must satisfy $\\cis=\\epspi/\\epsilon^*\\lesssim10^{-7}$. This means that the field trajectory is almost exactly parallel to the $\\chi$ axes initially and typically this requires a finely tuned initial condition. In the case of a product separable potential we then also require that the final value of the angle of the background trajectory lies in a narrow range much greater than the initial value (in relative terms) but still nearly parallel to the $\\chi$ axis. The analysis for a sum separable potential is very similar but more complicated to analyse. The initial background trajectory must be similarly fine-tuned to be nearly parallel to one of the axes of the inflaton fields, but the tuning at the end of inflation is not written so simply. We have also shown how the constraints on generating a large non-Gaussianity may be eased if we generalise the inflaton model to allow a non-canonical kinetic term.  We have presented two explicit models where a large non-Gaussianity can be generated during slow roll inflation. One is a product potential driven by a field with a quadratic potential, which ends inflation when it approaches the minimum of the potential. If the second field with an exponential potential has a sufficiently small initial value then for certain values of the model parameter, $\\lambda$, this field grows by the right amount to generate a large non-Gaussianity, $\\fnl\\sim100$. We also consider the sum separable model of hybrid inflation. We provide a few specific choices of the model parameters, which are effectively $\\etap$ and $\\etac$, for a suitably large ratio of the initial field values this model generates a large non-Gaussianity. We find this is possible either if the $\\eta$'s have the opposite sign (so that inflation takes place near a saddle point) or when both of the $\\eta$'s have the same sign. However we have also found that if one of the field values is too small then quantum fluctuations may perturb the background trajectory to an extent that the classical trajectory in field space is no longer valid. There is also the possibility that the large scale loop correction which we have not generally considered in this paper may not be negligible. We intend to return to these issues and make a more thorough investigation of the hybrid model in a future publication."
    },
    "0807/0807.1337_arXiv.txt": {
        "abstract": "We investigate the origin of the intrinsic scatter in the correlation between black hole mass (\\mbh) and bulge luminosity (\\lbul) in a sample of 45 massive, local ($z$ \\lax\\ 0.35) type~1 active galactic nuclei (AGNs).  We derive \\mbh\\ from published optical spectra assuming a spherical broad-line region, and \\lbul\\ from detailed two-dimensional decomposition of archival optical {\\it Hubble Space Telescope}\\ images.  AGNs follow the \\mlb\\ relation of inactive galaxies, but the zero point is shifted by an average of $\\Delta \\log M_{\\rm BH} \\approx -0.3$ dex.  We show that the magnitude of the zero point offset, which is responsible for the intrinsic scatter in the \\mlb\\ relation, is correlated with several AGN and host galaxy properties, all of which are ultimately related to, or directly impact, the BH mass accretion rate.  At a given bulge luminosity, sources with higher Eddington ratios have lower \\mbh. The zero point offset can be explained by a change in the normalization of the virial product used to estimate \\mbh, in conjunction with modest BH growth ($\\sim$10\\%--40\\%) during the AGN phase.  Galaxy mergers and tidal interactions appear to play an important role in regulating AGN fueling in low-redshift AGNs. ",
        "introduction": "Early-type galaxies commonly contain a central black hole (BH) whose mass strongly correlates with the bulge luminosity (Kormendy \\& Richstone 1995; Magorrian et al. 1998) and stellar velocity dispersion (Gebhardt et al. 2000; Ferrarese \\& Merritt 2000).  Lower-mass BHs found in late-type spirals and spheroidal galaxies follow a similar \\msig\\ relation (Barth et al. 2005; Greene \\& Ho 2006a) but apparently a different \\mlb\\ relation (Greene et al. 2008).  The BH-host scaling relations suggest that BHs play an important role in galaxy formation and evolution (e.g., Granato et al. 2004; Di~Matteo et al. 2005; Robertson et al. 2006).  Understanding the mechanism by which BHs coevolve with their hosts impacts current models of cosmological structure formation (e.g., Bower et al. 2006; Croton et al. 2006).  A key, unanswered question is how the BH-host scaling relations originated. This issue can be addressed by extending the BH-host scaling relations to {\\it active}\\ galaxies---wherein, the BH, by selection, is currently still growing---and by tracking the scaling relations to higher redshift to see when and possibly how they were established.  BH masses in type~1 (broad-line, unobscured) active galactic nuclei (AGNs) now can be routinely estimated to reasonable accuracy ($\\sim 0.3-0.5$ dex), from the ``virial method'' using single-epoch ultraviolet or optical spectra (e.g., Kaspi et al. 2000; McLure \\& Dunlop 2001; Vestergaard 2002; Greene \\& Ho 2005).  More challenging to obtain are reliable measurements of the underlying host galaxy, particularly of the bulge component, which is maximally affected by the bright AGN core (e.g., McLure et al. 1999; Floyd et al. 2004; Nelson et al. 2004; Greene \\& Ho 2006b; Kim et al. 2007), although substantial progress has been made.  Recent studies present tantalizing evidence that the BH-host scaling relations for active galaxies evolve with redshift, even by $z \\approx 0.4$ (Woo et al. 2006; Treu et al. 2007) and as far back as $z \\approx 4$ (Peng et al. 2006a, 2006b; Shields et al. 2006; Ho 2007).  Compared to local, inactive systems, the sense of the evolution is that for a given host galaxy mass or gravitational potential, higher redshift AGNs have a larger BH mass than similar systems at low redshift.  Taken at face value, this suggests that the growth of the BH precedes, or at least outpaces, the growth of the galaxy at higher redshifts.  On the other hand, local ($z \\approx 0$) AGNs seem to behave quite differently. McLure \\& Dunlop (2002) studied a sample of 72 nearby AGNs and find that they roughly follow the same \\mlb\\ relation defined by inactive galaxies, albeit with a somewhat greater scatter.  In the \\msig\\ relation of AGNs, highly accreting AGNs seem to have a different normalization in the sense that they tend to have a lower \\mbh\\ for a given $\\sigma_\\star$ (Greene \\& Ho 2006b; Shen et al. 2008). A similar trend is seen by Ho et al. (2008), using H~I line widths to constrain the gravitational potential of the underlying host galaxy.  As a concrete step toward establishing a robust $z=0$ baseline for comparison with high-$z$ studies, this series of papers (Kim et al. 2007, 2008; M. Kim et al., in preparation) focuses on quantifying the local \\mlb\\ relation for \\psfig{file=fig1.ps,width=8.5cm,angle=0} \\figcaption[fig1.ps] {Distribution of \\mbh\\ vs. Eddington ratio. Radio-loud and radio-quiet objects are denoted by filled and open symbols, respectively. \\label{fig1}} \\vskip 0.3cm \\noindent active galaxies.  Among nearby inactive galaxies, BH mass correlates only marginally less tightly with bulge luminosity or mass than with bulge stellar velocity dispersion (Marconi \\& Hunt 2003; H\\\"aring \\& Rix 2004; Novak et al. 2006). Moreover, the \\mlb\\ relation for active galaxies shows no large systematic differences from that of inactive galaxies (McLure \\& Dunlop 2002).  This suggests that the \\mlb\\ relation can be used as a useful substitute for the \\msig\\ relation, being an especially effective observational tool to track the cosmological evolution of the BH-galaxy connection (Peng et al. 2006a, 2006b; Treu et al. 2007).  Whereas stellar velocity dispersions are difficult, if not impossible, to measure for distant quasars, for example, photometric measurements of quasar hosts continue to be feasible even out to high redshifts, either through direct imaging (e.g., Kukula et al. 2001; Ridgway et al. 2001) or through strong lensing magnification (Peng et al. 2006b).  In a companion paper, Kim et al. (2008) demonstrate that the bulge luminosity of type~1 AGNs can be measured to a reasonable accuracy ($\\sim 0.5$ mag) in {\\it Hubble Space Telescope (HST)}\\ images, even in the regime when the active nucleus far outshines the galaxy.  Instead of characterizing the full \\mlb\\ relation for AGNs, this paper restricts itself to only one important aspect: the origin of the intrinsic scatter. By choosing a sample  for which we can estimate reliable BH masses and for which we can derive robust measurements of bulge luminosity from \\hst\\ images, our objective is to quantify the true intrinsic scatter of the relation and to characterize possible variations of the scatter with physical properties of the AGN, host galaxy, or environment.  By elucidating the physical drivers that influence the scatter of the \\mlb\\ relation, we hope to gain insights on how the BH-host galaxy relations were established.  This paper is structured as follows. We describe the sample selection in \\S{2}. We present the image-fitting procedure for measuring bulge luminosities and our image decomposition results in \\S{3}. We investigate the \\mlb\\ relation for type~1 AGNs in \\S 4. Finally, Section 5 discusses the origin of \\psfig{file=fig2.ps,width=8.5cm,angle=0} \\figcaption[fig2.ps] {Distribution of redshifts for our sample.  The open histogram shows the total sample; the hatched histogram shows the objects with low Eddington ratio (\\edd $\\leq 0.1$); the dashed histogram shows the radio-loud objects. \\label{fig2}} \\vskip 0.3cm \\noindent the intrinsic scatter in the \\mlb\\ relation, ending with a summary in \\S6. Throughout we adopt the following cosmological parameters: $H_0 = 100\\,h = 71 \\, $\\kms~Mpc$^{-1}$, $\\Omega_m = 0.27$, and $\\Omega_{\\Lambda} = 0.75$ (Spergel et al. 2003). ",
        "conclusions": "\\subsection{Which is the Primary Variable?}  We have assembled a sample of local massive type~1 AGNs with reliable spectroscopic data and host galaxy photometric measurements to investigate the origin of \\psfig{file=fig10.ps,width=8.5cm,angle=0} \\figurenum{10} \\figcaption[fig10.ps] {Similar to Figure 6, except that here we show the dependence on redshift: low-redshift sources are marked with red circles and dashed line, and high-redshift sources are marked with blue stars and solid line.  The \\mlb\\ relation for inactive galaxies is shown by the thick hatched line. \\label{fig10}} \\vskip 0.3cm \\noindent the intrinsic scatter in the correlation between BH mass and bulge luminosity.  Assuming a geometrical factor of $f = 0.75$ for the BLR, we find that the AGNs in our sample lie systematically below the \\mlb\\ relation of inactive galaxies by an average offset of $\\Delta \\alpha \\approx -0.3$ dex. Moreover, we have shown that the magnitude of the offset correlates with secondary parameters connected with the AGN (Eddington ratio and degree of radio-loudness) and host galaxy (redshift, bulge-to-disk ratio, and signs of morphological disturbance).   Among the several variables that correlate with the offset in the \\mlb\\ relation, the only one that appears unphysical is that related to redshift. As we noted in \\S4.5, the apparent dependence on redshift most likely reflects the selection effect that higher redshift sources tend to be biased toward higher Eddington ratios (e.g., Boyle et al.  2000) and more luminous, earlier Hubble types.  If we divide the sample into two at $z = 0.15$, the \\mlb\\ relation for both the nearby and distant halves continue to exhibit the dependence on \\edd\\ and $B/T$ that we see for the full sample.  Still, among the rest of the variables that correlate with the zero point offset in the \\mlb\\ relation, which is the primary one?  Given that many galaxy and AGN parameters are mutually correlated, this is not a trivial question to answer. We propose that the primary physical driver is the mass accretion rate, as reflected in the Eddington ratio.  We argue that the host morphology and degree of tidal disturbance directly affect the AGN accretion rate, and that the accretion rate, in turn, is linked to the radio-loudness parameter.  Although our 2-D fits indicate that the hosts of many of our AGNs have a non-zero disk component apart from a bulge, it is important to recognize that, with few exceptions (Fairall 9,  HE 0306$-$3301, MS 1059.0+7302, PG 2130+99), most of the sources in our sample do not have regular, normal disks.  The vast majority of the sample---by construction when we imposed the \\mbh\\ cut---is decisively bulge-dominated.  Many of the features that we attribute to a ``disk,'' in fact, are simply diffuse, extended features above and beyond the dominant bulge component, which we have parameterized using a single S\\'ersic function.  There is no a priori reason why the bulge should be defined in such a manner, that it cannot have a more complex light distribution, especially at large radii.  In other instances, the extra-bulge component is highly disturbed and almost certainly of tidal origin.  With few exceptions (Bennert et al. 2008), these features have generally never been measured before quantitatively in AGN host galaxies.  However, it is entirely debatable whether any of these structures truly belongs to or will ever settle into a normal disk component.   Instead, we surmise that many of the tidal tails and extended, distorted features, in fact, should be considered as part of the {\\it bulge in formation}.  They are reminiscent of morphological signatures attributed to the late, advanced stages of gas-rich mergers (e.g., Barnes \\& Hernquist 1996; Lotz et al. 2008) or possibly even gas-poor (``dry'') mergers (e.g., van~Dokkum 2005; Naab et al. 2006).  Plausible examples of this phenomenon in our sample include [HB89] 0316$-$346 (Fig. 15), HE 1434$-$1600 (Fig. 19), and PG 1012+008 (Fig. 32).  In support of this hypothesis, Figure 11{\\it b}\\ illustrates that the morphological segregation seen in the \\mlb\\ relation (Fig. 7) essentially disappears when the bulge luminosity is replaced with the {\\it total}\\ luminosity of the host.  The scatter also goes down slightly, from 0.40 dex to 0.36 dex in the \\mlt\\ relation.  The dependence on Eddington ratio, however, remains (Fig. 11{\\it a}).  The above interpretation offers a plausible explanation for the apparent link between morphological type and accretion rate, which otherwise is somewhat perplexing.  Within our sample it is the apparently earlier-type, more bulge-dominated systems that actually have {\\it higher}\\ accretion rates. This runs counter to the trend normally seen in nearby AGNs (e.g., Heckman et al.  2004; Greene \\& Ho 2007a) and the general tendency for present-day early-type galaxies to be more gas-poor than late-type galaxies.  However, if high-mass, luminous AGNs result from the aftermath of gas-rich galaxy-galaxy mergers (e.g., Sanders et al. 1988; Hopkins et al. 2006), our results imply that it is during the most advanced stages of the merger that accretion on the central BH attains its maximum rate.  Some of the most highly accreting objects (those with large \\edd) in our sample, which coincide among those with the largest offset in the \\mlb\\ relation, also happen to be among the ones that show the most conspicuous morphological signatures of tidal perturbations, as measured by the  Fourier parameter $a_1$ (see Fig. 8).  These include HE 0354$-$5500 (\\edd = 0.57; $a_1 = 0.33$), PG 1012+008 (\\edd = 0.32; $a_1 = 0.32$), PG 1613+658 (\\edd = 0.44; $a_1 = 0.40$), and PG 1700+518 (\\edd = 0.71; $a_1 = 0.39$).  However, not every object with a high \\edd\\ has a large value of $a_1$.  This may imply that not all quasar episodes are triggered by major galaxy interactions, or that accretion can proceed at a substantial rate even after the tidal features have disappeared.   To the extent that the peak star formation rate in the merger has already subsided during this phase, our scenario offers an additional explanation for why type~1 AGNs contain intermediate-age stars (Kauffmann et al.  2003) but generally not much concurrent star formation (Ho 2005; Kim et al. 2006). \\begin{figure*} \\psfig{file=fig11.ps,width=18.5cm,angle=0} \\figurenum{11} \\figcaption[fig11.ps] {Correlation between BH mass and absolute $R$-band magnitude for the {\\it total}\\ emission from the host galaxy.  ({\\it a}) The symbols are the same as in Figure 6; the dependence on Eddington ratio still remains. ({\\it b}) The symbols are the same as in Figure 7; the dependence on galaxy morphology is much weaker. The relation between \\mbh\\ and bulge luminosity for inactive galaxies is denoted by the thick hatched line. \\label{fig11}} \\end{figure*} \\noindent  Within our sample, four AGNs have compact sources---plausibly small accreted companions---projected close to the primary host galaxy.  These may be examples of minor mergers.  None has a large value of $a_1$. Two of the four lie exactly on the \\mlb\\ relation of inactive galaxies (Fig. 8) and have relatively low Eddington ratios (\\edd\\ = 0.05 for PG 1426+015 and \\edd\\ = 0.06 for PKS 2135$-$14).  The other two lie offset below the relation, but only PG 1302$-$102 has a large \\edd\\ (0.4); PHL 1093 has \\edd\\ = 0.03.  Thus, within our limited statistics, we have no evidence that minor mergers play a significant role in elevating the accretion rate in AGNs.  Within this backdrop, we can offer a tentative explanation for the zero point difference between radio-loud and radio-quiet sources (Fig. 9), one that ultimately links the generation of powerful radio jets to the BH accretion rate and/or host galaxy morphology.  But first we should clarify some terminology.  There are two popular definitions of ``radio-loud'' AGNs in the literature, and it is important not to confuse them. One common usage of this term refers to sources that are classified solely by their radio-to-optical flux ratio ($R$) as defined by Kellermann et al. (1989)\\footnote{Terashima \\& Wilson (2003) advocated a closely related radio-loudness parameter $R_{\\rm X}$ based on the radio-to-X-ray flux ratio.}, regardless of their radio power.  On this basis, the vast majority of AGNs in the local Universe (Ho 2008), most with extremely low luminosities, qualify as being radio-loud, with the degree of radio-loudness increasing with {\\it decreasing}\\ \\edd\\ (Ho 2002; Terashima \\& Wilson 2003; Greene et al. 2006).  The radio emission in most of these low-power sources is largely dominated by a compact core, and any jet-like features are confined to sub-galactic scales.  The host galaxies encompass all morphological types, including spiral galaxies (Ho \\& Peng 2001).  The Milky Way's Sgr~A$^\\star$ is a familiar example.  The second commonly used definition of radio-loudness is less clear-cut, but it involves some combination of relative ($R \\gg 10$) and absolute ($P_{\\rm rad}$ \\gax\\ $10^{23-24}$ W Hz$^{-1}$) measures of radio power.  When detected, the radio jets have super-galactic dimensions and are highly collimated.  The radio-loud sources in this study, with median $R = 820$ and $P_{\\rm 6cm}^{\\rm tot} = 7\\times10^{25}$ W Hz$^{-1}$, belong to this second category.  Strong radio sources of this variety invariably reside in early-type galaxies (e.g., McLure et al. 1999) and are associated with {\\it high}\\ accretion rates (e.g., Maccarone et al. 2003; K\\\"{o}rding et al. 2006).  Nevertheless, only a minority of highly accreting AGNs in massive, early-type host galaxies are radio-loud.  There is no universally accepted explanation for this longstanding quandary.  One possibility is that a necessary ingredient for the generation of powerful, collimated jets is the existence of a BH with a large spin (e.g., Sikora et al. 2007), which is more easily attained in merger-driven accretion events during the formation of giant elliptical galaxies than in disk or spiral galaxies (Volonteri et al. 2007).  \\subsection{Physical Interpretation}  The principal conclusions of our study are that the zero point of the \\mlb\\ relation for AGNs is offset from that of inactive galaxies and that the magnitude of the offset correlates with several physical properties of the AGN and host ultimately connected to the accretion rate.  These effects account for the bulk of the observed intrinsic scatter in the AGN \\mlb\\ relation.  Interestingly, other studies of nearby AGNs using different probes of the host galaxy have independently arrived at very similar conclusions.  Onken et al. (2004; see also Nelson et al. 2004), analyzing a small sample of type~1 AGNs with available stellar velocity dispersions and BH masses determined through reverberation mapping, concluded that the active systems lie offset below the \\msig\\ relation of inactive galaxies by $\\sim 0.2$ dex if one assumes a spherical BLR.   Greene \\& Ho (2006b) obtained an almost identical result using a much larger sample of type~1 AGNs with measured $\\sigma_\\star$ and \\mbh\\ estimated from single-epoch spectra.  These conclusions have been reaffirmed by Shen et al. (2008) in their analysis of composite AGN spectra derived from the Sloan Digital Sky Survey.  A different approach was taken by Ho et al. (2008) through 21~cm H~I observations. Using the maximum rotational velocity and total dynamical mass of the galaxy as new variables to represent the gravitational potential of the host, they find that both quantities strongly correlate with \\mbh.  In qualitative agreement with the results from this study, Ho et al. find that the zero point in the scaling relations depends primarily on the accretion rate in the sense that, at a given galaxy rotation velocity and especially dynamical mass, AGNs with higher Eddington ratios have systematically lower \\mbh\\ than those with lower Eddington ratios.  These trends can be interpreted in one of two ways.  In the first instance, we note that AGNs, by selection, have actively growing BHs.  We can envision that, at a fixed bulge potential (velocity dispersion, luminosity, mass), active galaxies have less massive BHs than inactive galaxies {\\it if}\\ the bulk of the star formation precedes, and is not well synchronized with, a major accretion event.  This particular time sequence, which is required in order to imprint a net negative offset in the \\mbh\\ versus host galaxy relations, seems to be supported by the prevalence of post-starburst signatures (Canalizo \\& Stockton 2001; Kauffmann et al. 2003) as well as the low levels of ongoing star formation (Ho 2005; Kim et al. 2006) found in type~1 AGNs.  If the AGN phase lasts, say, for $\\sim 10^8$ yrs, which is near the upper end of the currently estimated lifetimes (Martini 2004), then a $10^8$ \\solmass\\ BH would increase its mass by a factor of 2 ($\\sim 0.3$ dex) if it is radiating at \\edd\\ = 0.5 with a radiative efficiency of 0.1.  This example is merely illustrative.  In reality, the AGNs in our sample span a wider range of \\edd\\ and the AGN lifetimes may be shorter.  Nevertheless, on average we expect luminous AGNs to lie systematically below inactive galaxies on the \\mlb\\ relation.  Moreover, our study, as do those of Ho et al. (2008) and Shen et al. (2008), further shows that the magnitude of the zero point offset depends on the Eddington ratio: the BH masses of high-Eddington ratio AGNs have more catching up to do than the BH masses of lower-Eddington ratio AGNs.  Such a systematic trend can only come about if the accretion rate directly relates to the evolutionary phase of the accretion event.  This seems plausible, in light of the apparent association between accretion rate and the degree of tidal perturbation.  Instead of the BH being undermassive, perhaps it is actually the bulge that is overluminous, by  $\\Delta M_{R,{\\rm bul}} \\approx 0.5-0.6$ mag.  First, we dismiss the possibility that the luminosity enhancement could be due to contamination from nebular emission from the narrow-line region.  Although the spatial extent of the narrow-line region in quasars can reach several kpc (Bennert et al. 2002), substantially overlapping with the bulge, the typical \\oiii\\ luminosities in our sample ($\\langle L_{\\rm [O III]} \\rangle = 10^{42}$ ergs s$^{-1}$) contribute less than 3\\% to the luminosity of the bulge ($\\langle$\\rbulge$\\rangle=-22.75$ mag).  Given the evidence outlined in \\S 5.1 that the most extreme outliers in the \\mlb\\ seem to have undergone a recent merger or tidal interaction, a more likely possibility is that the bulge luminosity may be moderately enhanced by the latest episode of central star formation.  Indeed, for a small sample of reverberation-mapped Seyfert 1 galaxies with stellar velocity dispersion and bulge luminosity measurements, Nelson et al. (2004) have shown that these objects are somewhat brighter ($\\sim 0.4$ mag) than inactive galaxies at a given velocity dispersion.  They attributed this offset in the Faber-Jackson (1976) relation to younger stellar populations in AGNs.  Boosting the $R$-band luminosity by 0.5 mag requires $\\sim 15$\\% of the stellar mass to come from a 1 Gyr population with solar metallicity (Nelson et al. 2004).  While this offers a plausible explanation for the offset in the \\mlb\\ relation seen in our sample, it cannot account for the fact that the most recent and largest samples of nearby type 1 AGNs, statistically at least, show negative offsets (with respect to inactive galaxies) when \\mbh\\ is compared to {\\it all}\\ bulge or host galaxy parameters (stellar velocity dispersion, bulge luminosity, total galaxy dynamical mass).  The direction of the offset is the same (at a given galaxy parameter, \\mbh\\ is lower), and the magnitude of the offset is also similar ($\\sim 0.2-0.3$ dex).  Although this clearly needs to be verified with a large sample that has reliable measurements of both velocity dispersion and bulge luminosities, the most recent studies suggest, contrary to Nelson et al. (2004), that local type 1 AGNs actually do not depart from the standard Faber-Jackson (1976) relation.  Furthermore, for a large sample of sources with measurements of rotational velocity and total galaxy luminosity, Ho et al. (2008) show that type 1 AGNs show no obvious deviations from the Tully-Fisher (1977) relation of inactive galaxies.  In light of these considerations, we favor the view that the negative offset in the \\mlb\\ relation represents a deficit in \\mbh\\ rather than an excess in \\lbul.  Alternatively, we can assert that both active and inactive galaxies intrinsically should obey the {\\it same}\\ BH mass versus host scaling relations. From this standpoint, the zero point offset between AGNs and inactive galaxies, as well as the variations of the offset with \\edd, can be viewed as a systematic {\\it underestimate}\\ of the true value of \\mbh.  Recall that our BH masses are based a virial product assuming a spherical distribution of BLR clouds with isotropic velocities, for which the geometric factor is $f = 0.75$. If, for example, the BLR (or at least the portion of it that predominantly emits the Balmer lines) has a flattened, disk-like geometry with kinematics dominated by rotation, and on the scale of the BLR type~1 sources happen to be preferentially more face-on to our line of sight, then we systematically underestimate the deprojected rotation velocity and hence \\mbh.  Wu \\& Han (2001) invoked this line of reasoning to interpret the offset of type~1 AGNs on the \\msig\\ relation and concluded that on average their BLRs are inclined by $\\langle i \\rangle \\approx 36$\\deg.  The same argument can, in principle, be applied to the observed offset in the \\mlb\\ relation.  While the effect of inclination probably enters at some level (see also Collin et al. 2006), it cannot account for the fact that the magnitude of the offset depends on \\edd. The latter is not a trivial consequence of the mass being underestimated because sources with high \\edd\\ truly {\\it do}\\ exhibit characteristically distinct X-ray, optical, and radio properties (Boller et al. 1996; Boroson 2002; Greene et al. 2006).  Whatever the physical origin of the offset (Onken et al. 2004; Collin et al. 2006; Marconi et al. 2008), we can empirically adjust the normalization factor of the virial product by forcing the AGN sample to agree with the fiducial reference of inactive galaxies.  To remove the zero point offset of $\\Delta \\alpha \\approx -0.2$ to $-0.3$ dex, then, the normalization should be increased by a factor of $\\sim 1.6-2$, from $f = 0.75$ to $f \\approx 1.2-1.5$.  For the most extreme offsets of $\\Delta \\alpha \\approx -0.6$, $f$ increases to $\\sim 3$.  Without additional information, unfortunately, the above two alternative explanations---an undermassive \\mbh\\ versus an underestimated $f$ factor---are degenerate. It is easy to imagine that both effects must operate jointly. On the one hand, the BHs in AGNs are, after all, gaining mass.  On the other hand, as discussed in Collin et al. (2006), there are multiple reasons to believe that the BLR has a nonspherical geometry and that the Eddington ratio may influence its structure and dynamics.   The only way to resolve this degeneracy is to obtain independent estimates of \\mbh\\ for AGNs that do not rely on the BLR virial technique.  To date, efforts to apply resolved stellar dynamical techniques to reverberation-mapped AGNs have yielded very rough estimates of \\mbh\\ for only a couple of sources (NGC 3227: Davies et al. 2006; Hicks \\& Malkan 2008; NGC 4151: Onken et al. 2007), and thus attempting to cross-calibrate the two techniques is still far too premature.  BH mass estimators based on X-ray variability seem more promising.  Gierli\\'nski et al. (2008; see also Hayashida et al. 1998) find that, for accreting BHs in their hard spectral state, the amplitude of their high-frequency X-ray variability scales inversely with \\mbh\\ over a very wide range of masses.  For a small subset of nearby type~1 AGNs, the X-ray-derived masses show rough agreement with \\mbh\\ obtained through reverberation mapping assuming $f \\approx 1.2$.  A very similar conclusion was reached by Niko\\l ajuk et al. (2006).  From comparison of \\mbh\\ for reverberation-mapped sources with masses obtained using the X-ray excess variance method, these authors estimate $f = 1.06\\pm 0.26$, which, interestingly, lies in between the values of the $f$ factor for the two extreme alternatives discussed above.  Assuming that the X-ray-derived normalization factor truly does represent the correct normalization factor for the virial masses, then the inferred growth rates for \\mbh\\ in luminous AGNs are much more modest, from typically as little as 10\\%--40\\% (0.05--0.15 dex) to at most 280\\% (0.45 dex).  Still, we note that the tendency for BHs in local AGNs to be less massive than the BHs in inactive galaxies of similar type runs counter to the trend observed at higher redshift.  Already by $z=0.36$, type~1 AGNs begin to depart from the local \\msig\\ and \\mlb\\ relations (Woo et al. 2006; Treu et al. 2007), but in the {\\it opposite}\\ direction as that seen at lower redshifts: at a given $\\sigma_\\star$ or \\lbul, AGNs are offset compared to local inactive systems by $\\Delta \\log M_{\\rm BH} \\approx +0.5$ dex. This trend has now been extended by Woo et al. (2008) out to $z = 0.57$ using stellar velocity dispersion measurements.  At even higher redshifts, direct $\\sigma_\\star$ measurements are no longer feasible, but surrogate estimates of $\\sigma_\\star$ using narrow emission lines (Salviander et al. 2007) as well as probes of the host galaxy using imaging (Peng et al. 2006a, b) and CO emission lines (Shields et al. 2006; Ho 2007) support the notion that the growth of the BHs in AGNs have been decoupled from, and outpaced, the underlying host."
    },
    "0807/0807.0277_arXiv.txt": {
        "abstract": "The observational data from some black hole candidates suggest the importance of electromagnetic fields in the vicinity of a black hole. Highly magnetized disk accretion may play an importance rule, and large scale magnetic field may be formed above the disk surface. Then, we expect that the nature of the black hole spacetime would be reveiled by magnetic phenomena near the black hole. We will start to investigate the motion of a charged particle which depends on the initial parameter setting in the black hole dipole magnetic field. Specially, we study the spin effects of a rotating black hole on the motion of the charged particle trapped in magnetic field lines. We make detailed analysis for the particle's trajectories by using the Poincar\\'{e} map method, and show the chaotic properties that depend on the black hole spin. We find that the dragging effects of the spacetime by a rotating black hole weaken the chaotic properties and generate regular trajectories for some sets of initial parameters, while the chaotic properties dominate on the trajectories for slowly rotating black hole cases. The dragging effects can generate the fourth adiabatic invariant on the particle motion approximately. ",
        "introduction": "In the central region of AGNs, compact X-ray sources and GRBs, many black hole candidates are reported.  The observational data from their active regions would include the informations about the gravitational and electromagnetic fields.  However, there are no direct observational evidences for the existence of a black hole in the center of the compact active regions.  So, it is very important how to find the informations on the curved spacetime from observational data, and how to interpret their nature. Although the inner part of a gas disk around a supermassive black hole emits non-thermal radiation, this may be explained by synchrotron emission from ultra-relativistic electrons in strong magnetic field regions around a black hole.  In the central black hole system, however, the emission process is very little understood, although the consensus is in favor of emission due to plasma process.  The emission from such particles would give us the information about the black hole spacetime and/or the electromagnetic field.   In a realistic model of a black hole--disk system, we should consider interactions between a plasma particle and a magnetic field in the curved spacetime.  The plausible electromagnetic field around a black hole--disk system is still unknown, although some authors have discussed this problem in the frame work of magnetohydrodynamics \\citep[e.g.,][]{BZ77,Camenzind87,NTT91,TT01}.  In general, the task for solving the structure of magnetic fields in a black hole system is very hard.  Then, to make a comprehensive analysis, we will give a {\\it vacuum} external magnetic field around a rotating black hole, and investigate in detail the motion of a charged  particle confined in the magnetic field.  The stationary and axisymmetric vacuum electromagnetic fields around a black hole are discussed by \\cite{Pett75,CV75,Li00,Ghosh00,TT01}. We apply the black hole dipole magnetic field in Kerr geometry, which is a solution of the vacuum Maxwell's equations.  Although the dipole magnetic field could correspond to the intrinsic field of a compact object like a neutron star, in the case of black holes it has to results from current rings exterior to the event horizon but they can be very close to it \\citep{Prasanna78}. As known in the Earth's magnetosphere, the dipole magnetic field can trap charged particles within its magnetic bottle (i.e., the Van Allen belts).  The particle gyrates around the magnetic field line, and drifts in the toroidal direction \\citep[see, e.g.,][]{GP04}. Furthermore, in the poloidal plane, the particle oscillates along the magnetic field line.  In the black hole case, we can also expect similar situations.   In a black hole magnetosphere, by solving equations of motion of a charged particle moving in the electromagnetic field numerically \\citep[see][]{Prasanna78}, we can see the gyration around the magnetic field line, the drift in the toroidal direction and the multi-periodic motion in the poloidal plane.  Such a orbit is very complicated in both poloidal and toroidal directions. Although a stationary and axisymmetric magnetosphere is assumed, there are two conserved quantities; that is, the energy and angular momentum of the charged particle.  In addition to these, the rest mass of the particle is the third invariable quantity.  However, the fourth invariable quantity related to the azimuthal motion of the particle does not exist when the electromagnetic field is considered around a black hole.  Thus, the motion of a test charge in a black hole magnetosphere is not an integrable system.  Then, the chaotic motion is expected \\citep{NI93}. To examine the trajectory, we plot the Poincar\\'{e} map in the two-dimensional $r$-$p^r$ plane, which shows intersections of the particle's trajectories with the surface of section in phase space \\citep{LL92}. If the plotted points form a closed curve, the motion is regular (not chaotic).  This is because a regular trajectory moves on a torus in the phase space and the curve is a cross section of the torus.  On the other hand, if the plotted points are distributed randomly, the motion is irregular (chaotic).  From the distribution of the points in the Poincar\\'{e} map, we can judge whether or not the motion is chaotic. Then, we know chaotic behavior on the trajectory of a charged particle in a black hole magnetosphere.  Furthermore, we find the black hole's spin effects on that chaotic motion.   This study will be related to the source of radiation near a black hole. We will start this work as a basic step before considering the plasma as a magnetized fluid.  Although we find chaotic and/or regular motion in a black hole magnetosphere, the problem is how to find the phenomena in the observed spectrum.  At this stage, however, we can not show the relation with the observational data in this work. The main purpose of this study is to understand the black hole spin effects (the dragging effects by a rotating black hole) on a test charge, which will be considered as the source plasma of high energy radiation in future.   In \\S~\\ref{sec:basic-eqs}, we review the basic equations for the motion of a charged particle in a black hole magnetosphere, using the Hamilton-Jacobi formalism. To understand the nature of orbits of a charged particle in the black hole magnetosphere, we also consider in \\S~\\ref{sec:orbit-mag} the effective potential including the four-vector potential of the magnetic field. We present the motions of a test charge in the black hole dipole magnetic field for various values of black hole's spin parameters. In \\S~\\ref{sec:chaos} we make detailed analysis for its trajectories by using the Poincar\\'{e} map method, and show the chaotic properties that depend significantly on the black hole spin obviously. We see simultaneous presence of regular trajectories and regions of stochastic trajectories on the Poincar\\'{e} map.  Furthermore, we find that due to the dragging effects of a rapidly rotating black hole, the fourth constant of motion {\\it approximately}\\/ turns out to exit. Then, the motion of a test charge in a dipole magnetic field around a rotating black hole can behave like a integrable system. ",
        "conclusions": "In this paper, we have discussed the off-equatorial motion of a test charge in the black hole dipole magnetosphere.  The charged particle gyrates around a magnetic field line, drifts in the toroidal direction and oscillates between the northern and southern hemisphere, and then the orbits becomes very complicated.  So, the numerical study by the Poincar\\'{e} map is effective.  Then, we find the spin dependence of the trajectory on the Poincar\\'{e} map.  That is, for a slowly rotating black hole case, the particle's trajectory shows chaos. More interestingly, we also find that, for a rapidly rotating black hole case, the chaotic behavior of the trajectories weakens and the fourth invariance of the motion, which is an adiabatic invariant related to the longitudinal motion, can be approximately generated.       The chaotic motion in a black hole magnetosphere may be related to the origin of cosmic rays.  The magnetic mirror concepts play a prominent role in space and plasma astrophysics.  Examples are the Van Allen belt in the magnetosphere of the Earth, where high energy particles are trapped in a magnetic bottle and some kind of acceleration mechanisms are expected.  % In the case of a black hole magnetosphere, we may expect the formation of the magnetic bottle near the black hole.  When such a situation is realized, the particles trapped in the ``black hole Van Allen belt'' would emit high-energy radiation, or some particles may accelerate to ultra-relativistic velocity by some electromagnetic interactions (ex., the Fermi acceleration by the perturbed magnetic bottle). Furthermore, some kinds of waves (i.e., Alfv\\'en waves and/or fast waves) in the magnetosphere will give influence to the motion of a particle moving in a magnetized plasma. Our simple model with the black hole dipole magnetic field presented here could be a preliminary step toward the understanding the motion of charged particles around a black hole.   In this paper, we do not consider the interactions between charged particles, and the emission process from the chaotic/regular motion of a charged particle.  However, we can expect that the spectrum emitted from such charged particles in periodic motions in the inhomogeneous magnetic field carries the informations on the black hole spin and the strength and/or distribution of the electromagnetic field. These studies should be considered in future works."
    },
    "0807/0807.0794_arXiv.txt": {
        "abstract": "The next generation of neutrino and $\\gamma$-ray detectors should provide new insights into the creation and propagation of high-energy protons within galaxy clusters, probing both the particle physics of cosmic rays interacting with the background medium and the mechanisms for high-energy particle production within the cluster. In this paper we examine the possible detection of $\\gamma$-rays (via the GLAST satellite) and neutrinos (via the ICECUBE and Auger experiments) from the Coma cluster of galaxies, as well as for the $\\gamma$-ray bright clusters Abell 85, 1758, and 1914. These three were selected from their possible association with unidentified EGRET sources, so it is not yet entirely certain that their $\\gamma$-rays are indeed produced diffusively within the intracluster medium, as opposed to AGNs. It is not obvious why these inconspicuous Abell-clusters should be the first to be seen in $\\gamma$-rays, but a possible reason is that all of them show direct evidence of recent or ongoing mergers. Their identification with the EGRET $\\gamma$-ray sources is also supported by the close correlation between their radio and (purported) $\\gamma$-ray fluxes. Under favorable conditions (including a proton spectral index of $2.5$ in the case of Abell 85, and $\\sim 2.3$ for Coma, and Abell 1758 and 1914), we expect ICECUBE to make as many as $0.3$ neutrino detections per year from the Coma cluster of galaxies, and as many as a few per year from the Abell clusters 85, 1758, and 1914. Also, Auger may detect as many as 2 events per decade at $\\sim 10^{18}$ eV from these gamma-ray bright clusters. ",
        "introduction": "A population of high-energy cosmic-ray (CR) protons, pervading galaxy clusters, was first invoked (Dennisson 1980) to explain highly polarized radio luminosities ($L_r \\sim 10^{40}-10^{42}$ erg s$^{-1}$) now observed in over 30 clusters. Evidently the radio light is synchrotron emission from highly relativistic electrons ($\\gamma_e \\sim 10^6$) gyrating in large magnetic fields ($B \\sim 1 \\ \\mu$G), despite the fact that such electrons would lose their energy to radiation well before they could traverse the $\\sim$Mpc extent of observed emission. However, protons injected into the cluster diffuse throughout the intracluster medium, essentially unchanged and confined for cosmological times. In this environment they may interact with the background gas and create a decay cascade. Among the products of this decay are the relativistic electrons responsible for radio emission.  Since Dennisson's (1980) pioneering work, conclusive observations of this `secondary' process for electron production within clusters have been few. In Hercules A (Nulsen et al. 2005b), MS 0735 (McNamara et al. 2005), and Hydra A (Nulsen et al. 2005b), vacated areas of X-ray bremsstrahlung indicate the background of thermal electrons is being pushed aside by relativistic electrons injected from the axes of a central source, implying a phenomenally large power, $L_p \\sim 10^{44}$ erg s$^{-1}$. These `bubbles' coincide with sources of radio emission as the protons produce a cascade of synchrotron-emitting electrons, confirming that the secondary decay process does occur in clusters. The protons released during such an episode would be confined and built up within the cluster over cosmic times, totaling some $10^{61}$ erg of nonthermal hadrons---very near energy equipartition with the thermal background (Berezinsky, Blasi, and Ptuskin 1997; Hinton \\& Domainko 2006). This is, we find, the order of nonthermal energy required to be deposited in the nonthermal hadron population for proton-proton scatterings to produce $\\pi^0$ decay $\\gamma$-rays within the EGRET detection sensitivity.  The energy liberated by the infall of the galaxy group NGC~4839 toward the center of the Coma cluster, observed by XMM-Newton (Neumann et al. 2001), is an example of another potential source of high-energy protons. NGC~4839 achieves a velocity of $\\sim 1,400$ km s$^{-1}$, and since the sound speed corresponding to Coma's gas temperature of $\\sim 8$ keV is $\\sim 1,000$ km s$^{-1}$, the subcluster's supersonic motion is expected to produce shocks (which Neumann et al. claim to observe directly in the imaging of NGC~4839). It is likely that the shock compression in these regions energizes a fraction of the thermal particles within the intracluster medium (ICM) by first-order Fermi acceleration. However, how the protons are actually energized is not directly relevant to our analysis. For example, since the cosmic rays remain trapped within the cluster for over a Hubble time, it is likely that second-order acceleration by a turbulent distribution of Alfv\\'en waves may be even more effective at helping the protons attain their highest energies (see, e.g., Liu, Petrosian, and Melia 2004; Liu, Melia, and Petrosian 2006). Our principal goal is to determine the observational signatures such a population would produce, should it be present within the cluster. But cluster mergers do provide us with a viable mechanism for energizing the cosmic rays, regardless of how the individual particles are ultimately accelerated. In this regard, we note that the gravitational energy $\\sim 10^{64}$ erg released in a large relic merger could reasonably supply $\\sim 10^{62}$ erg in a nonthermal hadron population, again sufficient to supply detectable $\\gamma$-rays. In this case, the radio halo may be interpreted as a signature of the cosmological development of a cluster (the oldest known cluster being some 9 billion years old).  Aside from a radio halo produced by secondary electrons, conclusive evidence for the properties of CR protons within galaxy clusters rests on the detection of other products of the protons' decay cascade: $\\gamma$-rays and neutrinos. The closest active galaxy cluster, the Coma cluster, does not emit copiously enough to have been detected by EGRET given its $\\sim10^{-11}$ cm$^{-2}$ s$^{-1}$ flux sensitivity, although we will show that this does not preclude a possible detection by the GLAST satellite or a detectable neutrino flux. Our preliminary identification of clusters which \\emph{are} $\\gamma$-ray bright, and, therefore, plausible extrasolar neutrino candidates, has come from a systematic search for likely clusters among unidentified EGRET sources.  In a search for new radio galaxy clusters, the NRAO VLA (Giovannini et al. 1999) and Westerbork Northern (Kempner \\& Sarazin 2001) sky surveys were systematically cross referenced with X-ray bright objects from the ROSAT survey (Ebeling et al. 1998), providing a list of more than 30 known cluster halos or relics, presented in Table 1. While none of the clusters EGRET was hoped to detect produce a measurable $\\gamma$-ray flux (Reimer et al. 2003), three radio sources coincide with unidentified $\\gamma$-ray bright objects detected by EGRET (Colafrancesco 2002)---the cluster Abell 85 and the EGRET source 3EG J0038-0949, Abell 1914 and EGRET source J1424+3734, and Abell 1758 and EGRET source 3EG J1337 +502a. These three EGRET sources are not unequivocally the quoted clusters, but here we shall assume they are in fact the same objects.  The assumption is not certain. In a study by Reimer et al. (2003) it was shown that the probability that one of the 170 unidentified EGRET sources and a member of their list of 58 considered galaxy clusters coincided was 48.1\\%. Of their list of 58 clusters worthy of consideration, only Abell 85 and the unidentified source J0038-0949 showed a considerable overlap, which could therefore be explained as an entirely chance occurrence. However only six of their 58 considered clusters ---included for their proximity ($z<0.14$) and X-ray brightness--- were shown by the NRAO VLA or Westerbork Northern radio sky surveys to be radio halo/relic clusters, so that in only these six cases is there any reason to believe a significant nonthermal proton population may be present. Also, Abell 1758 and 1914 were not considered by Reimer et al. (2003), as they are above $z >0.14$.  The three candidates are not the most X-ray bright clusters, most of which were assigned only upper limits in the nine years of EGRET observations. It is not obvious why three inconspicuous Abell-clusters are the first to be seen in gamma rays when the closest and most well-studied clusters---known to have evidence for non-thermal particle populations from radio, EUV, or hard X-ray detections---are not. Perhaps a key is that each of the $\\gamma$-ray bright candidates shows direct evidence of recent or ongoing mergers, indicating (if confirmed) that this may outcompete the AGN outburst as a source of cluster cosmic rays. In Abell 85, a 4 Mpc filament (Durret et al. 2005) forms a chain of several groups of galaxies. In particular, an impact region (the `south blob') displays a temperature ($8.8$ keV) far in excess of the rest of the cluster ($2.7$ keV). In Abell 1914 (Govoni et al. 2004) a NE-SW arch-like hot region extends across the cluster center and may represent an ongoing shock. Abell 1758 is a double cluster (David \\& Kempner 2004): A1758N is in the late stages of a large impact parameter merger between two $7$ keV clusters, while A1758S is undergoing a smaller impact parameter merger between two 5 keV cores. In each case, the flux of $\\gamma$-rays correlates with the radio flux, as it would if both were linked to secondary decay products. Abell 85 also shows evidence of a hard X-ray excess (Lima-Neto et al. 2001), indicative of a nonthermal population of electrons.  In the end, however, we caution that while the identification of Abell clusters 85, 1758, and 1914 with three unidentified EGRET sources may be motivated by the reasons we have outlined here, it is not entirely clear that their $\\gamma$-rays are indeed produced diffusively within the ICM. Thus, one must accept the results presented in this paper, particularly the predicted neutrino fluxes, with this important caveat in mind.  As clusters confine CR protons over cosmological times, they provide a promising source for a possible first detection of high-energy ($\\gtrsim$ TeV) extra-solar neutrinos. Aside from probing high or ultra-high energy particle physics processes occurring within astrophysical sources, any such neutrino (or $\\gamma$-ray) detection would also establish the properties of the proton population within the cluster and settle the possible role of such populations in various cluster phenomena---e.g., generating the radio halo (Dennisson 1980), their contribution to the `cooling flow' problem (Silk 1995), or their acceleration within cluster merger shocks (Sarazin 2004).  In this paper we take the optimistic view that protons within the Coma cluster create a flux of $\\gamma$-rays which is within the detection limits of the GLAST satellite, without violating the EGRET null detection. In the same spirit of optimism, we identify the Abell clusters 85, 1758, and 1914 with their gamma-ray bright counterparts. The resulting predictions for neutrino and gamma-ray detections for these four clusters will determine whether this optimism was well-founded.  Work on the hadronic origin of one or more spectral components in the overall cluster emissivity has been quite extensive, and our investigation overlaps with several earlier treatments, though our model is unique in the treatment of the particle kinetics and proton-electron scattering events. Dolag and Ensslin (2000) considered the possible origin of the radio halos from a hadronic secondary electron injection, based on a detailed modeling of the cluster magnetic field (see also Blasi \\& Colafrancesco 2000, Jones 2004, Brunetti et al. 2007, Marchegiani et al. 2007). Subsequently, Pfrommer and Ensslin (2004) constrained the cosmic-ray population in several cooling flow clusters from their radio and $\\gamma$-ray emissivities (see also Reimer et al. 2004, and Sarazin 2007). Not surprisingly, the $\\gamma$-ray emission from proton-proton collisions in clusters may contribute a non-negligible fraction of the extragalactic background radiation field (Kuo et al. 2005). And though we will not be addressing $\\gamma$-ray line emission in this paper, cosmic rays may also collide with heavier nuclei in the ICM to produce measurable $\\gamma$-ray line intensities (Iyudin et al. 2004). An extensive review of the $\\gamma$-ray emissivity in clusters may be found in Blasi et al. (2007).  The issue of hadronic acceleration during large-scale structure formation has been considered by several authors, including Miniati et al. (2001), Berrington \\& Dermer (2003), Blasi (2004), Ensslin et al. (2007), Jubelgas et al. (2008), Ando and Nagai (2008), Nakar et al. (2008) and, in the context of secondary acceleration due to stochastic processes, Brunetti et al. (2004), Brunetti and Blasi (2005), and Petrosian and Bykov (2008). But cosmic rays may also play an important dynamical role throughout the cluster's lifetime. Collisions between energetic hadrons and the ICM have been invoked as a means of heating the gas condensing toward the middle of the cluster, thereby preventing an overly rapid cooling of the flow (Colafrancesco, Dar, and De R\\'ujula 2004, Guo \\& Oh 2008).  Our goal in this paper is specifically to calculate the neutrino flux expected from $\\gamma$-ray bright clusters, but our approach differs from earlier attempts in several important ways, which we will describe over the next several sections. Our model requires that the number of protons be consistent with a secondary explanation of the Coma cluster's radio halo, using magnetic fields ($\\sim 1 \\mu$G) consistent with observations of Faraday rotation, rather than those ($\\sim 0.1 \\mu$G) consistent with an inverse-Compton description of the hard X-ray excess detected by Beppo\\emph{SAX} and RXTE (see also Wolfe and Melia 2006). The only model-dependent characteristics of this work are then the distribution of gas and that of protons diffusing throughout the cluster, which---together with the details of calculating the various links in the decay cascade---we discuss in Section 2. Our technique for calculating the detected number of neutrinos, given the predicted flux of Coma, is discussed in Section 3. Finally in Section 4, we review the forecast for those clusters for which a positive $\\gamma$-ray detection has been proposed. ",
        "conclusions": "Observations of both $\\gamma$-rays ($\\gamma_{pp}$) and neutrinos (both $\\nu_{pp}$ and $\\nu_{p\\gamma}$) in clusters provide answers to a series of yes-or-no questions. If both $\\gamma_{pp}$ and $\\nu_{pp}$ are observed from a cluster, then proton-proton induced cascades are certainly at work and their contribution to the radio halo can be evaluated---as well as potential models for the creation of viable proton populations with total energies $>10^{61}$ erg. If $\\gamma_{pp}$ is detectable by either EGRET or GLAST and $\\nu_{pp}$ is not observed, then $\\gamma_{pp}$ is not likely to be caused by secondary decay (this is not a firm `no' as the source in question may occupy an unfavorable position in the night sky, or the proton population may experience a coincidental cutoff just after 100 MeV) and new models would need to be considered. An Auger $\\nu_{p\\gamma}$ detection would confirm our expectation that very energetic cosmic-ray photons interact with the CMB to produce debris particles, thereby losing their energy. This would add some support to the view that cosmic ray events above $6 \\times 10^{19}$ eV cannot be hadronic if the incident particles originate from beyond the GZK limit. Finally, if $\\nu_{pp}$ is observed without $\\nu_{p\\gamma}$, we may assume that the spectral index of protons is steeper than $2.1$, or that the cluster is incapable of accelerating protons to this energy.  The combination of X-ray, radio, $\\gamma$-ray, and neutrino observations on galaxy clusters would give a vital new perspective on the mechanisms producing high-energy particles in clusters. A simultaneous detection of $\\gamma$-rays and neutrinos from pp scatterings in a cluster would confirm conclusively that secondary production is at work, and provide a new probe for cosmic-ray production in the MeV and TeV ranges, as well as shock acceleration in clusters and astrophysically in general. A detection (or upper limit) of p$\\gamma$ neutrinos, meanwhile, would demonstrate whether clusters are capable of accelerating protons to ultra-high energies, and, indeed, how particle physics tested in colliders evolves towards these energies.  A confirmation via GLAST of Abell 85, 1758, or 1914, as $\\gamma$-ray bright would confirm them as strong neutrino sources, potentially detectable in the ICECUBE and, perhaps, Auger experiments. If a nonthermal proton population approaching energy equipartition is demonstrated, these then form a dominant source of pressure in clusters, affecting structure and evolution. Of course, the fact that neutrino detections are only possible near equipartition means that there may be only a narrow window of opportunity for actually measuring a neutrino flux, or maybe none at all. But we find it promising that an interesting prediction may be made even without requiring a super-equipartition proton population. And, anyway, a near-equipartition situation between the protons and the background electrons is suggested by a dynamic coupling between these two populations, as discussed in Wolfe and Melia (2008). Taken together with evidence of ongoing mergers, these observations would give new insight into the period of cluster formation, beginning some $>9$ billion years ago, and what role mergers play in cosmic-ray production."
    },
    "0807/0807.2334_arXiv.txt": {
        "abstract": "We present a new determination of the mass content of the Sculptor dwarf spheroidal galaxy, based on a novel approach which takes into account the two distinct stellar populations present in this galaxy. This method helps to partially break the well-known mass-anisotropy degeneracy present in the modelling of pressure-supported stellar systems. ",
        "introduction": "The determination of the mass content of galaxies is a fundamental step for our understanding of their formation and evolution. This is particularly true for small systems such as dwarf spheroidal galaxies (dSphs), whose total mass can determine their destiny as being fragile objects, strongly perturbed by supernovae explosions and/or interactions with the environment, or mostly undisturbed objects.  Part of the interest in determining the masses of dSphs has surely been driven by the fact that these galaxies, with mass-to-light (M/L) ratios up to 100s, appear to be the most dark matter (DM) dominated objects known to-date. This makes them potentially good testing grounds for different DM theories of galaxy formation; however at the moment it is still unclear if dSphs inhabit cusped DM profiles, as predicted by cold DM theories of galaxy formation, or cored profiles, typical of warmer kinds of DM. This distinction is often hampered by the degeneracies which are intrinsic to the method generally used to determine the mass of pressure-supported stellar systems, i.e. a Jeans analysis of the line-of-sight (l.o.s.) velocity dispersion obtained considering all stars as a single component embedded in an extended DM halo. This analysis is subject to the well known degeneracy between the mass distribution and the orbital motions of the individual stars in the system (mass-anisotropy degeneracy). Recent observations have shown that some dSphs host multiple stellar populations which have distinct spatial distribution, metallicity and kinematics. The presence of multiple stellar populations requires a modification in the way these systems are dynamically modelled to derive their mass content.  Here we describe the results of our determination of the mass content of the Sculptor (Scl) dSph by adopting a more detailed kinematic modelling than the traditional one, i.e. by considering Scl as two-(stellar) components embedded in an extended DM halo. We first briefly describe and discuss the observational properties of Scl derived from our work (\\cite[Tolstoy et al. 2004]{tolstoy2004}; \\cite[Battaglia 2007]{battaglia2007}; \\cite[Battaglia et al. 2008]{battaglia2008}). We then move to the results of the two-component modelling; we show a clear example of the mass-anisotropy present when modelling Scl as a one-stellar component system, and finally discuss our results. For a more detailed description we refer to \\cite[Battaglia (2007)]{battaglia2007} and \\cite[Battaglia et al. (2008)]{battaglia2008}. ",
        "conclusions": "The combined information of velocity and metallicity from a large sample of FLAMES spectra of individual stars in Scl has allowed us to separate the different kinematics of the two stellar populations in Scl, and to carry out a more detailed kinematic modelling which allows us to partially break the mass-anisotropy degeneracy.  Under the hypotheses explored for the behaviour of the velocity anisotropy $\\beta(r)$, the two-component modelling indicates that a cored profile is slightly favoured over a cusped profile, as the latter tends to over-predict the central values (at $R \\lesssim$ 0.1 kpc) of the MR velocity dispersion. It could be argued that the relatively large central dispersion predicted for the MR stars from the cusped profile could be lowered by allowing the velocity anisotropy of the MR population to be tangential at $R \\lesssim$ 0.1 kpc. However it should be noted that in order to reproduce the rapid decline of the MR velocity dispersion profile, the MR $\\beta(r)$ needs to become very radial already on the scale of 0.3-0.4 kpc, given the observed spatial distribution of MR stars. Observational determinations of $\\beta$, either from direct measurement or from distribution functions, would be the most desirable (and definitive) solution to this issue.  \\cite[Strigari et al. (2007)]{strigari2007} showed that, independently of $\\beta(r)$ and the total mass distribution of dSphs, the mass enclosed within 0.6 kpc (M$_{\\rm 0.6}$) can be derived more accurately than at any other distance from the centre. However, given that the evolution of dSphs is most likely to be governed by their {\\it total} mass and that dynamical analyses provide only the {\\it mass enclosed within the last measured point}, it is important to have reliable determinations of the mass enclosed as farther out as possible. Our analysis shows that, independently of the adopted mass profile, the best fits from our two-component modelling yield consistent values for the mass enclosed within the last measured point, at 1.8kpc.  \\begin{figure}[!ht] \\begin{center} \\includegraphics[width=0.5\\textwidth]{mv_vs_mrax.eps} \\includegraphics[width=0.45\\textwidth]{mass_w_r.eps} \\caption{Left: Mass within the last measured point ($r_{\\rm max}$) versus absolute $V$ magnitude, M$_V$, for MW dSphs. The masses are from \\cite[Walker et al. 2007]{walker2007} (squares), except for Scl which is from our two-components modelling (asterisk); M$_V$ is from \\cite[Irwin \\& Hatzidimitriou (1995)]{ih1995}. The arrows indicate that the masses measured for Sextans and Fornax are likely to increase given that these objects have been surveyed just to 0.3 and 0.6 tidal radii. The solid line indicates a mass of 4$\\times 10^7$ \\sm. Right: Mass enclosed within $r$ for NFW halos of virial masses 2,4,8,20,40$\\times 10^8$ \\sm (solid lines), i.e. the range found by \\cite[Walker et al. (2007)]{walker2007} to fit the observed l.o.s. velocity dispersion profile of their sample of 7 dSphs. Concentrations are derived using the formula in \\cite[Koch et al. (2007)]{koch2007}, extrapolated from \\cite[Jing \\& Suto (2000)]{jing2000}. The dashed lines enclose masses between 2-7$\\times 10^7$ \\sm; the dotted line is placed at a radius of 0.6 kpc. }\\label{fig:mass} \\end{center} \\end{figure}  Finally we would like to comment on the issue that dwarf galaxies might be embedded in haloes of similar mass ($\\sim 4 \\times 10^7$ \\sm) as has been claimed in the literature. Figure~\\ref{fig:mass}(left) shows a revisited version of the mass enclosed within the last measured point versus absolute magnitude plot. The masses plotted are the ones derived in \\cite[Walker et al. (2007)]{walker2007} for a sample of dSphs, because they are derived from the most extended velocity samples and from a homogeneous analysis. For Scl we use our own determination, which is the most accurate so far. The figure shows that the DM mass of Scl is now much larger than the typical DM mass derived so far for the other dSphs. Also Fornax deviates considerably from the relation. In the plot shown here, the masses of Fornax and Sextans were derived from data extending out to $\\sim$0.6 and 0.3 tidal radii, respectively. Given that Fornax and Sextans have been recently mapped much farther out (Fornax out to $\\sim$ 1 tidal radius: Battaglia et al. 2006; Sextans out to $\\sim$ 0.8 tidal radius: Battaglia et al. in preparation), their revised masses are likely to increase quite significantly, and introduce even larger dispersion in this plot. In some cases, M$_{\\rm 0.6}$ has been used in this kind of analysis, instead than the mass enclosed within the last measured point. However, Fig.~\\ref{fig:mass}(right) shows that M$_{\\rm 0.6}$ is very insensitive to the total mass of the dSph. In fact, NFW haloes of virial masses between 2-40 $\\times 10^8$ \\sm (i.e. in the range found by \\cite[Walker et al. 2007]{walker2007} to fit their samples of dSphs) give remarkably similar values of M$_{\\rm 0.6}$, probably because at such a small distance there is not much sensitivity to the total mass. Therefore there is no indication of a common mass scale amongst the dSphs in the Local Group, while perhaps one may draw the conclusion of the existence of a minimum mass for the formation of dSphs."
    },
    "0807/0807.2102_arXiv.txt": {
        "abstract": "We discuss the total mass and dark halo properties of the Small Magellanic Cloud (SMC) for reasonable $V$-band stellar-to-mass-to-light ratios ($M_{\\rm s}/L_{\\rm V}$) based on the high-resolution neutral hydrogen (HI) observations of the SMC. We adopt both the Burkert and the NFW profiles for the dark matter halo of the SMC in order to model the radial profile of the observed rotation curve. We show that the mass ($M_{\\rm dm}$) and the mean density (${\\rho}_{\\rm dm}$) of the dark matter  halo within the optical radius ($\\sim 3$ kpc) of the SMC can be significantly  low for $M_{\\rm s}/L_{\\rm V}=2-4$ reasonable for a galaxy dominated by older stellar populations. For example, the estimated $M_{\\rm dm}$ and ${\\rho}_{\\rm dm}$ are $7.5 \\times 10^7 {\\rm M}_{\\odot}$  and $6.7 \\times 10^{-4} {\\rm M}_{\\odot}$ pc$^{-3}$, respectively, for $M_{\\rm s}/L_{\\rm V} = 3.8$. The models with lower $M_{\\rm s}/L_{\\rm V}$ ($<1.0$) can show higher $M_{\\rm dm}$ and ${\\rho}_{\\rm dm}$ yet  have difficulties in reproducing the inner rotation curve profile both for the Burkert and the NFW profiles. The  Burkert profile with  a larger core radius ($>1$ kpc) (thus a low density) and a large mass (${\\rm M}_{\\rm dm} > 3 \\times 10^9 {\\rm M}_{\\odot}$)  can be better fit to the observed slowly rising rotation curve than the NFW one for a reasonable  $M_{\\rm s}/L_{\\rm V} \\sim2$. This means that most of the dark halo mass can be initially located outside the optical radius in the SMC and thus that the dark halo would have already lost a significant fraction of its original mass due to the strong tidal interactions with the Galaxy and possibly with the Large Magellanic Cloud (LMC). We suggest that the dark mater halo of the SMC  is likely to have the initial total mass and core radius as large as, or larger than,  $6.5 \\times 10^9 {\\rm M}_{\\odot}$ and $3.2$ kpc, respectively. We discuss limitations of the present model in estimating the total mass of the SMC. ",
        "introduction": "Observational studies of the neutral hydrogen (HI) gas in  the Magellanic Clouds (MCs) have played a key role in the understanding of many aspects  of their formation, such as the formation of the Magellanic Stream (MS; Mathewson et al. 1974; Putman et al. 1998), gaseous kinematics of the Large Magellanic Cloud (LMC; Staveley-Smith et al. 2003),  and the formation of the Magellanic Bridge (MB; Muller et al. 2004). These observations furthermore have attracted much attention from theoretical studies of the roles of LMC-Small Magellanic Cloud (SMC)-Galaxy interactions in the formation of the MS (e.g., Murai \\& Fujimoto 1980), the importance of the last strong LMC-SMC interaction in the formation of the peculiar HI morphology (Bekki \\& Chiba 2007), and the origin of the bifurcated kinematics in the MB (e.g., Muller \\& Bekki 2007). Recent observations of the structural and kinematical properties of the HI gas in the SMC have provided valuable information on the mass distribution  of the SMC (Stanimirovi\\'c et al.  2004, S04).  \\begin{table*} \\centering \\begin{minipage}{185mm} \\caption{Parameter vales for the representative models.} \\begin{tabular}{cccccccc} {Model \\footnote{The mass model for the dark matter halo of the SMC: ``NDM'', ``B'', and ``NFW'' represent a model with no dark matter (or with no specific dark matter profiles considered),  the Burkert profile for the radial density profile of the dark matter halo, and NFW one, respectively.}} & {$i$ ($^{\\circ}$) \\footnote{The inclination angles of the gas disk in the SMC.}} & {$M_{\\rm S}/L_{\\rm V}$ \\footnote{The $V$-band stellar-to-mass-to-light-ratio.}} & {$M_{\\rm dm}$ (${\\rm M}_{\\odot}$) \\footnote{The total  mass of the dark matter halo in units of $10^9 {\\rm M}_{\\odot}$.}} & {$r_{\\rm 0}$ (kpc) \\footnote{The scale radius in the Burkert profile.}} & {$r_{\\rm s}$ (kpc) \\footnote{The scale radius in the NFW profile.}} & {$c$ \\footnote{$c$ ($=r_{\\rm vir}/r_{\\rm s}$) parameter in the NFW profile.}} & Comments \\\\ NDM1 &  40 & 3.8  & - &  - & - &  -  & without dark matter halo \\\\ NDM2 &  30 & 3.8  & - &  - & - &  -  & lower inclination angle \\\\ NDM3 &  40 & 2.3  & - &  - & - &  -  & lower $M_{\\rm s}/L_{\\rm v}$ \\\\ B1 &  40 & 2.3  & 6.5 & 3.2 & - &  -  & Burkert profile \\\\ B2 &  40 & 2.3  & 79.3 & 9.5 & - &  -  & \\\\ B3 &  40 & 0.6  & 79.3 & 9.5 & - &  -  & very low $M_{\\rm s}/L_{\\rm v}$\\\\ B4 &  40 & 1.0  & 79.3 & 9.5 & - &  -  & \\\\ NFW1 &  40 & 2.3  & 6.5 & - & 2.5 &  10  & NFW profile\\\\ NFW2 &  40 & 2.3  & 6.5 & - & 5.1 & 5  &  \\\\ NFW3 &  40 & 2.3  & 79.3 & - & 17.7 & 5  &  \\\\ \\end{tabular} \\end{minipage} \\end{table*}    One of intriguing results by S04 is that the HI gas in the SMC has  rotational kinematics with the maximum circular velocity ($V_{\\rm c}$) of $\\sim 60$ km s$^{-1}$. This kinematics is in a striking contrast to the stellar kinematics with no or little  rotation (e.g., Harris \\& Zaritsky 2006), which recently have led Bekki \\& Chiba (2008) to suggest that the SMC have experienced a major merger event long time ago. Another intriguing result in S04 is that $V_{\\rm c} \\sim 60$  km s$^{-1}$, corresponding to the total mass of $2.4 \\times 10^9 {\\rm M}_{\\odot}$, can be explained by the mass distribution of baryonic components (i.g., gas and stars) without a dark matter halo for the central 3 kpc: this result however depends strongly on the adopted stellar distribution. These properties certainly make the SMC a highly unusual  dwarf galaxy. However, no detailed dynamical studies have been conducted to  explain the origin of such unusual dark-matter halo.  The purpose of this paper is thus to discuss  the existence of dark matter within the central 3 kpc of the SMC, based on the detailed models for the mass distributions of the dark matter halo and the baryonic components. We derive the possible total mass ($M_{\\rm dm}$) and mean density (${\\rho}_{\\rm dm}$) of the dark matter halo within the optical radius of the SMC required to reproduce the observed HI rotation curve. We discuss whether the modeled distributions  of   the dark mater halo in the SMC can be consistent with the observed slowly rising rotation curve. We also discuss a possible initial  mass of the SMC before the interaction with the Galaxy and the LMC.   This paper focuses on implications and interpretation of observational results presented in S04. We stress that the observed rotation curve of the SMC depends on several adopted parameters: inclination, position angle, the disk thickness and shape of the SMC. In addition, possible radial gas-flows in the SMC, caused by the LMC-SMC-Galaxy interactions, can result in non-circular motions of HI gas and make the estimate of the rotation curve even harder. Some of the adopted parameters can not be directly determined by observations,  because the 3D HI structure remains observationally unclear for the SMC. Since observational uncertainties involved in the derivation of the rotation curve have already been discussed in S04, we just briefly summarize them in this paper and explain limitations they impose on the derivation of the dark matter halo. We also point out that the total mass of the SMC derived in S04 agrees well with the estimate based purely on the stellar kinematics (Harris \\& Zaritsky 2006). ",
        "conclusions": "We have investigated the mass ($M_{\\rm dm}$) and the mean density  (${\\rho}_{\\rm dm}$) of the dark matter  halo within the optical radius ($\\sim 3$ kpc) of the SMC based on the high-resolution HI observations of the SMC. Although there are some limitations of the present models in estimating the total mass of the SMC solely from the observed HI rotation curve by S04, we have tried to derive the  mass and best possible dark matter models. The results and their implications are described as follows.  (1) The total mass, mean density, and  distribution of the dark matter halo depend strongly on the adopted $M_{\\rm s}/L_{\\rm V}$. For example, the estimated $M_{\\rm dm}$ and ${\\rho}_{\\rm dm}$ are $7.5 \\times 10^7 {\\rm M}_{\\odot}$ and $6.7 \\times 10^{-4} {\\rm M}_{\\odot}$ pc$^{-3}$, respectively, for $M_{\\rm s}/L_{\\rm V} =3.8$. This implies that if the SMC is dominated by older stellar populations, the SMC has a low-density dark matter halo described  by B95 rather than by the NFW.  (2) The models with lower $M_{\\rm s}/L_{\\rm V}$ ($<1$) can show higher $M_{\\rm dm}$ and ${\\rho}_{\\rm dm}$ yet have difficulties in reproducing the observed rotation curve, in particular,  the inner part ($R<1$ kpc) of the curve, both for the Burkert and the NFW profiles. This implies that the SMC is unlikely to be dominated by younger stellar populations.   (3) The Burkert profile can fit better the observed HI rotation curve of the SMC than the NFW profile for a reasonable $M_{\\rm dm}$ and $M_{\\rm s}/L_{\\rm V}$ for reasonable   $M_{\\rm s}/L_{\\rm V}$ ($\\sim 2$). However the NFW profiles with low-densities, small $c$ ($\\sim 5$), and large total masses ($M_{\\rm dm} >   10^{10} {\\rm M}_{\\odot}$) can be also fit well the observations, though such low-density NFW profiles are highly unlikely for low-mass dwarfs like the SMC. The SMC can be the closest example among dIrr's in which rotation curves are  difficult to fit to the mass profiles of dark matter halos predicted by hierarchical clustering models.  (4) The dark matter halo of the SMC is likely to have the initial total mass of ${\\rm M}_{\\rm dm} \\sim  6.5 \\times 10^9 {\\rm M}_{\\odot}$ and the core radius of $r_{\\rm 0} \\sim 3.2$  kpc, which means that the SMC could be much more massive than what was previously thought.   (5) Orbital evolution models of the SMC and the LMC,  on which formation models of the MS and the MB  are based, would need to be  revised by considering the derived much larger total mass of the SMC. The observed large concentration of the baryonic components in the SMC is due to its unique formation history and is very different from  those of other dIrr's."
    },
    "0807/0807.0017_arXiv.txt": {
        "abstract": "We report the spectrophotometric investigation of a planetary nebula sample located at the disk-bulge interface of the Milky Way. The main goal of this work was to determine the galactocentric distance where, according to the intermediate mass population, bulge and disk properties separate. In order to derive such distance, new abundances were derived for a sample of PN located at this region, and the results were combined with additional data from the literature. The abundance analysis indicates a chemical abundance distribution similar to that derived from bulge stars, as already pointed out by other authors.  Statistical distance scales were then used to study the distribution of chemical abundances across the disk-bulge connection. A Kolmogorov-Smirnov test was used to find the distance in which the chemical properties of these regions better separate, resulting in a best value of 2.9 kpc to define the inner limit of the disk. ",
        "introduction": "Chemical properties of the galactic bulge and inner disk should intersect in a given region. In this region bulge characteristics such as diversity in chemical abundances should meet disk properties such as the inner edge of the radial gradient of chemical abundances.  The main goal of this work is to separate intermediate mass populations from the bulge and inner disk. This was made through the analysis of the abundance distribution of planetary nebulae (PN) located in the inner Galaxy. New abundances for a PN sample were derived, then added to data from the literature for other objects. These results were combined and then statistically analyzed, in order to estimate a galactocentric distance where bulge and disk characteristics better separate. Ultimately, we intend to provide an additional constraint to chemical evolution models for the Galaxy. ",
        "conclusions": ""
    },
    "0807/0807.0684_arXiv.txt": {
        "abstract": "The QCD phase tansition has important consequences in the context of both the early universe as well as compact stars. Such transitions are being studied for high temperature and small chemical potential scenario in the laboratory. There are also plans to study systems with large chemical potential and small temperatures. Here we have reviewed the role of strange quark matter and the phase transition in all the above scenarios. ",
        "introduction": "\\label{sc.intro}  The study of strongly interacting system is an active area of research. The existence of quark-hadron phase transition, order of transition, signature and consequences of such transitions in different scenarios are some of the topics which are being studied both experimentally as well as theoretically. The fact that the QCD perturbative series shows poor convergence except for very small coupling at very high temperatures($ \\alpha_s < 0.5$ $T \\sim 10^{5} T_c$), makes the calculations more difficult. On the other hand, the fact that nature has provided us with two extreme scenario in the phase diagram, makes this area of research more intriguing as well as interesting.  One of the natural scenario is the large T with vanishingly small chemical potential - the scenario which is supposed to have existed in the Early Universe and a similar situation is supposed to arise in the ultrarelativistic heavy ion collisions - the central collision in the central rapidity region, ideal hydrodynamics and longitudinal expansion mimicking the hubble expansion.  Off course there are differences, for example, Hubble expansion time scale is much larger than strong interaction scale whereas in ultrarelativistic collision expansion rate is comparable to the strong interaction.  The second natural scenario is the large chemical potential and small temperature limit which is expected to exist inside neutron stars. At present we do not know whether quark matter does exists inside the neutron stars. But with so many observational programmes going on it is really worth to study the possibilities of exsistence of such matter in the universe and especially inside the neutron star. This may lead us to that unique signature, the observational detection of which will resolve our doubt.  At present there exist large number of observational data on mass-radii of neutron stars. But, except very soft equation of states most of the other equation of states can explain these static properties within tthe error bars \\cite{schafer}. So it is necessary to study the different dynamical features of quark matter in its various forms. ",
        "conclusions": "QCD transition is one of the mose interesting areas of research. We have discussed the impact of the QCD transition in natural scenario of neutron stars and early universe. Though, an observational signature of phase transition and strange quark matter in nature is still elusive, the study of these phenomena is extremely important for a better understanding of the physics of strongly interacting matter. Moreover, there are many questions remain to be answered before we can make any conclusions."
    },
    "0807/0807.0151_arXiv.txt": {
        "abstract": "We discuss the problem of ultra high energy nuclei propagation in astrophysical backgrounds. We present a new analytical computation scheme based on the hypothesis of continuos energy losses in a kinetic formulation of the particles propagation. This scheme enables the computation of the fluxes of ultra high energy nuclei as well as the fluxes of secondaries (nuclei and nucleons) produced by the process of photo-disintegration suffered by nuclei. ",
        "introduction": "A clear understanding of the composition features of the UHECR spectrum passes through a detailed study of the propagation of UHE protons and nuclei in astrophysical backgrounds. In particular, UHE nuclei interacting with such backgrounds suffer two main process: (i) photo-disintegration, giving rise to the production of secondary nucleons (hereafter protons) and lighter nuclei, (ii) pair production, that produces a depletion in the nucleus energy. The study of the photo-disintegration process of UHE nuclei has already attracted the attention of several authors that have studied the propagation of these particles with the aim of Monte Carlo (MC) simulations \\cite{Stecker76},\\cite{Epele98},\\cite{Stecker99},\\cite{Bertone02},\\cite{Yamamoto04},\\cite{Allard06} (for a detailed list of references see also \\cite{NoiPRD}). One of the pioneering works in this field was performed in \\cite{Stecker76}, where the authors introduced, for the first time, a very useful parameterization of the photo-disintegration cross-section, this result was up-dated in \\cite{Stecker99}. Recently, a more detailed parameterization of the photo-disintegration cross-section was presented in \\cite{Allard05}, this computation confirms the findings of \\cite{Stecker76} and \\cite{Stecker99} at energies below $10^{20}$ eV, showing some deviations at the extreme energies. The studies performed so far have shown the important role that the Infrared/Optical Background (IRB) plays in nuclei propagation, this is an important result that, on the other hand, complicates the understanding because of the less precise determination of such background respect to the well known Cosmic Microwave Background (CMB).  In this paper we will briefly present a new analytical computation scheme that, based on a kinetic approach in the approximation of continuos energy losses, enables the computation of UHE nuclei propagation in any background, with the determination of the expected fluxes on earth of the primary injected particles as well as of all secondaries (nuclei and nucleons) produced in the photo-disintegration process. Following \\cite{NoiPRD}, in order to explain the method proposed, we will restrict the discussion to the case of the CMB radiation alone. ",
        "conclusions": ""
    },
    "0807/0807.3963_arXiv.txt": {
        "abstract": "The critical star formation rate density required to keep the intergalactic hydrogen ionised depends crucially on the average rate of recombinations in the intergalactic medium (IGM). This rate is proportional to the clumping factor $C \\equiv \\langle \\rho_{\\rm{b}}^2 \\rangle_{\\rm IGM} / \\langle \\rho_{\\rm{b}} \\rangle^2$, where $\\rho_{\\rm{b}}$ and $\\langle \\rho_{\\rm{b}} \\rangle$ are the local and cosmic mean baryon density, respectively and the brackets  $\\langle \\rangle_{\\rm IGM}$ indicate spatial averaging over the recombining gas in the IGM. We perform a suite of cosmological smoothed particle hydrodynamics simulations that include radiative cooling to calculate the volume-weighted clumping factor of the IGM at redshifts $z \\ge 6$. We focus on the effect of photo-ionisation heating by a uniform ultra-violet background and find that photo-heating strongly reduces the clumping factor because the increased pressure support smoothes out small-scale density fluctuations. Photo-ionisation heating is often said to provide a negative feedback on the reionisation of the IGM because it suppresses the cosmic star formation rate by boiling the gas out of low-mass halos. However, because of the reduction of the clumping factor it also makes it easier to keep the IGM ionised. Photo-heating therefore also provides a positive feedback which, while known to exist, has received much less attention. We demonstrate that this positive feedback is in fact very strong. Using conservative assumptions, we find that if the IGM was reheated at $z \\gtrsim 9$, the observed population of star-forming galaxies at $z \\approx 6$ may be sufficient to keep the IGM ionised, provided that the fraction of ionising photons that escape the star-forming regions to ionise the IGM is larger than $\\sim 0.2$. ",
        "introduction": "The absence of a Gunn-Peterson trough in the majority of the observed absorption spectra towards high-redshift quasars suggests that the reionisation of intergalactic hydrogen was completed at a redshift $z \\gtrsim 6$ (see \\citealp{Fan:2006} for a recent review) and that it remained highly ionised ever since. Current observational estimates of the ultra-violet (UV) luminosity density at redshifts $z \\lesssim 6$ (e.g.~\\citealp{Stanway:2003}; \\citealp{Lehnert:2003}; \\citealp{Bunker:2004}; \\citealp{Bouwens:2004}; \\citealp{Yan:2004}; \\citealp{Sawicki:2006}; \\citealp{Bouwens:2006}; \\citealp{Mannucci:2007}; \\citealp{Oesch:2008}; \\citealp{Bouwens:2008}), on the other hand, may imply star formation rate (SFR) densities several times lower than the critical SFR density required to keep the intergalactic medium (IGM) ionised (but see, e.g., \\citealp{Stiavelli:2004}; \\citealp{Malhotra:2005}; \\citealp{Panagia:2005}). Taken at face value, these low SFR densities pose a severe challenge to commonly employed theoretical models in which the observed population of star-forming galaxies is the only source of ionising radiation in the high-redshift Universe. \\par There are, however, large uncertainties associated with both the observationally inferred (see, e.g., the comprehensive analysis of ~\\citealp{Bouwens:2007}) and the critical SFR densities. The critical SFR density, \\begin{eqnarray} \\dot{\\rho}_* &\\approx&  0.027 \\Msun \\invyr \\invMpc \\nonumber \\\\ &\\quad& \\times {f_{\\rm{esc}}}^{-1} \\left ( \\frac{C}{30} \\right )  \\left ( \\frac{1+z}{7} \\right )^3  \\left ( \\frac{\\Omega_{\\rm{b}}h_{70}^2}{0.0465} \\right )^2, \\label{Eq:CriticalSfr} \\end{eqnarray} here rescaled to match the most recent WMAP estimate for the cosmic baryon density (\\citealp{Komatsu:2008}), has been derived by \\cite{Madau:1999} using an early version of the \\cite{Bruzual:2003} population synthesis code, assuming a Salpeter initial stellar mass function (IMF) and solar metallicity. It results from simply equating the spatially averaged rate at which ionising photons are emitted into the IGM to the spatially averaged rate at which the intergalactic gas recombines. Eq.~\\ref{Eq:CriticalSfr} is therefore incapable of addressing a number of potentially important physical effects. Some ionising photons will, for instance, redshift below the ionisation threshold before ionising and some ionising photons will have been emitted longer than a recombination time ago upon impact with a neutral atom, so that equating instantaneous rates is not appropriate and one even may have to take source evolution into account. It is therefore important to keep in mind that Eq.~\\ref{Eq:CriticalSfr} is likely only accurate within factors of a few. \\par The critical SFR is inversely proportional to the escape fraction $f_{\\rm{esc}}$, i.e. the fraction of ionising photons produced by star-forming galaxies that escape the interstellar medium (ISM) to ionise the IGM, and proportional to the average recombination rate in the IGM. The latter is parametrised using the dimensionless clumping factor $C \\equiv \\langle \\rho_{\\rm{b}}^2\\rangle_{\\rm IGM}  / \\langle \\rho_{\\rm{b}}\\rangle^2$, where $\\rho_{\\rm{b}}$ is the baryon density, $\\langle \\rho_{\\rm{b}}\\rangle$ is the mean baryon density of the Universe and the brackets $\\langle \\rangle_{\\rm IGM}$ indicate spatial averaging over the gas constituting the recombining IGM. Under the assumption of a uniformly ionised IGM, the clumping factor expresses the spatially averaged number of recombinations occurring per unit time and unit volume in the ionised IGM, relative to that in gas at the cosmic mean density $\\langle \\rho_{\\rm{b}}\\rangle$. A larger escape fraction implies a smaller critical SFR density, as more photons are available to ionise the IGM. On the other hand, a larger clumping factor implies a larger critical SFR density since more ionising photons are required to compensate for the increased number of recombinations. \\par Most observational studies that compare\\footnote{Note that although the critical SFR density is sensitive to the IMF, this comparison is insensitive to the IMF provided the same IMF is used to compute the critical and observationally derived SFR densities. This is because the UV luminosity density is dominated by the same massive stars that are responsible for the emission of ionising photons with energies $>13.6 \\eV$ (\\protect\\citealp{Madau:1999}).} the SFR density derived from estimates of the UV luminosity density at redshift $z \\approx 6$ to the critical SFR density assume an escape fraction $f_{\\rm esc} \\lesssim 0.5 $ and a clumping factor $C = 30$. While a variety of both observational and theoretical studies (e.g. \\citealp{Inoue:2006} and references therein; \\citealp{Razoumov:2006}; \\citealp{Gnedin:2008a}) have ruled out larger escape fractions, the estimate for the clumping factor comes from a single cosmological simulation performed more than 10 years ago (\\citealp{Gnedin:1997}). It is on the basis of these values for the escape fraction and the clumping factor that the observed population of galaxies has been found to be incapable of keeping the intergalactic hydrogen ionised, forming massive stars at a rate which is up to an order of magnitude lower than required by Eq.~\\ref{Eq:CriticalSfr}. \\par It has been pointed out that this discrepancy between the inferred and critical SFR densities could be resolved if the employed clumping factor were too high (e.g.~\\citealp{Sawicki:2006}; see also the discussion in \\citealp{Bouwens:2007}). Indeed, in most (but not all) of the more recent theoretical theoretical studies (e.g.~\\citealp{Valageas:1999}; \\citealp{Miralda:2000}; \\citealp{Gnedin:2000a}; \\citealp{Haiman:2001}; \\citealp{Benson:2001}; \\citealp{Chiu:2003}; \\citealp{Iliev:2007}; \\citealp{Srbinovsky:2007}; \\citealp{Kohler:2007}; \\citealp{Bolton:2007}; \\citealp{Maio:2007}; \\citealp{Furlanetto:2008}) significantly lower clumping factors were derived. On the other hand, it is sometimes emphasised that simulations underestimate the clumping factor, due to a lack of resolution (see, e.g., \\citealp{Madau:1999}). In this work we perform a set of cosmological Smoothed Particle Hydrodynamics (SPH) simulations that include radiative cooling and photo-ionisation by a uniform UV background in the optically thin limit to study the clumping factor of the IGM. \\par We focus on the effect of photo-ionisation heating on the evolution of the clumping factor. Previous investigations of the impact of photo-heating on the reionisation of the IGM have almost exclusively come to the conclusion that it acts as to provide a negative feedback. Photo-heating boils the gas out of the potential wells of dark matter (DM) halos with virial temperatures $T_{\\rm vir} \\lesssim 10^4\\K$ (e.g.~\\citealp{Thoul:1996}; \\citealp{Navarro:1997}; \\citealp{Barkana:1999}; \\citealp{Kitayama:2000}; \\citealp{Gnedin:2000b}; \\citealp{Dijkstra:2004}; \\citealp{Shapiro:2004}; \\citealp{Hoeft:2006}; \\citealp{Crain:2007}; \\citealp{Mesinger:2008}; \\citealp{Okamoto:2008}; \\citealp{Pawlik:2008}). This inhibits the formation of stars in these low-mass halos and thus decreases the ionising emissivity, which makes it more difficult to reionise the Universe. The same mechanism that reduces the number of ionising photons that are emitted into the IGM does, however, also affect the evolution of the clumping factor (e.g.~\\citealp{Haiman:2001}; \\citealp{Oh:2003}; \\citealp{Kuhlen:2005}; \\citealp{Wise:2005}; \\citealp{Furlanetto:2006}; \\citealp{Ciardi:2007}). \\par In this paper we demonstrate that photo-heating significantly lowers the clumping factor and hence the average recombination rate in the IGM.  While photo-ionisation heating undoubtedly impedes the production of ionising photons, our results imply that it also makes it much easier to keep the IGM ionised. \\par The paper is structured as follows. In Section~\\ref{Sec:Simulations} we give a detailed description of our set of simulations. In Section~\\ref{Sec:Results} we use our simulations to compute the clumping factor of the IGM. Finally, in Section~\\ref{Sec:Conclusions}, we discuss our results and their implications for the value of the critical SFR density. ",
        "conclusions": "\\label{Sec:Conclusions} Several observational studies have claimed that the star formation rate (SFR) density at redshift $z \\approx 6$ is smaller than the critical SFR density required to keep the intergalactic medium (IGM) ionised. In the absence of a large population of unseen sources of ionising radiation, this discrepancy between the two SFR densities would be in direct conflict with the high degree of ionisation inferred from the non-detection of a Gunn-Peterson trough in the majority of the line-of-sight spectra towards $z \\lesssim 6$ quasars. \\par The critical SFR density is inversely proportional to the spatially averaged fraction of ionising photons that escape into the IGM per unit time and proportional to the clumping factor $C \\equiv \\langle \\rho_{\\rm{b}}^2 \\rangle_{\\rm IGM} / \\langle \\rho_{\\rm{b}}\\rangle^2$, a measure for the average recombination rate in the IGM. One may therefore ask whether the discrepancy between the observed and critical SFR densities could be resolved by changing the assumptions about the values of either of these two quantities. In this work we considered the hypothesis that most observational studies overestimate the critical SFR density because they employ a clumping factor that is too large. \\par We re-evaluated the clumping factor, analysing the gas density distributions in a set of cosmological smoothed particle hydrodynamics simulations that include radiative cooling and photo-ionisation by a uniform UV background in the optically thin limit. The clumping factor of the IGM depends critically on the definition of which gas is considered to be part of the IGM. Following \\cite{Miralda:2000}, we assumed that all gas with densities below a threshold density constitutes the IGM and computed the clumping factor as a function of this threshold density. In addition, we introduced two physically well-motivated definitions, $C_{100}$, the clumping factor of gas with overdensities $\\Delta < 100$ and $C_{-1}$, the clumping factor of gas with proper densities below  $n_{\\rm H} = n_{\\rm{H}}^* \\equiv 10^{-1}\\cmci $, our threshold density for the onset of star formation. \\par By comparing simulations that include photo-ionisation by a uniform UV background to one that does not, we showed that photo-heating strongly influences the evolution of the clumping factor of the IGM. Photo-ionisation heating expels the gas from within halos of virial temperatures $T_{\\rm vir} \\lesssim 10^4 \\K$ and prevents its further accretion by raising the Jeans mass in the IGM. By suppressing the formation of stars in these low-mass halos, photo-heating from reionisation decreases the rate at which ionising photons are emitted into the IGM and is therefore correctly said to exert a negative feedback on the reionisation process. The fact that photo-heating also leads to a decrease in the clumping factor and hence provides a strong positive feedback by making it easier to keep the IGM ionised, is however often overlooked (but see, e.g., \\citealp{Haiman:2001}; \\citealp{Oh:2003}; \\citealp{Kuhlen:2005}; \\citealp{Wise:2005}; \\citealp{Furlanetto:2006}; \\citealp{Ciardi:2007}). \\par At redshift $z = 6$, we find that $C_{-1} \\approx 6$ and $C_{100} \\approx 3$ and that these values are insensitive to the redshift $z_{\\rm r}$ at which the UV background is turned on, as long as $z_{\\rm r} \\gtrsim 9$. These values for $C_{-1}$ and $C_{100}$ are at least three times smaller than they would be in the absence of photo-heating. We demonstrated that our default simulation is converged at $z = 6$ with respect to the employed resolution. It is converged with respect to changes in the box size for threshold overdensities $\\log_{10}\\Delta_{\\rm thr} \\lesssim 2$. In Appendix~\\ref{Sec:Appendix} we provide fits to the evolution of the clumping factor for various (over-)density thresholds and reheating redshifts. There we also compare the probability density function (PDF) of the gas densities at $z = 6$ obtained from our default simulation to the widely used fit provided by \\cite{Miralda:2000}. We update their fitting parameters to best fit the PDF from our default simulation. Finally, we compared our results for the clumping factor of the IGM to those obtained in previous works. \\par Since even our most conservative estimate for the clumping factor ($C_{-1} \\approx 6$) is five times smaller than the clumping factor that is usually employed to determine the capacity of star-forming galaxies to keep the $z = 6$ IGM ionised, our results may have important implications for the understanding of the reionisation process. Setting $C = 6$ in Eq.~\\ref{Eq:CriticalSfr}, the critical SFR density becomes $\\dot{\\rho}_* = 0.005~f_{\\rm esc}^{-1}\\Msun \\invyr \\invMpc$. This is smaller than recent observational estimates for the SFR density at $z \\approx 6$, $\\dot{\\rho}_* = 0.022 \\pm 0.004\\Msun \\invyr \\invMpc$ (integrated to the observed $z \\approx 6$ faint-end limit $L > 0.04 ~L_{z=3}^\\ast$ and dust-corrected; \\citealp{Bouwens:2007}), provided that $f_{\\rm esc} \\gtrsim 0.2$. \\par Our study thus suggests that the observed population of star-forming galaxies may be capable of keeping the IGM ionised, relaxing the tension between observationally inferred and critical SFR density in view of the observation of a highly ionised IGM at redshifts  $z \\lesssim 6$. We note that at $z \\approx 7$, the SFR density is estimated to be $\\dot{\\rho}_* = 0.004 \\pm 0.002 \\Msun \\invyr \\invMpc$ (integrated to the observed $z \\approx 7$ faint-end limit $L > 0.2 ~L_{z=3}^\\ast$ and dust-corrected; \\citealp{Bouwens:2008}), whereas the critical SFR density is $\\dot{\\rho}_* = 0.008~f_{\\rm esc}^{-1} \\Msun \\invyr \\invMpc$ (using $C = 6$). The observed population of star-forming galaxies at $z \\approx 7$ is therefore not able to keep the IGM ionised. If the Universe were ionised by this redshift, then the sources that were responsible remain to be discovered. \\par We caution the reader that the comparison of the critical and observed SFRs is subject to considerable uncertainty. First, the SFR inferred from UV galaxy counts probably underestimates the true SFR, because these counts miss UV galaxies fainter than the faint-end limit implied by their sensitivities. These galaxies may, however, significantly contribute to the UV luminosity density if the faint-end slope of the UV luminosity function is sufficiently steep. Complementary estimates of the high-redshift star formation rate based on measurements of the high-redshift ($z = 4-7$) gamma ray burst rate (\\citealp{Yuksel:2008}) and measurements of the Lyman-alpha forest opacity at redshifts $z \\sim 3$ (\\citealp{Faucher:2008}) indeed suggest $z\\sim 6$ SFRs that exceed those inferred from UV galaxy counts by factors of a few. \\par Second, the expression for the critical SFR (Eq.~\\ref{Eq:CriticalSfr}) is only approximate. As already mentioned in the introduction, this expression is based on equating the rate at which ionising photons escape into the intergalactic medium to the rate at which the intergalactic gas is recombining, both averaged in space. It neglects effects like the cosmological redshifting of photons below the ionisation threshold energy and evolution of the ionising sources during a recombination time. Because the cross-section $\\sigma_{\\rm HI} \\sim \\nu^{-3}$ for absorption of ionising photons by neutral hydrogen decreases with increasing photon frequency $\\nu$, these effects may become important for photons whose mean free path is comparable to the cosmic horizon. Our limit on the escape fraction required to keep the Universe at $z \\approx 6$ ionised may therefore only be accurate within a factor of a few. \\par Our simulations demonstrate that radiation-hydrodynamical feedback due to photo-ionisation heating plays a key role in shaping the properties of the IGM at redshifts $z \\gtrsim 6$.  We have studied the impact of photo-ionisation heating on the clumping factor of the IGM assuming a uniform ionising UV background in the optically thin limit. In reality the reionisation process will, however, be more complex. We have demonstrated the robustness of our conclusions with respect to uncertainties in the temperature of the IGM resulting from our simplified treatment of the reionisation heating, but there are other factors whose importance is more difficult to assess. \\par Our use of the optically thin approximation neglects self-shielding, a radiative transfer effect due to which halos that would otherwise be completely photo-evaporated could keep some of their gas (e.g., \\citealp{Kitayama:2000}; \\citealp{Susa:2004}; \\citealp{Dijkstra:2004}). Since the self-shielded gas remains neutral, it should be excluded when computing the clumping factor. Self-shielding becomes important for $N_{\\rm HI} \\gtrsim 10^{18} \\cmsi$, which for self-gravitating gas clouds corresponds to densities (\\citealp{Schaye:2001}) $n_{\\rm H} \\gtrsim 10^{-2} \\cmci (\\Gamma / 10^{-12} \\invs)^{-1}$, where $\\Gamma$ is the HI photo-ionisation rate.  In our discussion of the critical SFR we have conservatively adopted the clumping factor $C_{-1}$, defined using a threshold density $n_{\\rm H} = 10^{-1} \\cmci$, which is an order of magnitude larger than the critical density for self-shielding. \\par Self-shielding might affect our results because it may lower the speed with which halos are photo-evaporated. The work by \\cite{Iliev:2005b} (in combination with the work by \\citealp{Shapiro:2004}) shows that photo-evaporation times obtained in simulations that employ an optically thin UV background may differ by factors of a few from those obtained in detailed radiation-hydrodynamical simulations. However, if any halos would resist photo-evaporation much longer than predicted by our approximate treatment of photo-heating, then the clumping factor of the IGM would be even lower, because self-shielding locks the gas that would otherwise contribute to the clumping factor of the IGM in its neutral state. \\par If absorption by optically thick (self-shielded) clouds becomes important, then the mean free path of ionising photons may be set by the mean distance between these clouds (\\citealp{Zuo:1993}) instead of by the opacity of the diffuse IGM. In this case it might be appropriate to supplement the \\citealp{Miralda:2000} model for the computation of the average recombination rate with a more direct account for these clouds as discrete photon sinks (e.g., \\citealp{Iliev:2005a}; \\citealp{Ciardi:2006}). The effect of absorption by optically thick clouds is, however, largely degenerate with the ionising efficiency of the population of star-forming galaxies, as explained in \\cite{Iliev:2007}. If, on the one hand, these clouds are ionising sources themselves, then their contribution to the average recombination rate can be described by their escape fractions. If, on the other hand, these clouds do not host ionising sources, then their contribution to the average recombination can be accommodated by changes in the ionising efficiency of star-forming galaxies because of their biased clustering around these galaxies. The effects of optically thick clouds can therefore approximately be accounted for by adjusting the properties of the ionising sources. \\par Self-shielding is only one example of the physical effects that we are ignoring. The inclusion of metals and molecules, for instance, would increase the ability of the gas to cool, which may lead to an increase in the clumping factor (e.g.~\\citealp{Maio:2007}). In the presence of photo-heating and for the threshold densities we employ to compute the clumping factors we however expect this effect to be very small. A more efficient cooling due to metals and molecules may also enable star formation and associated kinetic feedback in low-mass haloes with virial temperatures $T_{\\rm vir} < 10^4 \\K$. Both are processes that we have ignored but which are likely to affect the clumping factor evolution (e.g.~\\citealp{Wise:2008}). \\par The evolution of the clumping factor will also depend on the morphology of the reionisation transition. Our treatment implicitly assumed that all gas with overdensities smaller than a threshold overdensity is uniformly ionised, while all gas with larger overdensities is fully neutral. This picture probably only applies to the late stages of reionisation, when individual ionised regions start to overlap and the only neutral gas that remains to be ionised is locked up in regions of high gas overdensities. Before overlap, other reionisation models may be more useful for the description of the clumping factor evolution. For example, \\cite{Furlanetto:2005} point out that if the large-scale dense regions are ionised first, the clumping factor may be somewhat larger than one would otherwise expect, because the photons are initially confined to these dense regions. Moreover, \\cite{Furlanetto:2008} show that in fossil ionised regions, that is, regions in which the gas freely recombines, the clumping factor will generally be smaller than for regions in photo-ionisation equilibrium, because the densest gas which contributes most to the clumping factor becomes neutral first. \\par The clumping factor is an important ingredient of (semi-)analytic treatments of reionisation. It would therefore be highly desirable to evaluate the approximations we have employed in our simplified treatment of the photo-heating process using large high-resolution radiation-hydrodynamical simulations of reionisation that include cooling by metals and molecules and feedback from star formation. At the moment such simulations are, however, not yet feasible. \\par In fact, current state-of-the-art radiative transfer simulations typically make use of clumping factors in their ``sub-grid'' modules because they lack the resolution to resolve the clumpiness of the gas directly. In addition, they typically do not include the effect of photo-heating. In fact, many radiative transfer simulations ignore hydrodynamics altogether and assume the gas to trace the dark matter. \\par With this work we hope to have presented a conservative assessment of the clumping factor of the post-reionisation IGM that may provide a useful input to future (semi-)analytic models and simulations of the reionisation process."
    },
    "0807/0807.3014_arXiv.txt": {
        "abstract": "In this paper, we study the feasibility of obtaining near-infrared spectra of bright extrasolar planets with the 2nd generation VLTI Spectro-Imager instrument (VSI), which has the required angular resolution to resolve nearby hot Extrasolar Giant Planets (EGPs) from their host stars. Taking into account fundamental noises, we simulate closure phase measurements of several extrasolar systems using four 8-m telescopes at the VLT and a low spectral resolution ($R = 100$). Synthetic planetary spectra from T. Barman are used as an input. Standard $\\chi^2$-fitting methods are then used to reconstruct planetary spectra from the simulated data. These simulations show that low-resolution spectra in the $H$ and $K$ bands can be retrieved with a good fidelity for half a dozen targets in a reasonable observing time (about 10 hours, spread over a few nights). Such observations would strongly constrain the planetary temperature and albedo, the energy redistribution mechanisms, as well as the chemical composition of their atmospheres. Systematic errors, not included in our simulations, could be a serious limitation to these performance estimations. The use of integrated optics is however expected to provide the required instrumental stability (around $10^{-4}$ on the closure phase) to enable the first thorough characterisation of extrasolar planetary emission spectra in the near-infrared. ",
        "introduction": "\\label{sec:intro}  %  Since the discovery by Mayor \\& Queloz\\cite{mayor95} of the first exoplanet around 51 Pegasi, the study of planetary systems has received an increasing attention, with the continuous development of new techniques. Among the direct detection techniques, interferometry is one of the most promising for the near future. It already provides the required angular resolution, but the dynamic range needs to be improved. The detection and characterization of extrasolar planets is one of the main science cases of the 2nd generation VLTI Spectro-Imager instrument (VSI)\\cite{VSIMalbet}.  At a distance of $a \\approx 0.05$ AU, most of the detected exoplanets, called hot Extrasolar Giant Planets (EGPs), receive from their parent star about $10^4$ times the amount of radiation intercepted by Jupiter from our Sun. So close to the star, they are also believed to be tidally locked such that half of the planet faces permanently the star while the other half stays in the dark. The fraction of incident light absorbed by the atmosphere -- parametrized by its albedo $A$ -- heats the planets. This heat is redistributed to the night side by strong wind and reradiated by the planet. Besides providing their heat source, the strong radiation illuminating hot EGPs also structures their atmosphere. It suppresses the convection to depths well below the photosphere, leading to a fully radiative photosphere across most of the day side\\cite{guillot96}. Whether silicate clouds can persist in the photosphere then results from the competition between the sedimentation and advective timescales. The sedimentation timescale of radiative photosphere is short and the winds are believed not to be strong enough to prevent dust from settling. Hence, the atmospheres of hot EGPs are pictured as being free of clouds. Finally to reproduce hot-Jupiter spectra, their chemical composition should be taken into account. As the opacity of each species regulates the emergent flux as a function of wavelength, their emergent spectra strongly depart from an ideal black body assumption. They display large molecular bands and spectral features (see Fig.~\\ref{fig:sudarsky}).  \\begin{figure}[tp] \\begin{center} \\begin{tabular}{cc} \\includegraphics[width=0.45\\textwidth]{sudarsky1.eps}& \\includegraphics[width=0.45\\textwidth]{sudarsky2.eps} \\end{tabular} \\end{center} \\caption[step] { \\label{fig:sudarsky} Model fo spectra for hot EGPs, with $T_{eq} = 1000$K and 1500K, orbiting a Sun-like star. Important changes in abundances and band strength are seen between both cases\\cite{sudarsky03} } \\end{figure}  The transit technique revolutionised the field of exoplanetology. It was a breakthrough to peer into hot EGPs' structure and to already give a glimpse about their composition. However, the error bars remain large to discriminate between different models and most of the planetary spectrum remains unknown. Particulary, the $1 - 2.4$ $\\mu$m spectrum is very rich in spectral information and may provide unprecedented constraints on our understanding of planetary atmosphere. VSI will have the ability to observe hot EGPs in $J$, $H$ and $K$ bands, with a low ($ R = 100$) or medium ($R = 1000$) resolution. Model fitting of low-resolution spectra gives a measurement of their albedo and test the cloud-free assumption. The phase dependance of the measured signal shall constrain the heat redistribution (through the temperature across the surface) and the weather conditions. At medium resolution, it will be possible to measure the abundance of CO and test the presence of CH$_4$. Contrarily to current characterisation techniques, VSI will not be restricted to transiting planets. The sample of favourable targets counts already 7 planets today and may count at least twice this number at the time of VSI will be in operation. Therefore, VSI will not only enhance our knowledge on few transiting planets, it will also allow statistical studies to be carried out and literally enable comparative exoplanetology.  The goal of this work is to study the feasibility of obtaining near-infrared emission spectra of bright extrasolar giant planets (EGPs) with VSI. In Sect.~\\ref{sec:meth}, we explain the method used for the simulation, from the choice of the targets to the simulations. Sect.~\\ref{sec:res} contains the results of the simulation, and Sect.~\\ref{sec:conclu} the discussion and the conclusions of the study. ",
        "conclusions": "\\label{sec:conclu}  From the results of the fit of the planetary spectra, it becomes evident that VSI will be a powerful tool to characterize hot EGPs. The simultaneous measurement of the contrast at various wavelengths provides an insight into the thermal, physical and dynamical structure of their atmospheres. In particular, the slope of the spectrum in $H$ and $K$ band directly informs on the presence of CH$_4$ in the planetary atmosphere, while the CO absorption feature around 2.3 $\\mu$m could also be detected from some of the selected targets. Furthermore, the repeated observations at various orbital phases provide an important information to constrain the heat distribution mechanisms by measuring the temperature and atmospheric composition around the planet.  We note that the simulated observations are more successful and constraining if the star-planet system is close to the observer ($ \\leq 20 pc$). For such targets, the VLTI angular resolution well matches the star-planet separation and the planet is bright enough to provide a good signal-to-noise ratio. It is thus recommended to choose the closest systems as the first targets of VSI.  In the light of these results, one can safely conclude that the prime criterion for the selection of additional targets is the magnitude of the host star: it drives the signal to noise ratio on the closure phases and therefore the quality of the fit of the planetary data. The semi-major axis of the planetary orbit, or more precisely, the temperature of the planetary companion, is of course another critical parameter, as well as its radius (which is generally unknown). From these observations, additional targets can be proposed for the VSI exoplanet sample such as HD 75286 b and HD 160691 d (from the hot Jupiter family) or 55 Cnc e (a hot Neptune). All give satisfactory results when repeating the above simulation procedure, yet with larger relative error bars on the measured planet/star contrast. A few other planets may be added to the list in the coming years.  Systematic errors, not included in our simulations, could be a serious limitation to these performance estimations. The use of integrated optics is however expected to provide the required instrumental stability (around $10^{-4}$) to enable the first thorough characterisation of extrasolar planetary spectra in the near-infrared.  The direct detection of hot EGPs is undoubtely one of the most challenging VSI programs. This program was already one of AMBER's goal. However VSI intrinsic design offers multiple improvements with respect to AMBER which rely on two main axes : \\begin{itemize} \\item Improvement in the observable signal-to-noise ratio and accuracy : with a combining core made of an integrated optics circuit in which the incoming beams are spatially filtered out and routed so that each of the 4UT beams are carefully interfered with each other, the intrinsic stability of VSI is much higher than a classical bulk optics solution like AMBER. Moreover, unlike AMBER, VSI has included in its study an internal fringe tracker located as close as possible to the science instrument in order to control the stabilisation of the fringes and allow cophasing. This fringe tracker significantly increases the signal-to-noise ratio on the closure phases\\cite{VSICorcione}. \\item Increasing the number of simultaneous points: using 4 telescopes allows 4 closure phases to be measured simultaneously while AMBER  only permits one such measurement. This simultaneity reduces time dependant drifts, improves the calibration and permits to constrain the flux ratio by 4 points in a single measurements instead of 1 for AMBER, improving dramatically the quality of the fit. \\end{itemize}"
    },
    "0807/0807.4402_arXiv.txt": {
        "abstract": "The interaction of a shock with a cloud has been extensively studied in the literature, where the effects of magnetic fields, radiative cooling and thermal conduction have been considered.  In many cases, the formation of fully developed turbulence has been prevented by the artificial viscosity inherent in hydrodynamical simulations. This problem is particularly severe in some recent simulations designed to investigate the interaction of a flow with multiple clouds, where the resolution of individual clouds is necessarily poor. Furthermore, the shocked flow interacting with the cloud has been assumed to be completely uniform in all previous single-cloud studies. In reality, the flow behind the shock is also likely to be turbulent, with non-uniform density, pressure and velocity structure created as the shock sweeps over inhomogenities upstream of the cloud (as seen in recent multiple cloud simulations). To address these twin issues we use a sub-grid compressible $k$-$\\epsilon$ turbulence model to estimate the properties of the turbulence generated in shock-cloud interactions and the resulting increase in the transport coefficients that the turbulence brings.  A detailed comparison with the output from an inviscid hydrodynamical code puts these new results into context.  Despite the above concerns, we find that cloud destruction in inviscid and $k$-$\\epsilon$ models occurs at roughly the same speed when the post-shock flow is smooth and when the density contrast between the cloud and inter-cloud medium, $\\chi \\ltsimm 100$. However, there are increasing and significant differences as $\\chi$ increases. The $k$-$\\epsilon$ models also demonstrate better convergence in resolution tests than inviscid models, a feature which is particularly useful for multiple-cloud simulations.  Clouds which are over-run by a highly turbulent post-shock environment are destroyed significantly quicker as they are subject to strong ``buffeting'' by the flow. The decreased lifetime and faster acceleration of the cloud material to the speed of the ambient flow leads to a reduction in the total amount of circulation (vorticity) generated in the interaction, so that the amount of vorticity may be self-limiting.  Additional calculations with an inviscid code where the post-shock flow is given random, grid-scale, motions confirms the more rapid destruction of the cloud.  Our results clearly show that turbulence plays an important role in shock-cloud interactions, and that environmental turbulence adds a new dimension to the parameter space which has hitherto been studied. ",
        "introduction": "\\label{sec:intro} Circumstellar, interstellar, and intergalactic environments are inhomogeneous, with clouds of various densities and temperatures embedded in a hotter, more tenuous, substrate.  This substrate is often turbulent, in part due to the interaction of shocks, shells, winds and jets with these clouds.  The nature of such interactions is interesting, because the evolution and morphology of large-scale flows can ultimately be regulated by objects of much smaller size.  In the interstellar medium (ISM), for instance, the interaction of supernova shock waves with interstellar clouds creates a continuous interchange of mass and energy between various thermal phases \\citep{Cox:1974,McKee:1977}. Several outcomes are possible.  The shocks may destroy the clouds, and mix their material into their surroundings.  Alternatively, the shocks may trigger the collapse of the clouds and the formation of new stars, thereby removing material (at least temporarily) from the ISM \\citep{Elmegreen:1977}.  The shocks themselves will slow and material behind them will cool to form thin dense shells. These shells may then fragment and form new clouds. Clouds which survive the passage of the shell subsequently find themselves either inside a hot, low density bubble (if the bubble is energy-conserving), or exposed to a fierce, high Mach number wind (if the shell is momentum-driven). An important process is the generation of vorticity as the diffuse medium flows around the clouds. Understanding the interaction between shocks/shells/winds/jets and interstellar clouds is therefore a key step in studies of the structure and evolution of the ISM \\citep[see the recent reviews by][]{Elmegreen:2004,Scalo:2004}.  A similar interaction occurs in massive, early-type, stellar systems where each star blows a powerful, clumpy, wind which collides with the other.  The impact of the clumps causes the wind-wind collision region to become highly turbulent \\citep{Pittard:2007a}, with implications for particle acceleration, the timescales for equilibrium ionization and electron heating, and the physical mixing of the winds.  In this work we consider the adiabatic interation of a shock with a cloud.  An extensive literature of analytical and numerical investigations of shock-cloud interactions now exists \\citep[see, e.g.,][and references therein]{Nakamura:2006}. The effects of ordered magnetic fields, radiative cooling, and thermal conduction have all been considered, but not simultaneously until recently \\citep{Orlando:2008}.  High-power laser experiments of shock-cloud interactions \\citep[e.g.,][]{Klein:2003} have complemented this literature. While thermal conduction acts to prevent hydrodynamic instabilities, radiative cooling enhances them. The effect of ordered magnetic fields is more complicated: instabilities are prevented in cases where a magnetic field provides a high tension at the surface of the cloud, but the effects depend on the strength, orientation and scale of the magnetic field. Instabilities and vortical motions may be prevented in some directions, but not necessarily in others. Interestingly, in three-dimensional simulations a magnetic field may actually enhance the fragmentation of a cloud \\citep*{Gregori:1999,Shin:2008}, though the actual mixing of the cloud and ambient medium remains hindered.  Even if the magnetic fields are not strong enough to directly affect the dynamics, they can significantly reduce the effects of thermal conduction in directions normal to the field lines, thus allowing instabilities to develop \\citep{Orlando:2008}.  Given the fact that in shock-cloud interactions the Reynolds number is typically high, and the above predisposition for instabilities to develop, the fluid velocity field around the cloud is expected to vary significantly and irregularly in both position and time. This ``turbulence'' transports and mixes the cloud material much more effectively than a comparable laminar flow. Turbulence is also effective at ``mixing'' the momentum of a fluid (i.e. accelerating material ripped off the cloud to the ambient flow speed), and at transferring heat.  In turbulent flows of high Reynolds number there is a separation of scales - large-scale motions will be strongly influenced by the geometry of the cloud, and control the transport and mixing of cloud material, while the behaviour of small-scale motions is determined almost entirely by the rate at which they receive energy from the large scales, and by the viscosity.  Hence these small-scale motions have a universal character, independent of the flow geometry.  While much insight has been gained from previous numerical investigations of shock-cloud interactions, the artificial viscosity inherent in all such simulations has the potential to prevent the formation of fully developed turbulence, to limit the turbulent mixing of cloud material into the surrounding medium, and to hinder the destruction of the cloud.  Furthermore, all previous simulations are highly idealized in the sense that the flow behind the shock is assumed to be perfectly smooth and uniform (i.e. laminar). In reality, random inhomogenities in the ambient medium upstream of the cloud will deform the shock, and will cause velocity, density and pressure structures to develop in the post-shock flow. The post-shock flow will then also be ``turbulent'', and it is expected that the destruction of the cloud will be more rapid in such conditions.  In this work we address whether fully developed turbulence has been prevented in all previous numerical works. For high Reynolds number flows, the only tractable method is a statistical approach i.e. to describe the turbulent flow, not in terms of a velocity field ${\\bf u}(x,t)$, but in terms of some statistics. A model based on such statistics can lead to a tractable set of equations, because statistical fields vary smoothly (if at all) in position and time.  To do this we use a sub-grid turbulent viscosity model, where an attempt is made to calculate the properties of the turbulence and the resulting increase in the transport coefficients. The most widely used is the so-called $k$-$\\epsilon$ model, where the properties of the turbulence are described by two variables, the turbulent energy per unit mass, $k$, and the turbulent dissipation rate per unit mass, $\\epsilon$. The addition of viscous and diffusive terms in the fluid equations simulates the turbulent mixing of the cloud and intercloud medium due to shear instabilities. In this way the subgrid turbulence model emulates a high Reynolds number flow. Incorporating a $k$-$\\epsilon$ model into shock-cloud simulations should produce more realistic results than those from inviscid codes, where the viscosity is purely numerical and the size of shear instabilities is determined by the resolution of the numerical grid. A turbulent viscosity model differs from simply adding physical viscosity to the grid, because the turbulent viscosity is largest in shear layers and essentially vanishes in regions with little shear, whereas models with grid viscosity have the same viscosity everywhere.  The structure of this paper is as follows. The key physics of a shock-cloud interaction is reviewed in Section~\\ref{sec:review}. Section~\\ref{sec:setup} introduces the investigation and the numerical method used. The results are presented in Section~\\ref{sec:results}, where the effects of turbulence on the cloud evolution are described.  A discussion of the relevance of our results to shock-cloud and wind-cloud observations is given in Section~\\ref{sec:discussion}. Section~\\ref{sec:summary} summarizes the conclusions of this work, and possible future work is noted in Section~\\ref{sec:future}. ",
        "conclusions": "\\label{sec:summary} This is the first in a series of papers investigating the turbulent destruction of clouds. Here we have investigated the destruction of a cloud by an adiabatic shock using a hydrodynamical code which incorporates a subgrid $k$-$\\epsilon$ turbulence model, in which an attempt is made to calculate the properties of the turbulence and the resulting increase in the transport coefficients. The results are compared against those from a hydrodynamical code which solves only the inviscid Euler equations of fluid motion. The motivation for this study is that fully developed turbulence is prevented by the artificial viscosity in all numerical codes, though it is expected in astrophysical environments where the Reynolds number of the shock-cloud interaction is $\\gtsimm 10^{5}$.  The effect of a highly turbulent environment sweeping over the cloud is also investigated - all other works studying the interaction of a shock with a single cloud have assumed that the post-shock environment is completely smooth, though this is clearly not the case when there are upstream inhomogeneties. Our main results are summarized thus:  \\begin{enumerate} \\item The evolution and destruction of clouds with inviscid and $k$-$\\epsilon$ models occurs at roughly the same speed when the post-shock flow is smooth. This is because the turbulence takes some time to be generated and for its effects to be subsequently felt. However, it is clear that there are increasing differences between such models as the density contrast, $\\chi$, between the cloud and inter-cloud medium increases.  We also show that the $k$-$\\epsilon$ model results are far less dependent on the spatial resolution than calculations with an inviscid code. This behaviour is attractive given the recent interest in simulations with multiple clouds.  \\item Turbulence is mainly generated at the slip surface around the cloud, though a small amount is also generated behind shocks. The set-up time to form a turbulent wake behind the cloud is $\\sim t_{\\rm cc}$. There is slower growth of turbulence around clouds with a smooth density profile.  The opening angle of the turbulent mixing layer is consistent with experimental results. The fraction of energy in small-scale sub-grid turbulent motions initially increases with time as turbulence continues to be generated by the velocity shear of flow past the cloud, but then decays as turbulent energy subsequently dissipates as heat.  The peak value of the turbulent energy fraction increases with $\\chi$ due to the longer drag-time of denser clouds and the higher velocity shear at the cloud surface, and decreases as the boundary of the cloud becomes smoother. The interaction of shocks with smooth clouds is generally milder, as previously reported.  \\item Clouds which are subject to a highly turbulent post-shock environment are destroyed significantly quicker than those within a smooth flow, due to the enhanced transport and diffusion coefficients. This effect increases with $\\chi$, since high density clouds are subject to a longer period of ``buffeting''.  Clouds with small density contrasts (e.g., $\\chi=10$) are destroyed so quickly after the initial shock passage that they experience very little ``buffeting''. Strong environmental turbulence increases the momentum coupling between the cloud and its surroundings, which increases the acceleration of the cloud. The degree by which the destruction of the cloud speeds up should depend on the strength of the turbulence imposed on the post-shock flow, but here our intention is simply to demonstrate that it does, in fact, occur more rapidly.  \\item The rate of mass-loss from a cloud overrun by a Mach 10 shock due to hydrodynamic ablation is found to be broadly consistent with theoretical expectations, but shows large variations over the destruction period: the peak mass-loss rate is about 5 times higher than the time-independent value from a simplified theory. Clouds in a highly turbulent environment can be destroyed approximately twice as fast. However, significant differences between theoretical and numerically determined mass-loss rates exist at higher Mach numbers (Pittard, in preparation).  \\item Material stripped off the cloud is fragmented and irregular in structure when $\\chi \\ltsimm 10^{2}$. In contrast, the mass-loss from models with $\\chi=10^{3}$ better resembles a single tail-like feature. The length-to-width ratio of the tail increases with $\\chi$ (denser clouds live longer, allowing material to be more dispersed in the axial direction), but decreases with the level of environmental turbulence.  \\item The vorticity generated in the interaction is broadly similar in models with differing levels of post-shock turbulence when $\\chi \\ltsimm 10^{2}$. However, models with high environmental turbulence generate substantially less total circulation when $\\chi=10^{3}$ compared to models with a smooth post-shock environment, due to the more rapid destruction of the cloud. This may self-limit the total circulation that can be generated in such interactions.  \\item Confirmation of the general speeding up of cloud destruction in a highly turbulent post-shock environment is attained by additional calculations where strong, grid-scale motions and inhomogeneities are imposed in the post-shock flow of an inviscid calculation. Differences in the resulting evolution compared to those from the subgrid turbulence models with high post-shock turbulence reveal that the interaction is sensitive to the details of the turbulence, such as the maximum eddy size. The degree by which the destruction speeds up should again depend on the strength and properties of the turbulence imposed on the post-shock flow. Future simulations with a non-smooth inter-cloud medium would also be of interest. \\end{enumerate}"
    },
    "0807/0807.4919_arXiv.txt": {
        "abstract": "The rate of the hydrogen-burning carbon-nitrogen-oxygen (CNO) cycle is controlled by the slowest process, $^{14}$N(p,$\\gamma$)$^{15}$O, which proceeds by capture to the ground and several excited states in $^{15}$O. Previous extrapolations for the ground state contribution disagreed by a factor 2, corresponding to 15\\% uncertainty in the total astrophysical S-factor. At the Laboratory for Underground Nuclear Astrophysics (LUNA) 400\\,kV accelerator placed deep underground in the Gran Sasso facility in Italy, a new experiment on ground state capture has been carried out at 317.8, 334.4, and 353.3\\,keV center-of-mass energy. Systematic corrections have been reduced considerably with respect to previous studies by using a Clover detector and by adopting a relative analysis. The previous discrepancy has been resolved, and ground state capture no longer dominates the uncertainty of the total S-factor. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0807/0807.4634_arXiv.txt": {
        "abstract": "Population III stars, the first generation of stars formed from primordial Big Bang material with a top--heavy IMF, should contribute substantially to the Universe reionization and they are crucial for understanding the early metal enrichment of galaxies. Therefore it is very important that these objects, foreseen by theories, are detected by observations. However PopIII stars, searched through the HeII 1640\\AA\\ line signature, have remained elusive. We report about the search for the HeII line in a galaxy at z=6.5, which is a very promising candidate. Unfortunately we are not yet able to show the results of this search. However we call attention to the possible detection of PopIII stars in a lensed HII dwarf galaxy at z=3.4, which appeared in the literature some years ago, but has been overlooked. ",
        "introduction": "Theoretical models foresee that the first generation of stars, forming from primordial Big Bang material, should have unusually massive stars (a {\\it top--heavy} IMF), with masses up to about 500 $M_\\odot$ (\\cite[Schneider et al. 2002]{schn02}). These stars are called Population III (PopIII) stars, and at least some of them are expected to quickly release metals in the interstellar medium (ISM); therefore later generations would soon be metal enriched. The metallicity threshold in the ISM for producing PopIII stars should be around $10^{-5}$--$10^{-4} Z_{\\odot}$. The very massive PopIII stars would then produce a short phase of unusually hard and strong UV radiation ($T_{eff} \\sim 100000 K$), resulting in a specific line emission signature (\\cite[Shaerer 2002]{sch02}). The most prominent and unique emission line is expected to be HeII 1640\\AA, which can become as strong as 1/3 of Lyman $\\alpha$, making it observable with current 8--10m telescopes in the highest redshift galaxies known (\\cite[Scannapieco et al. 2003]{sca03}).  The detection of PopIII stars would be extremely important for understanding the early chemical evolution of galaxies, and because they should be crucial contributors to the reionization of the Universe (\\cite[Ciardi et al. 2003]{cia03}). PopIII stars have however remained elusive: searches for HeII 1640\\AA\\ through stacking of spectra of Lyman $\\alpha$ emitting galaxies (\\cite[Dawson et al. 2004]{daw04} and \\cite[Ouchi et al. 2008]{ouc08}), through deep spectroscopy of an individual galaxy (\\cite[Nagao et al. 2005]{nag05}), or through dual (Ly$\\alpha$ and HeII) narrow--band imaging (\\cite[Nagao et al. 2008]{nag08}) have failed. \\cite[Jimenez \\& Haiman (2006)]{jim06} have ascribed to PopIII stars the HeII 1640\\AA\\ line detected in the composite spectrum of $\\sim$1000 Lyman break galaxies at $z\\sim 3$, but the line is only about 1/10 of Lyman $\\alpha$ and is considerably broader: it could therefore be attributed to a stellar wind feature associated with massive WR stars (\\cite[Shapley et al. 2003]{sha03}). The upper limits to the star formation rate (SFR) for PopIII stars, which can be set from these negative results are now quite close to the rates expected from the models of \\cite[Tornatore et al. (2007)]{tor07} (see the contribution by T. Nagao to these Proceedings).  We report about the search for the HeII 1640\\AA\\ line in a galaxy at z=6.5 and about the possible detection of PopIII stars in a lensed dwarf HII galaxy at z=3.4, which has appeared in the literarure (\\cite[Fosbury et al. 2003]{fos03}), but has so far been overlooked. ",
        "conclusions": ""
    },
    "0807/0807.3717.txt": {
        "abstract": "{} {This paper describes the design goals and engineering efforts that led to the realization of AMBER (Astronomical Multi BEam combineR) and to the achievement of its present performance.} {On the basis of the general instrumental concept, AMBER was decomposed into modules whose functions and detailed characteristics are given. Emphasis is put on the spatial filtering system, a key element of the instrument. We established a budget for transmission and contrast degradation through the different modules, and made the detailed optical design. The latter confirmed the overall performance of the instrument and defined the exact implementation of the AMBER optics.} {The performance was assessed with laboratory measurements and commissionings at the VLTI, in terms of spectral coverage and resolution, instrumental contrast higher than 0.80, minimum magnitude of 11 in K, absolute visibility accuracy of 1\\%, and differential phase stability of 10$^{-3}$rad over one minute.} {} ",
        "introduction": "AMBER \\citep{petrovd} is the near infrared focal beam combiner of the Very Large Telescope Interferometric mode (VLTI). A consortium of French, German, and Italian institutes is in charge of the specification, design, construction, integration, and commissioning of AMBER (http://amber.obs.ujf-grenoble.fr). This instrument is designed to combine three beams coming from any combination of Unit 8-m Telescopes (UT) or Auxiliary 1.8-m Telescopes (AT). The spectral coverage is from 1 to 2.4$~\\mu$m  with a priority to the K band.  AMBER is a general user instrument with a very wide range of astrophysical applications, as is shown in \\citet{richichi} and \\citet{petrovd}. To achieve the ambitious programs, \\citet{petrovb} presented the associated specifications and goals, such as: \\newline \\indent- Spectral coverage: J, H, and K bands, from 1.0 $\\mu$m to 2.3 $\\mu$m (goal: up to 2.4 $\\mu$m). \\newline \\indent- Spectral resolutions: minimum spectral resolution 30$<$R$<$50, medium spectral resolution 500$<$R$<$1\\,000, and highest spectral resolution 10\\,000$<$R$<$15\\,000.\\newline \\indent- Minimum magnitude: K=11, H=11 (goals: K=13, H=12.5, J=11.5).\\newline \\indent- Absolute visibility accuracy: 3$\\sigma_V$=0.01 (goal: $\\sigma_V$=10$^{-4}$).\\newline \\indent- Differential phase stability: 10$^{-3}$rad (goal: 10$^{-4}$rad) over one minute integration. This allows us to compute phase closure that is necessary in the search of brown dwarfs and extra solar planets \\citep{segr}.\\newline  These specifications have been the starting point of a global system analysis \\citep{malbet} initiated by a group of interferometrists from several French institutes and completed by the Interferometric GRoup (IGR) of AMBER. This work led to the current definition of the AMBER instrument whose broad outlines are recalled in the paragraphs below.  To reach a sufficient sensitivity, ESO provided a 60 actuator curvature sensing system MACAO (Multi-Application Curvature Adaptive Optics) \\citep{arseno} specified to deliver at least a 50\\% Strehl ratio $@$ 2.2$\\mu$m for on-axis bright sources (V=8) under median seeing conditions (0.65\") and a 25\\% Strehl ratio $@$ 2.2$\\mu$m for faint sources (V=15.5) under the same seeing conditions. The required high accuracy of the absolute visibility measurements implies the use of spatial filters \\citep{mege,tatub} with single mode fibers based on the experience of other smaller, successful interferometers such as IOTA/FLUOR \\citep{coude}. The atmospheric noise is reduced to photometric fluctuations, which can be monitored, and to Optical Path Difference (OPD) fluctuations between the different pupils, which can be frozen by very short exposures or adaptively corrected by a fringe tracker.  The simultaneous observation of different spectral channels is ensured by dispersed fringes. This very significantly increases the number and the quality of the measurements and subsequently the constraints imposed on the atmospherical models. The modularity of the concept was a strong argument in favor of the multi-axial scheme, as carried out on the Grand Interferom\\`etre \\`a 2 T\\'elescopes (GI2T) of the Plateau de Calern \\citep{mourard}. In addition, it was demonstrated that the instrument must correct the atmospheric transversal dispersion in J and H \\citep{tallon}. The need of an image cold stop was assessed by \\citet{malbet1} to reduce the thermal background coming from the blackbody emission of the fiber heads, which can be greater than the detector RON, especially for long time exposures in the K-Band. A pupil mask also acts as a cold stop. To perform the data reduction, the ABCD algorithm \\citep{chelli,millour}, as used with co-axial configurations with a temporal coding, was chosen. The associated complete data calibration procedure was then fully defined \\citep{hofman,chelli,tatuc}.  On the basis of the general instrumental concept resulting from the global system analysis, we defined the main modules and necessary accessories (such as the alignment units) of AMBER. We established a budget for throughput and contrast degradation through the different modules, made the detailed optical design, and performed a complete optical analysis. The latter confirmed the expected overall performance of the instrument in terms of signal-to-noise ratio (SNR). This entire process is described in the present paper. The procedure described here to allocate the specifications of the different modules of an interferometer could be used, after some changes, for the design of other interferometrical instruments, such as, for example, the VLTI second generation instruments. ",
        "conclusions": "To achieve the ambitious astrophysical program of AMBER, it is necessary to obtain visibility measurements with an accuracy better than 10$^{-2}$ (goal: 10$^{-3}$) on sources of magnitude K relatively smaller than 11 and to reach an instrumental stability of the differential phase of 10$^{-3}$ rad (goal: 10$^{-4}$ rad) over one minute. It was shown that these general specifications are driven by the necessity for the instrument to provide interferometric requirements that consist mainly of contrast level and accuracy, optical throughput and stability. To assess the AMBER performance, a tolerance procedure was defined: with the scientific specifications, the interferometric specifications were determined, in terms of error budgets for the fiber coupling degradation factor due to misalignments and for the instrumental contrast. Both allowed us to define the degrees of freedom of each optical element, respectively located before and after the optical fibers, and represented inputs for the detailed tolerance analysis performed with a ray-tracing tool.  The complete tolerance study, confirmed by the observations, assessed the feasibility of the requirements, i.e., an overall instrumental contrast level higher than 80\\% in the K-band and the 11th K-magnitude with the lowest resolution. After separate integration and tests in the institutes in charge of the warm optics \\citep{roba}, SPG \\citep{lisi}, and DET, the whole instrument was integrated and tested in LAOG in 2003. AMBER successfully passed the Preliminary Acceptance in Europe in November 2003, resulting in the validation of the instrument laboratory performance, of the compliance with the initial scientific specifications \\citep{rouss}, and of the acceptance of ESO for AMBER to be part of the VLTI. After the transportation of the instrument to Paranal, Chile in January 2004, the Assembly Integration and Verification phase occurred mid-March with a successful first fringe observation of bright stars with the VLTI siderostats \\citep{robb}.  Since then four commissioning runs have been performed, strengthening the performance. The astrophysical results presented in this issue of A\\&A is the best evidence that AMBER is today a very powerful tool allowing unequalled astrophysical knowledge. More results are expected once VLTI has completely solved vibration problems and installed the fringe sensor. %________________________________________________________________"
    },
    "0807/0807.0401_arXiv.txt": {
        "abstract": "It has been argued that the spacetime of our universe can be accurately described by a perturbed conformal Newtonian metric, and hence even large density inhomogeneities in a dust universe can not change the observables predicted by the homogeneous dust model.  In this paper we study a spherically symmetric dust model and illustrate conditions under which large spatial variations in the expansion rate can invalidate the argument. ",
        "introduction": "In cosmology, one models the evolution and observables associated with an inhomogeneous universe of density $\\rho(\\vec{x})$ and expansion rate $H(\\vec{x})$ by employing a Friedmann-Lema\\^{\\i}tre-Robertson-Walker (FLRW) homogeneous/isotropic model of density $\\rho = \\langle\\rho(\\vec{x})\\rangle$, where $\\langle\\cdots\\rangle$ denotes some suitably defined spatial average. One then assumes that the expansion rate and cosmological observables are those obtained in the corresponding FLRW model.  One of the proposals to explain ``dark energy'' calls into question this long-standing (86-year) procedure.  The idea is that the expansion rate and cosmological observables of a suitably inhomogeneous universe containing only dust, if analyzed within the framework of a homogeneous model, seems to behave as if the stress tensor also contains a fictitious negative-pressure fluid (\\textit{i.e.,} dark energy).  Although this proposal is conservative in the sense that it does not involve a cosmological constant of incredibly small magnitude (presumably rationalized on some anthropic basis), a scalar field of unbelievably small mass, or an entirely unmotivated modification of general relativity, it is rather revolutionary because it implies that there is no dark energy and the expansion of the universe does not actually accelerate (at least, not in the usual sense).  At present, the idea that the backreaction\\footnote{By the backreaction idea we mean both that inhomogeneities influence cosmological observables (weak backreaction) and that inhomogeneities influence the average/background scale factor (strong backreaction). Weak and strong backreactions are, however, likely connected if we assume that averages should be performed along the light cone. } of inhomogeneities accounts for the observational evidence usually attributed to dark energy is more of a concept than a predictive model.  However, it is generally agreed that if the proposal is to be relevant, nonlinearities are required.  There have been many criticisms of this approach.  One of them \\cite{Ishibashi:2005sj,Gruzinov:2006nk} is based on the claim that even in the presence of highly nonlinear density perturbations ($\\delta\\rho/\\rho\\gg1$) the metric for our universe can everywhere be written  as a perturbed conformal Newtonian metric of the form\\footnote{We will refer to a metric written in the form of Eq.\\ (\\ref{newtpert}), satisfying the stated conditions, as the \\textit{perturbed conformal Newtonian} metric.} \\begin{equation} ds^2=a^2(\\tau)\\left[-(1+2\\psi)d\\tau^2 + (1-2\\psi)\\gamma_{ij} dx^idx^j \\right]  , \\label{newtpert} \\end{equation} where $d\\tau = dt/a$ is conformal time, $\\gamma_{ij}$ is a metric of a three-space of constant curvature, and  $\\psi$ satisfies the Newtonian conditions $|\\psi| \\ll 1$,  $\\left| \\partial\\psi/\\partial t \\right|^2 \\ll a^{-2}D^i\\psi D_i\\psi$, and $\\left(D^i\\psi D_j\\psi\\right)^2\\ll \\left(D^iD^j\\psi\\right)D_iD_j\\psi$.  The covariant derivative with the metric $\\gamma_{ij}$ is denoted by $D_i$.  The usual statement is that in the dust case one is allowed to use the perturbed conformal Newtonian metric either in the linear regime (\\textit{i.e.,} perturbations of every quantity being small) or in the weak-field (Newtonian) regime.\\footnote{But already at second order in perturbation theory the spatial part of the metric differs from the temporal part of the metric simply from terms quadratic in the peculiar velocities (\\textit{i.e.,} an effective anisotropic stress term is in the stress-energy tensor).}  The claim is that if the metric can be written in the above form and satisfies the above conditions, even in the presence of large inhomogeneities, any cosmological observable will be the same as  the cosmological observable calculated with $\\psi=0$, \\textit{i.e.,} in the homogeneous/isotropic model. This has been described as a ``no-go'' theorem that backreactions can not \\textit{in principle} account for the observations.  While it is impossible to consider the most general inhomogeneous solutions, there are spherically symmetric inhomogeneous dust solutions, which are not perturbations of Einstein-de Sitter, that can be constructed to give observables similar to $\\Lambda$CDM models.  These models serve as a counterexample to the no-go argument.  In this paper we will show why these models can not be described in terms of a conformal Newtonian metric perturbed about a spatially flat background, and attempt to understand the implications for the backreaction proposal.  Indeed, while it may turn out that backreactions are not the answer, we argue that assuming the results of measurements of the luminosity distance as a function of redshift usually interpreted as an accelerated expansion, the metric describing our universe can not be written in the form of a perturbed conformal Newtonian metric where $a(t)$ is calculated from the homogeneous dust model. In other words, if the expansion history of the universe is well described by the $\\Lambda$CDM model, then perturbing about an Einstein-de Sitter model by the perturbed conformal Newtonian metric of Eq.\\ (\\ref{newtpert}) is inappropriate, because Einstein-de Sitter would be the wrong background. This is because of large peculiar velocities with respect to the background Einstein-de Sitter space.  So if inhomogeneities are responsible for the observables usually attributed to dark energy, the universe can not be obtained by small perturbations of the Einstein-de Sttter model.  In other words, the reason we interpret the observations as evidence for dark energy and acceleration of the universe is that we are comparing the observables to observables computed in the wrong background.  As we will discuss, the reason is that the proper meaning of ``peculiar'' is ``after subtraction of a background Hubble flow term.'' We will argue that large peculiar velocities must be present if the backreaction program works, and the peculiar velocities are not related to ``local'' departures from the Hubble flow that would show up as large velocity dispersions.  As an explicit example, consider the argument of Ref.\\ \\cite{Kolb:2005da}. They propose that the backreaction of inhomogeneities in a dust universe modifies the evolution of the effective volume-averaged scale factor and results in an evolution of the volume expansion that resembles a $\\Lambda$CDM model, rather than the unperturbed spatially flat dust model. If one would write the metric for such a perturbed universe in terms of a perturbed conformal Newtonian metric, then one would have to use $a(t)$ described by a $\\Lambda$CDM model, \\textit{not} the $a(t)$ from an unperturbed spatially flat dust model.  If one would attempt to express the metric in terms of a perturbed metric with $a(t)$ described by a spatially flat dust model, then there would be enormous peculiar velocities in the Hubble flow.  We explore the issue by considering the spacetime of a particular Lema\\^{\\i}tre-Tolman-Bondi spherically symmetric model. In the model we study, the observer is centered in a large (Gpc) underdense region in an asymptotically Einstein-de Sitter universe.  This model has large, but not highly nonlinear, density and curvature inhomogeneities.  We show by explicit calculation that the metric \\textit{can not} be expressed as a perturbed conformal Newtonian metric during the entire evolution.  We also show that the reason for this can be traced to large peculiar velocities that develop with respect to the Einstein-de Sitter background.  In the next Section we describe the relevant features of the Einstein-de Sitter  and Lema\\^{\\i}tre-Tolman-Bondi models.  In Section \\ref{LTBPG} we develop the formalism to express the original Lema\\^{\\i}tre-Tolman-Bondi metric in the form of a perturbed conformal Poisson metric. In Section \\ref{Discussion} we discuss the reason for the numerical results of the previous Section.  Finally, in Section \\ref{Conclusions} we present our conclusions. ",
        "conclusions": "}  The LTB model analyzed in this paper is a very useful toy model to explore the role of large inhomogeneities in the determination of cosmological observables such as the luminosity distance as a function of redshift. We employed the toy model to demonstrate where and why the no-go argument discussed in the Introduction fails. To this end we have performed a linear gauge transformation in order to understand the problems we have to face if we want to express the LTB model in the form of a perturbed conformal Newtonian metric, Eq.\\ (\\ref{newtpert}).  We would like to stress that we regard the AAG model as a tool useful to show a general feature needed to evade the no-go argument: we focused on the \\textit{general} physical cause of the nonlinear behavior of the AAG model more than on its particular initial conditions.  We came to the conclusion that the way the no-go argument is proved is by assuming, as a starting point, the validity of the argument itself. The way to escape the no-go argument is by allowing large ``background peculiar velocities,'' where these velocities have nothing to do with, \\textit{e.g.,} velocity dispersions in clusters or anything that can be measured as a local effect. If indeed inhomogeneities are responsible for the ``apparent'' acceleration of the universe, then the evolution of the expansion rate must resemble a $\\Lambda$CDM model.  If that is true, then there will naturally be ``peculiar'' velocities if one describes the evolution in terms of perturbations of an EdS model.  One cannot simply assume peculiar velocities are small and then use that as a basis to show inhomogeneities can have no effect. In the calculation of peculiar velocities it is necessary to specify a background solution about which to calculate velocities.  The peculiar velocities can be small in the $\\Lambda$CDM background but not in the EdS background. The velocities appear as a departure from EdS of the Hubble flow, preventing the metric from being written in the conformal Newtonian form.  In other words, to evade the no-go argument we need to free the inhomogeneous models of dynamical restrictions (small background peculiar velocities) which impose the existence of a global background.  Another way to see how the no-go argument of Eq.\\ (\\ref{newtpert}) fails is through the equations of motion: in order to have a model, \\textit{e.g.,} an LTB model, explain the observed departure of $d_L(z)$ from the EdS result requires nonlinear equations of motion, and this is because of the large background peculiar velocities which can generally occur. By this, we mean that in order to end up with a large spatial variation of the background $H$ requires a nonlinear evolution like the one of the LTB model discussed in this paper. And this can occur even if it is possible, as in the AAG model, to describe an inhomogeneous universe by means of the perturbed conformal Newtonian metric from (let's say) $z=2$ to present time.  This is shown here by explicit calculation.  Therefore, one must reconsider very carefully the statement that inhomogeneities can not have a significant effect on cosmological observables such as the luminosity distance as a function of redshift if arguments are based on the \\textit{assumption} that the metric describing our universe can be written in the perturbed conformal Newtonian form of Eq.~(\\ref{newtpert}).  Now we enumerate some of our findings.  \\begin{enumerate}  \\item  In the model we consider, even at late times when the density contrast is linear and the velocities are small, one can not express the LTB metric in the perturbed conformal Newtonian form of Eq.\\ (\\ref{newtpert}) by means of a linear gauge transformation. The breakdown of the linear transformation clearly shows the importance of background velocities.  The linear gauge transformation breaks down mainly because of the large $\\Delta r = \\beta'$.   This is caused by the different expansion rates ($\\Delta H$) of the void and the EdS solution, as can be seen in Fig.\\ \\ref{h}, which shows the expansion rates as a function of $v$ for several values of $x$. The result is that the two metrics do not remain close enough for a linear transformation to connect them.  As discussed in the previous section, the basic problem is that the linear gauge transformation neglects terms like ($\\nabla^2 \\beta)^2$ compared to terms like $\\phi$. Summarizing, the cause of the break down of the linear gauge transformation is the large $\\Delta H$, and the underlying cause of that is the nonlinearity of the equations of motion. The failure described above can be stated as the fact that the linearized equations of motion work only for a limited range of times.  \\item  It is possible to find a nonlinear gauge transformation to overcome this problem with a perturbed conformal Newtonian metric description of the LTB metric whose potential follows nonlinear equations \\cite{VanAcoleyen:2008cy}. These nonlinear equations (which are actually the exact Einstein equations) at some time will typically give large potentials and make the perturbed conformal Newtonian gauge description inconsistent (large $\\Delta t = a\\zeta$ in the LTB model examined) because at some point the peculiar velocities will be larger than $c$ if we do not restrict the dynamics, as shown by the AAG model.  \\item Although at the present time, in the model we consider, the magnitude of the potentials are small, they are not equal at the level of the linearized theory, \\textit{viz.}\\ $\\psi \\neq \\phi$.  This could na\\\"{\\i}vely be interpreted as some anisotropic stress or something else funny with the stress tensor.  This is exactly the point that has been made in, \\textit{e.g.,} Ref.\\ \\cite{Kolb:2005da}: Analyzing an inhomogeneous model in the framework of a homogeneous model introduces spurious components and behavior to the stress tensor (\\textit{possibly even a negative-pressure fluid!}).  \\item If we start with a perturbed EdS model, then the growth of perturbations will lead to the formation of large voids, filamentary structure, and other nonlinearities.  Perhaps the LTB models are useful toy models to study the effect of that phenomenon.  The model explored in this paper is a good example because it is capable of dealing with nonlinearities within a large timespan.  \\item Finally, another point worth mentioning is the physical nature of super-Hubble modes of $\\phi$ and $\\psi$.  At the initial time, within each Hubble radius patch, $\\phi$ and $\\psi$ are very nearly constant, but possibly larger than one (with a value of about $20$ near the origin).  That \\textit{does not} mean that we can simply subtract different values, because we are only allowed to do this once.  If we do it at $x^{-1}=1100$, then at $x^{-1}=1$ we would have $\\phi$ and $\\psi$ near the origin of $-20$.  \\end{enumerate}   Summarizing, inhomogeneities, in particular voids not dynamically restricted, undermine the existence of a global background.  The interesting question for future consideration is if backreactions\\footnote{We do not really distinguish between weak or strong backreaction because we think that what matters are the observables. Assuming, indeed, that an inhomogeneous model gives a $\\Lambda$CDM-like luminosity distance, then it is not straightforward to understand which category of backreaction (weak or strong) the model belongs to. The strong backreaction indeed is about a change of background, but this is also the case for the weak backreaction because, as in the case of this example, it ``defines'' an effective background, the $\\Lambda$CDM model. Moreover, the change of background of the strong backreaction should be related to observable data on the light cone as it is for the weak backreaction.} of cosmological perturbations can substantially distort derived conclusions about the energy-momentum tensor of the universe without allowing dramatic violations of \\textit{statistical} homogeneity. The present work addresses the key issue that backreaction effects cannot be assumed \\textit{in principle} to be small on the basis of the no-go argument explained in the Introduction and therefore opens a wide range of possibilities to tackle the study of the inhomogeneities in the universe."
    },
    "0807/0807.2404_arXiv.txt": {
        "abstract": "We observed the nearby, low-density globular cluster M71 (NGC 6838) with the Chandra X-ray Observatory to study its faint X-ray populations. Five X-ray sources were found inside the cluster core radius, including the known eclipsing binary millisecond pulsar (MSP) PSR J1953+1846A. The X-ray light curve of the source coincident with this MSP shows marginal evidence for periodicity at the binary period of 4.2 h. Its hard X-ray spectrum and luminosity resemble those of other eclipsing binary MSPs in 47 Tuc, suggesting a similar shock origin of the X-ray emission. A further 24 X-ray sources were found within the half-mass radius, reaching to a limiting luminosity of $1.5\\times10^{30}$ \\ergl (0.3-8 keV). From a radial distribution analysis, we find that $18\\pm6$ of these 29 sources are associated with M71, somewhat more than predicted, and that $11\\pm6$ are background sources, both galactic and extragalactic. M71 appears to have more X-ray sources between $L_X=10^{30}$--$10^{31}$ ergs s$^{-1}$ than expected by extrapolating from other studied clusters using either mass or collision frequency. We explore the spectra and variability of these sources, and describe the results of ground-based optical counterpart searches. ",
        "introduction": "}  Globular cluster X-ray sources are of interest for many reasons. Dense globular clusters bring stars into close dynamical encounters that lead to the production of X-ray binaries (e.g. Hut, Murphy \\& Verbunt 1991), and studies of globular clusters with different structural parameters can elucidate the details of these mechanisms. Globular clusters provide concentrations of faint X-ray sources for study of X-ray populations at a known distance, age and metallicity. This is of interest even in the least dense globular clusters, where dynamically formed X-ray binaries may be fewer than those descended from primordial binaries. Observations of globular clusters may provide science unique to individual X-ray sources that is enabled by knowledge of the X-ray source's distance, reddening, and other properties. An ensemble of such results for many clusters may shed light on the evolution of globular clusters and their binary populations.  Bright X-ray sources, associated with accreting neutron stars, have long been understood to be produced in globular clusters (Clark 1975). Fainter X-ray sources, composed of combinations of accreting neutron stars in quiescence, cataclysmic variables (CVs), millisecond radio pulsars (MSPs), and chromospherically active binaries (ABs), were known in the 1980s (Hertz \\& Grindlay 1983) and have been resolved with Chandra's high spatial resolution (e.g. Grindlay \\et\\ 2001a, Verbunt \\& Lewin 2006). A number of globular clusters have been observed with Chandra to fairly low X-ray luminosities ($<10^{32}$ \\ergl). Studies of dense clusters have identified large numbers of accreting neutron stars, CVs, ABs and MSPs (e.g. Grindlay et al. 2001a, 2001b, Pooley et al. 2002, Becker et al. 2003). Two relatively sparse nearby clusters have been carefully studied with Chandra: M4 (Bassa et al. 2004) and NGC 288 (Kong et al. 2006). Those studies have indicated that ABs are prevalent in sparse clusters, and that CVs are few in number, but possibly larger in number than predicted by empirical extrapolations with the density and mass of the cluster core.  M71 (NGC 6838) is of particular interest due to its close proximity to Earth (4 kpc). A moderately short Chandra study can identify X-ray sources down to a few $10^{30}$ \\ergl (0.5--2.5 keV), probing the populations of faint CVs, MSPs, and ABs. This cluster is of moderate central density ($\\rho_c = 10^{3.05} L_{\\odot}$/pc$^3$) and shows no evidence for core collapse (its central concentration parameter (Djorgovski 1993) is 1.15). It is moderately reddened (E(B-V)=0.25) and has a globular cluster metallicity ([Fe/H]=-0.73, Harris 1996, updated 2003) slightly higher than average. The core, half-mass radius, and tidal radii are $r_c = 0.63\\arcmin$, $r_h = 1.65\\arcmin$, and $r_t = 8.96\\arcmin$, respectively (Harris 1996, updated 2003). Neither {\\sl Einstein} (Hertz \\& Grindlay 1983) nor \\ROSAT\\ (Verbunt 2001) detected X-ray sources obviously associated with the cluster, with the 0.5--2.5 keV \\ROSAT\\ HRI upper limit being $2.6 \\times 10^{31}$ \\ergl\\ (for an assumed bremsstrahlung spectrum with temperature 0.9 keV). The \\ROSAT\\ HRI did detect 10 sources in its field of view, but none of these were within the cluster half-mass radius (Verbunt 2001).  Ransom \\et\\ (2003, 2005) and Hessels \\et\\ (2007) reported the presence of a binary MSP in M71, with $P_{PSR} = 4.89$ ms and $P_{orb} = 4.24$ h, and the presence at 20 cm of eclipses which last roughly 20\\% of the orbital period.  In this paper, we report on \\Chandra\\ X-ray Observatory observations of M71. We discuss the observations, data processing, and source detection, providing source lists, in \\S\\ref{s:obs}, and the radial distribution of sources with respect to the nominal cluster center in \\S\\ref{s:dists}. In \\S\\ref{s:powerlaw} we present the results of power-law spectral fits to the brighter detected X-ray sources, and in \\S\\ref{s:colorlum} we present and discuss the X-ray color-color diagram for these same sources. We then discuss candidate counterparts in \\S\\ref{s:counterparts}, including those extracted from the 2MASS, USNO B1.0, and TYCHO-2 catalogs (\\S\\ref{ss:catalogs}), a recent tabulation of variable stars (\\S\\ref{ss:variables}), X-ray source catalogs (\\S\\ref{ss:xrays}), and the MSP in M71 (\\S\\ref{ss:msp}). Comparison with Hubble Space Telescope (HST) observations of M71 (\\S\\ref{ss:hst}) will be reported in Huang \\et\\ (2008, in preparation). In \\S\\ref{s:m71individual} we discuss spectral and temporal properties of the brightest detected sources. A summary of our results is given in \\S\\ref{s:discussion}.  \\section {Observations, Data Reprocessing, and Source Detection \\label{s:obs}}  We obtained a 52.4-ks \\Chandra\\ observation (ObsID 5434) of M71 (nominal center of cluster at J2000 RA  $19^h \\ 53^m \\ 46.1^s$, DEC $18\\arcdeg \\ 46\\arcmin \\ 42\\arcsec$) on 20--21 December 2004 using the Advanced CCD Imaging Spectrometer (ACIS) in very faint (VF), timed-exposure mode, with 3.141-s frame time. Starting from the standard \\Chandra\\ X-ray Center (CXC) processing (ASCDS version number 7.6.7.1 and CALDB 3.2.1; third reprocessing) level 1 files, we reprocessed the data without applying pixel randomization.  The reprocessing included applying the current charge transfer inefficiency correction; selection of the standard ASCA grades 0, 2, 3, 4 and 6; and application of the good time filter.  We column-cleaned the data using a variant of a method developed at Pennsylvania State University\\footnote{See http://www.astro.psu.edu/xray/acis/recipes/index.html}.  As is appropriate for data taken in VF mode, we used CLEAN55\\footnote{See http://cxc.harvard.edu/cont-soft/software/clean55.1.0.html} to reduce the background and remove cosmic ray afterglows. In analyzing data and unless otherwise specified, we utilized events in pulse-invariant channels corresponding to 0.3 to 8 keV.  We searched for X-ray sources in the observed field employing techniques described in Tennant (2006) which use a circular-Gaussian approximation to the point spread function (PSF). That author gives a rather detailed account of the method in an Appendix, and there shows that for the \\Chandra\\ Deep Field-North (Brandt \\et\\ 2001) it gives results consistent with those obtained using the CIAO script wavdetect. Using a 30 ks subset of the much longer Deep Field-North observation made it possible to know exactly which sources should be found. The method has been used previously in other published work (\\eg\\ \\Chandra\\ observations of the globular cluster M28 reported in Becker \\et\\ 2003). As a result of operating in VF mode, background levels were low throughout the observation (\\eg\\ $\\sim 1.4 \\times 10^{-7}$ counts/s/pixel over 0.3-8.0 keV for $r_{M71} \\le r_h$). Therefore within twice the M71 half-mass radius ($r_h$ = 1.65$\\arcmin$), we set the signal-to-noise threshold, $S/N$, for detection to 2.0, but also required the number of source counts to be at least 5 times the statistical uncertainty in the local background estimate. The empirical relation derived by Tennant, C$_{min} = (S/N)^2/0.81$ then implies a point-source sensitivity limit of about 4.9 counts for $r_{M71} \\le 2 r_h$ and in the energy band 0.3--8.0 keV. Because of the increase in PSF size with off-axis distance and the associated increase in background within a detection cell, for $r_{M71} > 2 r_h$ we set the S/N threshold for detection to 2.4, and again required the number of source counts to be at least 5 times the statistical uncertainty in the local background estimate. The point-source sensitivity limit rises to about 7 counts. Following Tennant (2006) we expect a completeness limit of about 10 counts. We discuss the luminosities corresponding to these count sensitivity limits in \\S\\ref{s:dists} and \\S\\ref{s:powerlaw}.  The positions of X-ray sources found in this manner inside $2 r_h$ are listed in Table~\\ref{t:1}, which is divided into those inside $r_h$ and those with $r_h < r_{M71} \\le 2 r_h$.  The positions of both sets of sources are displayed in Figure~\\ref{f:1}. As can be seen in Figure~\\ref{f:1}, the circle with radius $2 r_h$ extends slightly off the S3 CCD toward S2, which is a front-side illuminated CCD.  We find one source, {\\it s55}, inside $2 r_h$ on S2 and include it in Table~\\ref{t:1}. In this table, columns 2--5 give right ascension RA (J2000), declination DEC (J2000), detect cell radius, $r_{ext}$, and the approximate number of X-ray source counts, C$_x$, detected from the source in the 0.3--8.0 keV band. Column~6 lists the radius about the source position for inclusion of the source with 68\\% confidence. The corresponding radii for inclusion with 95\\% or 99\\% confidence, $r_{95}$ and $r_{99}$, are found by multiplying column~6 by 1.62 or 2.01, respectively. Column~7 gives the distance, $r_{M71}$, of the source from the nominal center of M71. The remaining columns have to do with candidate counterparts and are discussed in \\S\\ref{ss:catalogs}.  Uncertainties in the aspect solution for ACIS-S\\footnote{See http://cxc.harvard.edu/cal/ASPECT/celmon/, upper left panel of the first figure} imply $\\sigma_{\\rm sys}\\approx 0\\farcs2$ (radial, $\\sim0.4\\arcsec$ at 90\\% confidence); to be conservative, we set $\\sigma_{\\rm sys}=0\\farcs2$ per axis.",
        "conclusions": ""
    },
    "0807/0807.0895_arXiv.txt": {
        "abstract": "{% We study the spatial structure and sub-structure of regions rich in {\\em Hipparcos} stars with blue $B_T-V_T$ colours. These regions, which comprise large stellar complexes, OB associations, and young open clusters, are tracers of on-going star formation in the Galaxy. The DBSCAN (Density-Based Spatial Clustering of Applications with Noise) data clustering algorithm is used to look for spatial overdensities of early-type stars. Once an overdensity, ``agglomerate'', is identified, we carry out a data and bibliographic compilation of their star member candidates. The actual membership in agglomerate of each early-type star is studied based on its heliocentric distance, proper motion, and previous spectro-photometric information. We identify 35 agglomerates of early-type {\\em Hipparcos} stars. Most of them are associated to previously known clusters and OB associations. The previously unknown P~Puppis agglomerate is subject of a dedicated study with Virtual Observatory tools. It is actually a new, nearby, young open cluster ($d \\sim$ 470\\,pc, age $\\sim$ 20\\,Ma) with a clear radial density gradient. We list P~Puppis and other six agglomerates (including NGC~2451~A, vdBH~23, and Trumpler~10) as new sites for substellar searches because of their youth, closeness, and spatial density. We investigate in detail the sub-structure in the Orion, CMa-Pup and Pup-Vel OB~complexes (``super-agglomerates''). We confirm or discover some stellar overdensities in the Orion complex, like the 25~Ori group, the Horsehead region (including the $\\sigma$~Orionis cluster), and the $\\eta$~Orionis agglomerate. Finally, we derive accurate parallactic distances to the Pleiades, NGC~2451~A, and IC~2391, describe several field early-type stars at $d <$ 200\\,pc, and discuss the incompleteness of our search.} ",
        "introduction": "\\label{introduction}  The works by Pannekoek (1929), Ambartsumian (1947) and Ruprecht (1966) were, according to de~Zeeuw et~al. (1999), among the most important ones for the understanding of the OB associations in the early pre-{\\em Hipparcos} era. We also quote the comprehensive review by Blaauw (1964). OB associations are very young regions in the Galaxy containing high-mass O- and B-type stars. The ultraviolet radiation injected into the intra-association medium by these luminous blue stars play an important r\\^ole on the subsequent low-mass star formation process. The study of OB associations is, therefore, necessary for answering critical questions in Astrophysics, like `how does the fragmentation of a primordial molecular cloud take place', `which is the shape of the initial mass function', `how is the dynamical evolution at the first stages of a recently-born open cluster', `where are the spiral arms of our Galaxy', or `do early-type stars inhibit or facilitate the formation of substellar objects'.  After the {\\em magnus opus} by de~Zeeuw et~al. (1999), who carried out a comprehensive census of the stellar content of nearby ($d <$ 1\\,kpc) OB associations, based on {\\em Hipparcos} positions, proper motions, and parallaxes, and prior to launch of the European Space Agency (ESA) mission {\\em GAIA}, very few ``fresh'' discoveries can be achieved on OB associations. The work by de~Zeeuw et~al. (1999) is not, however, the only search for star clusters and OB associations from the {\\em Hipparcos} data: Platais, Kozhurina-Platais \\& \\linebreak van~Leeuwen (1998) and Robichon et~al. (1999) had previously done it (see also a posterior work by Baumgardt, Dettbarn \\& Wielen~[2000]).  The superior capabilities and results of the ESA astrometry mission {\\em Hipparcos} (Perryman et~al. 1997) have allowed astronomers to go on the characterization of the Galactic stellar populations in general, and of OB associations in particular. There have been other works with the {\\em Hipparcos} catalogue aimed at studying particular regions containing early-type stars (Baumgardt 1998; de Bruijne 1999; Subramaniam \\& Bhatt 2000) or the Galactic structure attending to the spatial distribution of such stars (Comer\\'on, Torra \\& G\\'omez 1998; Ma\\'{\\i}z-Apell\\'aniz 2001; Schr\\\"oder et~al. 2004). We all place our trust in the ``fresh'' discoveries on OB associations that will arise in the near future with the overwhelming {\\em GAIA} dataset. It is expected, however, that the {\\em GAIA} mission will be launched in the second half of~2011.  In the interim, we can still explore the OB associations with the {\\em Hipparcos} catalogue and different aims, tools, and reductions of the raw data. On the one hand, in the current work, we have used the late re-reduction of the {\\em Hipparcos} data by van~Leeuwen (2007a), who has obtained improvements by up to a factor 4 in the astrometric accuracies for particular bright stars (see also details on the validation of the new {\\em Hipparcos} reduction in van~Leeuwen [2007b]). On the other hand, we have applied a novel (in Galactic astronomy) clustering analysis to part a dataset into subsets (``clusters'' or ``agglomerates''). Through {\\em data clustering}, objects with a common distinguishing feature are classified into different groups. In our case, the common quality is just the spatial location of {\\em Hipparcos} stars with blue colours (i.e. with early spectral types). That is, our aim is to study the structure (and super- and sub-structure) of Galactic OB associations, and derive some basic properties. ",
        "conclusions": "\\subsection{Two new clusters? $\\eta$~Orionis and P~Puppis}  \\subsubsection{$\\eta$~Orionis} \\label{etaOri}  Among the 35 identified agglomerates, two of them have never been proposed to be real clusters (i.e. their stars are gravitationally bound). They are the $\\eta$~Orionis overdensity and the P~Puppis cluster candidate.  As already mentioned in Section~\\ref{escorial4}, the centre of the $\\eta$~Orionis overdensity does not coincide with the massive (double) star $\\eta$~Ori~AB. In order to avoid the limits of the {\\em Hipparcos} catalog and our colour selection, we carried out a new Virtual Observatory analysis of the region. In particular, we loaded with Aladin all Tycho-2 stars at less than 2\\,deg to the central coordinates in Section~\\ref{escorial4} ($N$ = 734), and selected all of them with colours $B_T - V_T <$ +0.25\\,mag and magnitudes $V_T <$ 10.50\\,mag ($N$ = 96). We checked that all the selected stars have very low proper motions ($\\mu \\le$ 10\\,mas\\,a$^{-1}$) and follow a relatively narrow secuence in the $B_T$ vs. $B_T - V_T$ colour-magnitude diagram, consistent with its membership in Orion, and that redder non-selected stars of similar brightess have colours (and parallaxes) typical of field dwarfs in the foreground.  We counted the number of selected stars in coronae centred on both $\\eta$~Ori~AB and the survey area central coordinates, and did not find a clear evidence of a {\\em dense} star agglomeration as has been found in other sites in Ori~OB1 (like in the Orion Nebula Cluster and the $\\sigma$~Orionis cluster). Besides, the elongated shape of the overdensity gets more patent with the Tycho-2 data, which contrasts with the symmetrical, radial density gradient found in well-characterized clusters (Cartwright \\& Whitworth 2004; Caballero 2008a). To sum up, we confirmed that a spread overdensity of bright, early-type stars does exist in the area of $\\eta$~Ori~AB in the Ori~OB1a association. However, we failed to corroborate the existence of a (gravitationally bound) cluster in that region.   \\subsubsection{P~Puppis} \\label{PPup}  \\begin{figure*} \\centering \\includegraphics[width=0.48\\textwidth]{thePPupBTBTVT.eps} \\includegraphics[width=0.48\\textwidth]{thePPupiiKs.eps} \\includegraphics[width=0.48\\textwidth]{thePPupdeltaalpha.eps} \\includegraphics[width=0.48\\textwidth]{thePPupfrr.eps} \\caption{Some diagrams illustrating the follow-up of the P~Puppis cluster. {\\em Top left:} $B_T$ vs. $B_T - V_T$ colour-magnitude diagram showing Tycho-2 cluster member candidates at less than 45\\,arcmin to P~Pup~AC (filled --red-- stars) and remaining objects at up to 90\\,arcmin to the early-type binary (small --blue-- dots). {\\em Top right:} the faintest, reddest portion of the $i$ vs. $i-K_{\\rm s}$ diagram showing DENIS/2MASS cluster member candidates at less than 45\\,arcmin to P~Pup~AC (filled --red-- circles) and remaining objects at the same angular separation (small --blue-- dots). Solid (red) lines are for the 5, 10, and 20\\,Ma-old NextGen98 isochrones (from top to bottom) shifted at $d$ = 470\\,pc, while dashed (red) lines are for the previous isochroness shifted at $d$ = 400 (above) and 570\\,pc (below). {\\em Bottom left:} spatial distribution of the Tycho-2 (filled stars) and DENIS/2MASS (filled circles) cluster member candidates. {\\em Bottom right:} normalized cumulative number, $f(r)$, of cluster member candidates, in solid (red) line, and theoretical power-law distributions for $f(r) \\propto r^{1/2}$, $r^1$, and $r^2$ (from top to bottom). Compare with figures in Caballero~(2008a).} \\label{thePPup} \\end{figure*}  We followed a similar procedure to that of the $\\eta$~Orionis overdensity in the previous section for studying the possible cluster nature of the agglomerate Escorial~25/P~Puppis. First, we selected all Tycho-2 stars satisfying the colour and magnitude restrictions listed above. In this case, we only loaded the stars in a 90\\,arcmin-radius circle centred on P~Pup~AC. The photometric sequence in the \\linebreak colour-magnitude diagram in top left panel of Fig.~\\ref{thePPup} is very sharp, as already noticed in Section~\\ref{escorial25} with the {\\em Hipparcos} data. Remarkably, the three brightest stars in the agglomerate (i.e. the two resolved components in P~Pup~AC and the subgiant QS~Pup) lie on a circle with a radius of only 45\\,arcmin, which subtends and area four times smaller than the 90\\,arcmin-radius circle. Far from expecting four times less blue, bright stars, approximately half of them are within the smaller radius (i.e. there are twice more blue bright stars in the inner circle than expected -- the deviation is at the level of five times the square root of the Poissonian variance), suggesting a {\\em dense} star agglomeration.  We probed the very low-mass stellar population (close to the substellar boundary) in the 45\\,arcmin-radius circle using the Dee Near Infrared Survey of the Southern Sky (DENIS; Epchtein et~al. 1997) and the Two-Micron All Sky Survey (2MASS; Skrutskie et~al. 2006) and the NextGen98 theoretical isochrones of Baraffe et~al. (1998). We proceed as in Caballero \\& Solano (2008), by loading DENIS and 2MASS data and cross-matching the sources. We filtered the faintest ($i >$ 10\\,mag), reddest ($i-K_{\\rm s}$) sources with the highest quality photometry ($\\delta i$, $\\delta J$, $\\delta H$, $\\delta K_{\\rm s} <$ 0.1\\,mag). Since the Virtual Observatory selection of P~Puppis cluster member candidates cannot be carried out through the definition of a lower envelope of spectroscopically confirmed young stars (as in the $\\sigma$~Orionis; Caballero 2008c) or of a redder envelope that both maximizes the number of cluster members and minimizes the number of contaminants (as in the Alnilam/Mintaka region; Caballero \\& Solano 2008), we opted for the classical selection criterion based on theoretical isochrones. As can be seen in the colour-magnitude diagram in top left panel in window Fig.~\\ref{thePPup}, the theoretical P~Puppis sequence, considering all possible uncertainties in the cluster heliocentric distance ($d$ = 470$\\pm$70\\,pc) and age ($10^{+10}_{-5}$\\,Ma), is quite broad and overlaps with the location of a lot of (foreground) field M-type dwarfs. Other colour-magnitude diagrams, performed with the same technique in regions of exactly the same area and galactic latitude, and of approximately the same galactic longitude, are quite similar, which indicates that many of the selected cluster member candidates (between the 5\\,Ma, 400\\,pc and the 20\\,Ma, 540\\,pc NextGen isochrones) are actually late-type interlopers. Although some elaborated de-contamination technique could be applied (e.g. Caballero, Burgasser \\& Klement 2008), a central overdensity of selected cluster members is already obvious (see the spatial distribution of selected sources in the bottom left panel in Fig.~\\ref{thePPup}). Since there are 425 DENIS/2MASS selected sources in the 45\\,arcmin-radius circle, one should expect $\\sim$100 such objects in the four-times smaller area defined by a 22.5\\,arcmin-radius circle centred on P~Pup~AC. In contrast, we found more than 150 of them. In other words, there are at least $\\sim$50 (young) low-mass stars at less than 20--25\\,arcmin to the B0 giant that probably form part of the P~Puppis cluster.  Other way of illustrating the central overdensity of \\linebreak P~Puppis is the diagram of the normalized cumulative number of cluster member candidates, $f(r) = N(r) / N(r_{\\rm max})$, as a function of angular distance to the the centre, $r$. A wealth of details on how to compute and interpret it are in Caballero (2008a). In short, a power-law distribution $f(r) \\propto r^2$ corresponds to a uniform distribution of objects in the survey area, while $f(r) \\propto r^1$ corresponds to a volume density proportional to $r^{-2}$, which is consistent with the collapse of an isothermal spherical molecular cloud (e.g. Burkert, Bate \\& Bodenheimer 1997). A distribution with an intermediate index, as that one observed in the Mintaka cluster (Caballero \\& Solano 2008) and in P~Puppis from $r$ = 0 to $\\sim$25\\,arcmin (bottom right panel in Fig.~\\ref{thePPup}) corresponds to the superposition of a radially-concentrated cluster with $f(r) \\propto r^1$ and a homogeously-distributed population of (late-type, field dwarf) contaminants.  To sum up, we discover a star cluster surrounding the young, massive P~Pup~AC giant. However, a dedicated spectroscopic follow-up will be necessary to disentangle \\linebreak between the stellar populations of the cluster and of the field.   \\subsection{New sites for substellar searches} \\label{browndwarfs}  Some of the agglomerates identified in our work are valuable ``hunting grounds'' for the search of faint brown dwarfs and planetary-mass objects. In particular, the first brown dwarf, \\object{Teide~1}, was found in the Pleiades cluster (Rebolo et~al. 1995), while the Orion Nebula and the $\\sigma$~Orionis clusters harbour the great majority of the planetary-mass objects with spectroscopy (Lucas \\& Roche 2000; Zapatero Osorio et~al. 2000). Other listed agglomerates that have held substellar searches are the Orion Belt (surrounding Alnilam and Mintaka), NGC~2264, NGC~2547/$\\gamma$~Velorum, and IC~2391. However, there are agglomerates in our list that share properties (youth, closeness, low extinction) that also make them ideal sites for substellar searches.  After discarding the possible spurious agglomerates and the groupings associated to Collinder~121 and Vel~OB2, that are too far away and/or too sparse for efficient searches, we retained 14 agglomerates for a follow-up inspection of their parameters. Of them, five seem to be at prohibitive heliocentric distances larger than 650\\,pc (i.e. NGC~2353, Collinder~132, NGC~2362, Pismis~4, and IC~2395). Besides, M~47, although being located at $d \\sim$ 470\\,pc, has a Pleiades-like age and, therefore, their hypothetical brown dwarfs will be about 3\\,mag fainter.  We maintain eight agglomerates as new sites for substellar searches. Two of them are obvious choices, because they resemble their ``elder brothers'' (in contraposition to ``fraternal twins'') in Orion: 25~Orionis and $\\eta$~Orionis. Even in the case that $\\eta$~Orionis is not a real cluster, the large observed stellar surface density will facilitate forthcoming studies (stellar and substellar objects are expected to have proportional surface densities; Caballero 2008a). \\linebreak NGC~2451~A, with an age of $\\sim$45\\,Ma and a very close heliocentric distance of $d$ = 180--200\\,pc (see Section~\\ref{distance.ic2391}) is other apparent option.  If compiled basic data for vdBH~23 (Escorial~27) and Trumpler~10 (Escorial~33), with distances and ages in the approximate intervals 320--360\\,pc and 1--4\\,Ma, are correct, then they might be new cornerstones for the substellar \\linebreak searches in the near future. Finally, there remain in order of interest, NGC~2232 (Escorial~8; $d \\sim$ 320\\,pc, age $\\sim$ 30\\,Ma), Collinder~140 (Escorial~19; $d \\sim$ 370\\,pc, age $\\sim$ 20\\,Ma), and our new cluster P~Puppis (Escorial~25; $d \\sim$ 470\\,pc, age $\\sim$ 10\\,Ma). The first two clusters in the trio seem to be IC~2391 twins at larger distances, while P~Puppis resembles the stellar groupings in the Orion~Belt.   \\subsection{Sub-structure of super-agglomerates} \\label{sub-structure}  We find evidence of spatial sub-structure within classical OB associations. First, the ``isolated cluster'' Escorial~35 could be a physical grouping within Lac~OB1 \\linebreak (Section~\\ref{escorial3435}). Secondly, our three super-agglomerates \\linebreak (Orion, CMa-Pup, and Pup-Vela) are, by definition, the \\linebreak combination of a number of individual agglomerates. There can be sub-sub-structure (i.e. sub-structure within the agglomerates), as described in, for example, Sections~\\ref{escorial9} (NGC~2264), \\ref{escorial6} (Orion Belt), and~\\ref{escorial2427} (NGC~2451~AB), but it was not our aim to detect~it.    The Orion super-agglomerate is composed of five agglomerates, whose boundaries do {\\em not} coincide with classical divisions within the Ori~OB1 complex. The Orion Sword (Escorial~5) is the superposition of 23 stars previously ascribed to the sub-associations Ori~OB1c[1--4] and \\linebreak Ori~OB1d, and includes the brightest stars of, e.g., the Orion Nebula Cluster. The Orion Belt (Escorial~6) coincides with Ori~OB1b, except for the fact that it does not contain the stars in the nearby $\\sigma$~Orionis cluster, that form the largest spatial overdensity in a new agglomerate of early-type stars close to the Horsehead Nebula (Escorial~7). Finally, we \\linebreak identify two agglomerates surrounding the very bright stars 25~Ori (Escorial~3) and $\\eta$~Ori (Escorial~4) in the Ori~OB1a sub-association. While the 25~Ori agglomerate was identified very recently, we report here for the first time the existence of a possible open cluster surrounding $\\eta$~Ori.  The classical view of sequential formation of subgroups in OB associations by Elmegreen \\& Lada (1977) is reinforced by our results. As shown in Fig.~\\ref{theaggregates.orion}, the stars in the Orion super-agglomerate are aligned in the north-south direction. Noticeably, the oldest (or less young) agglomerate is located northward (25~Ori; $\\sim$10\\,Ma), the youngest one southward (Orion Sword; $\\sim$1\\,Ma), and those ones with intermediate ages (Orion Belt, Horsehead, and, possibly, \\linebreak $\\eta$~Ori) in the middle. Extrapolating backward, it might happen that the $\\lambda$~Orionis cluster (to the north of the Orion super-agglomerate) is older than 25~Ori. On the other hand, the ionization and front shocks produced in the Orion Sword might also initiate in the near future another cycle of OB star formation to the south of Orion, where large concentrations of molecular gas have been detected by the {\\em IRAS}~satellite.    We have broken up the CMa-Pup super-agglomerate \\linebreak into Collinder~132 (Escorial~17) and Collinder~140 (Escorial~19) to the south and Collinder~121 to the north. \\linebreak Collinder~121 has been splitted, in its turn, into ten agglomerates. One of them, Escorial~10, would be the original \\linebreak Collinder~121 compact group that gives the name to the extended, loose OB association. The few dedicated spectro-photometric analyses in the region do not allow to conclude if there are younger and older regions within the \\linebreak super-agglomerate. The spatial sub-structure suggests, however, that Collinder~121 is an analog to the Ori~OB1 association at about 1\\,kpc to the~Sun.    The Pup-Vel super-agglomerate displays a much larger spatial distribution complexity. Indeed, the large spread on heliocentric distances, from $\\overline{d}$ = 147$\\pm$6\\,pc of IC~2391 to the $d \\sim$ 800\\,pc of IC~2395 indicates that not all of them are physically bounded. NGC~2451~A, at $\\overline{d}$ = 181$\\pm$9\\,pc, would be another open cluster in the foreground of the extended Vel~OB2 association. However, the remaining agglomerates in the super-agglomerate could actually have a common origin. All of them except Pismis~4 ($d \\sim$ 500--600\\,pc) lie in a relatively narrow interval of heliocentric distances: ($i$) vdBH~23 at $\\overline{d}$ = 320$\\pm$30\\,pc; ($ii$) NGC~2451~B at $\\overline{d}$ = 340$\\pm$30\\,pc; ($iii$) Trumpler~10 at $d \\sim$ 360\\,pc; ($iv$) $\\gamma$~Velorum, the core of the Vel~OB2 association, NGC~2547, and V468~Pup at $d$ = 400--450\\,pc; and ($v$) P~Puppis at $\\overline{d}$ = 470$\\pm$70\\,pc. The mean age of these star clusters is about 30\\,Ma, with only vdBH~23 having been reported to be significatively younger (age $\\sim$ 4\\,Ma). Although a careful investigation of the parallactic distances to the stars in these agglomerates will not be carried out until the {\\em GAIA} mission is operative, the resemblance between the ages and distances of $\\gamma$~Velorum and NGC~2547 (that coincide with the trio of agglomerates Escorial~26, 28, and~29) suggests that an important fraction of the Vel~OB2 association is actually the superposition of different star-forming regions {\\em \\`a~la}~Orion.   \\subsection{Early-type stars at $d \\le$ 200\\,pc}  A bit less than 10\\,\\% of the investigated OB-type stars are at relatively short heliocentric distances, at $d \\le$ 200\\,pc. Most of them belong to the nearby clusters Pleiades (Escorial~1), NGC~2451~A (Escorial~24), and IC~2391 (Escorial~31), but there are also possible contaminants in the foreground of our agglomerates (listed in last column in Table~\\ref{agglomerates}), and errors in the {\\em Hipparcos} catalogue. On the one hand, cluster stars with accurate parallax measurement are important to derive cluster distances, and compare its spectro-photometric sequence with theoretical models. On the other hand, some of the foreground early-type stars may belong to the rare type of bright B-type dwarfs unbound to known open clusters, which are generally subject of intensive follow-ups in the literature. Next, we describe the identified early-types stars at $d \\le$~200\\,pc.   \\subsubsection{Pleiades} \\label{distance.pleiades}  Because of its importance for the ``cosmic distance ladder'', the disagreement on the actual heliocentric distance to the Pleiades is still an open problem in Astrophysics (see recent works by Zwahlen et~al. [2004], Soderblom et~al. [2005], Southworth, Maxted \\& Smalley [2005], and Percival, \\linebreak Salaris \\& Groenewegen [2005]). Far from taking back again the dispute between the main-sequence fitting and the {\\em Hipparcos} distance moduli, we have computed an additional measurement of the Pleiades heliocentric distance by \\linebreak weight-averaging the re-reduced {\\em Hipparcos} parallaxes of \\linebreak the ten stars in our agglomerate Escorial~1. The computed distance to the Pleiades, $\\overline{d}$ = 120$\\pm$5\\,pc, is still shorter than the adopted value of 134--139\\,pc derived from isochrone fitting, but slightly larger than the previous (probably incorrect) {\\em Hipparcos} distance of 118\\,pc. The two {\\em Hipparcos} measurements are, however, consistent with each other within error bars. There is no especial improvement in our $\\overline{d}$ if we also add Asterope and Calaeno, classical Pleiads with {\\em Hipparcos} parallax that were not in our input catalogue (Section~\\ref{escorial1}).   \\subsubsection{NGC 2451 A} \\label{distance.ngc2451a}  In Section~\\ref{escorial2427}, we identified seven {\\em Hipparcos} stars as \\linebreak members of the NGC~2451~A cluster. Our weight-averaged heliocentric distance to the cluster, $\\overline{d}$ = 181$\\pm$9\\,pc, is consistent with (although slightly lower than) the value derived by Robichon et~al. (1999) from the original {\\em Hipparcos} data ($d$ = 188.7$^{+7.0}_{-6.5}$\\,pc). Carrier et~al. (1999) derived, however, $d$ = 197$\\pm$12\\,pc using Geneve photometry and isochrone fitting. Nevertheless, the difference between this value and ours, $\\Delta d$ = 16$\\pm$15\\,pc, could be lower if state-of-the-art theoretical tracks and improved metallicity and age were used (as in Platais et~al. [2007] for IC~2391 -- see below). As a result, there is no hint of NGC~2451~A  displaying a Pleiades-like disagreement between  heliocentric and isochrone distances.   \\subsubsection{IC 2391} \\label{distance.ic2391}  IC~2391 (Section~\\ref{escorial30313233}), with an age of about 40\\,Ma, is probably the youngest open cluster in the solar vicinity (see, however, Barrado y Navascu\\'es et~al. [2004] and references therein). Platais et~al. (2007) have provided the most complete discussion on the problem of the distance to IC~2391. They derived an isochrone distance of 159.2\\,pc, while the classical {\\em Hipparcos} distance is $d \\sim$ 146\\,pc. From the \\linebreak weight-average of the parallaxes of the seven stars in our Escorial~31 agglomerate, we derive a new {\\em Hipparcos} distance of $\\overline{d}$ = 147$\\pm$6\\,pc. The distance controversy is, therefore, still~standing.   \\subsubsection{Early-type field dwarfs} \\label{fielddwarfs200pc}  \\begin{table*} \\caption[]{Early-type field dwarfs at $d \\le$ 200\\,pc.} \\label{fielddwarfs} $$ \\begin{tabular}{c l l c c c} \\hline \\hline \\noalign{\\smallskip} HIP\t\t& Name\t\t\t& Spectral\t& $d$\t\t& $M_V$\t& $B-V$\t\\\\ & \t\t\t& type\t\t& [pc]\t\t& [mag] & [mag] \\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} 25469\t\t&        {HD 35716}   \t& B9\t\t& 180$\\pm$30\t& 2.21\t&--0.03\t\\\\ % 25592\t\t&        {HD 35926} AB  & B7IV+\t\t& 180$\\pm$40\t& 2.08\t&--0.07\t\\\\ % 25600\t\t&        {HD 35957}\t& B8\t\t& 190$\\pm$40\t& 2.07\t&--0.05\t\\\\ % 35083\t\t&        {HD 56342}   \t& B3V\t\t& 193$\\pm$8\t&--1.10\t&--0.14\t\\\\ % 37173\t\t&        {m Pup} A(C)B  & B8IV++\t& 190$\\pm$9\t&--1.72\t&--0.09\t\\\\ % 37229\t\t&      {k$^{01}$~Pup} A & B6V\t\t& 110$\\pm$8\t&--0.79\t&--0.16\t\\\\ % &      {k$^{01}$~Pup} B & B5IVn\t\t& 110$\\pm$8\t&--0.58\t&--0.11\t\\\\ % 37304\t\t&        {HD 61687} AB \t& B6V+\t\t& 193$\\pm$18\t& 0.34\t&--0.11\t\\\\ % 39716\t\t&        {HD 67704}  \t& A0V\t\t& 200$\\pm$14\t& 0.45\t&--0.06\t\\\\ % \\noalign{\\smallskip} \\hline \\end{tabular} $$ \\end{table*}  In Table~\\ref{fielddwarfs}, we list eight foreground agglomerate contaminants at less than $d \\le$ 200\\,pc (one binary was actually resolved by Tycho-2 in its two components). Apart from shorter heliocentric distances than to the corresponding agglomerates, they display different proper motions and, in most cases, are located in abnormal positions in the colour-magnitude diagrams. Six of them are B6--A0 dwarfs at distances $d \\approx$ 180--200\\,pc and are not of especial interest. \\linebreak However, the two remaining stars stand~out.  As described in Section~\\ref{escorial1719}, {HD~56342} (B3V, $V_r$ = +33.4$\\pm$2.8\\,km\\,s$^{-1}$, $v \\sin{i}$ = 26\\,km\\,s$^{-1}$) is a contaminant of the poorly known Collinder~132 cluster. The differences between the HD~56342 and cluster heliocentric distances, \\linebreak $\\Delta d \\approx$ 450\\,pc, and proper motion, $\\Delta \\mu \\approx$ 11\\,mas\\,a$^{-1}$, make unlikely its membership in Collinder~132. The star has been spectroscopically investigated in detail by Lyubimkov et~al. (2002) and subsequent papers. HD~56342 has typical effective temperature, mass, surface gravity, and helium abundance for its spectral type and class. Furthermore, they derived $d$ = 209$\\pm$36\\,pc based solely on spectro-photometric parameters, which is consistent with our parallactic distance ($d$ = 193$\\pm$8\\,pc). % Savage et~al. (1985) and Berghoefer, \\linebreak Schmitt \\& Cassinelli (1996) also derived independent distances at  $d \\approx$ 240 and 236\\,pc from ultraviolet interstellar extinction with the Astronomical Netherlands Satellite and from hydrogen column density with {\\em ROSAT}. With age and mass of 54$\\pm$9\\,Ma and 5.3$\\pm$0.3\\,$M_\\odot$, respectively (Lyubimkov et~al. 2002), HD~56342 is one of the very few early B dwarfs in the solar neighbourhood that do not belong to a known~cluster.  The case of {k$^{01}$~Pup}~AB (CD--26~4707) is even more extreme. Its proper motion, of $\\mu \\sim$ 29\\,mas\\,a$^{-1}$, is inconsistent with membership in the Escorial~23 agglomerate, that probably forms part of the Collinder~121 complex ($d \\sim$ 0.5--1.0\\,kpc). {k$^{01}$~Pup}~AB is a binary system of B5--6 stars (HD~61555 and HD~61556: $\\rho$ = 9.913$\\pm$0.003\\,arcsec, $\\theta$ = 318\\,deg, $\\Delta H_P$ = 0.21$\\pm$0.01\\,mag; Perryman et~al. 1997) located at only $d$ = 110$\\pm$8\\,pc. This distance agrees with that estimated by Lindroos (1985), at $d \\sim$ 125\\,pc. He also derived a very young age, of about 12\\,Ma. The evolved nature of the secondary (class IV), that is a helium variable star (Rivinius et~al. 2003), may indicate a slightly older age. In any case, {\\em if} the {\\em Hipparcos} parallactic distance is correct (binarity may have affected the parallax measurement), {k$^{01}$~Pup}~AB could be the closest massive very young star (age $\\le$ 50\\,Ma, $M \\ge$ 5\\,$M_\\odot$), even closer than the OB-type stars in Upper~Scorpius, R~Coronae~Australis, $\\rho$~Ophiuchi, Chamaeleon~I+II, and IC~2391 ($d$ = 130--150\\,pc)\\footnote{The most massive nearby star is probably \\object{$\\beta$~Pic} (A6V, $d$ = 19.27$\\pm$0.19\\,pc, age $\\sim$ 12\\,Ma). {k$^{01}$~Pup}~AB, roughly contemporary, would be five times further, but would also be much more massive.}. We are ignorant of the star-forming region where {k$^{01}$~Pup}~AB was born or if the binary has an associated young moving~group.   \\subsection{Missing agglomerates? Missing stars?}  Obviously, our {\\em clustering} algorithm is not able to identify all kinds of star {\\em clusters}, but is mostly sensitive to agglomerates with at least six bright blue stars separated by less than $R_\\epsilon$ (in our case, $R_\\epsilon$ = 0.8\\,deg). Missing cluster types are globular (e.g. $\\omega$~Cen, M~13) and super star clusters (e.g. Westerlund~1). None of their stellar components were in the input catalogue described in Section~\\ref{theinputcatalogue} because of their faintness or reddening in the optical $B_T V_T$ bandpasses (due, in their turn, to very large heliocentric distances, evolution towards the red giant and asymptotic giant branches, or high interstellar extinction).  Some well-known open clusters are also missing. These lacking open clusters can be divided into two classes: ($i$) clusters with ages in the approximate interval 100--600\\,Ma (e.g. $\\alpha$~Persei, NGC~2516, M~35, Prasepe, Hyades), and ($ii$) clusters with ages in the approximate interval 1--10\\,Ma (e.g. Serpens, MBM~12, $\\rho$~Ophiuchi, IC~348, $\\lambda$~Orionis, $h+\\chi$~Persei). There exist also explanations for their non-detection: for example, there is no Hyades member with {\\em Hipparcos} colour bluer than $B_T-V_T$ = 0.0\\,mag. The brightest Hyads, Aldebaran and $\\theta^{02}$~Tau, are K5III and A7III giants, while the bluest Hyad with reliable {\\em Hipparcos} photometry, 68~Tau~AB, is an A2IV subgiant. In other words, because of stellar evolution, the Hyades turn-off point is redwards of our colour selection criterion in Section~\\ref{theinputcatalogue}. Something similar happens to the other 100--600\\,Ma open~clusters.  In contrast, the effect of the post-main sequence evolution is barely detectable in the 1--10\\,Ma-old open clusters. However, the early-type stars in the youngest ones (e.g. $\\rho$~Ophiuchi, IC~348) also have red colours for their respective spectral types because of intra-cluster extinction (there are abundant populations of Class~0 and~I objects --recently born stars surrounded by thick shells-- in many of these star-forming regions). The remaining young open clusters with relatively low extinction (e.g. $\\lambda$~Orionis) have looser, less abundant, early-type stellar populations than our 35 agglomerates. Our clustering algorithm is not sensitive, either, to the detection of sparse very young associations, like Taurus-Auriga and TW~Hydrae, and low-density Galactic OB associations, like many of those listed in Garmany \\& Stencel~(1992).   \\subsection{Effect of the parametrization on the results} \\label{effectoftheparametrization}  We have investigated how a different choice of parameters of the DBSCAN algorithm (in particular, $R_\\epsilon$) would have affected our results. On the one hand, Table~\\ref{agglomerates0.6} shows the early-type {\\em Hipparcos} stars that belong to agglomerates \\linebreak when we use $R_\\epsilon$ = 0.6\\,deg instead of 0.8\\,deg. With the new value, there appear two basic differences: ($i$) there are only 21 agglomerates, and ($ii$) the agglomerates contain no more than 11 stars. For example, Escorial~1 (the Pleiades) now contains only six stars, while many other agglomerates, including the young open clusters Collinder~132, M~47, and IC~2395, are not identified at all. Besides, Escorial~3 (25~Ori) is splitted into two different sub-agglomerates (Escorial~3a and 3b), which may indicate the presence of \\linebreak sub-sub-structure in Ori~OB1a. Finally, HIP~26727 \\linebreak (Mintaka), that is one of the three bright supergiants in Escorial~6 (Orion Belt), is classified now as a member of Escorial~7 (Horsehead and $\\sigma$~Orionis). To sum up, $R_\\epsilon$ = 0.6\\,deg gives an incomplete sampling of our aggregates.  On the other hand, further increasing $R_{\\epsilon}$ to 1.0\\,deg results in the fusion of some aggregates, the detection of some new clusters, and the growth of nearby agglomerates \\linebreak (Fig.~\\ref{thegalaxy}). In~particular:  \\begin{itemize} \\item 6 agglomerates remain unchanged with respect to the $R_\\epsilon$ = 0.8\\,deg parametrization (Escorial~1 [Pleiades], 2 [spurious], 4 [$\\eta$~Ori], 9 [NGC~2264], 27 [vdBH~23], and~32 [IC~2395]).  \\item 14 agglomerates ``capture'' nearby bright blue stars and increase in size. Ten of the agglomerates accrete only three or less stars. There are, however, three agglomerates that drastically grow up, capturing 8 (Escorial~3 --25~Ori--), 9 (Escorial~16 --NGC~2232--), and up to 12 bright blue stars (Escorial~26 --$\\gamma$~Velorum and \\linebreak NGC~2547--).  \\item 15 agglomerates are fused into five larger entities (with additional capture of nearby bright blue stars). The resultant groupings are the combination of Escorial~6 \\linebreak and~7 (Orion Belt and Horsehead in the Orion super-agglomerate), 10--15~and~17, 18~and~20, 22~and~23 (all agglomerates associated to Collinder~121 along with \\linebreak Collinder~132), and 28~and~29 (in the dense trio in \\linebreak Vel~OB2; Section~\\ref{escorial262829}). As a result, using $R_{\\epsilon}$ = 1.0\\,deg, the CMa-Pup super-agglomerate is splitted into only \\linebreak four groupings: Collinder~140 (Escorial~19, the only agglomerate in CMa-Pup that maintain its independence) and the three fusions in Collinder~121. The new Orion division gives the same number of groupings (25~Ori, $\\eta$~Ori, Orion Sword, and ``Orion Belt+Horsehead''). Remarkably, the large amassment Escorial~10--15~and~17 contains 135 bright blue stars, a value that does not differ very much from the number of such stars in \\linebreak de~Zeeuw et~al. (1999)'s Collinder~121 asociation.  \\item 19 agglomerates are new. All of them, except three, contain six or seven stars and have surface densities that are $\\pi 0.8^2 / \\pi 1.0^2$ = 0.64 times smaller than our spurious agglomerate Escorial~2. Many of them are probably spurious agglomerates as well, and they would require a careful follow-up, as that carried out for the original parametrization. Two of the remaining agglomerates lie on CMa-Pup-Vel region, with N$_\\star$ = 8 and 9, and may belong to the sparse stellar population of the local spiral arm of the Galaxy (Section~\\ref{results}). Finally, there is a new isolated agglomerate with a relative large number of stellar components, N$_\\star$ = 10, and separated from the large overdensities in Fig.~\\ref{thegalaxy}. It includes the B0Vp-type star \\object{$\\theta$~Car}, that is the brightest one of the IC~2602 open cluster (age $\\sim$ 30\\,Ma, $d$ = 135$\\pm$9\\,pc -- Randich et~al. 1995; Stauffer et~al.~1997).  \\end{itemize}  \\begin{table*} \\caption[]{Agglomerates of early-type {\\em Hipparcos} stars with the alternative parameters $N_{\\rm MinPts}$ = 6 and $R_\\epsilon$ = 0.6\\,deg.} \\label{agglomerates0.6} $$ \\begin{tabular}{l l l} \\hline \\hline \\noalign{\\smallskip} No.\t& Original \t& {\\em Hipparcos} stars\t\t\t\\\\ & agglomerate \t& (HIP)\t\t\t\t\t\\\\ & (Escorial) \t& \t\t\t\t\t\\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} I\t& 1    \t\t& 17499, 17527, 17531, 17573, 17608, 17702 \\\\ II\t& 2    \t\t& 23279, 23287, 23295, 23328, 23473, 23508 \\\\ III\t& 3a   \t\t& 25163, 25235, 25288, 25302, 25340, 25469 \\\\ IV\t& 3b   \t\t& 25241, 25378, 25411, 25533, 25567, 25582, \\\\ &      \t\t& 25592, 25648, 25655, 25751, 25752 \\\\ V\t& 4    \t\t& 25293, 25394, 25480, 25552, 25557, 25600 \\\\ VI\t& 5    \t\t& 26197, 26199, 26241, 26314, 26345, 26427 \\\\ VII\t& 6    \t\t& 26106, 26213, 26311, 26319, 26334, 26405, \\\\ &      \t\t& 26439, 26464, 26508, 26683 \\\\ VIII\t& 7    \t\t& 26549, 26551, 26579, 26656, 26694, 26713, \\\\ &      \t\t& [26727] \\\\ IX\t& 8    \t\t& 30580, 30660, 30700, 30758, 30761, 30772, \\\\ &      \t\t& 30789 \\\\ X\t& 9    \t\t& 31917, 31951, 31955, 31978, 32030, 32053 \\\\ XI\t& 10   \t\t& 33062, 33070, 33087, 33165, 33208, 33215, \\\\ &      \t\t& 33276, 33309, 33410 \\\\ XII\t& 11   \t\t& 33586, 33666, 33695, 33721, 33796, 33841 \\\\ XIII\t& 12   \t\t& 33935, 33970, 34041, 34048, 34167, 34227 \\\\ ...\t& 13   \t\t& ... \\\\ ...\t& 14   \t\t& ... \\\\ ...\t& 15   \t\t& ... \\\\ ...\t& 16   \t\t& ... \\\\ ...\t& 17   \t\t& ... \\\\ XIV\t& 18   \t\t& 35267, 35370, 35412, 35461, 35503, 35539, \\\\ &      \t\t& 35597 \\\\ XV\t& 19   \t\t& 35700, 35761, 35795, 35822, 35855, 35905, \\\\ &      \t\t& 36045 \\\\ ...\t& 20   \t\t& ... \\\\ ...\t& 21   \t\t& ... \\\\ ...\t& 22   \t\t& ... \\\\ ...\t& 23   \t\t& ... \\\\ ...\t& 24   \t\t& ... \\\\ XVI\t& 25   \t\t& 37926, 37953, 38020, 38028, 38159, 38164 \\\\ XVII\t& 26   \t\t& 39873, 39919, 40011, 40016, 40024, 40059 \\\\ XVIII\t& 27   \t\t& 40218, 40255, 40268, 40274, 40321, 40324 \\\\ ...\t& 28   \t\t& ... \\\\ XIX\t& 29   \t\t& 40662, 40742, 40825, 40851, 40872, 40921 \\\\ ...\t& 30   \t\t& ... \\\\ XX\t& 31   \t\t& 42400, 42459, 42504, 42535, 42536, 42715, \\\\ &      \t\t& 42726 \\\\ ...\t& 32   \t\t& ... \\\\ XXI\t& 33   \t\t& 43055, 43085, 43182, 43209, 42340, 43326, \\\\ &      \t\t& 43392 \\\\ ...\t& 34   \t\t& ... \\\\ ...\t& 35   \t\t& ... \\\\ \\noalign{\\smallskip} \\hline \\end{tabular} $$ \\end{table*}   \\subsection{Possible errors in the input catalogue} \\label{possiblehipparcoserrors}  Apart from the large number of {\\em Hipparcos} stars without parallax determination (indicated with ellipses ``...'' in the $d$ column in Table~\\ref{theaggregates}), we have also identified two stars that suffered from systematic errors during the {\\em Hipparcos} reduction. Both of them display proper motions with very large error bars and improbable distances of a few tens AU (indicated with square brackets ``[~]'' in the $d$ column in Table~\\ref{theaggregates}). They are the binary HIP~23279 + HIP~23287 (\\object{HD~32039} + \\object{HD~32040}, B9Vn, in Escorial~2; $d$ = 34$\\pm$9\\,pc) and HIP~35503 (\\object{HD~57281}~AB, B5V, in Escorial~18; $d$ = 70$\\pm$30\\,pc). We have not taken them into account in the analysis.  Finally, it seems that there was a misunderstanding in the Bonner Durchmusterung and/or the {\\em Hipparcos} \\linebreak catalogues between two 10--11th-magnitude stars separated by $\\sim$57\\,arcsec. The actual K2-type star BD+05~1825 \\linebreak (TYC~189-1314-1), presented in Section~\\ref{theinputcatalogue}, is the easternmost and brightest one of the pair at passbands \\linebreak $V_T R I J H K_{\\rm s}$. The westernmost star (HIP~38575, TYC-189-1700-1), that is brighter than BD+05~1825 only at passband $B_T$, is the only {\\em Hipparcos} star in the area. HIP~38575 has typical colours of late B or early A~dwarfs."
    },
    "0807/0807.3896_arXiv.txt": {
        "abstract": "{From a macroscopic  theory of the quantum vacuum in terms of conserved relativistic charges (generically denoted by $q^{(a)}$ with label $a$), we have obtained, in the low-energy limit, a particular type of $f(R)$ model relevant to cosmology. The macroscopic quantum-vacuum theory allows us to distinguish between different phenomenological $f(R)$ models on physical grounds.}  \\begin{document} ",
        "introduction": "The development of high-energy physics and cosmology over the last years has led to the realization that, most likely, Einstein's theory of gravity needs to be modified. At the moment, we are only able to calculate higher-order curvature corrections to the Einstein action coming from scalar, spinor, and vector quantum fields propagating over a fixed classical background (see, e.g., Refs.~\\cite{BirrellDavies1982,CopelandSamiTsujikawa2006}). The full theory of the modified-Einstein action requires the underlying microscopic theory  -- what is usually called \\emph{quantum gravity theory}. Since the latter theory is not yet established, phenomenological models are needed. Among these are the so-called $f(R)$ models (see, e.g., Refs.~\\cite{CopelandSamiTsujikawa2006,SotiriouFaraoni2008} for two recent reviews) which have \\emph{ad hoc} powers of the curvature invariants added to the linear term of the Einstein--Hilbert action~\\cite{Weinberg1972}.   We have proposed another approach to modified Einstein gravity, which is based on the treatment of the Lorentz-invariant quantum vacuum as an extended  self-sustained system characterized by a conserved relativistic charge  $q$~\\cite{KlinkhamerVolovik2008a,KlinkhamerVolovik2008b}. Here, $q$ is a microscopic variable describing the physics of the deep (ultraviolet) vacuum, but its thermodynamics and dynamics are described by macroscopic equations, because $q$ is a conserved quantity. This quantity $q$ is similar to the particle density in liquids, which describes a microscopic quantity -- the density of atoms -- but obeys the macroscopic equations of hydrodynamics, because of particle-number conservation. Different from known liquids, the quantum vacuum is Lorentz invariant. The quantity $q$ must, therefore, be Lorentz invariant. This treatment has allowed us to discuss both the \\emph{thermodynamics} and the \\emph{dynamics} of a Lorentz-invariant quantum vacuum.  The thermodynamic approach~\\cite{KlinkhamerVolovik2008a} assumes that the quantum vacuum  is a stable self-sustained equilibrium state, which is described by compressibility and other characteristics of the response to external perturbations. In this approach, we have found that the vacuum energy density appears in two forms.  First, there is the microscopic vacuum energy density which is characterized by an ultraviolet energy scale $E_\\text{UV}$, so that $\\epsilon(q)\\sim E_\\text{UV}^4$. Most likely, $E_\\text{UV}$ corresponds to the standard Planck energy scale defined in terms of Newton's gravitational constant $G_\\text{N}$, $E_\\text{Planck} \\equiv \\sqrt{\\hbar\\, c^5/G_\\text{N}} \\approx 1.2 \\times 10^{28}\\:\\text{eV}$. But, here, we prefer to stay as general as possible and keep $E_\\text{UV}$ distinct from $E_\\text{Planck}$.  Second, there is the macroscopic vacuum energy density which is determined by a particular thermodynamic quantity, \\mbox{$\\widetilde{\\epsilon}_\\text{vac}(q) \\equiv \\epsilon - q\\,d\\epsilon/dq$,} and it is this type of energy density which contributes to the effective gravitational field equations at low energies. For a self-sustained vacuum in full thermodynamic equilibrium and in the absence of matter, the effective (coarse-grained) vacuum energy density $\\widetilde{\\epsilon}_\\text{vac}(q)$ is automatically nullified (without fine tuning) by the spontaneous adjustment of the vacuum variable $q$ to its equilibrium value $q_0$, so that $\\widetilde{\\epsilon}_\\text{vac}(q_0)=0$. This implies that the effective cosmological constant $\\Lambda$ of a perfect quantum vacuum is strictly zero, which is consistent with the requirement of Lorentz invariance for zero external pressure.  The dynamic approach~\\cite{KlinkhamerVolovik2008b} demonstrates how, in a flat Friedmann--Robertson--Walker universe, the vacuum energy density $\\widetilde{\\epsilon}_\\text{vac}$ (effective cosmological ``constant'') relaxes from its natural Planck-scale value at early Planckian times to a naturally small value at late times.  In this Letter, we show that the macroscopic theory of the quantum vacuum, when applied to cosmology, gives rise to a specific class of $f(R)$ models of modified gravity. ",
        "conclusions": "\\label{sec:Conclusion}  In $q$--theory~\\cite{KlinkhamerVolovik2008a,KlinkhamerVolovik2008b} with chemical potentials $\\mu^{(a)}$ at their equilibrium values  $\\mu^{(a)}_0$, the vacuum energy density $\\rho_\\text{V}$ has been found to relax, according to \\eqref{eq:ocillations}, from its  natural Planck-scale value at early times when the system is far from equilibrium to a naturally small value at late times when the system is close to equilibrium. This solves the main cosmological problem: the present cosmological constant is small compared to Planck-scale values simply because the universe happens to be old compared to Planck-scale times. The remaining problem is to understand why the cosmological constant does not completely relax to zero as $t\\rightarrow \\infty$ or, in other words, to determine the origin of the small residual part of the vacuum energy which remains (almost) constant during the present epoch.  The suggested $f(R)$ model in which the cosmological constant appears just at the latest stage \\cite{Starobinsky2007} does not follow from the macroscopic quantum-vacuum approach. That is why one needs to find another explanation for the observed value of $\\Lambda$. In order to produce a nonzero $\\Lambda$, one must find, according to \\eqref{eq:lambda}--\\eqref{eq:LambdaGeneral}, processes which shift the chemical potentials $\\mu^{(a)}$ away from the equilibrium mean-field values $\\mu^{(a)}_0$.  There has been the suggestion to relate the present value of $\\Lambda$ to quantum fluctuations of the vacuum energy, for the case that the vacuum energy itself is nullified (see, e.g.,  Ref.~\\cite{Padmanabhan2008}). This approach could be helpful in $q$--theory, where the vacuum energy is necessarily zero in full equilibrium. The observed nonzero value of $\\Lambda$ would then correspond to quantum or thermodynamic fluctuations of the chemical potentials $\\mu^{(a)}$ compared to the equilibrium mean-field values $\\mu^{(a)}_0$. Still, there could be other processes which are able to shift $\\mu^{(a)}$ somewhat away from the mean-field values $\\mu^{(a)}_0$, one example being the backreaction of matter and another example being matter production by the rapid oscillations which accompany the relaxation of the vacuum energy \\eqref{eq:ocillations}. It remains to be seen which of these processes, if any,              % is the dominant one for the inferred small but nonzero value of the vacuum energy density from observational cosmology.    \\vspace*{0mm}"
    },
    "0807/0807.1541_arXiv.txt": {
        "abstract": "{We present the detection of a low surface-brightness stellar counterpart to an enormous (190 kpc) ring of neutral hydrogen (\\HI) gas that surrounds the nearby radio galaxy B2~0648+27. This system is currently in an evolutionary stage between major merger and (radio-loud) early-type galaxy. In a previous paper we investigated in detail the timescales between merger, starburst and AGN activity in B2~0648+27, based on its unusual multi-wavelength properties (large-scale \\HI\\ ring, dominating post-starburst stellar population and infra-red luminosity). In  this {\\sl Research Note} we present deep optical B- and V-band imaging that provides further evidence for the merger origin of B2~0648+27. The host galaxy shows a distorted optical morphology and a broad tidal arm is clearly present. A low surface-brightness stellar tail or partial ring curls around more than half the host galaxy at a distance of up to 55 kpc from the centre of the galaxy, following the large-scale, ring-like \\HI\\ structure that we detected previously around this system. The gas and stars in this ring are most likely tidally expelled material that slowly fell back onto the host galaxy after the merger event. There also appear to be sites of star formation in the \\HI\\ ring that may have formed within the gaseous tidal debris after the merger. We argue that the observed properties of the gas and stars in B2~0648+27, as well as the apparent time-delay between the merger and the starburst event, may be the logical result of a merger between two gas-rich disk galaxies with a prominent bulge, or of a merger between an elliptical and a gas-rich spiral galaxy. There also appears to be a significant time-delay between the merger/starburst event and the current episode of radio-AGN activity.} ",
        "introduction": "\\label{sec:intro}  \\begin{figure*}[htb] \\centering \\includegraphics[width=0.8\\textwidth]{0046fig1_lowres.eps} \\caption{Deep optical B+V band image of radio galaxy B2~0648+27. {\\sl Left:} high-contrast plot of B2~0648+27, highlighting the low surface-brightness stellar tail or partial ring. {\\sl Right:} low-contrast plot (logarithmic intensity scale) of the inner region of B2~0648+27.} \\label{fig:optical} \\end{figure*}  Gas-rich galaxy mergers and interactions are often invoked to be the trigger of starburst and Active Galactic Nuclear (AGN) activity. Gravitational and hydrodynamical perturbations that are exerted on the inter-stellar medium in merging galaxies induce bursts of star formation and may also remove enough angular momentum from the gas for it to be deposited onto a super-massive black hole in the centre of the merging system \\citep{bar91,bar96,mih94,mih96,spr05,kap05,dim07}. After a luminous infra-red phase, in which starburst activity is heavily obscured \\citep[e.g.][]{san96}, feedback from stellar winds and AGN activity may clear the merger product \\citep[e.g.][]{hec90,rup02,dim05,spr05,hop05}. What remains is often a newly formed early-type galaxy. The presence of an active nucleus in such a system may depend critically on the triggering time-scale and life-cycle of the AGN.  In a recent paper \\citep[][$\\ $hereafter \\paperI]{emo06} we investigated in detail the timescales of merger, starburst and AGN activity in the nearby radio galaxy B2~0648+27. Currently, the AGN in B2~0648+27 consists of an emission-line AGN \\citep{emo06thesis} plus compact radio source \\citep{gir05a}. Although the AGN is not as powerful as often observed in high-redshift quasars and QSOs, the proximity of B2~0648+27 ($z = 0.0412$) allows a detailed study of the host galaxy. In \\paperI\\ we presented the detailed analysis of a large-scale ring of neutral hydrogen (\\HI) gas ($M_{\\rm HI} \\approx 8.5 \\times 10^{9} M_{\\odot}$; diameter $\\approx$ 190 kpc) that is present around the host galaxy and which was initially discovered by \\citet[][]{mor03}. We argued that the \\HI\\ ring formed from a major merger event $\\gtrsim 1.5$ Gyr ago, when gas that was tidally expelled during the merger fell back onto the host galaxy and settled in regular rotation. This merger scenario was supported by the detection of a 0.3-0.4 Gyr young post-starburst stellar population that dominates the starlight throughout the host galaxy and which could have given B2~0648+27 the appearance of an (Ultra-) Luminous Infra-Red Galaxy at the time of the starburst event. The off-set in time between the initial merger event and the starburst activity indicates that the starburst was triggered in an advanced stage of the merger. The spectral age of the radio source in B2~0648+27 has been estimated by \\citet{gir05a} to be about 1 Myr, which suggests that there is also a significant time-delay between the starburst event and the triggering of the current episode of radio-AGN activity.  In order to unambiguously verify the merger scenario that we proposed for B2~0648+27 in \\paperI\\ and to investigate whether stars are present in the large-scale \\HI\\ ring that surrounds the host galaxy, we obtained deep optical B- and V-band imaging of this system. Previous optical imaging of B2~0648+27 by \\citet{hei94} already showed that -- while the host galaxy appears to contain smooth but very elongated elliptical isophotes -- low surface-brightness structures (a faint tail/plume and shells/arcs) appear to be present, although the limited sensitivity of their data only allowed an ambiguous classification of the optical morphology of the host galaxy. A detailed analysis of the radial light-profile of B2~0648+27 from V-band surface photometry by \\citet{gon00} revealed that the host galaxy can be fitted by an r$^{1/4}$-law, although isophote twisting is present. This indicates that B2~0648+27 already has a primary optical characteristic of an early-type galaxy, possibly in the process of formation. The new, much deeper optical imaging presented in this {\\sl Research Note} clearly shows some remarkable stellar features, which provide additional evidence that a major merger indeed formed this radio galaxy and which allow us to verify our conclusions from \\paperI.\\\\ \\vspace{0mm}\\\\ Throughout this paper we will use the name B2~0648+27 for both the radio source and the host galaxy. We assume $H_{\\circ} = 71\\ {\\rm km~s}^{-1}\\ {\\rm Mpc}^{-1}$, which puts B2~0648+27 at a distance of 174 Mpc and 1 arcsec = 0.84 kpc. ",
        "conclusions": "\\label{sec:conclusions}  We presented new, deep optical B+V band imaging of the nearby radio galaxy B2~0648+27. The distorted optical morphology of the host galaxy and the presence of a broad stellar tidal arm provide further evidence for the scenario proposed in \\paperI\\ that a major merger event formed this system. A low surface-brightness stellar tail or partial ring follows the large-scale \\HI\\ ring that surrounds the host galaxy at a radius of up to 55 kpc from the centre. The gas and stars in this ring are most likely tidally expelled material that slowly fell back onto the host galaxy after the merger event, although there also appear to be sites of star formation that may have formed within the gaseous tidal debris after the merger. The observed properties of the gas and stars, as well as the time-delay between the merger and starburst event in B2~0648+27, are in agreement with simulations of a merger between two disk galaxies with a prominent bulge, or between an elliptical and a gas-rich spiral galaxy. Apparently, the triggering of the current radio source in this system was delayed with respect to the merger and starburst event."
    },
    "0807/0807.4687_arXiv.txt": {
        "abstract": "We further the recent discussion on the relation between the star-formation rate (SFR) of a galaxy and the luminosity of its brightest star-cluster (SFR vs. M$_{V}^{\\rm brightest}$).  We first show that the observed trend between SFR vs. M$_{V}^{\\rm brightest}$ is due to  the brightest cluster in a galaxy being preferentially young ($\\leq 15$~Myr - for a constant SFR) and hence a good tracer of the current SFR, although we give notable exceptions to this rule.  Archival HST imaging of high-SFR galaxies, as well as additional galaxies/clusters from the literature, are used to further confirm the observed trend.   Using a series of Monte Carlo simulations we show that a pure power-law mass function with index, $\\alpha=2$, is ruled out by the current data.  Instead we find that a Schechter function (i.e. a power-law with an exponential truncation at the high mass end) provides an excellent fit to the data.  Additionally, these simulations show that bound cluster formation (in \\msun/yr) represents only $\\sim8\\pm3$\\% of the total star-formation within a galaxy, independent of the star-formation rate.  From this we conclude that there is only a single mode of cluster formation which operates over at least six orders of magnitude in the SFR.  We provide a simple model of star/cluster formation feedback within dwarf galaxies (and star-forming complexes within spirals) which highlights the strong impact that a massive cluster can have on its surroundings.  Using this relation, we can extrapolate backwards in time in order to estimate the peak SFR of major merger galaxies, such as NGC~7252, NGC~1316, and NGC~3610.  The derived SFRs for these galaxies are between a few hundred and a few thousand solar masses per year.  The inferred far infrared luminosity of the galaxies, from the extrapolated SFR, places them well within the range of Ultra-luminous galaxies (ULIRGs) and for NGC~7252 within the Hyper-luminous infrared galaxy regime.  Thus, we provide evidence that these post merger galaxies passed through a ULIRG/HLIRG phase and are now evolving passively.  Using the current and extrapolated past SFR of NGC~34, we infer that the ULIRG phase of this galaxy has lasted for at least 150~Myr. ",
        "introduction": "Young massive star clusters, which often surpass the globular clusters in the Galaxy in terms of brightness, mass, and density, are seen to result from intense episodes of star-formation in galaxies. However, star clusters are also found in relatively quiescent, low star-formation rate (SFR) galaxies, albeit at much lower masses (e.g. Larsen \\& Richtler~1999, 2000).  This difference in the types (mass) of clusters produced in various galactic environments has been suggested to be caused by size-of-sample effects, in which galaxies with high SFRs form proportionally more clusters, and hence are able to sample the cluster mass function out to higher masses (Larsen~2002).  This effect has been quantitatively observed through the use of the \\relation~relation (Larsen~2002), where M$_{V}^{\\rm brightest}$ is the brightest cluster in V-band absolute magnitude, in the sense that the most luminous clusters in galaxies with high SFRs are brighter.  This trend, along with the similar log N vs. $M_{V}^{\\rm brightest}$ relation (where N is the number of clusters brighter than a certain magnitude limit; Whitmore~2003), have been used to argue for a universality of cluster formation, i.e. stochastic sampling from a universal underlying mass function.  Size-of-sample effects, together with cluster population synthesis models (e.g.~Gieles et al.~2005) have become a common means to investigate the properties of clusters and cluster systems.  For example, Hunter et al.~(2003) used the relation of the most massive cluster per logarithmic age bin in the LMC and SMC in order to estimate the exponent of the cluster initial mass function ($\\alpha$).  This proceedure was recently revisited by Gieles \\& Bastian~(2008) who used the same relation to rule out mass independent, long duration ($>10$~Myr) cluster disruption models.  Gieles et al.~(2006a) used the log N vs. $M_{V}^{\\rm brightest}$ relation to constrain $\\alpha$, and found a value of $\\sim2.4$, which is similar ($2.3$) to that derived by Weidner, Kroupa, \\& Larsen~(2004) using the \\relation~ relation.  This is significantly steeper than that derived from direct measurements of the mass/luminosity function of galaxies, namely $\\alpha=2.0\\pm0.1$ (e.g. de Grijs et al.~2003).  This discrepancy will be addressed in \\S~\\ref{sec:discussion}.  Wilson et al.~(2006) have tested whether the above relations still hold in the extreme environment of galaxy merger starbursts.  They studied the Ultra-Luminous Infrared Galaxy (ULIRG), Arp~220, and found that despite its high SFR ($\\sim240$~\\msun/yr), being an order of magnitude higher than any of the galaxies in the previous samples, falls nicely on the extrapolated fit to the more quiescent star-forming galaxies.  Weidner et al. (2004) used the \\relation~relation to constrain cluster formation scenarios, namely the timescale over which clusters form, which they estimate to be on the order of a crossing time.  They further suggest that a cluster population formation epoch (i.e. the timescale where a statistically full population of clusters is formed) is on the order of 10~Myr.  However, their analysis was based on the assumption that within a \"cluster population formation epoch\" the brightest cluster of a galaxy is also the most massive, hence that the \\relation~ trend is simply reflecting a relation between the SFR of a galaxy and the most massive cluster within it.  Observationally, it appears that this assumption is not valid, as the brightest cluster in a galaxy tends to be young, and more massive, older clusters may appear less luminous due to stellar evolution (Gieles et al.~2006a).  In this work our goals are threefold.  The first is to test the claim by Weidner et al.~(2004) that the brightest cluster within a galaxy is also the most massive.  This naturally leads to a discussion as to why the observed \\relation~relation holds.  The second is to investigate the implications of the observed relation, paying particular attention to the cluster initial mass function, and the implied connection between the cluster and star formation rates within a galaxy.  Thirdly, using the observed trend, combined with a correction for stellar evolutionary fading, to estimate the SFR in a sample of post-starburst merger galaxies.  This, in turn, allows us to place limits on the duration of the starburst phase of ULIRGS as well as trace their subsequent evolution.  In \\S~\\ref{sec:data} we present archival observations of two ongoing galaxy mergers and a collection of data taken from the recent literature.  \\S~\\ref{sec:why} presents a series of Monte Carlo simulations of cluster populations in order to investigate why the observed \\relation~ relation holds.  In \\S~\\ref{sec:discussion} we investigate the implications for the underlying cluster initial mass function,  the relation between star and cluster formation, and use the observed relation to derive the peak SFR of post-starburst galaxies.  Our conclusions are presented in \\S~\\ref{sec:conclusions}. ",
        "conclusions": "\\label{sec:conclusions}  Using the empirically derived relation between the star-formation rate (SFR) and the luminosity of the brightest star cluster within the galaxy (Larsen 2002) we have investigated the conditions (SFR and infrared luminosity) that were present during the peak star-forming epoch of galactic mergers.    We began with a series of simple Monte Carlo simulations of cluster populations with different assumptions on the underlying cluster mass function, and showed that, independent of the model construction, the youngest clusters ($<10$~Myr) were predominantly the brightest.  This, along with size-of-sample effects, can explain the observed \\relation~ relation, as the young clusters should be a good representation of the current SFR.  Using data from the literature as well as archival HST imaging, we have further tested the high SFR regime of the relation.  We found, in agreement with the results of Wilson et al.~(2006), that the relation continues to hold up to SFRs of a few hundred solar masses per year.  We have also presented a schematic model (following on Weidner et al.~2004) to explain why some dwarf (post) starburst galaxies lie on the left of the observed relation.  The basic model is that the formation of a massive cluster is enough to disturb a substantial part of the galaxy, and cause an ISM blowout (as seen in NGC~1569) which effectively terminates further star-formation.  This same process is expected to take place within star-forming complexes in spiral galaxies.  By generating a large sample of Monte Carlo simulations of cluster populations with varying underlying initial mass functions and cluster-formation rates, we have shown that the observed \\relation~ relation is best fit by a Schechter mass function with \\mstar$=1-5\\times10^{6}~\\msun$.  This rather high value of \\mstar~ makes it extremely difficult to detect a difference between a pure power-law or a Schechter type mass function directly (i.e. through binned cluster mass distributions).  However, size-of-sample effects which rely heavily on the most massive or brightest object in a sample are much more sensitive to small differences between these two functions at the high mass end.  We find that the data are inconsistent with the pure power-law, $\\alpha=2$ case at the $>8\\sigma$ level.  Using this set of simulations, we also showed that a constant fraction of star-formation goes into clusters.  This fraction, which we term $\\Gamma$, is $0.08\\pm0.03$, meaning that optically selected clusters represent only 8\\% of the total star-formation in a galaxy.  This is similar to that found for the solar neighborhood by Lada \\& Lada~(2003) and Lamers \\& Gieles~(2007),  and it is also in fair agreement with that derived for the SMC, namely $2-4$\\% (Gieles \\& Bastian~2008).  $\\Gamma$ appears to be independent of the star-formation rate over 6 orders of magnitude.  From this we conclude that cluster formation does not have 'multiple modes' and that the cluster formation history of a galaxy should accurately reflect its star-formation history (once a proper accounting of the cluster formation rate and variance in the stellar fields is taken into account).  Using the fit to the \\relation~ relation by Weidner et al.~(2004), and extrapolating to higher SFRs, we have estimated the peak SFR and infrared luminosity of seven (post-starburst) galactic mergers by estimating the absolute magnitude of the brightest star cluster when corrected for stellar evolutionary fading.  We estimate SFRs between 50 and $>4000$~\\msunyr~ and infrared luminosities between $10^{11.5} - 10^{13.4}$~\\lsun.  These are lower limits as we have not corrected for cluster mass-loss (fading) due to tidal fields and internal dynamics. Additionally, the infrared luminosities are also lower limits as  they do not take into consideration of any additional contribution to the \\lir~ by an AGN component.  These results further demonstrate that many mergers pass through a LIRG/ULIRG/HLIRG phase during their evolution.  NGC~34, which currently has a high SFR ($>75$~\\msunyr), also hosts a population of bright massive clusters with ages of $\\sim150-600$~Myr (Schweizer \\& Seitzer~2007).  Using these clusters we have estimated a peak SFR of $\\sim600$~\\msunyr~ which corresponds to an infrared luminosity of $10^{12.5}$~\\lsun, well within the ULIRG regime.  Thus we have placed a lower-limit to the duration of the LIRG/ULIRG phase of 150~Myr for this galaxy."
    },
    "0807/0807.3402_arXiv.txt": {
        "abstract": "The formation of a cool disk in the innermost regions of black hole X-ray transient systems in the low hard state is investigated.  Taking into account the combined cooling associated with the Compton and conductive energy transport processes in a corona, the radial structure of a disk is described for a range of mass accretion rates.  The mass flow in an optically thick inner region can be  maintained by the condensation of matter from a corona with the disk temperature and luminosity varying continuously as a function of the accretion rate.  Although such a disk component can be present, the contribution of the optically thick disk component to the total luminosity can be small since the mass flow due to condensation in the optically thick disk underlying the corona can be significantly less than the mass flow rate in the corona. The model is applied to the observations of the low quiescent state of the black hole source GX 339-4 at luminosities of $\\sim 0.01 L_{Edd}$ and is able to explain the temperature of the thermal component at the observed luminosities.  Since conductive cooling dominates Compton cooling at low mass accretion rates, the luminosity corresponding to the critical mass accretion rate above which a weak thermal disk component can be present in the low hard state is estimated to be as low as $0.001 L_{Edd}$. ",
        "introduction": "The class of objects known as transient X-ray binary systems containing black holes is one of the few examples of astrophysical laboratories available for investigating the nature of accretion flows and the structure of accretion disks surrounding compact objects.  Their luminosities provide us with an opportunity to investigate a given system over a wide range of mass accretion rates onto the central black hole. Indeed, these systems have been observed in distinct states generally categorized as low/hard and high/soft states (see McClintock \\& Remillard 2006).  In this categorization, the high/soft state is characterized by a strong thermal component corresponding to a color temperature of order 1 keV.  These states are observed during the outburst phase in which the luminosity can approach the Eddington luminosity, $L_{\\rm Edd} \\sim 10^{39}$ ergs s$^{-1}$ for a $10 \\msun$ black hole. On the other hand, the low/hard state is typically described by a power law with a photon power law index $\\sim 1.7$.  This state is observed during the quiescent stages, where the luminosities can be as low as $10^{-7} L_{\\rm Edd}$ (Corbel et al. 2006). Hence, it is generally accepted that these states are distinguished by the spectrum and a range of luminosity with a transition at $\\sim1-4\\%$ of the Eddington luminosity (Maccarone 2003) from the low/hard state to the high/soft state.  Given these observational diagnostic clues, progress in our theoretical understanding has advanced in parallel. Specifically, the spectrally soft state is believed to represent the emission from a cool optically thick and geometrically thin accretion disk extending to the innermost stable circular orbit ($3 R_S$ for a non-rotating black hole, where $R_S = 2GM/c^2$ is the Schwarzshild radius) as described in the seminal study by Shakura \\& Sunyaev (1973, hereafter SS).  In contrast, the spectrally hard state is thought to arise from a geometrically thick, hot optically thin coronal or advectively dominated accretion flow (ADAF; Narayan \\& Yi 1994; 1995a, b) located in the innermost regions of the disk near the black hole.  Here, the SS accretion disk is truncated at a distance of several hundred Schwarzschild radii from the black hole, leading to a transition from a disk dominated by a geometrically thin outer structure to a geometrically thick inner structure Esin et al. (1997, 1998, 2001).  Although a complete description of the truncation of the optically thick disk has yet to be worked out in detail, the mechanisms suggested to facilitate this transition have been investigated in the studies by Honma (1996) and Manmoto \\& Kato (2000) based on a radial conductive energy transport process and by Meyer et al. (2000), Spruit \\& Deufel (2002), and Dullemond \\& Spruit (2005) based on a vertical evaporative process.  Recently, additional observational clues concerning the nature of the accretion flow and geometry in the low quiescent state have been found pointing to the possible existence of an inner cool optically thick disk structure in black hole transient X-ray binary systems.  Although the observational signatures are difficult to extract because of the low X-ray count rates, the evidence for such a structure has been inferred from the source GX 339-4 (see Miller et al. 2006a, b; Tomsick et al. 2008) where a weak, soft, thermal emission component ($\\lesssim 0.3$ keV) has been detected and blurred reflection and iron features are indicated in fits to the X-ray spectrum for luminosities as low as $\\sim 0.01 L_{\\rm Edd}$. One of the goals of current observational studies has been to extend such studies to other sources (e.g., J 1817-330, see Rykoff et al. 2007; and J1753.5-0127, see Miller et al. 2006a; Ramadevi \\& Seetha 2007; Soleri et al. 2008) and to determine the luminosity range for which such a thermal component exists so that a full description of the accretion flow geometry in the various states of transient X-ray binary systems can be constructed.  Given the existence of this weak thermal component, the accretion disk structure in the quiescent state would be described by three distinct regions.  Specifically, a small inner cool, optically thick disk, a geometrically thick hot coronal region at intermediate radial scales, and a very extensive geometrically thin disk nearly filling the Roche lobe of the black hole would contribute to the weak thermal X-ray component, the power law X-ray component, and the infrared and optical emission respectively. Such a configuration was already contemplated in early work by R\\`o\\.za\\`nska \\& Czerny (2000) and more recently by Liu et al. (2006), Meyer et al. (2007) and Mayer \\& Pringle (2007). To maintain such a remnant cool disk in the innermost regions (in the presence of mass transport by turbulent viscous stresses) gas must condense from an overlying corona.  In this study we focus on the general description of the structure and properties of the disk in its innermost regions during the quiescent state of a black hole X-ray transient binary system, taking into account the effects of both thermal bremsstrahlung and comptonization.  This study differs from our recent study (Liu et al. 2007) where only one such process was considered dominant for any given structure. In \\S 2, we describe the theoretical model and the modifications required in considering more than one cooling process in any given region of the disk.  The solutions to the modified set of equations are presented for a range of mass accretion rates and viscosity parameters in \\S 3. The numerical results are compared to the observations of GX 339-4 in \\S 4 and their implications are discussed in the last section. ",
        "conclusions": "The structure of the inner disk region of accretion disks have been investigated for black hole X-ray transient systems in their low luminosity quiescent state within the context of an $\\alpha$ viscosity model.  Particular attention has been focused on the conditions under which a weak, soft thermal component would be present in this state. By generalizing the study of Liu et al. (2007) to include the cooling associated with both conduction and Compton processes together, we have examined the properties of an optically thick inner disk fed by the condensation of gas from a corona for a range of mass accretion rates.  It has been shown that such a disk component can be present with its spatial extent, temperature, and luminosity varying continuously and decreasing with a decreasing mass accretion rate. Despite its existence and the higher conversion efficiency of gravitational energy into radiation in the cool disk, the contribution of the disk component to the total luminosity can be very low due to the low condensation rate.  Throughout the quiescent state, the mass flow rate to the central black hole from the corona always exceeds the mass flow rate in the inner cool disk. Hence, the disk structure remains geometrically thick even in the innermost regions although a soft thermal component is present in the spectrum.  In our analysis, we have neglected the effect of X-ray irradiation to facilitate comparisons with our earlier work.  Although the qualitative features of our model will not change, the parameters describing the quantitative comparisons will be modified.  In particular, the inclusion of X-ray irradiation can affect the fitting parameters $\\alpha$ and $\\dot m$ since the observed disk temperature will not solely be determined by the mass flow in the inner disk. For example, for ADAF emission from the innermost region at height $H=R_{\\rm ISCO}=3R_s$, the flux irradiating the disk is given by \\begin{equation} F_{\\rm ir}(r)={{1\\over 2}L_c\\over 4\\pi R^2}{H\\over R}=1.716\\times 10^{26} {\\rm erg\\,s^{-1}\\,cm^{-2}} m^{-1}r^{-3}{L_c\\over L_{\\rm Edd}}. \\end{equation} Adopting a mean photon energy of 50keV, about one tenth of the irradiation flux (the ratio of the mean photon energy to electron rest mass energy) is absorbed by the disk through Compton scattering collisions with electrons in the disk.  Although the K-shell absorption of the metals is negligible at 50 keV, it can become significant at lower energies since the absorption cross section is proportional to $(m_ec^2/h\\nu)^{7/2}$. For a photon distribution of $n_\\nu\\propto \\nu^{-1.6}$, the flux absorbed by K-shell absorption is found to be $\\sim 30\\%$ of the total irradiation flux. Thus, about 40\\% of the total irradiation flux is absorbed by the K-shell absorption and Compton scattering. The re-radiation of the absorbed flux at the inner disk surface yields an effective temperature of \\begin{equation} T_{\\rm ir}(r)=\\left({0.4F_{\\rm ir}\\over \\sigma}\\right)^{1/4}=2.86keVm^{-1/4}r^{-3/4} \\left({L_c\\over L_{\\rm Edd}}\\right)^{1/4}. \\end{equation} At the distance of maximal temperature, $r_{tmax}=3{49\\over 36}$, the effective temperature caused by irradiation alone for spectrum\\#1 is 0.25keV and for Spectrum\\#2 is 0.19keV (see \\S4). These temperatures are slightly higher than the temperatures generated by disk accretion and could contribute to the observationally inferred temperature.  We note that this theoretical estimate is very approximate since an accurate calculation depends on the region of origin of the ADAF flux and the size of the coronal gap region. From Eqs.\\ref{L_cmp1} and \\ref{L_cmp2} it can be seen that the corona/ADAF radiation is dominant from a region outside the ISCO. In this case, the irradiation flux incident on the innermost disk decreases with distance of the source, $F_{\\rm ir} \\propto r^{-2}$. For instance, assuming the ADAF radiation is dominant from $5R_s$ rather than $3R_s$, the irradiation flux absorbed by the disk around the ISCO is 7\\% of the total flux. On the other hand, many coronal photons escape without interception by the disk for a large coronal gap, thereby, reducing the thermal effects associated with the irradiation. In the case for Spectrum \\#1 and \\#2 of GX339-4, the reflection covering factor inferred from observations is $\\sim 20\\%$ of the half sky, indicating that, indeed, a very large fraction of the coronal photons escape, in contrast to the case of Haardt \\& Maraschi (1991).  Hence, these estimates indicate that our results are indicative and that additional freedom exists in determining $\\dot m$ and $\\alpha$ when irradiation is taken into account.  For sufficiently low mass accretion rates, the cool inner disk vanishes.  An estimate of the mass accretion rate for which this occurs depends on whether condensation occurs in the vicinity of the ISCO. If condensation does not occur, accretion and/or evaporation can deplete the inner disk within a viscous time.  In the case of cooling dominated by conduction,  the critical condensation radius is given by $r_d=0.815\\alpha^{-28/3}\\dot m^{8/3}\\ge 3$.  Thus, the lower limit to the accretion rate is \\begin{equation}\\label{llimit1} \\dot m = 0.006 \\({\\alpha\\over 0.2}\\)^{7/2}. \\end {equation} For an accretion rate just above this lower limit, both the size of the inner disk and the condensation rate are very small, leading to a very low disk temperature. In this case, the Compton cooling is negligible, and the lower limit predicted by the conduction model is self-consistent.  For disk temperatures exceeding 0.14 keV, on the other hand, Compton cooling dominates in a corona around a $10 M_\\odot$ black hole.  If this temperature is produced by the accretion of condensed gas, the condensation rate would be greater than 0.002 $\\dot M_{\\rm Edd}$. However, this occurs at high accretion rates and is greater than the lower limit given by conduction model for $0.1 \\le \\alpha\\le 0.4$ (King et al. 2007).  Thus, the lower limit to the mass accretion rate for maintenance of an inner disk is set by the conduction model, i.e.  equation (\\ref{llimit1}). Notwithstanding the possible importance of irradiation of the inner disk by the corona in providing some latitude in disk parameters, $\\alpha$ and $\\dot m$ (see above), when fitting the observational data, a determination of the mass flow rate below which the cool disk component vanishes could provide a constraint, within the framework of the condensation model, on the $\\alpha$ viscosity parameter during the quiescent state.  In the future, global calculations of the full radial and vertical disk structure should be considered taking into account the effect of X-ray irradiation of the hot coronal gas on the inner disk and of the cooling associated with both the Compton and conductive processes in order to quantitatively confirm the picture presented in this paper.  With such models in hand, a detailed comparison of the spectra from such disks with observations can be undertaken."
    },
    "0807/0807.1625_arXiv.txt": {
        "abstract": "{The quiet Sun presents magnetized plasma whose field strengths vary from zero to about $2$~kG. The probability density function of the magnetic field strength $B$ efficaciously describes the statistical properties of the quiet Sun magnetic field.} {We simulate the dynamics and the evolution of quiet Sun magnetic elements to produce a probability density function of the field strengths associated with such elements.} {The dynamics of the magnetic field are simulated by means of a numerical model in which magnetic elements are driven passively by an advection field characterized by spatio-temporal correlations that mimick the granulation and mesogranulation scales observed on the solar surface. The field strength can increase due to an amplification process that occurs where the magnetic elements converge. Starting from a $\\delta$-like probability density function centered on $B=30$~G, we obtain magnetic field strengths of up to $2$~kG (in absolute value). To derive the statistical properties of the magnetic elements, several simulation runs are performed.} {Our model is able to produce kG magnetic fields in a time interval of the order of the granulation timescale. The mean unsigned flux density and the mean magnetic energy density of the synthetic quiet Sun reach values of $\\langle B \\rangle \\simeq 100$~G and $\\langle B^2 \\rangle^{1/2}\\simeq 350$~G respectively in the stationary regime. The derived probability density function of the magnetic field strength decreases rapidly from $B=30$~G to $B \\sim 100$~G and has a secondary maximum at $B=2$~kG. From this result, it follows that magnetic fields $\\ge700$~G dominate the unsigned flux density and magnetic energy density, although the probability density function of the field strength reaches a maximum at $B \\sim 10$~G.} {A photospheric advection field with spatio-temporal correlations, driving the magnetic elements, and reduced magnetic amplification rules are able to create a realistic probability density function of the quiet Sun magnetic field. It has been found that they naturally produce an excess of magnetic fields around $2$~kG if an upper limit is imposed on the field strength.}  \\titlerunning{Quiet Sun PDF: a simplified model approach} \\authorrunning{F.~Berrilli, D.~Del Moro \\and B.~Viticchi\\'{e}} ",
        "introduction": "\\label{intro} When observing the Sun in full disk magnetograms, magnetic concentrations define the network pattern associated with the supergranular motions of photospheric plasma, while the interior of the network pattern appears to be unmagnetized.\\\\ We know that weak polarization signals are measurable almost everywhere in the interior of network cells  \\citep{livhar75,smi75}. These signals are associated with the internetwork component of the solar magnetic field, also called the quiet Sun magnetic field, which covers more than 90\\% of the solar surface, independent of the solar activity \\citep{harang93}. Due to its large diffusion onto the solar surface, the quiet Sun magnetic field may be useful for understanding the physics of the photospheric magnetism. In fact, it could hold a large fraction of the unsigned flux density and magnetic energy density of the solar photosphere \\citep[e.g.][]{unn59,ste82,san98,san04,schtit03}.\\\\ The weak polarization signals measured in quiet Sun regions can be interpreted as the result of the linear combination of the polarization emerging from discrete magnetic flux tubes, which are typically smaller than the angular resolution, that fill the solar atmosphere with a complex topology \\citep[e.g.][]{emocat01a,san03}.\\\\ Knowledge of the quiet Sun has been enhanced significantly by improvements in instrumentation \\citep[e.g.][]{san96,linrim99,lit02,dom03a,dom03b,kho03}, in diagnostic techniques \\citep[e.g.][]{ste82,fausch93,fausch95,lan98,ree00,sanlit00,soc01,socsan02,san04,tru04,man04,san05,mar06b}, and numerical simulations \\citep[e.g.][]{cat99,emocat01b,stenor02,vog03th,vog05, stenor06,vogsch07}.\\\\ These improvements were crucial in characterizing the quiet Sun magnetic field and its different observational parameters. One parameter, the fraction of solar surface occupied by a certain field strength, is a very powerful tool. In fact, this can be trivially translated into the probability density function (PDF) of the magnetic field strength, $P(B)$, which is useful in describing the statistical properties of a quiet Sun region. For example, calculating the first and the second momentum of $P(B)$ provides information about the mean unsigned flux density and the mean magnetic energy density (apart from a constant factor $1/8\\pi$), respectively. Using a PDF assumes implicitly a continuum for values of the magnetic field strength. The limit imposed by the pressure equilibrium between the magnetic flux tubes and photospheric plasma also defines the variability range for quiet Sun magnetic fields to be between $0$~G to about $2$~kG.\\\\ \\citet{lin95} computed the PDF in active and quiet Sun regions using \\fei{15648} and \\fei{15652} lines, while \\citet{lit02} obtained it from the analysis of network regions and their surroundings using \\feis{} lines at $6300$~\\AA{}. Both authors found a bimodal distribution with a dominant contribution at $\\simeq 150$~G and a secondary peak at $\\simeq 1.5$~kG. In contrast, \\citet{col01}, studying quiet Sun regions using the same infrared lines as \\citet{lin95}, found an exponential decline in the magnetic flux density histogram from a peak at $\\simeq 200$~G and no counts for magnetic flux density $\\ga 1$~kG. \\citet{dom06} developed a set of PDFs for the quiet Sun magnetic fields that was consistent with observations and magnetoconvection simulations by using information from both Hanle and Zeeman observations; their study created substantial interest because the functions increased for strong fields, just before a cut-off at $B\\simeq 1.8$~kG. This characteristic implies that strong fields are important because, even when present in a small fraction of the photosphere, they dominate the unsigned flux density and magnetic energy density. On the other hand, the magnetic field distributions proposed by \\citet{mar06a,mar06b} and \\citet{mar07}, which were derived from simultaneous infrared and visible observations of \\feis{} lines, do not present any rise for strong fields. The inversion of the \\fei{6301} and \\fei{6302} profiles observed by the \\textit{HINODE} satellite \\citep{oro07} confirms this result.\\\\ Magnetoconvection simulations provide a wide variety of quiet Sun magnetic field PDFs to be compared with those derived from observations. Simulations show, for example, exponential decays \\citep{cat99,ste03,stenor06,vogsch07}, bimodal behaviors \\citep{vog03,vogsch03}, or an ``exponential-like core and a modest peak'' for $B \\sim 1$~kG \\citep{ser03}.\\\\ An important step is to ascertain the presence or absence of a bump for magnetic fields $\\sim 1$~kG in quiet Sun magnetic field strength PDFs; since the magnetic flux and magnetic energy both scale as powers of the field strength, kG magnetic fields should play a crucial role in defining the global properties of quiet Sun magnetic fields. The connectivity between the lower and the upper solar atmosphere may indeed be defined by the larger field strength concentrations in the photosphere \\citep{dom06}.\\\\ \\citet{san07} indicated that a bump in the PDF for quiet Sun fields $\\sim$~kG is to be expected if ``a magnetic amplification mechanism operates in the quiet Sun''. The author derived an equation for the temporal evolution of the PDF, whose shape can be modified by magnetic amplification and photospheric convection. In the stationary state, the model is characterized by a $P(B)$ with a bump for strong fields.\\\\ In our model, we implement the dynamics of a high number of photospheric magnetic elements, driven by a photospheric velocity field in which a large-scale organization (mesogranulation) is naturally derived from the cooperation of small-scale flows (granulation). The model also takes into account simplified processes that are able to amplify and limit the magnetic field strength. In \\S~\\ref{model}, we describe the model, in \\S~\\ref{pdfs}, the results of the simulation are presented, while in \\S~\\ref{disc}, they are discussed. Finally, in \\S~\\ref{conc}, the conclusions are outlined. ",
        "conclusions": "\\label{conc} We have presented a model that reproduces the dynamics and evolution of photospheric magnetic elements. These elements are passively driven by an horizontal velocity field characterized by spatio-temporal correlations that agree with those observed in the solar photosphere. The description of the velocity field follows \\citet{ras98,ras03}. In their evolution, the magnetic elements interact and increase their magnetic field strengths through an amplification process assumed to work on a timescale comparable with the granular timescale and able to recreate kG concentrations. The injection scale, constrained to be $B_{in}=30$~G, and the upper limit at $B_{lim}=2$~kG are both derived from observations and physical constraints \\citep[e.g.][]{tru04,par78}.\\\\ We investigated the statistical properties of the field strengths associated with the magnetic elements using a series of numerical simulations. The main conclusions are: \\begin{enumerate} \\item The model produces kG fields in a time interval of the order of the mesogranulation timescale. These strong fields are organized in a pattern superimposed on the mesogranulation pattern of the advection field. \\item In the stationary regime, the PDF of the magnetic field strength has a stable shape and has a secondary maximum that corresponds to the upper limit set at $2$~kG. \\item The associated mean unsigned flux density and mean magnetic energy density in the stationary regime are $\\langle B \\rangle_{sat} \\simeq 100$~G and $\\langle B^2 \\rangle_{sat}^{1/2}\\simeq350$~G, respectively. \\item The magnetic element system described by this PDF is strongly dominated by kG magnetic fields: the fractions of unsigned flux density and magnetic energy density for $B\\ge700$~G are $\\simeq 65\\%$ and $\\simeq 97\\%$, respectively. \\end{enumerate} The quiet Sun is certainly more complex than the model presented here. However, a more simplified approach, based on the advection field proposed by \\citet{ras03}, reproduced the observed distributions for the released magnetic energy and waiting times of nanoflare events \\citep{vit06}. In the present work, we showed how the same approach can also be applied to reproduce some of the statistical properties of the quiet Sun magnetic field."
    },
    "0807/0807.1763_arXiv.txt": {
        "abstract": "We report high spatial resolution 11.2 and 18.1 $\\mu$m imaging of V838 Monocerotis obtained with Gemini Observatory's Michelle instrument in 2007 March.  Strong emission is observed from the unresolved stellar core of V838 Mon in our Gemini imagery, and is confirmed by Spitzer MIPS 24 $\\mu$m imaging obtained in 2007 April.  The 2007 flux density of the unresolved mid-infrared emission component is $\\sim$2 times brighter than that observed in 2004.  No clear change in the net amount of 24$\\mu$m extended emission is observed between the 2004 and 2007 epoch Spitzer imagery. We interpret these data as evidence that V838 Mon has experienced a new circumstellar dust creation event. We suggest that this newly created dust has condensed from the expanding ejecta produced from V838 Mon's 2002 outburst events, and is most likely clumpy.  We speculate that one (or more) of these clumps might have passed through the line-of-sight in late 2006, producing the brief multi-wavelength photometric event reported by \\citet{bon06} and \\citet{mu07b}.  We detect no evidence of extended emission above a level of $\\sim$1 mJy at 11.2 $\\mu$m and $\\sim$7 mJy at 18.1 $\\mu$m over radial distances of 1860 - 93000 AU (0$\\farcs$3-15$\\farcs$0) from the central source.  Using the simple assumption that ejecta material expands at a constant velocity of 300-500 km s$^{-1}$, this gap of thermal emission suggests that no significant prior circumstellar dust production events have occurred within the past $\\sim$900-1500 years. ",
        "introduction": "\\label{introduction}  V838 Monocerotis experienced three photometric outbursts in early 2002, beginning with an initial event on 2002 January 6.6 \\citep{bro02} and continuing with events in 2002 February and 2002 March \\citep{kim02,mun05}.  The star's spectrum cooled through the F, G, and K spectral types throughout February \\citep{mun07}, transitioned to a late M-type giant by 2002 April 16 \\citep{rau02}, and had cooled to a L-type supergiant appearance by 2002 October \\citep{eva03}.  The circumstellar ejecta from the outburst events was not spherically symmetric, as revealed by the intrinsic polarization detected until 2002 February 13 \\citep{wi03a,des04}; \\citet{wis07} estimated that ejecta located interior to the region producing H$\\alpha$ and Ca II emission was flattened by a \\textit{minimum} of 10\\% from a spheroidal shell.  A renewed intrinsic polarization component observed on 2002 October 22 might be attributable to a new source of asymmetrical scatterers in V838 Mon's circumstellar envelope or to a change in the opacity of the existing envelope \\citep{wi03b}.  V838 Mon produced a spectacular light echo first detected on 2002 February 17 by \\citet{hen02}, which evolved dramatically over time \\citep{bon03}. \\citet{ban06} imaged V838 Mon with the \\textit{Spitzer Space Telescope} at 24, 70, and 160$\\mu$m, and attributed the observed excess of thermal infrared (IR) emission from the unresolved central source as evidence of the formation of circumstellar dust produced by the 2002 outbursts.   An IR excess was also observed from the spatially resolved light echo, and \\citet{ban06} concluded that the source of this emitting dust could have both interstellar and circumstellar components.  V838 Mon resides in a young open cluster \\citep{afs07} at a distance of 6.2 $\\pm$1.2 kpc \\citep{spa08}.  It has a known, albeit unresolved, B3V binary companion \\citep{mun05}.  The nature of V838 Mon prior to outburst is a subject of active debate (c.f. \\citealt{mun05, tyl05}); a wide variety of mechanisms have been proposed to explain the eruptive events in V838 Mon, including a shell thermonuclear event \\citep{mun05}, a born-again star which experienced an accretion event \\citep{law05}, a stellar merger \\citep{tyl06}, and a star which engulfed several planets \\citep{ret06}.  Determining the amount and spatial distribution of dust within the V838 Mon line of sight, and quantifying how much is circumstellar versus interstellar in nature, would help constrain whether its outburst mechanism was a variant of a repeatable phenomenon, like a thermonuclear runaway or nova-like explosion, or more likely a singular event, like a stellar merger. ",
        "conclusions": "\\label{disc}  We detect no evidence of extended 18.1 $\\mu$m emission above a level of $\\sim$7 mJy from radial distances of 1860 AU (0$\\farcs$3) to 93000 AU (15$\\farcs$0), assuming a distance of 6.2 kpc \\citep{spa08}.  Making the simple assumption that V838 Mon's circumstellar ejecta expands at a constant velocity, this gap of thermal IR emission suggests no circumstellar dust producing events have occurred within the past $\\sim$900 (if v$_{ejecta}$ $\\sim$300 km s$^{-1}$; \\citealt{wi03a}) to $\\sim$1500 years (if v$_{ejecta}$ $\\sim$500 km s$^{-1}$; \\citealt{mun02}).  Our analysis of 2007 imagery of V838 Mon from both Gemini's Michelle and Spitzer's MIPS indicates the presence of significantly enhanced 10-70$\\mu$m emission in the unresolved central stellar core in 2007 as compared to 2004.  We interpret our 2007 data as evidence that a \\textit{new} circumstellar dust production event has occurred since 2004.  Spectroscopic evidence of V838 Mon's binary companion disappeared and H$\\alpha$ returned to emission in late 2006, corresponding with a temporary multi-color photometric fading \\citep{bon06,mu07b}. \\citet{mu07b} attributed these events to either a) an eclipse between the primary and secondary; or b) a dust cloud characterized by E$_{B-V}$ = 0.55 and R$_{V}$ = 3.1, gravitationally bound to the binary system, which eclipsed the B3V secondary.  We suggest alternate interpretations are possible if these photometric and spectroscopic events are correlated to the new mid-IR behavior we have found in data from early 2007.  A gravitationally bound dust cloud would radiate a constant amount of thermal IR emission throughout its orbit except during the brief time that it is occulted by the secondary; thus, \\textit{unless it was only recently formed}, such a cloud could not be singularly responsible for producing the observed enhanced mid-IR emission.  A simple eclipse between the primary and secondary also would not lead to enhanced thermal IR emission.  \\citet{bon06} suggested that the photometric and spectroscopic events of late 2006 were indicative of the 2002 ejecta interacting with and overtaking the B3V binary.  More recent observations indicate that V838 Mon is experiencing a prolonged period of reduced photometric brightness (H. Bond 2008, personal communication), further supporting this interpretation.  While the temporal resolution of our data do not allow us to quantitatively determine whether this proposed encounter is related to the new circumstellar dust creation event we report, we suggest that such a correlation is plausible.  We suggest that the expanding ejecta from the 2002 outbursts has condensed to form \\textit{new} circumstellar dust, producing the enhanced IR excess observed inside a radial distance of 1860 AU.  The dust condensation radius of $\\sim$3.5 AU for the B3V companion, which assumes T$_{cond}$ $\\sim$1000K, $R_{cond} \\sim\\frac{R_{star}}{2} (\\frac{T_{star}}{T_{cond}})^{2} $ \\citep{ire05}, and the estimated primary-secondary separation of $\\sim$34 AU derived from interferometry \\citep{lan05}, indicates that dust can survive within the radial distance constraint of $<$1860 AU implied by our observations.  This type of event is not altogether unexpected, as \\citet{lyn04} previously reported the detection of an expanding region of molecular gas which they noted was likely dust precursor material.  The enhanced flux at 18.1, 24, and 70 $\\mu$m we observed in 2007 suggests that the equally strong (2007 epoch) 11.2 $\\mu$m flux might not represent an additional enhancement of gas-phase molecular emission above that suggested by \\citet{lyn04}, but rather represents a warm thermal dust emission component.  Early post-outburst polarimetric observations \\citep{wi03a,des04} indicated this ejecta was asymmetric; hence, we suggest that it is likely that dust which condenses from this ejecta will be spatially non-uniform and/or clumpy.  We speculate that it is possible that one or more of these clumps eclipsed the line of sight in late 2006, producing the attenuation event reported by \\citet{bon06} and \\citet{mu07b}.  Unlike the dust cloud eclipse event invoked by \\citet{mu07b}, we believe that any eclipsing dust blob is more likely to located in the expanding ejecta which is condensing to form new dust.  Followup interferometric observations might reveal differences in the angular size of the system from that found in 2004 epoch observations by \\citet{lan05} if this ejecta has indeed recently reached the location of the B3V secondary star."
    },
    "0807/0807.3550_arXiv.txt": {
        "abstract": "We have detected over $400$ \\HI\\ clouds in the lower halo of the Galaxy within the  pilot region of the Galactic All-Sky Survey (GASS), a region of the fourth  quadrant that spans $18\\degr$ in longitude, $40\\degr$ in latitude and is centered on the Galactic equator. These clouds have a median peak brightness temperature of $0.6$~K, a median velocity width of $12.8$~km~s$^{-1}$, and angular sizes $\\lesssim 1\\degr$. The motion of these clouds is dominated by Galactic rotation with a random cloud-to-cloud velocity dispersion of $18$~km~s$^{-1}$. A sample of clouds likely to be near tangent points was analyzed in detail. These clouds have radii on the order of $30$~pc and a median \\HI\\ mass of $630 M_{\\odot}$. The population has a vertical scale height of $400$~pc and is concentrated in Galactocentric radius, peaking at $R=3.8$~kpc. This confined structure suggests that the clouds are linked to spiral features, while morphological evidence that many clouds are aligned with loops and filaments is suggestive of a relationship with star formation. The clouds might result from supernovae and stellar winds in the form of fragmenting shells and gas that has been pushed into the halo rather than from a galactic fountain. ",
        "introduction": "Neutral atomic hydrogen (\\HI) is ubiquitous throughout the Galaxy with a wide variety of morphologies and kinematics, and exhibits many complex structures, including worms \\citep{1992Koo}, sheets and filaments \\citep{1967Heiles,1990Dickey}, shells \\citep{1979Heiles,2002McClure-Griffiths}, and clouds \\citep{2002Lockman}. The Galactic \\HI\\ disk extends to Galactocentric radii, $R\\geq 30$~kpc and its thickness varies from $\\leq 100$~pc inside $R=3.5$ kpc to $\\sim 3$ kpc in the outer Galaxy, with a roughly uniform thickness of $230$~pc between $3.5$~kpc and the solar circle (see \\citealt{2001Ferriere} and references within). \\HI\\ is also known to extend far beyond the thin \\HI\\ disk as a layer into the Galactic halo \\citep{1984Lockman}. However, recent high angular resolution observations using the Green Bank Telescope (GBT) have revealed that this layer is not smooth but instead is composed of small \\HI\\ clouds with sizes on the order of a few tens of parsecs and masses of $50 M_\\odot$ \\citep{2002Lockman}; confusion may limit the detectability of such clouds at low heights, in which case they may not be confined to the halo (see, e.g., \\citealt{2006Stil}). These clouds follow Galactic rotation and are discrete clumps of \\HI\\ that are localized in space and velocity. Although they are sometimes related to larger structures, for example, embedded in filaments, each cloud appears to be a distinct object. The gross properties of the Milky Way's thick \\HI\\ layer may in fact be a consequence of the statistical properties of these \\HI\\ clouds and it is possible that thick \\HI\\ disks in other galaxies, once observed with sufficient resolution, would reveal a similar structure.  While the origin of the halo clouds is unknown, one possible explanation is a ``galactic fountain'' model, where hot gas produced by supernovae rises into the halo of the Galaxy, cools and condenses into \\HI\\ clouds, which then fall back to the plane \\citep{1976Shapiro,1980Bregman}. \\citet{1990Houck} predicted that, from a lower temperature fountain, gas in Galactic rotation could be formed at heights close to the plane. This scenario is supported by observations of intermediate velocity clouds (IVCs), such as those of cloud g1, whose location, kinematics and abundances match those expected \\citep{2008Wakker}. The abundances in the IV Arch \\citep{2001RichterA} and LLIV Arch \\citep{2001RichterB} also suggest a fountain origin; as their abundances are near solar it is likely that they originated from material enriched from the disk. The location of these IVCs also support this, as they are roughly $1$~kpc from the disk.  Another possibility is that the halo clouds originate in environments where supernovae and stellar winds disrupt the surrounding medium. Such events can result in the formation of a bubble, and models suggest that with a large enough energy source it is possible for these bubbles to expand beyond the thickness of the Galactic disk (e.g., \\citealt{1988Tomisaka,1990Heiles}). These bubbles are encompassed by an \\HI\\ shell as a result of either radiative cooling, which then accumulates more \\HI\\ as the bubble continues to expand, or solely the sweeping up of ambient material, depending on the wind speed \\citep{1992KooB}. The shell remains as a single entity until in some cases Rayleigh-Taylor instabilities cause it to fragment \\citep{1989MacLow}. Once this fragmentation has occurred, gas that has been shock heated is expelled outwards, mixing material from the disk with that in the halo \\citep{1989Norman} and the remaining fragments of the \\HI\\ shell may be the observed halo clouds \\citep{2006McClure-Griffiths}. It is also possible for the hot gas that has been expelled to cool and recombine, as seen in models of \\citet{2000deAvillez}, or perhaps energy from increased supernova activity in areas of active star formation has simply pushed disk gas into the halo, forming clumps of \\HI.  Large samples of halo clouds are required to constrain the properties and distribution of the population and provide insight into their origin. As the \\citet{2002Lockman} sample only has $38$ clouds in a small region of the Galaxy, the properties of the population that constitute the \\HI\\ layer are currently not well-determined. In this paper we present a catalog of over $400$ \\HI\\ clouds in the lower halo of the inner Galaxy, which we have detected in the Galactic All-Sky Survey pilot region, along with an analysis and discussion of their properties and distribution. We begin with an overview of the observations and data in \\S\\ref{observations} and present the observed properties of all clouds in \\S \\ref{clouds}. In \\S \\ref{tangentpoints} we determine the physical properties of a subset of these clouds that can be assumed to be located at tangent points, where the observer's line-of-sight is tangent to circles of constant Galactocentric radius. An analysis of the cloud properties is presented in \\S \\ref{discussion} and implications of these results are discussed in \\S \\ref{nature}. We summarize the results and discuss future work in \\S\\ref{conclusions}. ",
        "conclusions": "\\label{conclusions}  We have detected over $400$ \\HI\\ clouds in the lower halo of the Galaxy in the Galactic All-Sky Survey pilot region. These clouds have a median peak brightness temperature of $0.6$~K, a median velocity width of $12.8$~km~s$^{-1}$, and angular sizes $\\lesssim 1\\degr$. As these clouds follow Galactic rotation, a subset was selected that is likely to be located at tangent points of the inner Galaxy, allowing us to determine their distances and therefore their sizes and masses. The tangent point clouds have radii on the order of $30$~pc and a median \\HI\\ mass of $630 M_{\\odot}$. The properties of these clouds suggest that they belong to the same population of clouds discovered by \\citet{2002Lockman}.  We simulated the population of clouds to constrain their random cloud-to-cloud velocity dispersion, $\\sigma_{cc}$, and spatial distribution. We found that $\\sigma_{cc}=18$~km~s$^{-1}$, but if the clouds were left to evolve in the Galactic potential without any disruptions, these random motions would produce a scale height of $90$~pc, which is inconsistent with our derived scale height of $400$~pc. This suggests that the clouds do not belong to an equilibrium population. We detected clouds throughout the entire GASS pilot region, up to the latitude boundaries ($|b|\\lesssim 20\\degr$). Few clouds were observed at low latitudes due to confusion, which may have resulted in an underestimate of the number of clouds at low heights, and therefore an overestimate in the derived scale height.  Our large, homogeneously-selected sample has allowed us to determine the spatial distribution of these halo clouds for the first time and has revealed that although clouds were observed at all longitudes within the GASS pilot region, they do not appear to be uniformly distributed but instead are concentrated in radius, peaking at $R=3.8$~kpc. We analyzed this distribution and suggest that the clouds are related to the spiral structure of the Galaxy. In particular, the peak in the radial distribution is suggestive of a relation to the $3$~kpc arm. This relation to a specific spiral feature remains speculative until further analysis using a larger sample of clouds and better constrained spiral structure models can be performed. It is therefore unlikely that the halo clouds are a result of a standard galactic fountain, as radial enhancements would not be expected in this scenario \\citep{1980Bregman}. Instead, it appears that the clouds may be directly related to areas of active star formation, in the form of fragmenting \\HI\\ shells and \\HI\\ gas that has been pushed into the halo. This is further supported by the appearance of numerous clouds related to filaments and loops, whose structures may have resulted from stellar winds and supernovae \\citep{1990Dickey}. However, a comparison between the radial surface density distribution of \\HI\\ clouds and \\HII\\ regions provides conflicting evidence: if clouds are related to areas of star formation, a relationship between the clouds and \\HII\\ regions would be expected but is not observed.  The morphology of \\HI\\ at large positive latitudes in the GASS pilot region suggests that some of the clouds may be related to high velocity cloud complex L, whose lower height estimates of $4$~kpc and distance estimates of $8$~kpc place it in the vicinity of the clouds presented here. If they are related, it is possible that whatever process is responsible for the halo clouds, e.g., star formation, is also responsible for complex L. This is further supported by recent observations of IVCs, which suggest that the IVCs are related to energetic events in the Galactic disk and that they are likely linked to spiral structure \\citep{2006Kerton}. The majority of those IVCs have $V_{\\mathrm{dev}}\\leq 20$~km~s$^{-1}$, implying  a Galactic origin, and we suggest that they are similar to the cloud population presented here.  A future study of halo clouds within the entire inner Galaxy observed by the Galactic All-Sky Survey will provide a much larger sample of tangent point clouds, enabling more complete statistics for distribution studies and a better understanding of Galactic structure. These studies will include comparisons between the distribution of halo clouds and that of tracers of star formation, such as \\ion{H}{2} and $\\mathrm{H}_2$, and will allow us to better constrain the origin of these clouds."
    },
    "0807/0807.2884_arXiv.txt": {
        "abstract": "The Pierre Auger Cosmic-Ray Observatory uses the earth's atmosphere as a calorimeter to measure extensive air-showers created by particles of astrophysical origin.  Some of these particles carry joules of energy. At these extreme energies, test beams are not available in the conventional sense. Yet understanding the energy response of the observatory is important. For example, the propagation distance of the highest energy cosmic-rays through the cosmic microwave background radiation (CMBR) is predicted to be strong function of energy.  This paper will discuss recently reported results from the observatory and the use of calibrated pulsed UV laser \"test-beams\" that simulate the optical signatures of ultra-high energy cosmic rays. The status of the much larger 200,000 km$^3$ companion detector planned for the northern hemisphere will also be outlined. ",
        "introduction": "By the time this article appears in print, the base-line configuration of Pierre Auger Detector in Argentina (Auger South) will be completed. The tank commemorating completion of the surface detector was activated (Fig. \\ref{tlast}), two weeks after this conference concluded.  \\begin{figure}[h] \\begin{minipage}{16pc} \\includegraphics[width=16pc]{figures/last-tank.eps} \\caption{\\label{tlast}June 13th 2008: Auger South staff celebrate deployment of the commemorative final surface detector tank. The design specification of 1600 operation tanks has been achieved.} \\end{minipage}\\hspace{2pc}% \\begin{minipage}{22pc} \\includegraphics[width=22pc]{figures/auger-map.eps} \\caption{\\label{config}Auger South baseline configuration.} \\end{minipage} \\end{figure}  Roughly the size of Provincia di Pavia, Auger South is the largest detector of cosmic rays above $5\\times10^{18}$eV in the world. The instrument has been described in previous CALOR conferences \\cite{bianca06} and elsewhere \\cite{paoprotonim}. So this talk will focus instead on an aspect of this experiment that has not been presented in this conference series yet relates closely to its calorimetric theme. This is the use of atmospheric laser \\textgravedbl test-beams\\textacutedbl \\space that simulate many of the optical properties of an extensive air-shower (EAS). ",
        "conclusions": ""
    },
    "0807/0807.0553_arXiv.txt": {
        "abstract": "{} {We describe the discovery of an extremely wide pair of low-mass stars with a common large proper motion and discuss their possible membership in a Galactic halo stream crossing the Solar neighbourhood.} {In a high proper motion survey of the southern sky we used multi-epoch positions and photometry from the SuperCOSMOS Sky Surveys. New nearby ultracool dwarf and subdwarf candidates were selected among the faint and red high proper motion objects, and subsequently confirmed by low-resolution classification spectroscopy. The resulting spectroscopic distance estimates, approximate radial velocity measurements and improved proper motions involving additional epochs from the Two Micron All Sky Survey and from the DEep Near-Infrared Survey were used to compute Galactic space velocities.} {The late-type (M7) dwarf SSSPM\\,J2003$-$4433 and the ultracool subdwarf SSSPM\\,J1930$-$4311 (sdM7) sharing the same very large proper motion of about 860~mas/yr were found in the same sky region with an angular separation of about 6\\degr. From the comparison with other high proper motion catalogues we have estimated the probability of a chance alignment of the two new large proper motions to be less than 0.3\\%. From the individually estimated spectroscopic distances of about $38^{+10}_{-7}$~pc and $72^{+21}_{-16}$~pc, respectively for the M7 dwarf and the sdM7 subdwarf, and in view of the accurate agreement in their large proper motions we assume a common distance of about 50~pc and a projected physical separation of about 5~pc. The mean heliocentric space velocity of the pair $(U,V,W)=(-232, -170, +74)$~km/s, based on the correctness of the preliminary radial velocity measurement for only one of the components and on the assumption of a common distance and velocity vector, is typical of the Galactic halo population.} {The large separation and the different metallicities of dwarfs and subdwarfs make a common formation scenario as a wide binary (later disrupted) improbable, although there remains some uncertainty in the spectroscopic classification scheme of ultracool dwarfs/subdwarfs so that a dissolved binary origin cannot be fully ruled out yet. It seems more likely that this wide pair is part of an old halo stream. Higher-resolution spectroscopic observations are needed to measure accurate radial velocities of both components. Further, we suggest to check the M7 dwarf for an unresolved binary status, which would explain its shorter spectroscopic distance estimate, and to place both objects on a trigonometric parallax program.} ",
        "introduction": "Common proper motion (CPM) stars are typically wide binary stars for which the measuring errors of the proper motion are much larger than the effect of orbital motion. Wide binaries of different masses are interesting for stellar evolution studies since their components, believed to have the same age and metallicity, have evolved without interaction and can be studied independently and compared to each other (e.g. Seifahrt et al.~\\cite{seifahrt05}; Catal\\'{a}n et al.~\\cite{catalan08}). Halo wide binaries provide important traces of the formation of the Galaxy (Allen et al. \\cite{allen07}) and can be used to constrain Galactic halo dark matter (Chanam{\\'e} \\& Gould \\cite{chaname04}).  CPM stars are usually found among stars with large proper motions where the compatibility of the proper motions is easier to detect. The ratio between the angular separation and the proper motion can be used as a criterion to discriminate between physical and optical pairs (Halbwachs \\cite{halbwachs86}). But even for very large separations, a physical connection in terms of membership in a cluster or stellar stream (or a dissolved binary status) is obvious if the individual proper motions are much larger than their measuring errors and agree within these errors.  The LDS (Luyten Double Stars) catalogue of CPM stars (Luyten~\\cite{luyten4087}) lists 20 pairs with angular separations larger than 600~arcsec (maximum 9000~arcsec). Only four of these pairs have a proper motion above the lower limit of the Luyten Half Second (LHS; Luyten~\\cite{luyten79a}) catalogue of 500~mas/yr. For LDS entries with smaller proper motions and relatively large separations there is a higher probability of random pairings (see e.g. LDS\\,4041 and LDS\\,4999 in Scholz, Meusinger \\& Jahrei{\\ss} \\cite{scholz05}). The catalogue of wide binaries by Allen, Poveda \\& Herrera (\\cite{allen00}) includes a few very wide CPM pairs with angular separations up to about 800~arcsec and with a maximum expected semi-major axis of about 125\\,000~AU. At least the primary components are typically bright Hipparcos stars so that the projected physical separations given in that catalogue rely on trigonometric parallax measurements.  Extremely wide (angular separation $>$1\\degr) binaries and multiple systems with pre-Hipparcos trigonometric parallaxes were presented in Gliese \\& Jahrei{\\ss} (\\cite{gliese88}). In their list of 32 systems with projected tangential separations exceeding 5000\\,AU, there are 9 systems with angular separations between 1\\degr\\ and 16\\degr, corresponding to minimum physical separations of 10\\,000\\,AU up to more than a million AU (their proper motions range between 0.2 and 3.8~arcsec/yr). Four of the latter were discussed by the authors as gravitationally unbound due to their large separations and differences in their parallaxes or radial velocities. The Hipparcos parallaxes (ESA~\\cite{esa97}) have shown that only 3 out of the above mentioned 9 systems, Gl~881+879, Gl~803+799AB, and Gl~559AB+551, can be considered as physical since their parallaxes agree to within 5\\% and their projected linear separations are 54\\,250, 47\\,180, and 10\\,380~AU, respectively\\footnote{Three other systems (Gl~331ABC+332AB, Gl~48+22AB, Gl~1255ABC+HD~199476) are questionable physical systems with projected linear separations between 100\\,000 and 350\\,000~AU and parallax differences of 10-15\\%.}. It is remarkable that for the first two of these systems, which are located at distances of 8 and 10~pc, respectively, all proper motion components do also agree to within about 10~mas/yr. This is not the case for Gl~559AB+551 ($= \\alpha$CenAB+Proxima), where the distance of only 1.3~pc and the orbital motions result in differences up to a few 100~mas/yr between the three components of this bound system (Wertheimer \\& Laughlin \\cite{wertheimer06}).   All previously known extremely wide systems include at least one star with $>$0.5~$M_{\\odot}$. In this paper, we present the first extremely wide CPM pair consisting of two very low-mass ($\\approx$0.1~$M_{\\odot}$) stars. In Section~\\ref{sect_cpm} we describe the astrometric measurements and show how accurately the large proper motions of the two faint red stars separated by about 6\\degr\\ agree with each other. In Section~\\ref{sect_spec} we present a rather surprisingly different spectroscopic classification of the components of this extremely wide CPM pair, and in Section~\\ref{sect_halo} we discuss their possible membership in a Galactic halo stream. ",
        "conclusions": "\\label{sect_halo}  The assumed distance interval leads to projected physical separations between 3.5 and 8~pc so that we can rule out a wide binary. Therefore, we investigate the kinematic connection between the two objects as possible members in a stellar stream.  The proper motion error of less than 10~mas/yr corresponds with the assumed distance of 50~pc to an error in the tangential velocity of less than 2~km/s. However, the large uncertainty in the distance to our pair translates with a proper motion of 860~mas/yr to a wide tangential velocity interval of $204^{+102}_{-61}$~km/s. A radial velocity (RV) measurement is available for only one of the components, SSSPM\\,1930, with $-262\\pm25$~km/s from Scholz et al. (\\cite{scholz04}). Therefore, the errors in the $UVW$ space velocity components are dominated by the uncertain distance estimates and by the (partly lacking) RV measurements, whereas the proper motion errors play only a minor role.  \\begin{figure} \\centering \\includegraphics[angle=-90,width=130mm]{9797fig4_res150.ps} \\caption{$UVW$ velocities of 35682 CRVAD-2 stars (dots). Four stars with velocity components within 50~km/s (squares) of the mean velocity of SSSPM\\,1930 and SSSPM\\,2003 are marked (circles). } \\label{uvw}% \\end{figure}   We have computed possible heliocentric Galactic space velocities $UVW$ (Johnson \\& Soderblom \\cite{johnson87}), positive towards the Galactic Centre, Galactic rotation, and North Galactic Pole, separately for the two objects using their positions, the individual proper motions (Table~\\ref{tab_pm}, run \\#4), and a wide parameter space for the parallaxes (from 10 to 40~mas in steps of 0.5~mas) and RVs (from $-$300 to $+$300~km/s in steps of 5~km/s). An exact agreement, to within 5~km/s for all three $UVW$ velocity components, could only be achieved for positive RVs, where with distances of the order of 50~pc the RVs were required to be larger than $+$140~km/s. Generally, the agreement in the space velocities was achieved with slightly larger parallaxes ($+$1~mas) and RVs ($+$5~km/s) for SSSPM\\,2003 compared to those of SSSPM\\,1930.  However, if we consider the large negative RV of SSSPM\\,1930 measured by Scholz et al. (\\cite{scholz04}) to be approximately correct, then we can only achieve an agreement in all three space velocity components to within 25~km/s. In this case, SSSPM\\,2003 would lie at $\\approx$32~pc but SSSPM\\,1930 would be much too close ($<$27~pc) compared to our spectroscopic estimate. If we allow for differences in the $UVW$ components of up to 30~km/s, then we get the following solution with individual distances close to our preferred common distance of about 50~pc:\\\\  SSSPM\\,1930: $d$=44~pc, RV=$-$235~km/s \\\\ $(U,V,W) = (-245 \\pm  25, -157 \\pm  51,  +59 \\pm  16)$~km/s\\\\  SSSPM\\,2003: $d$=50~pc, RV=$-$220~km/s \\\\ $(U,V,W) = (-219 \\pm  24, -183 \\pm  66,  +89 \\pm  15)$~km/s\\\\  Opposite to our spectroscopic distances, SSSPM\\,2003 turns out to be more distant and having a smaller negative RV value than SSSPM\\,1930 in all solutions with the condition of similar $UVW$ components. An unresolved equal-mass binary status of SSSPM\\,2003 would lead to a corrected spectroscopic distance of $\\approx$54~pc and help to explain the above mentioned discrepancy.  Although we are aware of the uncertainties in our input data (distances and RVs) and in the assumption of a common space motion of the pair, we can further speculate on the consequences of the computed $UVW$. According to Chiba \\& Beers (\\cite{chiba00}) the above listed large $U$ and $V$ and moderately large $W$ values are more typical of Galactic halo rather than of thick disk stars. To our knowledge, no halo stream with comparable velocity components has been identified so far. The Catalogue of Radial Velocities with Astrometric Data (CRVAD-2; Kharchenko et al.~\\cite{kharchenko07}) lists the RVs of about 55\\,000 bright stars together with their proper motions and trigonometric parallaxes if available. We have computed $UVW$ for all CRVAD-2 stars whose parallaxes are more than three times larger than their errors (35682 stars) and searched for stars with velocity components similar (within 50~km/s) to the mean $UVW$ of our wide pair (Fig.~\\ref{uvw}). Only four stars (marked on Fig.~\\ref{uvw}) were found to satisfy these criteria: Hip 10449, 12294, 18742, 81528. Two of them are in the equatorial zone, the others have large declinations with different signs. The nearest (Hip 10449) is at $\\approx$60~pc, the others lie between 150~pc and 420~pc. Three of these F8- to K0-type stars are widely discussed in the literature as thick disk/halo stars. However, their metallicities are different (-2.5 $<$ [Fe/H] $<$ -0.7 dex, according to the Simbad database).     The similarity of the two subdwarfs LHS\\,377 and SSSPM\\,1930 and their clearly different spectra compared to the normal dwarfs (including SSSPM\\,2003) is well seen in Fig.~\\ref{5spec}. However, if we apply the procedure proposed by L{\\'e}pine et al.~(\\cite{lepine07}) to derive revised metallicity classes from the measured spectral indices CaH2, CaH3 and TiO5 of LHS\\,377 (taken from Gizis \\cite{gizis97}) and SSSPM\\,1930 (from Scholz et al. \\cite{scholz04}) we get a $\\zeta_{\\text{TiO/CaH}}$ of 0.994 and 0.918, respectively. Consequently, the two objects would be classified as normal dwarfs (dM) instead of subdwarfs (sdM), M7.2 and M6.9, respectively. Unfortunately, their indices fall exactly in the nonvalidated region (0.2$<$TiO5$<$0.7, 0.5$<$CaH2+CaH3$<$0.9) of the classification diagram used by L{\\'e}pine et al.~(\\cite{lepine07}), where they mentioned a different slope of the iso-metallicity lines obtained from model spectra. Therefore, we prefer our sdM7 classification of SSSPM\\,1930, hinting at a lower metallicity compared to the M7 dwarf SSSPM\\,2003, until better measurements and classification schemes will be available. The metallicity problem does not affect our crude estimates for the distance and space velocity of the pair.  The two objects described in this paper form the widest-known pair of (very low-mass) stars with a large proper motion in common. More detailed investigations of this extreme pair are certainly needed, in particular accurate RV measurements of both components. Trigonometric parallaxes will further improve the velocity components and help to explain the physical association of the objects with each other and their possible membership in a Galactic halo stream."
    },
    "0807/0807.3850.txt": {
        "abstract": "{}{}{}{}{} % 5 {} token are mandatory  \\abstract % context heading (optional) % {} leave it empty if necessary {The major challenges for a fully polarized radiative transfer driven approach to Zeeman-Doppler imaging are still the enormous computational requirements. In every cycle of the iterative interplay between the forward process (spectral synthesis) and the inverse process (derivative based optimization) the Stokes profile synthesis requires several thousand evaluations of the polarized radiative transfer equation for a given stellar surface model.} % aims heading (mandatory) {To cope with these computational demands and to allow for the incorporation of a full Stokes profile synthesis into Doppler- and Zeeman-Doppler imaging applications as well as into large scale solar Stokes profile inversions, we present a novel fast and accurate synthesis method for calculating local Stokes profiles.} % methods heading (mandatory) {Our approach is based on artificial neural network models, which we use to approximate the complex non-linear mapping between the most important atmospheric parameters and the corresponding Stokes profiles. A number of specialized artificial neural networks, are used to model the functional relation between the model atmosphere, magnetic field strength, field inclination, and field azimuth, on one hand and the individual components $(I,Q,U,V)$ of the Stokes profiles, on the other hand.} % results heading (mandatory) {We performed an extensive statistical evaluation and show that our new approach yields accurate local as well as disk-integrated Stokes profiles over a wide range of atmospheric conditions. The mean rms errors for the Stokes $I$ and $V$ profiles are well below 0.2 \\% compared to the exact numerical solution. Errors for Stokes $Q$ and $U$ are in the range of 1 \\%. Our approach does not only offer an accurate approximation to the LTE polarized radiative transfer it, moreover, accelerates the synthesis by a factor of more than 1000.} % conclusions heading (optional), leave it empty if necessary {} ",
        "introduction": "The successful application of Doppler- and Zeeman-Doppler imaging (DI \\& ZDI) in the last two decades has enormously contributed to our current knowledge about stellar surface activity and magnetism \\citep{Strass98, Donati99b, Pisk02, Strass02, Donati03, Koch04, Donati06}. Though it has its limitation, Zeeman-Doppler imaging is the only available method that allows to estimate the surface distribution of stellar magnetic fields. Doppler- and Zeeman-Doppler imaging rely on the complex and computational expensive interplay between the forward process of spectral line synthesis on the one hand and the inverse process of a regularized optimization or fitting procedure on the other hand. Based on a highly discretized surface model Doppler- and Zeeman-Doppler imaging iteratively seeks the best possible solution of the surface temperature and/or magnetic field surface distribution, which is compatible with the observed rotationally-modulated line profiles.  The spectral synthesis plays a decisive role here and requires a good knowledge of accurate atomic data, continuous as well as line opacities, broadening mechanisms, model atmospheres etc. The usual way to cope with the synthesis in DI applications, which allows for the calculation of disk-integrated line profiles in a reasonable time frame, is the compilation of a precalculated database of local line profiles for large number of different atmospheric models and lines-of-sight (LOS) \\citep{Rice00}. But unlike DI, the radiative transfer in magnetized media requires the solution of the coupled set of equation for all four Stokes parameters \\citep{Rees87} and moreover, since the magnetic field is a vector quantity, a proper precalculation quickly becomes prohibitive for ZDI applications in the general case where no simple global magnetic field organization is present.  Due to these complexities and enormous demands many applications of ZDI (but also DI) still rely on approximate methods or even completely avoid a radiative transfer modeling approach. Some use an ad-hoc line synthesis where Stokes $I$ profiles are used from template stars or modeled by simple Gaussian profiles \\citep{Donati97a}, from which the corresponding Stokes $V$ profiles are derived by means of the weak-field approximation \\citep{Donati97b,Donati99a}. In this respect, one is able to entirely bypass radiative transfer modeling.  In an effort to provide the necessary tools for a fully polarized radiative transfer driven approach to ZDI as well as to provide the diagnostic capabilities for the next generation of a high-resolution spectropolarimeter (PEPSI) at the 8.4 m Large Binocular Telescope (LBT) \\citep{Strass03,Strass08}, we developed the Zeeman-Doppler imaging code \\emph{iMap} \\citep{Carroll07}. To accelerate the synthesis process, our ZDI code incorporates, besides the conventional numerical implementation of the polarized radiative transfer, a novel fast and accurate method for calculating local and disk-integrated synthetic Stokes profiles under local thermodynamic equilibrium (LTE). This approach uses artificial neural networks (ANNs) to approximate the polarized radiative transfer calculations. Artificial neural networks have already proven to be a fast alternative for solar Stokes profile inversions and Zeeman tomography \\citep{Carroll01, Socas05a, Socas05b, Lites07, Carroll08}. In the here presented approach ANNs are trained on the basis of the numerical solution of the polarized radiative transfer equation to find an approximate representation of the underlying complex mapping between atmospheric quantities like temperature, magnetic field, velocity, etc., and the corresponding Stokes profiles. A detailed assessment and evaluation of the ANN synthesis confirms the accuracy of our novel approach and moreover demonstrates that the ANN approach can accelerate the synthesis process by three orders of magnitude.  This paper is organized as follows: Since our method strongly depends on the accuracy of the underlying numerical method used for the calculation of synthetic Stokes spectra, we first perform an in-depth analysis and assessment of our polarized radiative transfer code in Sect. \\ref{Sect:2}. For this purpose, we made a detailed comparison between our code \\emph{iMap} and the C{\\small OSSAM} synthesis code \\citep{Stift00}, which was also part of the benchmark test and inter-agreement analysis of \\citet{Wade01}. In Sect. \\ref{Sect:3}, we give a short introduction into the ANN models we used in this work and describe the calculation of the training database for different atmospheric models, magnetic field configurations and Zeeman sensitive spectral lines. In Section \\ref{Sect:4} we evaluate the accuracy and performance of the trained ANNs relative to the exact numerical solution. Sect. \\ref{Sect:5} presents a benchmark test that demonstrates the exceptional speed of the new synthesis method compared to the conventional numerical procedure. Finally, we conclude in Sect. \\ref{Sect:6} with a discussion and summary of the presented fast Stokes profile synthesis.   %__________________________________________________________________ ",
        "conclusions": "\\label{Sect:6} In this paper we have introduced, described and evaluated a novel LTE Stokes spectrum synthesis method based on ANNs. Artificial neural networks are used to approximate the LTE polarized radiative transfer by modeling the functional relationship between the most important atmospheric parameters and the corresponding local Stokes spectra. The ANN approach with MLPs is based on an extensive training database of calculated synthetic Stokes profiles. Because of this crucial dependence for the network training we first assessed the accuracy of our polarized radiative transfer module integrated in our ZDI code \\emph{iMap} by a detailed comparison with the existing synthesis code C{\\small OSSAM} in Sect. \\ref{Sect:3}. Both codes exhibit congruent results for several composite values like the continuous absorption coefficient and broadening parameters. Moreover, both codes show a very good agreement in all calculated Stokes profiles. By verifying the agreement between \\emph{iMap} and C{\\small OSSAM}, we also indirectly established an inter-agreement between \\emph{iMap} and I{\\small NVERSE}10 as well as Z{\\small EEMAN}2, as derived from the inter-agreement study of \\citet{Wade01}.  In Sect. \\ref{Sect:4}, we introduced the concept of our proposed method and described the data preparation with a PCA method as well as the training process for the ANNs. A detailed evaluation of our ANN synthesis revealed the remarkable accuracy with which the ANN approach is able to approximate the polarized radiative transfer process. Stokes profiles of the spectral line \\ion{Fe}{i} $\\lambda$ 6173 \\AA\\ are calculated with high accuracy for different configurations of the atmospheric input parameters (effective temperature, magnetic field strength, field inclination, field azimuth, and LOS angle). The ANNs have clearly been able to \\emph{learn} the complex non-linear mapping between the atmospheric input parameters and the resulting Stokes profiles. Although we have restricted in this work the number of variable input parameters to 5, there is no principle limitation to extend the number of atmospheric input parameters. Moreover, the limits that we have set for the parameter ranges (see Table \\ref{Table:1}) are also not a principle restriction of the ANN approach but rather driven by the constraints to cope with the training and evaluation database in this study.  The ANN approach to spectral line synthesis is also not limited to one spectral line, the ANN method can easily be extended to a wider wavelength range to cover several neighboring spectral lines or successive regions of interest in the wavelength domain. Another interesting possibility, offered by the ANN approach is the modeling of co-added line profiles as used for polarized spectral line extraction and reconstruction algorithms like LSD \\citep{Donati97c} or by the PCA method  \\citep{Carroll07,Marian08}. Instead of synthesizing a huge number of individual Stokes profiles to model, the co-added (and therefore noise reduced) mean line profile, an ANN synthesis can be used to directly calculate this co-added line profile. This would offer a vast improvement in speed and accuracy.  It is also important to note that since the ANN approach is based on the accurate numerical modeling of line profiles, blends can be fully accounted for with this approach. Even heavily blended line profile can be accurately synthesized as long as the training database properly accounts for these blends.  One of the main reasons to approximate or even to entirely bypass a full radiative transfer approach in DI or ZDI applications, is the enormous requirement for massively calculating local Stokes profiles during the inversion process. In our benchmark tests, we could show that our ANN approach is three orders of magnitudes faster then the conventional numerical integration scheme. This enormous acceleration and the impressive accuracy makes the ANN synthesis approach a viable and promising alternative for polarized radiative transfer calculations and facilitates a complete radiative transfer driven approach to DI and ZDI as well as for large scale Stokes profile inversions.  %__________________________________________________________________"
    },
    "0807/0807.0623_arXiv.txt": {
        "abstract": "Systematic search for multiperiodicity in the LMC Cepheids (Moskalik, Ko{\\l}aczkowski \\& Mizerski 2004) has led to discovery of low amplitude nonradial modes in a substantial fraction of overtone pulsators. We present detailed discussion of this new type of multimode Cepheid pulsators and compare them to similar nonradial pulsators discovered among RR~Lyrae stars. Finally, we show first detections of secondary nonradial modes in FU/FO double-mode Cepheids. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0807/0807.3831_arXiv.txt": {
        "abstract": "New information on the relations between the Galactic disks, the halo, and satellite galaxies is being obtained from elemental abundances of stars having metallicities in the range $-1.5 <$ [Fe/H] $< -0.5$. The first results for a sample of 26 halo stars and 13 thick-disk stars observed with the ESO VLT/UVES spectrograph are presented. The halo stars fall in two distinct groups: one group (9 stars) has [$\\alpha$/Fe]= $0.30 \\pm 0.03$ like the thick-disk stars. The other group (17 stars) shows a clearly deviating trend ranging from [$\\alpha$/Fe]= 0.20 at [Fe/H]= $-1.3$ to [$\\alpha$/Fe]= 0.08 at [Fe/H]= $-0.8$. The kinematics of the stars are discussed and the abundance ratios Na/Fe, Ni/Fe, Cu/Fe and Ba/Y are applied to see if the ``low-alpha'' stars are connected to the thin disk or to Milky Way satellite galaxies. Furthermore, we compare our data with simulations of chemical abundance distributions in hierarchically formed stellar halos in a $\\Lambda$CDM Universe. ",
        "introduction": "As discussed by Venn et al. (2004), the chemical signatures of stars in Milky Way dwarf spheroidal (dSph) galaxies are different from the majority of Galactic thin-disk, thick-disk and halo stars. The dSph stars have lower values of $\\alpha$/Fe, where $\\alpha$ refers to the abundance of typical alpha-capture elements like Mg, Si, Ca and Ti, than the Galactic stars. Furthermore, the Ba/Y ratio is significantly higher in the dSph stars. Thus, Venn et al. conclude that the main components of our Galaxy have not been formed through the merging of dwarf galaxies similar to present-day dSphs. This could be seen as an argument against hierarchical structure formation as predicted in Cold Dark Matter cosmologies. Font et al. (2006) have, however, made $\\Lambda$CDM simulations of the chemical abundance distributions in the the Galactic halo, which show that early accreted dwarf galaxies have chemical properties different from those of present-day satellite galaxies. For the metal-poor part of the Galactic halo, they predict enhanced values of [$\\alpha$/Fe] corresponding to the ratio calculated for Type II supernovae, but for the metal-rich end of the halo a decline of [$\\alpha$/Fe] is predicted due to an increasing contribution of iron from Type Ia SNe.  In order to test these recent simulations, we have started a survey of abundance ratios of about 100 halo stars having $-1.5 <$ [Fe/H] $< -0.5$. This metallicity range also contains thin-disk and thick-disk stars, and hence a comparison of abundances may give new information on the relations between the various Galactic populations.  The survey is a continuation of a more limited study by Nissen \\& Schuster (1997), who determined abundances of 13 halo stars and 16 thick disk stars in the metallicity range $-1.3 <$ [Fe/H] $< -0.4$. Interestingly, 8 of the halo stars turned out to have significantly lower values of [$\\alpha$/Fe]\\footnote{Defined as [$\\alpha$/Fe] = $\\frac{1}{4}$\\,\\,([Mg/Fe]\\,+\\,[Si/Fe]\\,+\\,[Ca/Fe]\\,+\\,[Ti/Fe])} than the other halo stars. Considering the small number of halo stars studied it is, however, unclear from the work of Nissen \\& Schuster if the distribution of [$\\alpha$/Fe] in the halo is continuous or bimodal. One of the aims of the present study is to answer this question by observing a larger sample of stars. ",
        "conclusions": "Figures  1 and 2 show the derived abundances of the four alpha-capture elements, Mg, Si, Ca and Ti, relative to Fe. As seen, the halo stars fall in two distinct groups: $i)$ the ``high-alpha'' stars that have very near the same [$\\alpha$/Fe] as the thick-disk stars and a remarkably small scatter, $\\pm 0.03$\\,dex, around the ``plateau'' value of [$\\alpha$/Fe] and $ii)$ the ``low-alpha'' stars that show a decreasing trend of [$\\alpha$/Fe] as a function of increasing [Fe/H]. In the cases of Mg and Si the separation of the two groups of halo stars is seen for the whole range $-1.5 <$ [Fe/H] $< -0.7$ with a maximum separation of $\\sim 0.25$\\,dex in [Mg/Fe] and $\\sim 0.20$\\,dex in [Si/Fe]. For Ca and Ti the two groups tend to merge at [Fe/H] $\\simeq -1.2$, and the maximum separation is only about 0.12\\,dex for Ca and about 0.17\\,dex for Ti.  \\begin{figure}[ht] \\begin{center} \\includegraphics[width=13.0cm]{nissen.fig3.eps} \\caption{Spectra of two stars with nearly the same atmospheric parameters $T_{\\rm eff}$, log$g$ and [Fe/H].  The spectrum of the ``low-alpha\" star CD\\,$-45\\,3283$ ($T_{\\rm eff}$\\,=\\,5603\\,K, log$g$\\,=\\,4.57, [Fe/H]\\,=\\,$-0.89$, [$\\alpha$/Fe]\\,=\\,0.08) is shown with a thick (red) line, and that of the ``high-alpha\" star G\\,159-50 ($T_{\\rm eff}$\\,=\\,5648\\,K, log$g$\\,=\\,4.39, [Fe/H]\\,=\\,$-0.92$, [$\\alpha$/Fe]\\,=\\,0.29) with a lighter (green) line.} \\label{fig3} \\end{center} \\end{figure}  \\begin{figure}[ht] \\begin{center} \\includegraphics[width=13.0cm]{nissen.fig4.eps} \\caption{[Ni/Fe] and [Cu/Fe] vs. [Na/Fe] with the same symbols as in Fig. 1. In the case of Cu the deviation of [Cu/Fe] from the relation defined by the thick-disk stars is plotted.} \\label{fig4} \\end{center} \\end{figure}   A possible explanation of these trends is that the ``high-alpha'' and the thick-disk stars have been formed in regions with a relative high star-formation rate such that only Type II SNe have contributed to the chemical evolution up to [Fe/H] $\\sim -0.5$. The ``low-alpha'' stars, on the other hand, come from regions with a slower star-formation rate, where Type Ia SNe have started to contribute with iron at a metallicity [Fe/H] $\\simeq -1.5$. The reason for the smaller separation in [Ca/Fe] compared to [Mg/Fe], may be that Mg is almost exclusively produced in Type II SNe, whereas about 25\\% of Ca originates in Type Ia SNe according to the chemical evolution models of Tsujimoto et al. (1995).  The abundance differences between the two halo groups can be seen directly from the observed spectra. Fig. 3 shows two spectral regions for a ``high-alpha'' and a ``low-alpha'' star with similar $T_{\\rm eff}$, log$g$ and [Fe/H] values. As seen, the Fe\\,{\\sc i} lines of the two stars have nearly the same strength, whereas the Mg and Si lines are weaker in the ``low-alpha'' star. The same is the case for Na, Ni and Cu lines and, as shown in Fig. 4, [Ni/Fe] and [Cu/Fe] are well correlated with [Na/Fe]. The reason for these correlations may be that the yields of Na, Ni and Cu depend on the neutron excess in supernovae and that this excess is affected by the $\\alpha$/Fe ratio.   \\begin{figure}[ht] \\begin{center} \\includegraphics[width=10.0cm]{nissen.fig5.eps} \\caption{The ``Toomre'' diagram for the three groups of stars discussed in the present paper. Regions, where stars are most likely to belong to the halo, the thick disk and the thin disk, respectively (see Venn et al. 2004, Fig. 1), are separated by dashed circles corresponding to $V_{\\rm total}$\\,=\\,180 and 70\\,km\\,s$^{-1}$.} \\label{fig5} \\end{center} \\end{figure}   \\begin{figure}[ht] \\begin{center} \\includegraphics[width=7.0cm]{nissen.fig6.eps} \\caption{[Ba/Y] vs. [Fe/H] with the same symbols as in Fig. 1} \\label{fig6} \\end{center} \\end{figure}   The kinematics of the stars are shown in Fig. 5. As seen from this ``Toomre'' diagram, both groups of halo stars have an average Galactic rotation velocity close to zero in contrast to the thick-disk stars. Furthermore, the velocity dispersion for the ``low-alpha'' group is higher than that of the ``high-alpha'' group. This larger dispersion is mainly caused by larger $U$ velocities, which means that the ``low-alpha'' stars tend to move on higher energy radial orbits. A larger sample of halo stars is, however, needed before any definitive conclusions about the kinematics of the two groups can be made.   The distribution of the ``low-alpha'' stars in Figs. 1 and 2 shows a smooth continuation of the $\\alpha$/Fe trend for thin-disk stars, which according to several studies (Reddy et al. 2003, Bensby et al. 2005) stops with [$\\alpha$/Fe] $\\simeq 0.1$ at [Fe/H] $\\simeq -0.7$. The thin-disk stars do not share, however, the low [Na/Fe], [Ni/Fe] and [Cu/Fe] values that are being found for the ``low-alpha'' group. In addition, the ``low-alpha'' stars have enhanced Ba/Y ratios as shown in Fig. 6. In all these abundance deviations, the ``low-alpha'' stars resemble stars in dSph galaxies (Venn et al. 2004, Sbordone et al. 2007, Koch et al. 2008) and in the LMC (Pomp\\'{e}ia et al. 2007), although these satellite galaxies tend to have larger abundance offsets from disk stars than the ``low-alpha'' stars.  The trend of [Mg/Fe] for the ``low-alpha'' halo stars in Fig. 1 agrees remarkably well with the trend predicted by Font et al. (2006, Fig. 9) from simulations of abundance distributions for a hierarchically formed stellar halo in a $\\Lambda$CDM Universe. The fact that the ``low-alpha'' stars do not have quite as low [$\\alpha$/Fe] values as present-day satellite galaxies also agrees with their predictions. On the other hand, the ``high-alpha'' halo stars are not predicted from the simulations. The existence of this group suggests that the formation of the Galactic halo is more complicated than predicted from the $\\Lambda$CDM simulations. It remains to be seen if this group of ``high-alpha'' halo stars can be explained as due to the merger of an exceptionally large satellite or if it is a dissipative component of the Galaxy as suggested by Gratton et al. (2003). In the latter case one would expect the group of ``high-alpha'' stars to have some net rotation. Another problem is that the ``low-alpha'' halo stars tend to move on high-energy radial orbits that plunge into the Galactic central regions from the outer regions of the halo. As discussed by Gilmore \\& Wyse (1998), this requires possible parent satellite galaxies to have a very high mean density to provide ``low-alpha'' stars on orbits with such small perigalactic distances.  In order to obtain more insight into these problems, the elemental abundances of a larger sample of metal-rich halo stars should be studied so that better knowledge of the kinematics and the relative frequency of halo stars belonging to the ``high-alpha'' and ``low-alpha'' groups can be obtained. As mentioned earlier, such data is being obtained with the Nordic Optical Telescope.  \\medskip  {\\em Acknowledgement}. This work has been financially supported from CONACyT project 49434-F."
    },
    "0807/0807.3000_arXiv.txt": {
        "abstract": "We investigate the emergence of Active Region NOAA 10790 by means of time--distance helioseismology. Shallow regions of increased sound speed at the location of increased magnetic activity are observed, with regions becoming deeper at the locations of sunspot pores. We also see a long-lasting region of decreased sound speed located underneath the region of the flux emergence, possibly relating to a temperature perturbation due to magnetic quenching of eddy diffusivity, or to a dense flux tube. We detect and track an object in the subsurface layers of the Sun characterised by increased sound speed which could be related to emerging magnetic flux and thus obtain a provisional estimate of the speed of emergence of around $1 \\, {\\rm km \\ s^{-1}}$. ",
        "introduction": "\\label{sec:Intro}  Local helioseismology provides a tool for studying the emergence and evolution of active regions and complexes of solar activity. Several theoretical scenarios for the formation of active regions on the solar surface have been suggested within the context of the dynamo theory explaining the formation of the ``sunspot belt'' at low latitudes.   Recently developed ``flux-transport'' dynamo models include the magnetic flux transport by meridional circulation \\cite{Wang91, Choudhuri95, Dikpati99, NC02, Kriv05}, flux transport by an interface \\cite{parker93, charb97} and overshoot dynamo with a positive radial shear beneath the surface \\cite{rudiger95}, or by a distributed dynamo with a negative shear \\cite{roberts72}.  All of the models agree that the most favourable and promising place for the $\\alpha \\Omega $ dynamo is the region which includes the deepest layers of the convection zone, the convective overshoot layer and the tachocline. Nevertheless, it is difficult to keep flux tubes with strong magnetic field (B >100 G) in the interior for a long time against magnetic buoyancy \\cite{parker79} and then deliver them to the surface within the sunspot belt.  In a flux-transport model it is possible to calculate the velocities of macroscopic diamagnetic transport as a function of depth in the convective zone by taking into account the magnetic advection processes. \\inlinecite{Kriv05} estimated that in the subsurface layers the transport velocity increases from $0.5 {\\rm \\ km\\ s^{-1}}$ at a depth of 30 Mm to around $0.9 {\\rm \\ km\\ s^{-1}}$ at 10 Mm depth. However, the way in which the flux travels to the surface is not fully understood. There are ongoing arguments in favour of each model, and they can only be tested by comparison with the observations of active regions and sunspots on the surface and their appearance in the solar interior by means of helioseismology.   Time--distance helioseismology  \\cite{Duval93} is a fast-growing area of helioseismology. The data used for time-distance inversions are the estimated wave travel-times between different points on the surface of the Sun. The travel-time estimates are obtained by cross-correlating and averaging the Doppler-velocity signal at different locations at the Sun's surface and then obtaining the cross-correlation function. The travel times are then computed from cross-correlations either via Gabor wavelet fitting \\cite{Kosovichev} or by comparing with a reference cross-correlation function (see \\opencite{GB2005} and references therein). Because of the stochastic nature of solar oscillations, substantial spatial and temporal averaging of the data is required to measure the frequencies and travel times accurately.  The properties of the solar interior are then determined by, first, establishing the relations between travel-time variations and internal properties (variations of the sound speed, flow velocity, magnetic field). This is followed by the inversion of these relations, which are typically cast as linear integral equations, to infer the internal properties.  \\inlinecite{Chang99} applied the method of acoustic imaging to study the emergence of NOAA 7978, and reported the detection of the signature of upward-moving magnetic flux in the solar interior. The first time-distance investigation of an emerging active region was by  \\inlinecite{Kosovichev2000}, wherein an active region which emerged on the solar disk in January 1998, was studied with SOHO/MDI for eight days, both before and after its emergence at the surface. The results have shown a complicated structure of the emerging region in the interior, and suggest that emerging flux ropes travel very quickly through the depth range of our observations, estimating the speed of emergence at about $1.3 {\\rm \\ km\\ s^{-1}}$.  In Section \\ref{Sec:DataMethod} of this paper the method and the data used in this work are described. In Section \\ref{Sec:Discussion} the results of sound-speed inversions are presented, together with the discussion. This is followed by conclusions in Section \\ref{Sec:Conclusions}. ",
        "conclusions": "\\label{Sec:Conclusions} We have investigated the emergence of the Active Region NOAA 10790. The region has a complex magnetic structure and can be viewed as an amalgamation of two magnetic region as described in Section \\ref{sec:DisData}. We see shallow regions of increased sound speed at the location of increased magnetic activity, with regions becoming deeper at the locations of sunspot pores. It appears that a subsurface sound-speed perturbation at depth is associated with the emergence on the surface of a sunspot and sunspot pores. We observe an object with increased sound speed which may be related to the subsurface flux prior to its emergence in the photosphere at the second location. We also observe a loop-like structure of the decreased sound-speed perturbation located below the investigated region, which is present on the averages of the inversions. It would be interesting to see if similar structures characterised by a drop in sound speed can be detected in other instances of emerging active regions.  Based on our results, one can speculate that magnetised regions without any sunspots or sunspot pores lead to increased subsurface sound speed at shallow depths but no perturbation detected at greater depths. On the other hand, the areas associated with sunspot appearance have increased finger-like sound-speed perturbations which extend to depths of 10\\,--\\,20 Mm. Then by tracking the times of the emergence of magnetic flux in the subsurface layers of the Sun one can estimate the speed at which flux emerges to be about $1 {\\rm\\, km \\, s^{-1}}$.  We also note that the precise timing of the signature of emergence in the seismic imaging is very limited by the temporal resolution of  the sound-speed inversions presented here, since each set is a representative of eight hours of solar oscillations.  In view of the dynamo models discussed in Section \\ref{sec:Intro} (\\opencite{NC02}, \\opencite{Kriv05}), it could be interesting to compare such results for a variety of emerging active regions at different latitudes and stages of the solar cycle."
    },
    "0807/0807.4766_arXiv.txt": {
        "abstract": "We have used the Australia Telescope Compact Array (ATCA) and the Swedish-ESO Sub-millimetre Telescope (SEST) to map the large-scale atomic and molecular gas in the nearby (4~Mpc) Circinus galaxy. The ATCA \\HI\\ mosaic of Circinus exhibits the warps in position angle and inclination revealed in the single-pointing image of \\citet{jkeh98}, both of which appear to settle beyond the inner 30 kpc which was previously imaged. The molecular gas has been mapped in both the CO $J=1\\rr0$ and $J=2\\rr1$ transitions down to a column density of $N_{{\\rm H}_{2}}\\ga10^{21}$ \\scm ~($3\\sigma$), where we derive a total molecular gas mass of $M_{\\rm H_2} \\approx2 \\times 10^9$\\Mo. Within a radius of 3 kpc, i.e. where CO was clearly detected, the molecular fraction climbs steeply from $\\approx0.7$ to unity (where $N_{{\\rm H}_{2}}=4\\times10^{22}$ \\scm, cf. $N_{\\rm HI}=10^{21}$ \\scm) with proximity to the nucleus. Our \\HI\\ mosaic gives an atomic gas mass of $M_{\\rm HI} \\approx6 \\times 10^9$\\Mo, which is 70\\% of the fully mapped single dish value. Combining the atomic and molecular gas masses gives a total gas mass of $M_{\\rm gas}\\equiv M_{\\rm HI} + M_{\\rm H_{2}}\\approx1\\times10^{10}$\\Mo, cf. the total dynamical mass of $\\approx3\\times10^{11}$\\Mo\\ within the inner 50 kpc of our mosaiced image.  The total neutral gas mass to dynamical mass ratio is therefore 3\\%, consistent with the SAS3 classification of Circinus. The high (molecular) gas mass fraction of $M_{\\rm H_2}/M_{\\rm dyn}\\approx50$\\% found previously \\citep{cjrb98}, only occurs close to the central $\\approx0.5$ kpc and falls to $\\lapp10$\\% within and outwith this region, allaying previous concerns regarding the validity of applying the Galactic $N_{\\rm H_{2}}/I_{\\rm CO}$ conversion ratio to Circinus. The rotation curve, as traced by both the \\HI\\ and CO, exhibits a steep dip at $\\approx1$ kpc, the edge of the atomic/molecular ring, within which the star-burst is occurring. We find the atomic and molecular gases to trace different kinematical features and believe that the fastest part ($\\gapp130$ \\kms) of the sub-kpc ring consists overwhelmingly of molecular gas. Beyond the inner kpc, the velocity climbs to settle into a solid body rotation of $\\approx150$ \\kms\\ at $\\gapp10$ kpc. Most of the starlight emanates from within this radius and so much of the dynamical mass, which remains climbing to the limit of our data ($\\ga50$ kpc), must be due to the dark matter halo. ",
        "introduction": "% \\label{sec:intro}  Circinus is a nearby, highly inclined spiral galaxy located behind the Galactic Plane (at $l,b$ = 311\\fdg3, --3\\fdg8), where the high stellar density and dust extinction (\\AB\\ = 6.3 mag, \\citealt{sfd98}) make it difficult to study its optical properties and environment. At a distance of only $4.2 \\pm 0.8$ Mpc \\citep{fkl+77}, Circinus represents one of the closest examples of a star-burst galaxy with a Seyfert\\,2 nucleus \\citep{mo90,osmm94}. Low-resolution \\HI\\ observations of Circinus with the 64-m Parkes radio telescope revealed a vast galaxy with \\HI\\ 21-cm emission extending to a half-width of 32\\arcmin\\ \\citep{fkl+77}. Its enormous extent has also been noted in both the shallow and deep \\HI\\ surveys of the Zone of Avoidance (HIZSS, \\citealt{hse+00} and HIZOA; \\citealt{jsk+00}, respectively) and, most recently, in the \\HI\\ Parkes All-Sky Survey (HIPASS, \\citealt{ksk+04}).  High resolution, single-pointing ATCA \\HI\\ observations of Circinus \\citep{jkeh98} reveal a complex velocity field and gas distribution out to a diameter of $\\sim80$\\arcmin\\ ($\\sim100$ kpc), with strong kinematical warping (also evident between the large-scale position angles and inclinations, Table \\ref{info}). \\begin{table} % \\caption{General properties of the Circinus galaxy.\\label{info}} \\begin{center} \\begin{tabular}{lllc} \\hline central position    & $\\alpha_{\\rm J2000}$ & $14^{\\rm h}\\,13^{\\rm m}\\,09\\fs95$ ($\\pm$0\\fs02) & 1 \\\\ & $\\delta_{\\rm J2000}$ & --65\\degr\\,20\\arcmin\\,21\\farcs2 ($\\pm$0\\farcs1) &     \\\\ type                &                & SAS3                          & 2 \\\\ adopted distance    & $D$            & 4.2 Mpc                       & 3 \\\\ optical extent$^*$  & $D_{24.5}$     & $11\\farcm9 \\times 4\\farcm6$   & 3 \\\\ & $D_{26.6}$     & $17\\farcm2 \\pm 1\\farcm7$    & 3 \\\\ position angle      & $PA_{\\rm opt}$ &  220\\degr                     & 2 \\\\ & $PA_{\\rm CO}$  & $214\\degr \\pm 4\\degr$         & 4 \\\\ & $PA_{\\rm HI}$  & $210\\degr \\pm 5\\degr$         & 3 \\\\ inclination         & $i_{\\rm opt}$  &  64\\degr                      & 2 \\\\ & $i_{\\rm CO}$   & $78\\degr \\pm 1\\degr$  ($<30$\\arcsec) & 4 \\\\ & $i_{\\rm HI}$   & $65\\degr \\pm 2\\degr$          & 3 \\\\ systemic velocity   & $v_{\\rm HI}$   & $439 \\pm 2$\\kms\\              & 3 \\\\ \\HI\\ flux density   & \\FHI\\          & $1910 \\pm 130$ Jy\\kms         & 3 \\\\ \\HI\\ mass           & \\MHI\\          & $7.9\\pm0.5\\times10^{9}$\\Mo    & 3 \\\\ dynamical mass      & $M_{\\rm dyn}$  & $1.3\\pm0.2\\times10^{11}$\\Mo   & 3 \\\\ $B$-band luminosity & \\LB\\           & $7  \\pm2  \\times10^{9}$\\Lo    & 3 \\\\ $B$-band extinction & \\AB\\           & 6.3 mag                       & 5 \\\\ \\hline \\end{tabular} \\end{center} Notes: Freeman et al. (1977) derived extinction-corrected optical diameters for the Circinus galaxy at the 24.5th blue magnitude, $D_{24.5}$, and, through extrapolation, the 26.6th blue magnitude, $D_{26.6}$ (the Holmberg diameter).\\\\ References: (1) \\protect\\citet{gbe+03}, (2) \\protect\\citet{ddc+91}, (3) \\protect\\citet{fkl+77}, (4) \\protect\\citet{cjrb98}, (5) \\protect\\citet{sfd98}. \\end{table} The atomic gas is distributed in a $\\sim$10~kpc radius ring enclosing a bar, in turn terminating in a 1 kpc radius ring \\citep{jkeh98}. This inner \\HI\\ ring may be an outer component of the 600 pc CO ring \\citep{cjrb98}, where the star formation is dominant \\citep{mmoo94}. This itself may feed the nucleus via a nuclear bar nestled within the molecular ring \\citep{maa+00}, thus providing a means of transporting gas from the outer galaxy to the central engine (e.g. \\citealt{sfb89}).  The notion of a continuous structure from the kpc to sub-pc scale is supported by the VLBI observations of the \\WAT\\ maser emission \\citep{gbe+03} and VLTI near infrared observations \\citep{tmj+07}, which both show that the sub-pc accretion disk/obscuring torus shares a similar close to edge-on orientation to the larger-scale molecular gas \\citep{cjrb98}.  The inner edge of the warped maser disk shares a similar position angle to the large-scale gas disk ($PA = 209\\degr \\pm 3\\degr$), although this increases to $236\\degr \\pm 6\\degr$ at the outer edge \\citep{gbe+03}. Furthermore, a second population of masers are found to trace a wide angle outflow which coincides with the $>$100 pc-scale ionisation cone (Marconi et al. 1994, \\citealt{vb97}) and molecular outflow (\\citealt{crjb98}), as well as sharing the same position angle as the radio lobes \\citep{hwr+90,ehj+95}. The lobes, cone and outflow are all directed along the minor axis of the galaxy, coincident with the rotation axis of the molecular ring \\citep{crjb98,cur01}.  Previously, \\citet{jkeh98} obtained very sensitive, single-pointing \\HI\\ data of the Circinus galaxy with the Australia Telescope Compact Array (ATCA). Since the ATCA primary beam is $\\sim$33\\arcmin, these data were not sufficiently sensitive to detect the outer edges of the large gaseous disk of Circinus. Here we present and analyse an \\HI\\ mosaic of Circinus, obtained with the 375-m configuration of the ATCA, covering the whole galaxy. This is complemented by the most extensive CO~$1\\rr0$ and $2\\rr1$ maps of the galaxy to date. ",
        "conclusions": "%  We have studied the atomic and molecular emission in the Circinus galaxy through large-scale maps of the \\HI\\ and CO $1\\rr0$ \\& $2\\rr1$ transitions.  Our \\HI\\ mosaic represents the largest high resolution image of this galaxy to date, in which we recover 70\\% of the total \\HI\\ flux, while resolving the detailed gas distribution. We have mosaiced a field of $\\approx\\pm1\\degr\\times\\pm1\\degr$, thus recovering $\\approx30\\%$ more flux than the primary beam limited image of \\citet{jkeh98}, giving the atomic gas distribution and dynamics to a radius of 50~kpc.  The CO maps also represent the largest to date, spanning $\\pm260\\arcsec$ and $\\pm100\\arcsec$ along the major and minor axes, respectively, at a sensitivity of $\\approx10$ mK per each 10 \\kms\\ channel. From our observations: \\begin{enumerate}  \\item We confirm that the atomic gas in Circinus is distributed in a $\\ga50$ kpc irregular disk, which is a factor of $\\ga5$ times as extended as the optical and infrared emission. A tilted ring model of the \\HI\\ velocity field agrees very well with the known warping in position angle and inclination, found by \\citet{jkeh98}, out to $\\approx30$ kpc. Beyond this (to 50 kpc), the warping appears to dampen significantly.  \\item We obtain clear detections of CO at a galactocentric radius of 3~kpc. Although the position angles and inclinations yielded by the tilted ring model of the CO velocity field may be unreliable, by assuming circular orbits for the molecular gas we find that the molecular disk also undergoes a ($\\approx20\\degr$) warp within the inner 3~kpc, a radius at which it shares a similar inclination to the atomic gas. Note, however, that the apparent warp may be the result of non-circular motions, which are a distinct possibility due to the presence of the atomic bar on these scales \\citep{jkeh98}.   \\item The inner part of each gas structure also appears to maintain the same position angle on all scales; sub-pc (\\WAT), sub-kpc (CO) and $\\lapp10$~kpc (\\HI). This may suggest a continuous transportation of material to the active galactic nucleus, while ``kinks'' are introduced when the outer diffuse gas encounters a bar.  \\item We confirm the finding of the ring model of \\citet{cjrb98}, that there is more molecular (and less atomic) gas in the receding (SW) half of the ring.  \\item We find that the CO profile is wider than that of the \\HI\\ in both the full-width half maximum and the total velocity spread. A comparison of the \\HI\\ and CO position-velocity diagrams confirms that the molecular gas has different kinematics to the atomic gas.  We believe that this indicates that the fastest rotating gas, which occurs within the sub-kpc ring, is indeed primarily molecular.   \\item We also investigate the global integrated atomic and molecular gas intensities, and find:  \\begin{enumerate} \\item Circinus has a total molecular gas mass of $M_{\\rm H_{2}} \\approx 2\\times 10^{9}$\\Mo, cf. the atomic gas mass of $M_{\\rm HI}\\approx8 \\times 10^9$\\Mo\\ \\citep{jsk+00}. \\item Previously, \\citet{cjrb98} found a gas mass fraction of $M_{\\rm gas}\\approx M_{\\rm H_{2}}\\approx 0.5 M_{\\rm dyn}$, within the central kpc, which cast into doubt the validity of applying the standard Galactic conversion ratio of $N_{\\rm H_{2}}/I_{\\rm CO} \\approx 2 \\times 10^{20}$ \\scm/[K \\kms] \\citep{sbd+88}. However, over the global molecular gas emission we obtain a more canonical ($M_{\\rm gas}\\la0.1 M_{\\rm dyn}$, \\citealt{mb81}) ratio of $\\approx0.03$. This is much smaller than the value ($M_{\\rm gas}/M_{\\rm dyn}\\approx0.6$) obtained by \\citet{ekhj98}.  \\item Within the inner $\\approx$3 kpc, the gas remains predominantly molecular, with the molecular hydrogen column density ranging from $N_{\\rm H_{2}} \\approx 2\\times 10^{21}$ to $\\approx 4 \\times 10^{22}$ \\scm\\ cf. the atomic hydrogen column density of $N_{\\rm HI} \\la 10^{21}$ \\scm.  Furthermore, the CO $2\\rr1/1\\rr0$ intensity ratio of $\\approx$1 previously observed over the central beam is also seen over the whole of this region.  \\end{enumerate}  \\item From the rotational dynamics of the gas, the steepest gradient of enclosed mass, $dM/dr$, occurs from $\\approx$10 -- 400~pc, although there is a low abundance of gas in this region (again, $M_{\\rm gas}\\la0.1 M_{\\rm dyn}$), probably due to  consumption by the starburst. At $\\ga400$~pc, which is approximately coincident with the molecular ring ($M_{\\rm gas}/M_{\\rm dyn}\\approx0.5$), there is a steep drop in velocity. This continues to $\\approx1$~kpc, the extension of the atomic ring (\\citealt{jkeh98}). Beyond these radii, velocities climb once more, until settling into solid body rotation at $\\approx150$ \\kms. This continues to the 50~kpc limit of our tilted ring model, although most of the blue, visible and infrared radiation emanates from the central 10~kpc, suggesting that the dynamics beyond this region are halo dominated.  \\end{enumerate}"
    },
    "0807/0807.1143_arXiv.txt": {
        "abstract": "We present phase-resolved low resolution infrared spectra of AM Her and ST LMi, two low-field polars that we observed with SPEX on the IRTF. Optical/NIR lightcurves are also published to help constrain the viewing geometry and brightness of the objects at the time they were observed. Currently, only limited IR spectra have been published for these objects, and none with the phase-coverage presented here. In both cases, the resulting spectra are dominated by emission from the secondary star in the NIR. However, the emission regions are also self-eclipsed, allowing us to isolate the cyclotron emission through subtraction of the dim-phase spectrum. We use a ``Constant Lambda'' prescription to model the changing cyclotron features seen in the resulting data. For AM Her, we find a best fit model of: B = 13.6 MG, kT = 4.0 keV, and log$\\Lambda$ = 5.0. The cyclotron derived accretion geometry is consistent with $i$ = 50$^{\\circ}$ and $\\beta$ = 85$^{\\circ}$. For ST LMi, B = 12.1 MG, kT = 3.3 keV, and log$\\Lambda$ = 5.7 with $i$ = 55$^{\\circ}$ and $\\beta$ = 128$^{\\circ}$. ",
        "introduction": "Polars are interacting binary systems containing a primary white dwarf (WD) and a late type secondary star. Material flows from the secondary, through the L$_{1}$ point and falls ballistically toward the WD. The WDs in polars are highly magnetized with magnetic field strengths that range between 10 - 240 MG. Eventually, the accreting material couples to the field lines of the WD and is transported to the magnetic pole(s) of the star where a dense, standing-shock is formed, with nominal temperatures of 2 - 20 keV, which cools by emitting bremsstrahlung and cyclotron radiation .  AM Her is the prototype polar (see Tapia, 1977a). It has an orbital period of 3.094 hrs and is nearby (78 pc; Thorensten et al., 2003). Despite being extensively studied many characteristics of the system remain uncertain. One point of discrepancy is the exact mass of the primary WD with estimates ranging from 0.39 M$_{\\odot}$ (Young et al., 1981) to 1.22 M$_{\\odot}$ (Cropper et al., 1998). Additionally, the geometry of the system remains unclear with orbital inclination estimates ranging from $i$ = 35$^{\\circ}$ (Brainerd \\& Lamb, 1985) to $i$ = 60$^{\\circ}$ - 80$^{\\circ}$ (Watson et al., 2003), although the self-eclipse observed in the X-Ray and UV requires that $i$ + $\\beta$ $\\geq$ 90$^{\\circ}$, where $\\beta$ is the magnetic co-latitude. Indeed, both Sirk $\\&$ Howell (1998) and Gansicke et al. (1998) found that $i$ + $\\beta$ = 105$^{\\circ}$. Later, Gansicke et al.(2001) determined that combinations of $i$ and $\\beta$ between ($i$ = 50$^{\\circ}$, $\\beta$ = 55$^{\\circ}$) and ($i$ = 35$^{\\circ}$, $\\beta$ = 70$^{\\circ}$) best modeled the high-state optical lightcurves. In other ways, however, AM Her is well characterized. The temperature of the primary WD has been well constrained. Gansicke et al. (2006) modeled low-state FUSE and STIS spectra from AM Her with TLUSTY/SYNSPEC (Hubeny \\& Lanz, 1995), finding that T$_{wd}$ = 19800 $\\pm$ 700 K. Also, the secondary star is spectroscopically determined to be M4 $\\pm$ 1 (Kafka et al., 2005b), although there is some evidence that it is irradiated by the WD and thus the spectral type of the secondary is orbitally modulated (Davey \\& Smith, 1992). The published photometry of AM Her is exhaustive. Orbital lightcurves show variability in every pass-band from the UV (Gansicke et al. 1998) out to the $K$-band (see below). Extensive AAVSO and automatic photometric telescope, ``APT'' (e.g. RoboScope) monitoring from 1990 - 2004, has revealed that AM Her is usually in one of two states: a ``high-state'', with large intrinsic variability: 13.0 $\\leq V \\leq$ 14.0, and a ``low-state'' with $V$ $\\simeq$ 15.5 (Kafka et al., 2005a). Finally, Bailey, Ferrario, \\& Wickramasinghe (1991; henceforth BFW91) used a Constant Lambda (``CL'') code to model the NIR low-state (V $\\sim$ 15.0) cyclotron spectrum, which were binned into bright-phase ($\\phi$ = 0.46 - 0.88) and dim-phase ($\\phi$ = 0.46 - 0.88) spectra. The dim-phase showed only emission from the secondary star, while strong cyclotron emission was observed for the duration of the bright phase. After subtracting a 3250 K model atmosphere with log $g$ = 4.75 to mimic the secondary spectrum, BFW91 found B $\\simeq$ 14.5 MG, and a shock temperature of kT = 8.5 keV.  ST LMi (= CW1103 +254) is a short period polar (P$_{orb}$ = 114 min) containing a 0.7 M$_{\\odot}$ primary (Ramsay et al. 2004) with a likely temperature of 11000 K (Araujo-Betancor et al., 2005; Sion, 1999) and a M5 - M6 secondary (Knigge, 2006; Harrison et al., 2005; Howell et al., 2000; Warner, 1995) at a distance of 115 - 138 pc (Araujo-Betancor et al., 2005; Kafka et al., 2007). It was classified as an AM Her object by Stockman et al. (1983) on the basis of its highly variable polarization. For 70 \\% of the orbit the object shows no significant polarization. Subsequently, a strong pulse is observed peaking near $\\phi$ = 0.00 at 12 \\% and -20 \\% in linear and circular light, respectively (Cropper et al., 1986). This observed bi-modality is echoed in optical and IR orbital lightcurves, which show a quiescent ``dim-phase'' for most of the orbit in each band that is followed by a significant jump in brightness (the ``bright-phase'') coincident to the peak in polarization. Peacock et al. (1992) obtained multi-band photometry of ST LMi, showing that $\\Delta$m $\\simeq$  0.6, 1.4, 2.0, 1.5, and 1.3 for the $BRIJH$ bands, respectively. Long-term $V$-band lightcurves were obtained with RoboScope from 1990 to 2003 (Kafka et al., 2005a). From 1992 - 1997, the system was in a protracted ``low-state'' with $<$$V$$>$ = 17.5 $\\pm$ 0.2. From 1997 - 2003, a more variable, slightly higher state was observed with $<$$V$$>$ = 16.0 $\\pm$ 1.5. Additionally, instances of ``extreme low-states'' have been observed. In Ciardi et al. (1998) the $K$-band spectra of ST LMi, showed no obvious emission lines and were modeled successfully with a $\\simeq$ 3000 K atmosphere suggesting that accretion had almost completely shut off. Kafka et al. (2007) has presented photometry of a similar extreme low-state showing that the system can be as faint as $V$ = 18.5. $JHK$ cyclotron spectra were previously modeled by Ferrario, Bailey \\& Wickramasinghe (1993; henceforth FBW93) finding that two spots were necessary to fully model their spectra: a primary region with B = 12.0 and kT = 12 keV, and a secondary, ``cool-spot'' with  kT = 5.0 keV. The accretion geometry of the primary emission region has also been previously determined, with 55$^{\\circ}$ $\\leq i \\leq$ 64$^{\\circ}$ and 140$^{\\circ}$ $\\leq \\beta \\leq$ 150$^{\\circ}$ (Schmidt et al., 1983; Potter, 2000).  Limited IR spectroscopy exists for these two low-field polars.  Below, we present and model new phase-resolved low-state infrared spectra as well as $JHK$ lightcurves for AM Her and ST LMi. In both instances, we show that variable cyclotron emission over the orbit is responsible for the spectroscopic and photometric behavior. Additionally, we present a second epoch data set for ST LMi that shows no cyclotron emission and must be in an extreme low-state similar to that seen by Kafka et al. (2007). In the next section, we describe the observations of each object, in section 3 we fit these data with cyclotron models, discuss our results in section 4, and draw our conclusions in section 5.  \\section {Observations}  AM Her and ST LMi were observed using SPEX (c.f. Rayner et al., 2003) on the Infrared Telescope Facility (IRTF). AM Her was observed once on 2005 September 1, whereas ST LMi was observed on two different epochs: 2005 Feb 7 and 2006 Feb 2. Both AM Her and ST LMi were found to be in low-states, although as we discuss below, the 2005 Feb 7 data found ST LMi in an extreme low-state. SPEX was used in low-resolution ``prism'' mode with a 0.3'' x 15'' slit. To remove background, each object was nodded along the slit. In its low-resolution mode SPEX produces R(=$\\lambda$/$\\Delta \\lambda$) $\\sim$ 250 spectra, with short enough exposure times to obtain phase resolved spectra of polars with $K \\leq 16.0$.  For ST LMi, we used 240 second exposure times, where shorter, 120 s, integration times were adequate for AM Her. Each of these spectra were then median combined with 2-3 other spectra to allow for cosmic-ray removal and to improve the S/N ratio. The spectra were reduced using the SPEXTOOL package (Vacca et al., 2003). A telluric correction was applied using an A0V star of similar airmass to our program objects. We use the Kafka et al. (2005b) ephemeris to phase all observations of AM Her. For ST LMi because of the large phase uncertainty ($\\Delta \\phi \\simeq$ 0.10) in the Howell et al. (2000) ephemeris, we phased our observations to the $J$-band minimum found in the photometry presented in this paper, which worked out to a phase-shift of $\\Delta \\phi$ = 0.15 from that ephemeris.  Because of the narrow slit size on IRTF/SPEX (0.3'') infrared photometry is required to calibrate the fluxes of the spectra. The $JHK$ photometry for each object was obtained with SQIID on the KPNO 2.1-m telescope (Ellis et al., 1992). AM Her was observed on 2002 September 26, and ST LMi on 2003 April 9. In addition, we obtained simultaneous $BVRIJHK$ photometry on 2005 May 20 for AM Her. The $JHK$ photometry was obtained with NIC-FPS\\footnote{see http://www.apo.nmsu.edu/arc35m/Instruments/NICFPS/nicfpsusersguide.html} on the Apache Point 3.5-m, while the optical data set was obtained with the NMSU 1-m (see Harrison et al., 2003). To aid the reader, we have collated all the observational specifics in Table 1. ",
        "conclusions": "We obtained a full orbit of phase-resolved IR spectra for both AM Her and ST LMi. We found both objects to be dominated by emission from the secondary in the IR. To remove this component, we utilized the fact that emission regions for both stars were self-eclipsed. Thus, at each phase we subtract a dim-phase or ``Stream-Emission'' spectrum. Because of the changing spectral type of the secondary, we found that medianing dim-phase spectra over $\\simeq$ 25 \\% of an orbit produced a better subtraction. For AM Her, we found a phase averaged model of :  B = 13.6$_{-0.8}^{+1.0}$ MG, kT = 4.0$_{-1.0}^{+1.5}$ keV (see Table 2 for specifics at each phase). Additionally, we found that the viewing angle varied in a manner consistent with expectations from a system with $i$ = 50$^{\\circ}$ and $\\beta$ = 85$^{\\circ}$. For ST LMi, we collected two datasets. The first had ST LMi in a low-state with $V$ $\\simeq$ 17.4, and displayed weak cyclotron harmonics that were difficult to decouple from the water vapor signatures leftover after SE-subtraction. We found a phase averaged model of :  B= 12.1 $\\pm$ 0.5 MG, kT = 3.3 $\\pm$ 1.8 keV in an accretion region consistent with $i$ = 55$^{\\circ}$ and $\\beta$ = 128$^{\\circ}$. For ST LMi, we include a second data set, taken when the object was in an ``extreme low-state'' showing no substantial cyclotron emission. The non-detection of cyclotron emission contrasts with our earlier data from both EQ Cet and MQ Dra both of which show cyclotron emission, even while being in extremely low-states."
    },
    "0807/0807.0659_arXiv.txt": {
        "abstract": "For neutrinos streaming from a supernova (SN) core, dense matter suppresses self-induced flavor transformations if the electron density $n_e$ significantly exceeds the neutrino density $n_\\nu$ in the conversion region. If $n_e$ is comparable to $n_\\nu$ one finds multi-angle decoherence, whereas the standard self-induced transformation behavior requires that in the transformation region $n_\\nu$ is safely above~$n_e$. This condition need not be satisfied in the early phase after SN core bounce. Our new multi-angle effect is a subtle consequence of neutrinos traveling on different trajectories when streaming from a source that is not point--like. ",
        "introduction": "\\label{sec:introduction}  The neutrinos streaming from a supernova (SN) core or from the accretion torus of coalescing neutron stars are so dense that the neutrino-neutrino interaction causes collective flavor transformations~\\cite{Pantaleone:1992eq, Samuel:1993uw, Qian:1995ua, Pastor:2002we, Sawyer:2005jk, Fuller:2005ae, Duan:2005cp, Duan:2006an, Hannestad:2006nj, Balantekin:2006tg, Duan:2007mv, Raffelt:2007yz, EstebanPretel:2007ec, Raffelt:2007cb, Raffelt:2007xt, Duan:2007fw, Fogli:2007bk, Duan:2007bt, Duan:2007sh, Dasgupta:2008cd, EstebanPretel:2007yq, Dasgupta07, Duan:2008za, Dasgupta:2008my, Sawyer:2008zs, Duan:2008eb, Chakraborty:2008zp, Dasgupta:2008cu}.  At the same time, the density of ordinary matter is also very large, suppressing normal flavor conversions unless the neutrinos encounter an MSW resonance~\\cite{Dighe:1999bi,Schirato:2002tg,Fogli:2003dw,Tomas:2004gr}. On the other hand, one of the many surprises of self-induced flavor transformations has been that dense matter barely affects them. They are driven by an instability in flavor space that is insensitive to matter because it affects all neutrino and antineutrino modes in the same way. Therefore, it can be transformed away by going to a rotating frame in flavor space~\\cite{Duan:2005cp, Duan:2006an} in a sense to be explained below. Various numerical simulations confirm this insight.  We here clarify, however, that the matter density can not be arbitrarily large before it affects collective flavor conversions after all. The matter term is ``achromatic'' only if we consider the time--evolution of a homogeneous (but not necessarily isotropic) neutrino ensemble on a homogeneous and isotropic matter background. If the matter background is not isotropic, the current-current nature of the neutrino-electron interaction already implies that different neutrino modes experience a different matter effect.  It is more subtle that even without a current, matter still affects different neutrino modes differently if we study neutrinos streaming from a source. The relevant evolution is now the flavor variation of a stationary neutrino flux as a function of distance. For a spherically symmetric situation, ``distance from the source'' is uniquely given by the radial coordinate $r$. Neutrinos reaching a certain $r$ have travelled different distances on their trajectories if they were emitted with different angles relative to the radial direction. Therefore, at $r$ they have accrued different oscillation phases even if they have the same vacuum oscillation frequency and even if they have experienced the same matter background. In other words, if we project the flavor evolution of different angular modes on the radial direction, they have different effective vacuum oscillation frequencies even if they have the same energy. The same argument applies to matter that modifies the oscillation frequency in the same way along each trajectory, but therefore acts differently when expressed as an effective oscillation frequency along the radial direction.  The neutrino-neutrino interaction, when it is sufficiently strong, forces different modes to reach a certain $r$ with the same oscillation phase. To achieve this ``self-main\\-tained coherence'' the neutrino-neutrino term must overcome the phase dispersion that would otherwise occur. Such a dispersion is caused not only by a spectrum of energies, but also by a matter background.  To clarify and quantify these ideas we revisit in Sec.~\\ref{sec:homogeneous} the time evolution of a homogeneous ensemble and introduce the notion of the ``rotation-averaged equations of motion.'' In Sec.~\\ref{sec:stream} we discuss the geometric modifications when we study neutrinos streaming from a source and quantify the multi-angle effect caused by ordinary matter. We conclude in Sec.~\\ref{sec:discussion} with a discussion of the implications of our findings for realistic situations relevant for supernova neutrinos. ",
        "conclusions": "\\label{sec:discussion}  We have identified a new multi-angle effect in collective neutrino transformations that is caused by a matter background. Previous numerical studies of multi-angle effects had used a matter profile that satisfies the condition Eq.~(\\ref{eq:condition}) in the critical region~\\cite{Duan:2006an, Fogli:2007bk}. In other multi-angle studies matter was entirely ignored~\\cite{EstebanPretel:2007ec} and otherwise, single-angle studies were performed. Therefore, the multi-angle matter effect discussed here had escaped numerical detection.  In many practical cases relevant for SN physics or in coalescing neutron stars, the density of matter is probably small enough so that this effect can be ignored. On the other hand, for iron-core SNe, during the accretion phase the matter density can reach values and can have such a profile that it is important. The possibility of such large matter densities had led some of us to speculate that the second-order difference between the $\\nu_\\mu$ and $\\nu_\\tau$ refractive effect could sometimes play an interesting role for three-flavor collective transformations~\\cite{EstebanPretel:2007yq}. It is easy to show, however, that the density requirement for this mu--tau effect to be important implies that the multi-angle matter effect can not be avoided. In this sense, one complicated effect caused by a large matter density annihilates another one.  If at early times the matter density profile is such that our multi-angle effect is important, this will not be the case at later times when the explosion has occurred and the matter profile contracts toward the neutron star. In principle, therefore, interesting time-dependent features in the oscillation probability can occur.  A large matter effect can be ``rotated away'' from the EOMs when it is identical for all modes. Here we have seen that even a perfectly uniform medium provides a multi-angle variation of the matter effect. We note that the matter fluxes would not be important, in contrast to our first example of a homogeneous ensemble, because the relevant quantity is the spread of the matter effect between different modes. Therefore, whenever a flux term would be important, the matter density term already provides a strong multi-angle effect.  Possible polarizations of the electron-positron background in the strong SN magnetic fields would produce an additional axial term in the matter potential, proportional to the scalar product of the neutrino momentum with the electron polarization direction~\\cite{Nunokawa:1997dp}.  However, we expect that also this effect should be negligible with respect to the matter density term.  In addition, the medium can have density variations caused by convection and turbulence~\\cite{Scheck:2007gw} that is known to affect the MSW resonance under certain circumstances~\\cite{Loreti:1995ae, Fogli:2006xy, Friedland:2006ta, Choubey:2007ga, Kneller:2007kg}. Density variations in the transverse direction to the neutrino stream lines may well cause important variations of the matter effect between different modes. It remains to be investigated in which way collective flavor transformations are affected."
    },
    "0807/0807.0235_arXiv.txt": {
        "abstract": "The increasing number of transiting planets raises the possibility of finding changes in their transit time, duration and depth that could be indicative of further planets in the system. Experience from eclipsing binaries indeed shows that such changes may be expected. A first obvious candidate to look for a perturbing planet is GJ 436, which hosts a hot transiting Neptune-mass planet in an eccentric orbit. Ribas et al. (2008) suggested that such eccentricity and a possible change in the orbital inclination might be due to a perturbing small planet in a close-in orbit. A radial velocity signal of a 5~M$_\\oplus$ planet close to the 2:1 mean-motion resonance seemed to provide the perfect candidate. Recent new radial velocities have deemed such signal spurious. Here we put all the available information in context and we evaluate the possibility of a small perturber to GJ~436~b to explain its eccentricity and possible inclination change. In particular, we discuss the constraints provided by the transit time variation data. We conclude that, given the current data, the close-in perturber scenario still offers a plausible explanation to the observed orbital and physical properties of GJ~436~b. ",
        "introduction": "The study of eclipsing binary systems was initiated by the discovery by John Goodricke in 1783 and subsequent proof by Edward C. Pickering in 1881 that the variable star Algol is indeed composed of two stars undergoing eclipses. In the last century, the study of eclipsing binary systems has been a major component to Stellar Astrophysics as a whole, by improving our understanding of a variety of phenomena in binaries, by providing valuable information on the structure and evolution of stars, and by serving as indicators of, e.g., age or distance (see, e.g., Andersen 1991; Guinan 1993; Hilditch 2001; Ribas 2006; Bonanos 2007). Eclipsing binaries still continue to play an important role in this respect. However, in recent years we have witnessed a rebirth of the study of eclipsing binary systems in the particular case of very unequal mass and brightness ratios, i.e., transiting planets. Being eclipsing binaries a field of long tradition (both observational and theoretical) it is convenient to use such outstanding background and adapt it to suit current needs.  A particular aspect of eclipsing binary research has been the study of detached systems that show variable light curves. In some cases, the origin of the variability can be traced to the appearance, migration, and disappearance of inhomogeneities (starspots) on the surface of one or both components but in some instances the observed changes are of a more fundamental nature. For example, the eclipsing binaries SS~Lac and SV~Gem stopped eclipsing some 60 years ago (Torres \\& Stefanik 2000; Guilbault et al. 2001), V906~Sco stopped eclipsing in 1918, then restarted in 1963 and stopped again in 1986 (Lacy et al. 1999), and IU~Aur shows a fast variation of eclipse depth (Drechsel et al. 1994). All these changes are thought to originate from the perturbations of a third star in the system (S\\\"oderhjelm 1975; Mazeh \\& Shaham 1976). Obviously, such large perturbations are not very common among eclipsing binaries but it is also true that not many systems have been observed intensively enough and with a sufficiently long time baseline to uncover slow light curve variations.  In the case of transiting planets, variations in light curve properties have been predicted to occur from various sources, most notably from the effect of perturbing further bodies in the system. Transit time variations have been the subject of intense attention and indeed been proposed as a way to detect smaller perturbing planets (Miralda-Escud\\'e 2002; Holman \\& Murray 2005; Agol et al. 2005). But not only the transit central time may suffer variations. The duration and depth of the transit (Schneider 1994; Miralda-Escud\\'e 2002; Laughlin et al. 2005) may also be modified because of changes in the orbital inclination, semi-major axis, eccentricity and argument of periastron.  The numerous ongoing surveys from both the ground and space are now producing new transiting planet discoveries at an ever increasing pace. The total tally of exoplanets undergoing transit events has already surpassed 50. With these increasing statistics and the eclipsing binary experience one may wonder if changes in the light curve are already observable in some cases. ",
        "conclusions": "New radial velocity data have ruled out the presence of a 5~M$_\\oplus$ planet with a period of 5.18 d in GJ~436. Although the new data set a more stringent limit to possible further planet in GJ~436, the scenario of a close-in perturber to GJ~436~b is still plausible. Strong proof should come from changes in the transit duration measured over the coming seasons. If such changes are there, this will nail the case for a perturbing object in a slightly non-coplanar orbit, much in the same way as for some eclipsing stellar binaries. This will give rise to the interesting concept of ``transient transits''. Still, it is possible that the duration is stable over time and this would rule out a perturber in mild non-coplanarity.  The most certain observational fact is the eccentricity of GJ~436~b's orbit. Besides the close-in perturber scenario, there are a number of other possible explanations. A distant perturber was proposed by M07 and also by Bonfils (this volume). This would explain the radial velocity trend and be responsible for the eccentricity pumping. However, the trend seems not to be confirmed by newer data and, further, the effect of GR precession may prevent the building of significant eccentricity. This is because the GR timescale for GJ~436~b is only 15\\,000 yr and any eccentricity pumping effect with a longer timescale will not be efficient. The other obvious scenario is that of a large value of the tidal dissipation constant $Q'_{\\rm p}$ for GJ~436~b. The value needed is 10$^{6-7}$, which is one to two orders of magnitudes larger than Neptune's and larger than that of any object in the Solar System. A large value of $Q'_{\\rm p}$, if it can be generalized, should also be made compatible with the distribution of eccentricities of close-in planets. Further, there is weak evidence of tidal heating on GJ~436~b (Deming et al. 2007), which advocates for a normal $Q'_{\\rm p}$. Finally, one may also think of a recent event (less than 100 Myr ago) that pumped up the eccentricity of the planet (as suggested by Zakamska \\& Tremaine 2004, in another context), such as the close passing of a star or massive object. However, at this point this is just mere speculation.  We believe that a close-in perturber is still the most likely scenario to explain the observations of GJ~436. Some of the small planets found in around mass stars seem to belong to multiple systems, such as Gl~876, Gl~581, or the recent discovery of HD 40307. Thus, given this remarkable parallelism, it would not be a surprise if GJ~436 hosts more than one planet. This would also be the case in the framework of the hypothesis of packed planetary systems of Barnes \\& Raymond (2004). As more transiting planets are discovered, the chances of observing variations in their transit properties with time increase. The case of near-grazing events is especially suitable because of their sensitivity to perturbers. In the coming years, this technique combined by intensive studies of transiting planets (ensured by the interest in the field) should provide us with new insight into the architecture of planetary systems."
    },
    "0807/0807.2230_arXiv.txt": {
        "abstract": "{} {To study the interaction between young AGN and their host galaxies based on their ionized gas and radio emission, and to analyze possible implications for the radio galaxy evolution.} {The \\otres line and 5-GHz radio properties are compared and studied on a large, representative sample of GPS and CSS (i.e., young) quasars and radio galaxies as well as large-scale sources using \\otres line and 5-GHz radio data from literature and our observations. } {Several correlations between the \\otres line and 5-GHz radio emission have been found. The main result is that the \\otres emission is strongly related to the GPS/CSS source size indicating that the \\otres emission is clearly enhanced by the jet expansion through the host galaxy ISM. Shocks are the most likely enhancing mechanism, although jet-induced star formation could also be, partly, responsible for the \\otres emission. The data also suggests a possible deceleration of the jet as it grows. In this case, however, the correlation is weak.} {} ",
        "introduction": "We still know little about how radio galaxies are born and how they subsequently evolve but it is generally accepted that the GHz Peaked Spectrum (GPS) and Compact Steep Spectrum (CSS) radio sources are young, smaller versions of the large-scale powerful radio sources \\citep[with a few exceptions, see e.g., ][]{Stanghellini05, Marecki06}.  The GPS and CSS sources are powerful but compact radio sources whose spectra are generally simple and convex with peaks near 1 GHz and 100 MHz respectively. The GPS sources are contained within the  extent of the optical narrow emission line region ($\\lesssim 1$ kpc) while the CSS sources are contained within the host galaxy \\citep[$\\lesssim 15$ kpc, see ][ for a review]{O'Dea98}.  Current models of the evolution of powerful radio galaxies suggest that these sources propagate from the $\\sim 10$ pc to Mpc scales at roughly constant velocity through an ambient medium whose density declines as $\\rho(R) \\propto R^{-2}$ while the sources decline in radio luminosity as $L_{rad} \\propto R^{-0.5}$ \\citep[e.g., ][ and references therein]{O'Dea02}. In this scenario, GPS sources would evolve into CSS sources and these into supergalactic-size sources\\footnote{Recently, the High Frequency Peakers have been added to the sequence, as possible progenitors of GPS sources \\citep[e.g., ][ and references therein]{Orienti07}.}. Such a scenario is consistent with the observed number densities of powerful radio sources as a function of linear size \\citep[e.g.][]{O'Dea97,Fanti01}. However, to match observations, the radio jets of the young sources must slow down as they cross the host galaxy ISM and dim faster than predicted. The most likely mechanism to produce these effects is interaction of the radio source with the host environment \\citep[see e.g.,][]{Young93,Carvalho94,Carvalho98}.  The characteristics (size, radio power and young age) of GPS and CSS sources make them excellent probes of interaction (and evolution). Furthermore, they have not completely broken through the ISM, so these interactions are expected to be more important than in the larger sources. Observations of UV, \\ion{H}{i} and, especially, of the ionized gas in GPS and CSS sources suggest the presence of such interactions \\citep[e.g.,][]{Labiano08, Holt06, Labiano05, Axon00, Vries99, Vries97}.   Although evidence of interactions was found, previous investigations dealt with small samples or even just a few sources and, until now, there was no study of interaction of a large representative sample of GPS and CSS sources. Moreover, the sparseness of samples used did not allow for a general study of the consequences such interactions.  In order to examine the interaction between gas clouds and the radio source in a statistically significant manner, I collected a sample of almost one hundred sources, including GPS and CSS galaxies and quasars, as well as large-scale sources. I studied the properties of the \\otres line and 5-GHz radio emission aiming to finding traces of interaction and the mechanisms responsible for it.  All calculations were made with H$_0$=71,  $\\Omega_{\\mathrm{M}}$ = 0.27, $\\Omega_\\Lambda = 0.73$, \\citep{Spergel03} ",
        "conclusions": "\\label{sec:results}  For the objects in the sample, the following properties of the \\otres line and 5-GHz radio emission were compared: FWHM, luminosity, asymmetry, kurtosis (\\otres line), as well as power, size and turnover frequency (radio emission).  The \\otres FWHM shows no correlation with radio power, turnover frequency, and \\otres luminosity. Therefore, the shock velocity \\citep[closely related to the FWHM, see e.g.,][]{Bicknell97} is independent of the radio source strength, and does not affect the luminosity of the ionized gas or the radio spectral properties of the source. However, a correlation between \\otres FWHM and radio source size (Fig. \\ref{FWHMLS}) could be present, suggesting a possible deceleration of the jet as it crosses the host galaxy:   \\begin{center} log FWHM $\\simeq  2.89(\\pm0.04) - 0.08(\\pm0.05) \\times$ log LS \\end{center}  This correlation is not very strong but could be real given that the deceleration of the jet is required to explain observations \\citep[e.g.,][]{O'Dea98} and the latest models predict deceleration in the jet (see e.g., Kawakatu \\& Kuno 2008 in prep.). Unfortunately, there is not much data available on \\otres FWHM to confirm any of such models.  \\begin{figure} \\centering \\includegraphics[width=\\columnwidth]{O3FWHM_LS_ref.eps} \\caption{Plot of the \\otres FWHM of the sources in \\cite{Gelderman94} (diamonds), \\cite{Vries00b} (triangles) our Kitt Peak (circles) and STIS (shaded diamonds) observations, and a linear fit to the data. The ionization and kinematics of \\object{3C~67}, \\object{3C~277.1}, and \\object{3C~303.1} were studied by \\cite{Labiano05} and \\cite{O'Dea02}.  \\label{FWHMLS}} \\end{figure}   \\begin{figure} \\centering \\includegraphics[width=\\columnwidth]{RP_LS_ref.eps} \\caption{Radio power at 5 GHz vs. linear size for GPS, CSS and large-scale radio sources. \\label{RPLS}} \\end{figure}  As expected \\citep[see e.g.][Fig. \\ref{RPLS}]{O'Dea97}, the sample also shows no correlation between radio power and radio size, showing that GPS, CSS and FR2\\footnote{FRx stands for Fanaroff-Riley class x, \\citep{Fanaroff74}.} sources are equally powerful (log P$_{\\mathrm{5 GHz}}\\sim10^{26-27}$ W/Hz) while FR1 sources tend to be fainter \\citep[see e.g.,][]{Baum95, Zirbel95}. However, for giant sources ($\\gtrsim 1$Mpc), the radio power seems to decrease with size \\citep[e.g.,][]{Ishwara99}.  It is also clear that quasars tend to be brighter in \\otres than radio galaxies, consistent with the unification scenario: some of the \\otres may be hidden by the torus in radio galaxies \\citep[e.g.,][]{Hes93}. Also, the difference in \\otres emission could be due to selection effects: quasars are usually found at higher redshifts. On the other hand, the sample could be missing fainter quasars with luminosities similar to radio galaxies.  \\begin{figure} \\centering \\includegraphics[width=\\columnwidth]{O3_RP_ref.eps} \\caption{Plot of the \\otres luminosity vs. radio power at 5 GHz. Data for the CSS and GPS sources from \\cite{Gelderman94,O'Dea98, Vries00b}, the 2-Jy sample and our Kitt Peak observations. Data for the large-scale radio sources from the 2-Jy sample. The IDs of the sources have been updated with \\cite{Zirbel95}. I use \"C/G\" to name those sources with no clear ID as CSS or GPS. \\label{O3RP}} \\end{figure}  \\begin{figure} \\centering \\includegraphics[width=\\columnwidth]{O3_RP_Zoom_ref.eps} \\caption{As Fig. \\ref{O3RP}, showing only the GPS and CSS sources. \\label{O3RPzoom}} \\end{figure}  The sample shows the known relation \\citep[e.g.,][]{Baum89,Rawlings91} between \\otres luminosity and radio power (Figs. \\ref{O3RP} and \\ref{O3RPzoom}) for powerful (log P$_{\\mathrm{5 GHz}}\\gtrsim10^{25}$ W/Hz) radio sources. This relation is usually explained as the AGN powering both the ionized gas and radio emission. The correlations for different groups of object in our sample are:  \\begin{flushleft} GPS: \\end{flushleft} \\begin{center}  log L$_{[\\ion{O}{iii}]}$ = 32($\\pm$4) + 0.4($\\pm$0.1) $\\times$ log P$_{\\mathrm{5 GHz}}$ \\end{center} CSS: \\begin{center}  log L$_{[\\ion{O}{iii}]}$ = 22($\\pm$4) + 0.8($\\pm$0.2) $\\times$ log P$_{\\mathrm{5 GHz}}$\\end{center} GPS + CSS: \\begin{center} log L$_{[\\ion{O}{iii}]}$ = 25($\\pm$3) + 0.7($\\pm$0.1) $\\times$ log P$_{\\mathrm{5 GHz}}$ \\end{center} GPS + CSS + FR: \\begin{center} log L$_{[\\ion{O}{iii}]}$ = 12($\\pm$2) + 1.13($\\pm$0.07) $\\times$ log P$_{\\mathrm{5 GHz}}$ \\end{center}   It should be noted that when combining NVSS and SDSS data (Best 2009, in preparation), this correlation is not present for fainter (log P$_{\\mathrm{1.4 GHz}} \\lesssim 24.5$ W/Hz) radio sources, suggesting that a different mechanism may be producing the radio and \\otres emission in these sources.  There is no evident relation between kurtosis or asymmetry of the \\otres line \\citep[as defined by][]{Whittle85, Heckman81} with other properties of the line. Most of the sources have kurtosis lower than 1, suggesting the presence of broad wings in the gas. \\cite{Whittle85} also finds the same for his sample of Seyfert galaxies. Most sources show asymmetry values close to zero in their \\otres profile, suggesting the cocoon widens or narrows in a symmetric way  or there are no major variations between both sides. However, the ground spectra may lack sufficient resolution to discern more complex structures that could change the profile of the line.  \\citet{O'Dea98} discovered that in the \\citet{Gelderman94} sample, GPS galaxies tend to have lower \\otres luminosity than CSS galaxies but it was not clear if the trend would be followed by a larger sample with different selection criteria, which was the case of quasars and large FR sources. Also, the implications of the existence of that trend on the radio source (and host galaxy) evolution are not clear\\footnote{A direct consequence is that it is more difficult to find optical counterparts of GPS sources than CSS sources.}. To asses these issues, I looked for this trend in the sample presented here, which included not only GPS and CSS radio galaxies but also quasars and large-scale sources. I found that the GPS and CSS sources (galaxies and quasars) showed a strong correlation between \\otres luminosity and size of the radio source\\footnote{Note that the trend is also suggested by Fig. \\ref{O3RPzoom} and the different correlations between radio power and \\otres luminosity for GPS and CSS sources.}:   \\begin{flushleft} GPS + CSS: \\end{flushleft} \\begin{center} log L$_{[\\ion{O}{iii}]}$ = 42.43($\\pm$0.09) + 0.46($\\pm$0.09) $\\times$ log LS$_{\\mathrm{5 GHz}}$ \\end{center}  \\begin{flushleft} GPS$^*$ + CSS:  \\end{flushleft} \\begin{center} log L$_{[\\ion{O}{iii}]}$ = 42.44($\\pm$0.08) + 0.4($\\pm$0.1) $\\times$ log LS$_{\\mathrm{5 GHz}}$ \\end{center}  GPS$^*$ means the complete sample of GPS without the smallest source (\\object{1718--649}), to test if it is dominating the correlation.  In principle, this correlation could be due to the AGN enhancing both the radio and emission gas luminosities through photons. However, for the same radio power, small sources (GPS) are systematically fainter in \\otres than larger (CSS) sources.  I propose a scenario where the expansion of the radio source through the host ISM is triggering and/or enhancing the \\otres emission through direct interaction. Some contribution from the AGN light must be present but AGN light alone would not produce a correlation with size. Furthermore, the fact that the correlation disappears for larger radio sources ($\\gtrsim 15-20$ kpc) supports this model: once the radio lobes leave the host galaxy, the \\otres luminosity drops (Fig. \\ref{O3LS}).  This scenario is also supported by previous observations providing evidence of strong interaction between the jet and surrounding ISM, as well as proof of shock-ionized \\otres emission \\citep[e.g.,][]{Labiano05, O'Dea03}. Jet--ISM interactions are also found through  \\ion{H}{i} studies \\citep[e.g.,][ and references therein]{Labiano06, Holt06} and predicted by jet expansion models \\citep[e.g.,][]{Jeyakumar05, Saxton05}.  Another interesting enhancing mechanism to consider is that the jet could, at least partly, boost \\otres emission through indirect mechanisms such as jet-induced star formation (Labiano et al. 2008, in prep.) but new data are needed to study the possible contribution of recently formed stars. The hosts of GPS and CSS sources are usually elliptical galaxies so it is unlikely that the average/normal stellar population of the host galaxy has a strong contribution to the \\otres emission. These ``normal'' stars would, however, not create a correlation with radio jet size (and jet-induced stars would).  The correlations are as follows:\\\\  GPS: \\begin{center} log L$_{[\\ion{O}{iii}]}$ = 42.8($\\pm$0.2) + 0.8($\\pm$0.2) $\\times$ log LS$_{\\mathrm{5 GHz}}$ \\end{center}  GPS$^*$: \\begin{center} log L$_{[\\ion{O}{iii}]}$ = 42.5($\\pm$0.4) + 0.5($\\pm$0.4) $\\times$ log LS$_{\\mathrm{5 GHz}}$ \\end{center}  CSS: \\begin{center} log L$_{[\\ion{O}{iii}]}$ = 42.6($\\pm$0.2) + 0.2($\\pm$0.3) $\\times$ log LS$_{\\mathrm{5 GHz}}$\\end{center}  The correlation tends to disappear when the sample is divided into different types of sources. This could be due to low statistics or to the smaller range of sizes covered by each type. GPS could show lower \\otres emission due to high obscuration. However, it is more likely that young compact sources are too small to strongly affect their environment. This effect has also been observed in star formation histories of GPS hosts \\citep{Labiano08}. The UV luminosities of GPS sources seem to be as high as those of CSS sources. Furthermore, the UV luminosity of GPS sources could be correlated with their radio power \\citep{Labiano08}. These two effects suggest that obscuration is not too strong in GPS sources or, at least, similar to that in CSS.   \\begin{figure} \\centering \\includegraphics[width=\\columnwidth]{O3LS_ref.eps} \\caption{Plot of the [\\ion{O}{iii}] luminosity vs. linear size of the radio source, showing the correlation for GPS and CSS sources. Data from the same references as Fig. \\ref{O3RP}. \\label{O3LS}} \\end{figure}  Concerning the overall scenario of radio source evolution, where GPS and CSS sources evolve into the larger FR sources, the visual inspection of Fig. \\ref{O3LS} suggests that CSS sources would evolve into FR2. Some authors also found a possible decreasing trend linking FR2 to FR1 (Best 2008, private communication) with increasing size. However, our sample may lack enough supergalactic-sized sources to thoroughly address evolution beyond $\\sim15-20$ kpc and the FR2 -- FR1 connection is beyond the scope of this letter. Extensive discussions on the FR2 -- FR1 connection can be found in the literature: e.g., \\citet{Muller04, Best05, Wold07} and references therein or classical papers such as \\citet{Baum95, Zirbel95}."
    },
    "0807/0807.1929_arXiv.txt": {
        "abstract": "The massive Galactic black hole and the stars around it are a unique laboratory for studying how relaxation processes lead to close interactions of stars and compact remnants with the central massive black hole, in particular those leading to the emission of gravitational waves. I review new results on the processes of strong mass segregation and loss-cone refilling by massive perturbers and resonant relaxation; describe observational evidence that these processes play a role in the Galactic Center and can be studied there; and discuss some of the implications for Extreme Mass Ration Inspiral event rates and their properties. ",
        "introduction": "The $\\Mbh\\sim4\\times10^{6}\\,\\Mo$ massive black hole (MBH) in the Galactic Center (GC) \\cite{eis+05,ghe+05} is special not only because it is the nearest and observationally most accessible of all MBHs, but also because it is an archetype of the extragalactic MBH targets for the planned Laser Interferometer Space Antenna (LISA) gravitational wave (GW) detector. This is a coincidence: current technology limits LISA's baseline to $5\\times10^{6}\\,\\mathrm{km}$, which happens to coincide with the typical GW wavelength emitted by an object on the last stable circular orbit around a $\\sim10^{7}\\, M_{\\odot}$ MBH. GW from more massive MBH have longer wavelengths, beyond the sensitivity of LISA. As shown below, LISA's focus on the lower range of MBH masses implies that LISA targets will typically lie in dynamically relaxed, high density nuclei. In such systems, relaxation controls the rate at which compact remnants on wide orbits near the MBH are scattered into inspiral orbits ({}``loss-cone'' orbits) and become extreme mass ratio inspiral (EMRI) GW sources. Thus, to understand and predict EMRI rates and orbital properties, it is necessary to understand the various dynamical relaxational processes that occur close to a MBH, and the degree of concentration of compact remnants around it. I describe here several newly (re-)discovered relaxation mechanisms beyond standard 2-body relaxation, which likely play an important role in EMRI dynamics, and discuss how these theoretical ideas are tested by observations of the GC. ",
        "conclusions": "\\label{s:summary}  The empirical $\\Mbh/\\sigma$ relation implies that LISA extragalactic targets, MBHs in the range $\\Mbh\\!\\lesssim\\!10^{7}\\,\\Mo$, lie in very high density relaxed stellar cusps. The Galactic MBH happens to be an archetype of LISA targets. In such systems, it is the slowest relaxation process, non-coherent 2-body relaxation, or possibly other more efficient processes, that determine the central concentration of compact EMRI candidates, and the rate at which they are captured in EMRI orbits. We described several new results on processes that may strongly affect the cosmological EMRI rates, and can be tested by observations of the GC: strong mass segregation, which drives heavy SBHs into unusually high concentrations near the MBH, and may have left an imprint in the strong observed central suppression of light giants around the Galactic MBH; accelerated loss-cone refilling by massive perturbers, which strongly affects the rate at which binaries are tidally separated by the MBH, leaving compact objects very close to the MBH. This mechanism very naturally explains the cluster of young stars around the Galactic MBH and the numbers of hyper-velocity stars observed on their way out of the Galaxy; and RR, which dominates EMRI dynamics and likely enhances their rate. RR can explain some of the systematic differences between the dynamical properties of the various stellar populations observed in the GC, and conversely, these observations can be used to test this process and constrain its properties.  The implications of all these processes on the EMRI rates have still to be fully mapped. However, it is already safe to say that GW astronomy will tell us as much about fundamental questions in stellar dynamics, as it will tell us about cosmology and GR.  \\begin{theacknowledgments} This work was supported by ISF grant 928/06, Minerva grant 8563 and a New Faculty grant by Sir H. Djangoly, CBE, of London, UK. \\end{theacknowledgments}"
    },
    "0807/0807.3601_arXiv.txt": {
        "abstract": "We qualitatively examine properties of artificial deformation in shapes of objects (galaxies and stars) induced by the pixelization effects (also called as the aliasing effects) using toy mock simulation images. Two causes of the effects have been recognized: One is a consequence of observing the continuous sky with discrete pixels which is called as the first pixelization. And the other, called the second pixelization, is a consequence of resampling a pixelized image onto an another pixel grid whose coordinates are not perfectly adjusted to the input grid. We pay a special attention to the latter because it might be a potential source of a systematic noise in a weak lensing analysis. In particular, it is found that resampling with rotation induces artificial ellipticities in object shapes having a periodic concentric-circle-shaped pattern. Our major findings are as follows. (1) Root-mean-square (RMS) of artificial ellipticities in object shapes induced by the first pixelization effect can be as large as RMS$\\gtrsim 10^{-2}$ if a characteristic size of objects (e.g., the FWHM) is smaller than twice of the pixel size. While for larger objects, it quickly becomes very small (RMS$\\lesssim 10^{-5}$). (2) The amplitude of the shape deformation induced by the second pixelization effect depends on the object size. It also depends strongly on an interpolation scheme adopted to carry out resampling and on the grid size of the output pixels. The RMS of ellipticities in object shapes induced by the second pixelization effect can be suppressed to well below $10^{-2~}$ if one adopts a proper interpolation scheme (implemented in popular image processing softwares). We also discuss an impact of the pixelization effects on a weak lensing analysis. ",
        "introduction": "A precise shape measurement is of fundamental importance in astronomical research, not only because the morphology of celestial bodies provides us with their physical information, but also because tiny deformation in shapes of distant galaxies caused by gravitational lensing effect allows us to explore foreground mass distribution (see Fort \\& Mellier 1994; Mellier 1999; Bartelmann \\& Schneider 2001 for reviews).  Analyses of weak gravitational lensing effect (e.g., the cosmic shear correlation) especially require very precise shape measurement, since its amplitude is of a few percents level (Refregier 2003 and references therein). Therefore, any artificial shape deformation which may arise during an observation as well as data reduction must be properly understood and must be controlled down to a sufficiently small level.  One known artificial shape deformation is caused by the pixelization of images, known as {\\it pixelization effects} (also called as the {\\it aliasing effect}). As pointed out by Rhodes et al. (2007), there are two causes of the pixelization effect (see Fig. 6 of Rhodes et al. 2007 for an illustration): One is a consequence of observing the continuous sky with discrete pixels, called as the {\\it first pixelization}, which is an unavoidable effect.  And the other, called as the {\\it second pixelization}, occurs when resampling a pixelized image onto an another pixel grid whose coordinate is not perfectly adjusted to the input grid. During resampling, one input pixel may be resampled onto several output pixels, obviously resulting in deformation in object shapes (see Figure \\ref{fig:illust} for an illustration). The second pixelization occurs several stages of image processing processes involving resampling, for example, the correction for a geometric  distortion, mosaicking images from multiple CCD chips to generate a combined image, and stacking multiple dithered images.  In this paper, we are concerned with the second pixelization effect taking two examples of resampling; rotation and correction for an axially symmetric optical distortion. We especially pay an attention to rotation which must be involved in the mosaic-stacking process of a mosaic CCD camera if multiple CCDs are not installed in perfectly parallel with each other. In the case of the Subaru Prime Focus Camera (Suprime-Cam), rotations of $0.025-0.17$ degrees are necessary for generating a properly mosaicked image. An important point to notice is that resampling with rotation induces artificial ellipticities in object shapes having a concentric-circle-shaped pattern (explained in detail in the following sections and see Figures \\ref{fig:illust} and \\ref{fig:demo} for demonstrations).  Therefore it give rise to artificial shear correlations which can potentially act as a systematic noise in the measurement of cosmic shear correlation functions.   We notice that the artificial shape deformation induced by the second pixelization effect cannot be generally corrected by the anisotropic point spread function (PSF) correction (which is one of the important procedures in weak lensing analyses [Kaiser et al. 1995]) simply because they originate from different causes. The shape deformation by the anisotropic PSF arises during an observation with a real (thus non-perfect) instrument, thus it is an unavoidable effect and one has to develop a reliable correction scheme (e.g., Heymans et al. 2006; Massaey et al. 2007). Whereas the second pixelization effect occurs during image processing, and can be minimized by adopting an optimal resampling scheme. In fact, Rhodes et al (2007) developed such a resampling scheme for the HST ACS data in an empirical manner by searching for optimal parameters (the interpolation kernel and output pixel size) of the image processing software {\\it MultiDrizzle}\\footnote{see MultiDrizzle web page:\\\\ {\\tt http://stsdas.stsci.edu/pydrizzle/multidrizzle/}}.   The purpose of this paper is two-fold: The first is to quantitatively examine the effect of the second pixelization effect to understand its properties. The second is to explore an optimum way to minimizing it.  To do these, we use simple image simulations which is described in \\S 2.  Then in \\S 3, we give some illustrative examples for a visual impression and for demonstrating the origin of the concentric-circle-shaped pattern induced by resampling with rotation.  Results are presented in \\S 4.  Finally, \\S 5 is devoted to a summary and discussion. ",
        "conclusions": "\\label{sec:summary}  We have qualitatively examined ellipticities in object shapes induced by the pixelization effects paying a special attention to the periodic concentric-circle-shaped pattern induced by resampling of pixels with rotation. Our major findings are summarized as follows. \\begin{itemize} \\item Artificial ellipticities induced by the first pixelization effect can be as large as $\\langle e^2 \\rangle^{1/2} \\gtrsim 10^{-2}$ if a characteristic size of objects (e.g., the FWHM) is smaller than twice of the pixel size. Whereas for objects with the characteristic size being larger than three times of the pixel size, the RMS becomes negligibly small ($\\langle e^2 \\rangle^{1/2} \\lesssim 10^{-5}$). \\item The second pixelization effect preferentially induces the $e_1$ component (parallel to the grids). The reason of this is that pixels are square shaped and so the pixelization effect is, in general, not axially symmetric but has some special directions. \\item The size (e.g., RMS of $e$) of the shape deformation caused by the second pixelization effect depends on the object size. It also strongly depends on the interpolation scheme for resampling and on the grid size of the output pixels. If we set an upper limit of the RMS ellipticies by $\\langle e^2 \\rangle^{1/2} < 5 \\times 10^{-3}$ for objects with FWHM$>2.5 \\times l_{\\rm pixel}$  (corresponding to FWHM$>0.5$ arcsec for the case of Suprime-Cam), the interpolation schemes passing the above condition are (see Figure \\ref{fig:thetag}) the 5th order polynomial, Lanczos3, Lanczos4 and sinc kernel (as far as among ones considered in this paper). Adopting the grid refinement makes a great improvement. Actually, if one adopts twice finer grid for output pixels, even the linear polynomial can pass the above condition. \\item Resampling of a pixelized image with rotation induces a periodic concentric-circle-shaped pattern of artificial ellipticities in object shapes. The scale of the pattern is related to the pixel size and the rotation angle, $\\phi$,  by $L_{\\rm pattern} = l_{\\rm pixel} / \\tan\\phi$. \\end{itemize}  Before closing this paper, we would like to make a comment on an impact of the second pixelization effect on the actual weak lensing analysis using Suprime-cam data presented in Miyazaki et al. (2007). Miyazaki et al. (2007) carried out resampling adopting the 3rd order polynomial interpolation scheme and combined typically 4 dithered images\\footnote{Combining dithered images reduces the RMS ellipticities roughly as $\\propto N^{-1/2}$ for $N$ dithered images (Rhodes et al. 2007), because an object falls onto a different sub-pixel position in different exposures as a consequence of dithered exposures which results in different ellipticities with basically random orientations. Combining those images can mitigate the pixelization effects.}, thus for the images with FWHM$\\gtrsim 0.6$ arcsec (the typical PSF size), the RMS of ellipticities induced by the second pixelization effect should be well below $10^{-2}$. Whereas, the RMSs measured from stellar images are about a few $\\times 10^{-2}$, therefore we may safely conclude that the second pixelization effect is suppressed sufficiently, and is not a major source of the artificial ellipticities in object shapes.  \\bigskip  We would like to thank Richard Massey for valuable comments and Nick Kaiser for making the IMCAT software available. We would like to thank the anonymous referee for valuable and constructive comments on the earlier manuscript which improve the paper. This research was supported in part by the Grants--in--Aid from Monbu--Kagakusho and Japan Society of Promotion of Science (15340065 and 17740116). Numerical computations presented in this paper were carried out on computer system at CfCA (Center for Computational Astrophysics) and at ADAC (Astronomical Data Analysis Center) of the National Astronomical Observatory Japan."
    },
    "0807/0807.1097_arXiv.txt": {
        "abstract": "We develop population models of high-mass X-ray binaries (HMXBs) formed after bursts of star formation and we investigate the effect of electron-capture supernovae (ECS) of massive ONeMg white dwarfs and the hypothesis that ECS events are associated with typically low supernova kicks imparted to the nascent neutron stars. We identify an interesting ECS {\\em bump} in the time evolution of HMXB numbers; this bump is caused by significantly increased production of wind-fed HMXBs 20--60\\,Myr post starburst. The amplitude and age extent of the ECS bump depend on the strength of ECS kicks and the mass range of ECS progenitors. We also find that ECS-HMXBs form through a specific evolutionary channel that is expected to lead to binaries with Be donors in wide orbits. These characteristics, along with their sensitivity to ECS properties, provide us with an intriguing opportunity to probe ECS physics and progenitors through studies of starbursts of different ages. Specifically, the case of the Small Magellanic Cloud, with a significant observed population of Be~HMXBs and starburst activity 30--60\\,Myr ago, arises as a promising laboratory for understanding the role of ECS in neutron star formation. ",
        "introduction": "Recently, electron-capture supernovae (ECS) of massive ONeMg white dwarfs have attracted renewed attention as an important mechanism of neutron star (NS) formation. ECS events probably operate at the low end of the NS progenitor mass range, when a strongly degenerate ONeMg core forms. If the progenitor exceeds a critical core mass of 1.38~M$_\\odot$, high pressure and density lead to electron captures onto $^{24}$Mg and $^{20}$Ne. The resulting decrease in electron degeneracy pressure causes a supernova (SN) collapse before an iron core is ever formed \\citep{1980PASJ...32..303M, 1984ApJ...277..791N, 1987ApJ...322..206N}. Numerical simulations indicate that ECS events are under-energetic compared to iron-core supernovae (ICC-SN\\footnote{Here we use the term ICC to refer to both iron-core collapse supernovae, as well as direct collapse black holes (BHs) which form without a SN explosion (i.e, all non-ECS compact object formation events). We use ICC-SN to refer to only those iron-core collapse systems which result in an energetic supernova explosion}) by at least an order of magnitude, However they are, at the same time, relatively fast, neutrino-driven, delayed explosions. The short time available for the propagation of the shock is expected to prevent a strong build-up of core asymmetries, leading to typical kicks of $\\lesssim 100$\\,km\\,s$^{-1}$,  lower than those in ICC-SN events \\citep[][Janka 2007, private communication]{2006ApJ...644.1063D, 2006A&A...450..345K}.  The mass range of ECS progenitors is highly uncertain. While \\citet{1984ApJ...277..791N, 1987ApJ...322..206N} predicted an initial mass range of 8-10~M$_\\odot$ at solar metallicity, recent simulations suggest much stricter ECS boundaries for single stars. For example, \\citet{2008ApJ...675..614P} find a single star mass range of 9-9.25~M$_\\odot$, which accounts for only $\\simeq$4\\% of SNe for a given population. \\citet{2007A&A...476..893S} finds a metallicity-dependent mass range which stays consistently narrow (a width of $\\lesssim$1.5\\,M$_\\odot$), and could vanish altogether depending on the metallicity dependence of the stellar winds of evolved stars.  \\citet{2004ApJ...612.1044P} proposed that binary evolution can result in a much wider ECS initial mass range of 8-11~M$_\\odot$ in the following way. Stripping of the hydrogen-rich stellar envelope due to mass transfer and/or common envelope evolution aborts the second dredge-up phase, normally seen in single stars with their hydrogen envelope intact. Consequently the mass of the helium core is preserved, leading to relatively massive ONeMg cores that experience electron-capture collapse. They find that, in the closest binaries, even the first dredge up phase may be aborted, enabling ECS activity in the even lower mass range of 6-8~M$_\\odot$.  Given the uncertainties involved, probing the ECS physical mechanism and its relevance to NS formation compared to ICC-SN events is a challenge. If ECS events are indeed associated with lower kicks \\citep{2006ApJ...644.1063D, 2006A&A...450..345K}, then close binaries containing a NS are good candidate probes as their formation rates and properties significantly depend on SN kicks and mass loss. \\citet{2002ApJ...574..364P} highlighted a sub-class of high-mass X-ray binaries (HMXBs) with low eccentricities that can be explained by low NS kicks. \\citet{2004ApJ...612.1044P} and \\citet{2004ESASP.552..185V, 2007AIPC..924..598V} postulated that the ECS mechanism could be responsible for these low kicks, although a very wide ECS progenitor mass range may be required to account for the observed HMXBs.  In this {\\em paper}, we present results that uncover an alternative probe for ECS events through studies of HMXB formation in starbursts of varying ages. We find that ECS NS formation associated with low kicks (lower than $\\sim$50\\,km\\,s$^{-1}$) dramatically increases the formation efficiency of Be X-ray binaries in starburst ages of $\\simeq20-60$\\,Myr. In what follows, we discuss this effect in detail and analyze the reasons for this unique dependence, with the hope that observations of starbursts can be used to assess ECS kick velocities and progenitor mass ranges. We note that, in the Small Magellenic Cloud (SMC), dominated by $\\simeq 40$\\,Myr-old stellar populations, a systematic overabundance of Be XRBs has been observed \\citep{2000A&A...359..573H, 2004ApJ...609..133M} potentially making the SMC a unique laboratory for ECS NS formation. ",
        "conclusions": "\\label{disc}  We have shown that, if NS formation through ECS events is associated with natal kicks smaller than $\\sim50$\\,km\\,s$^{-1}$, production of relatively bright wind-fed HMXBs (L$_X$ $>$ 1 x 10$^{32}$\\,erg\\,s$^{-1}$) is significantly favored between 20 and 60\\,Myr after delta function star formation. We call this increased HMXB formation, which appears above a background HMXB population formed through standard ICC compact object formation, the ECS {\\it bump}. The width in post-starburst age of this ECS bump and its amplitude relative to the rest of the ICC-HMXB population\\footnote{including a narrow peak of HMXBs with massive BHs that appears earlier at $\\sim5$\\,Myr}  depend primarily on two ECS factors: the mass range of ECS progenitors and the typical magnitude of ECS natal kicks (Figure~\\ref{ecsplot}). Moreover, we find that these ECS-HMXBs form through a specific evolutionary channel that avoids a CE phase before the SN and instead includes stable mass transfer from the primary and a mass ratio inversion leading to orbital expansion. As a result, ECS-HMXBs have predominantly wide orbits ($\\sim$500--2500\\,R$_\\odot$) and evolved donors. These characteristics, along with their sensitivity to ECS properties, provide us with an intriguing opportunity to probe ECS physics and progenitors through studies of starbursts at different ages.  The SMC is a promising candidate for a study probing ECS physics. Observations of the SMC have revealed a major burst of star formation 30--60 Myr ago \\citep{2001A&A...379..864M, 2004AJ....127.1531H}. Recent observations have found a surprisingly large population of HMXBs with Be donors concentrated in the SMC bar region \\citep{2000A&A...359..573H, 2004ApJ...609..133M, 2009ApJ...697.1695A}.  There are four reasons to prefer the hypothesis that Be HMXBs are primarily formed through ECS creation channels. First, the population of HMXBs in the SMC has been found to have wide orbital separations and low eccentricities, which would require fine tuning of the initial orbital parameters if strong kicks were used \\citep{2002ApJ...574..364P}. Second, we note the coincidence of age between the peak of ECS driven HMXB activity, and the observed age of star forming regions of the SMC bar. Third, the spatial distribution of SMC HMXBs is highly peaked in the bar region, with few HMXBs elsewhere in the galaxy. This spatial distribution is highly unlikely to result from high velocity natal kicks occuring between 10 and 40 Myr before the present. Lastly, \\citet{2005ApJS..161..118M} have recently argued that Be stars form through spin up due to mass transfer in binaries that experience mass ratio inversion and associated orbital expansion, which fits the evolution pathway found for the ECS HMXBs in our models.  \\begin{figure} \\plotone{f3.eps} \\caption{Distance travelled by HMXBs with $L_X>10^{32}$\\,erg\\,s$^{-1}$ detected between 30-60 Myr after star formation for ECS SN with natal kick velocities of 26.5 km s$^{-1}$. Distances are calculated ignoring any deceleration effects from the underlying galactic gravitational potential.} \\label{dist} \\end{figure}  In order to estimate the overconcentration of ECS sources in the SMC bar, we select a value of $\\sigma_{ECS}$~=~26.5~km~s$^{-1}$, and examine the number of HMXB with L$_X$ $>$ 1 x 10$^{32}$\\,erg\\,s$^{-1}$ between 30 and 60 Myr after star formation, taking snapshots of the systems position every 0.1 Myr. In Figure~\\ref{dist}, we plot the cumulative distribution function of three-dimensional distances traveled due to the SN natal kick of both ECS and ICC sources, normalized to their total number. We ignore any effects due to the galactic gravitational potential on the motion of HMXB sources. We note that while overall there are approximately twice as many ICC sources, only around 15\\% of ICC sources are found within 1 kpc of the location in which they were formed. Inside this region, ECS sources dominate by a factor of approximately 4. We note that the deep $Chandra$ observations concentrate on the SMC bar region, while regions further outside are either not targeted by observations or their X-ray luminosity sensitivity limit is much higher. Consequently a significant number of XRBs predicted by the models to have moved to distances beyond $\\sim10$ kpc could remain undetected.  Our above findings strongly depend on our choice to set natal kicks for direct collapse BHs to be 10\\% as strong as our Maxwellian natal kick distribution. If we set direct collapse natal kicks to be strictly zero, we find a large population of persistent sources in the SMC bar with luminosity $L_X>10^{35}$\\,erg\\,s$^{-1}$, containing CHeB donors and massive BH accretors, and having orbital periods on the order of thousands of days. As these systems are not only unobserved in the SMC, but throughout all HMXB populations, we enable small direct collapse kicks which split these loosely bound systems.  Lastly, while the Be star phenomena is not understood from first principles, making it difficult to determine the absolute number of systems with Be-driven luminosities above $L_X>10^{34}$\\,erg\\,s$^{-1}$, we note that our simulations are consistent with the lack of observed supergiants with spherically driven winds. For each million initial binaries, we find approximately 7.5 such sources with luminosity $L_X>10^{34}$\\,erg\\,s$^{-1}$ which have travelled less than 1 kpc. Using standard assumptions for initial mass function and star formation rate in the bar region, this is consistent with a Poisson distribution of supergiant systems.  These results demonstrate the promise of using the ECS process in order to correctly model the great overabundance of Be-HMXB sources in the SMC. Because of the low natal kicks imparted to ECS systems, HMXBs formed through the ECS process are able to replicate the long period of Be-HMXB sources and high concentration of these systems in the SMC bar. While several models exist which allow for low natal kick velocities, we show that the ECS process should be favored due to its ability to create a dominant population of systems at times matching SMC observations. In addition, the ECS mechanism produces HMXBs through the same physical process suspected to produce a Be-HXMB population \\citep{2005ApJS..161..118M}. Detailed modeling of the HMXB population in the SMC, taking into account the duration, age and spatial distribution of star formation activity, the motion of the star forming clusters, and the lower metallicity of the SMC is beyond the scope of this {\\em paper}, but is currently underway. We anticipate that such studies of the SMC and potentially other nearby galaxies with recent starburst activity will greatly contribute to efforts of understanding SNe and compact object formation."
    },
    "0807/0807.3747_arXiv.txt": {
        "abstract": "We present an analysis of the stellar populations of the gas-stripped outer disks of ten Virgo Cluster spiral galaxies, utilizing SparsePak integral field spectroscopy on the WIYN 3.5m telescope and GALEX UV photometry. The galaxies in our sample show evidence for being gas-stripped spiral galaxies, with star formation within a truncation radius, and a passive population beyond the truncation radius. We find that all of the galaxies with spatially truncated star formation have outer disk stellar populations consistent with star formation ending within the last 500~Myr. The synthesis of optical spectroscopy and GALEX observations demonstrate that star formation was relatively constant until the quenching time, after which the galaxies passively evolved.  Large starbursts at the time of quenching are excluded for all galaxies, but there is evidence of a modest starburst in at least one galaxy. For approximately half of our galaxies, the timescales derived from our observations are consistent with galaxies being stripped in or near the cluster core, where simple ram-pressure estimates can explain the observed stripping. However, the other half of our sample galaxies were clearly stripped outside the cluster core. Such galaxies provide evidence that the intra-cluster medium is not static and smooth. For three of the most recently stripped galaxies, there are estimates for the stripping timescales from detailed gas stripping simulations. For all three of these galaxies, our stripping timescales agree with those from the gas stripping simulations, suggesting that star formation is quenched near the time of peak pressure. While the stripping of star-forming gas in the outer disk creates a passive population in our galaxies, there is still normal star formation in the center of our sample galaxies.  It may be that Virgo is not massive enough to completely strip these spiral galaxies and, in a more dynamically active cluster or a cluster with a higher density ICM, such a process would lead to passive spirals and/or S0s. ",
        "introduction": "The morphology-density relationship \\citep{dressler80} implies that clusters are creating S0 galaxies at the expense of spirals.  The most salient difference between the two galaxy populations is the presence (spirals) or absence (S0's) of current strong star formation. \\citet{kauffmann04} find that the star formation - environment correlation is the strongest correlation of any galaxy property with environment, suggesting that environment plays a key role in the termination of star formation.  While there is intense debate as to the specific environmental processes that cause the transformation, stellar population simulations \\citep{shioya04} have shown that simple quenching of star formation in a spiral galaxy can form S0-type spectra. However, not all S0 galaxies are faded spirals; there is a significant fraction of S0's with larger bulges than spiral galaxies (e.g.  \\citealp{christlein04}), something that cannot be explained by simple disk fading. Ram pressure stripping is not the sole cause of the transformation of spirals to S0's, but it does appear to be a key part of the transformation for many galaxies.  A recently discovered population of cluster galaxies, ``passive spirals'' \\citep{dressler99}, have noticeable spiral structure but little, if any, ongoing star formation.  UV photometry \\citep{moran06} has shown that these galaxies have intermediate stellar population properties between spirals and S0's; passive spirals have no ongoing star formation, but have stopped forming stars only recently. These galaxies are consistent with an intermediate phase between spirals and S0s. \\citet{poggianti06} suggest that S0's may represent a heterogeneous population: ``primordial passive galaxies'' that formed at $z>2.5$ and environmentally-formed S0s, which may evolve through a passive spiral phase. \\citet{shioya04} demonstrate that such evolution from active spiral to passive spiral to S0 is naturally explained by passive stellar evolution after a stripping event and also suggest that there are at least two paths to S0 galaxies.  In the nearby Virgo cluster, there exists a population of spiral galaxies with normal star formation inside some radius, but little or no star formation beyond that ``truncation radius'' \\citep{kk04}. Many of these galaxies have relatively undisturbed stellar distributions, despite their severely truncated H$\\alpha$ disks. Such a disturbance in the gas with little or no apparent stellar disk disturbance strongly suggests a gas-gas interaction; if a gravitational interaction disturbed the star formation, the stellar light should also be disturbed. The appearance of these galaxies beyond the stripping radius (i.e. normal stellar disk with no current star formation) is strikingly similar to the passive spirals in higher redshift clusters. By studying the stellar populations of our sample galaxies, we aim to understand when and where in the cluster galaxies are stripped. Are these Virgo spirals with truncated star-forming disks just passive spirals that have not been fully stripped?  Our sample of galaxies are drawn nearly entirely from the work of \\citet{koopmann01}, an atlas of $R$-band and H$\\alpha$-band images of spiral galaxies in Virgo. They find that $\\sim 50\\%$ of spiral galaxies in Virgo have truncated H$\\alpha$ disks \\citep{kk04}: relatively normal star formation in their central regions, with little or no star formation beyond a well-defined radius which we call the gas truncation radius. By studying the stellar populations of the outer disks of these galaxies through optical spectroscopy and UV imaging, we aim to understand where in clusters galaxies can be stripped and what effect gas stripping has on the star formation properties of spiral galaxies. The location in the cluster of stripping tells us how effective stripping is at different distances from the cluster center. Simple models of ISM-ICM stripping suggest that stripping can only occur in the cluster core, where the ICM densities are highest. Galaxies stripped far from the center of the cluster suggest a non-static, non-smooth ICM or extreme conditions outside the cluster core.  In \\S \\ref{sec-obs}, we outline our observations and explain the observational setup for optical spectroscopy and GALEX UV imaging. In \\S\\ref{sec-abslines}, we discuss the relevant spectral features that we use to determine the stellar population properties of these galaxies. In \\S\\ref{sec-models}, we describe the models that we have used to interpret the stellar population properties of our observed sample. In \\S \\ref{sec-galaxies}, we detail our observations and results and discuss what the results imply about the quenching of star formation in the galaxies' outer disks. In \\S \\ref{sec-discussion}, we discuss the implications of our individual galaxy results for the cluster as a whole, compare our results with simulations of neutral gas, and briefly discuss what our observations of these galaxies in Virgo may imply for galaxy evolution in clusters. Finally, in \\S \\ref{sec-summary}, we summarize our results. ",
        "conclusions": "\\label{sec-discussion}   \\subsection{Ages and Star Formation Histories from GALEX Observations} \\label{sec-galex}  \\begin{figure}[t] \\plotone{fig16.eps} \\caption[GALEX Flux vs. H$\\delta$A]{Outer disk H$\\delta$A absorption lines strengths plotted against the GALEX flux ratio for the same outer disk region. Shown as lines on the figure are Z=0.02 (top) and Z=0.008 (bottom) Starburst99 models for a SSP (dashed line), a quenched model with a 2\\% starburst (dotted line) and a quenched model (dash-dot line). The solid lines are lines of constant age; from top to bottom: 100, 200, 300, 400, and 500 Myr.} \\label{fig-galexvshda} \\end{figure}  Because of the extraordinary sensitivity of GALEX UV observations to the presence of young populations, we can potentially distinguish between different star formation histories using our GALEX observations.  In particular, we can extract information about a galaxy's star formation history by combining analyses of UV emission, which traces a very young stellar population and stellar absorption lines, which give us information about star formation over a longer timescale. There is moderately good agreement between the ages determined by spectral analysis of the outer disk and the GALEX UV observations (Figure \\ref{fig-galexvshda}) when assuming a quenching model. There is a clear sequence from NGC~4522, which has strong Balmer lines (young ages) as determined by spectral measurements and a very hard UV color, to NGC~4580 and IC~3392, which have comparatively weak Balmer lines (older ages) and significantly softer UV colors. While there is a clear sequence, the galaxy measurements are clearly offset from the solar-metallicity models (top panel of Figure \\ref{fig-galexvshda}). As can be seen when comparing to the Z=0.008 models (bottom panel of Figure \\ref{fig-galexvshda}), this is apparently a metallicity effect, showing that there is some degeneracy between star formation histories (i.e. quenching, quenching plus a burst, SSP) and metallicities. Metallicity estimates from the [MgFe]' index indicate that most of the galaxies have a luminosity-weighted mean metallicity that is intermediate between Z=0.008 and Z=0.02, suggesting that all of the galaxies are consistent with simple quenching of star formation, with the exception of NGC~4522 (See \\citealp{crowl06}).   As discussed in \\citet{crowl06}, the outer disk of NGC~4522 is most consistent with a small starburst at the end of its outer-disk star formation, such that 2\\% of the total stellar mass is formed in a very short interval at the time of quenching. As shown in Figure \\ref{fig-4522sp}, the absorption line indices can not be explained with a quenched star formation history; the Balmer lines in the outer disk are too strong. It is through GALEX observations (shown in Figure \\ref{fig-galexvshda} and in \\citealp{crowl06}) that we are able to obtain a stripping timescale without making an assumption about the star formation history.  \\subsection{Comparison of Stellar Population Ages to Gas Stripping Simulations}  \\begin{figure}[t] \\plotone{fig17.eps} \\caption[$\\tq$ vs. Simulation Age]{$\\tq$ plotted against the time since peak pressure from gas stripping simulations ($t_\\mathrm{sim}$) for the three galaxies where both data exist. The errors for $\\tq$ are the errors determined from the stellar population models and the error shown in $t_\\mathrm{sim}$ for NGC~4388 is due to uncertainty in the appropriate model. A 1:1 line is shown for comparison.} \\end{figure}  \\subsubsection{Timescale Comparisons}  For three of the galaxies in our sample, we can compare our timescale estimates to estimates from a completely independent analysis. While we observe and model the stellar properties of the stars left behind after ISM-ICM stripping, it is also possible to measure and model the gas that has been stripped from these galaxies. For most galaxies in our sample, the lack of high column density gas near the disk of the galaxy means it is difficult to constrain the interaction parameters from gas modeling. However, for three galaxies, NGC~4522, NGC~4388, and NGC~4569, we are able to compare our results to those from gas stripping models.  We find a very young star formation quenching age for NGC~4522, $\\tq = 100\\pm50$ Myr. This implies that star formation has ended very recently and, indeed, we do see large amounts of extraplanar gas near the disk midplane indicating that stripping is ongoing. A comparison of HI distribution and kinematics with ram-pressure stripping simulations of this galaxy \\citep{vollmer06} imply that peak pressure was $\\sim 50$ Myr ago, in excellent agreement with our results.  For NGC~4388, the stellar population analysis indicates that this galaxy was stripped longer ago, $\\tq = 225$ Myr. A comparison of the HI distribution and kinematics with a ram pressure simulation of this galaxy \\citep{vollmer03} originally suggested that it was stripped $\\sim 120$ Myr ago. This is in reasonable agreement with our data, but shows some difference between the HI simulations and the stellar population ages.  However, recent discoveries of more HI much further from the galactic plane \\citep{oosterloo05} imply that the stripping must have happened longer ago. Specifically, considering how far the HI is from the galactic plane, the time since peak pressure should be revised upward to $\\sim$ 200-250 Myr (B. Vollmer, private communication). While these models are based on much less complete HI data than the models for NGC~4522 and NGC~4569, the stripping ages are still in agreement with our estimates from the stellar population analysis.  For NGC 4569, we find the oldest stripping age of the three: $\\tq\\sim300$ Myr. This is precisely the age that \\citet{vollmer04} derive from simulations of their HI observations of this galaxy. In fact, \\citet{vollmer04} conclude that the HI features can only be explained if this galaxy is significantly past the peak pressure phase and much of the observed HI is falling back onto the disk. Additionally, \\citet{boselli06} model the broadband color profile of NGC~4569 and find that the time since interaction is constrained to be less than 400 Myr ago.  \\subsubsection{Implications of Timescale Comparisons}  The agreement of our results with those from stripping simulations \\citep{vollmer03,vollmer04,vollmer06} is remarkable for several reasons. First, the agreement of our estimates for \\textbf{quenching timescales} with their stripping timescales is encouraging. While both our stellar population analysis and HI morphology and kinematic analysis are built on a large set of assumptions, the assumptions are generally {\\it different} for the two methods. While the stellar population analysis relies on assumptions about star formation history and stellar evolution models, the HI simulations rely on how the ICM interacts with the ISM, how gas is transported away from the disk and the ram pressure profile that a galaxy experiences while interacting with the ICM. In the case of these three galaxies, the agreement between the timescales from HI models and the timescales from our stellar population analysis suggests that our timescales are probably correct.  Secondly, the agreement between the ages derived from HI data and stellar population data suggest that the \\textbf{star formation history} we've assumed is basically correct. While the stellar population models are largely insensitive to how stars formed early on, if there is a large burst at the end of star formation, it would lead to artificially young estimates if modeled with a quenched star formation history.  In fact, this can be seen for NGC~4522, where 2\\% of the stellar mass has to have formed in a burst at the end of star formation in order to match the stellar population data. The overall agreement of ages (even including NGC~4522) indicates that these galaxies only experience, at most, modest bursts at the end of their star formation epoch.  Finally, the agreement of our quenching time and the timescales predicted by gas simulations imply that \\textbf{star formation is cut off when the HI is stripped.} The value of $\\tq$ tells us how long it's been since star formation has ended at the location where we measure the stellar populations. In the case of our sample of H$\\alpha$ truncated galaxies, we have measured $\\tq$ just beyond the spatial truncation radius. Because of its relative proximity to the galaxy center, the region just beyond the truncation radius was the most difficult to strip of any region already stripped by the ICM interaction. This implies that this region was stripped by the {\\it strongest} pressure the galaxy had yet experienced. In the case of galaxies that are past peak pressure (as we believe to be the case for all of our sample), the region just beyond the truncation radius was stripped when the galaxy was experiencing peak or near-peak ICM pressure. The agreement between the time elapsed since peak pressure (from the ram pressure simulations) and the quenching time of star formation ($\\tq$ from our work) means that star formation stops near the time of peak pressure. This agreement implies that gas stripping happens quickly and the stripping of the neutral gas (HI) is roughly coincident with the end of star formation.  \\subsection{Locations of Stripping in Virgo} \\label{sec-locstrip}  In a study of nearby clusters, \\citet{solanes01} find that, while HI deficiency extends out to 2$R_A$, the proportion of gas-poor spirals increases toward the center of the clusters. Additionally, \\citet{gh85} find that the HI deficient galactic fraction correlates with cluster X-ray luminosity in the sense that those clusters with higher X-ray luminosities (and presumably higher mass) have more HI-deficient galaxies.  It appears, therefore, that cluster processes are driving HI deficiency and pushing gas out of galaxies.  If it is the case that galaxies are only (or primarily) stripped in the core, this implies that there should be a relationship between galaxies' location in the cluster (Figure \\ref{fig-contourpos}) and time since star formation quenching. While projection effects limit our understanding of the true distance between a galaxy and M87 (at the core of Virgo), we can set lower limits to the clustercentric distance by using the {\\it projected} distance from M87 ($d_{87}$). These are only lower limits because the galaxies are also displaced from M87 by an unknown amount along the line of sight.   \\begin{figure}[t] \\plotone{fig18.eps} \\caption[Galaxy Locations Relative to X-Ray Emission]{The positions of the sample galaxies compared to locations of hard X-Ray emission (0.4-2.4 keV), as traced by images from the ROSAT X-Ray map of Virgo Cluster \\citep{bohringer94}. Two sample galaxies are not shown here; NGC~4064 is off of the figure to the north-west and NGC~4580 is off of the figure to the south.} \\label{fig-contourpos} \\end{figure}    \\begin{figure}[t] \\plotone{fig19.eps} \\caption[Quenching Time vs. $d_{87}$]{$\\tq$ for the sample of stripped spirals against projected distance from the central elliptical galaxy M87. Also shown (as a dashed line) is the position a galaxy would have if its star formation were halted in the core and it had been traveling 1500 km/s in the plane of the sky away from M87.}  \\label{fig-agevsd87} \\end{figure}   In Figure \\ref{fig-agevsd87}, we show the distance of the galaxies from M87 plotted against $\\tq$ for each galaxy. The dashed line in Figure \\ref{fig-agevsd87} shows the distance a galaxy could travel in a given time if its velocity was 1500~km/s in the plane of the sky. The average line-of-sight velocity of a galaxy in Virgo relative to the cluster is $\\sim 700$km/s, so 1500~km/s is a generous estimate for most galaxies. Any galaxies to the right of the line in the figure are either too far out to have been stripped in the cluster core or have unusually large velocities in the cluster. This is simply a matter of travel time. Assuming some reasonable velocity, it is impossible for the galaxies to the right of the line to have reached their current location in the time since star formation ended. Therefore, they can not have been stripped in the core. While several galaxies are consistent with being stripped in the core, there are three galaxies clearly to the right of the line and one galaxy marginally to the right of the line; we will take them case by case.   \\subsubsection{NGC~4522}  NGC~4522, despite being far from the cluster core, has an exceedingly young population that precludes it from being stripped in the core. In addition to its young population, the large amounts of extraplanar HI and H$\\alpha$, and enhanced polarized radio continuum emission and flat radio spectral index along the leading edge all point to ongoing strong ram pressure. Simple estimates of ram pressure are too low by a factor of 10 to adequately explain the observed stripping \\citep{kenney04}. However, the ``simple'' estimates of ram pressure assume a smooth, static ICM. The existence of galaxies like NGC~4522 suggest that the ICM is either dynamic or ``lumpy.'' Observed ICM shocks \\citep{shibata01} suggest that the ICM may be moving with significant velocity in the region of NGC~4522.  In \\citet{crowl06}, we include a more complete discussion of the stellar populations of this stripped spiral galaxy.   \\subsubsection{NGC~4580}  NGC~4580 is quite far from the cluster core but has an intermediate-aged outer stellar disk ($\\tq = 475$ Myr). In many ways, NGC~4580 is a prototypical stripped spiral, with normal star formation within a well-defined radius and no star formation beyond, but with an undisturbed stellar disk. One slightly peculiar aspect of this galaxy is the presence of stellar spiral arms beyond the truncation radius. These stellar spiral arms in the outer disk are stronger than in any of the other sample galaxies with comparable quenching times. Without gas, the spiral arms should dissipate relatively quickly, perhaps within a few galaxy rotation periods, but it appears these arms have persisted.   \\subsubsection{NGC~4064}  NGC~4064 is very far out in the cluster, the furthest from M87 of any galaxy in our sample. Despite this, its stellar population points to star formation quenching only 425 Myr ago. It is possible that this galaxy was stripped of its star forming gas by an interaction with the ICM, but if that is the case, it did not happen in the core of Virgo. Optical imaging, kinematics, and CO interferometry \\citep{cortes06} suggest that this galaxy has recently undergone a minor gravitational interaction. However, this by itself is likely to not be enough to remove the large amount of gas that this galaxy has apparently lost. The HI is truncated well within the stellar disk of NGC~4064, demonstrating that the gas was removed without significantly disrupting the stars.  \\subsubsection{NGC~4424}  NGC~4424 is modestly below the 1500 km/s line and has a relatively young stripping age. This, coupled with the large HI tail observed in this galaxy \\citep{chung07}, suggests that this galaxy recently underwent an interaction. Optical morphology and stellar kinematics \\citep{cortes06} demonstrate that NGC~4424 has recently experienced a significant gravitational interaction. Such interactions are typically more common in the cluster outskirts, where the relative velocities of galaxies are lower. However, the large HI tail pointing toward M49 \\citep{chung07} suggests that a collision with that galaxy may have played a role in the stripping of the HI from NGC~4424. While this is clearly not the only environmental effect this galaxy has experienced, it appears that a ram pressure interaction did play a role in quenching the star formation in NGC~4424.   \\subsubsection{Galaxies Stripped in the Core}  The remaining galaxies fall near the $v_t = 1500$ km/s line. This implies that, within the regions probed by our spectroscopic observations, star formation was interrupted on a recent core passage. If the galaxies had been stripped on a previous passage through the Virgo core, they would fall to the left of the line shown in the figure; that is, they would be closer to the cluster core than their stripping age and a travel speed of 1500 km/s would imply.  The observation that these galaxies are at the minimum age that their distances allow suggests that the current appearance of these galaxies, with truncated HI and H$\\alpha$ disks, may be a short-lived state.  One possibility is that these galaxies are stripped further on a subsequent core passage. Note that because $\\tq$ has only been determined just beyond the truncation radius, we cannot comment on how much of the outer disk was recently stripped; it's possible that only a narrow annulus of gas was stripped $\\tq$ ago. As a galaxy enters a cluster, much of its gas is likely stripped on its first core passage but models suggest that, in a moderate-sized cluster like Virgo, it's unlikely that {\\it all} of the gas will be stripped from an 0.5$L_*$ or 1.0$L_*$ galaxy on that first passage (e.g. \\citealp{vollmer01}). However, on subsequent orbits, gas at smaller and smaller radii may be stripped. This is particularly likely if cluster conditions and galaxy orbits change, which are both likely in an unrelaxed cluster. However, our results appear to show that some fraction of star-forming gas was truncated on these galaxies' last passage through the cluster core.   \\subsection{Are Stripped Spirals ``Failed'' Passive Spirals?}  With the exception of small star-forming disks in their centers, the observed stripped spirals in our sample are similar to the ``passive spirals'' observed at higher redshift \\citep{dressler99}. It may be that Virgo is not massive enough to entirely strip the ISM from the spiral galaxies we have observed.  However, in the largest clusters, the ICM density is higher by a factor of $\\sim 10$ and the relative velocities are higher by a factor of $\\sim 2$; this means that ram pressure is higher by a factor of $\\sim 40$ in these clusters. Therefore, it may be that the true passive spirals are only found in higher mass clusters such as those of \\citet{dressler99} and that only truncated spirals are found in more modest clusters like Virgo.  While our sample is not a complete sample of truncated spirals in Virgo, it is still interesting that none of the stripped spirals in our sample were stripped more than 500 Myr ago. If this observed state (truncated star-forming disk) were long-lived, we might expect to see stellar populations with a range of ages; from very newly stripped disks (i.e. NGC~4522 or NGC~4388) to disks that had been stripped several gigayears ago. Those galaxies with the most regular, symmetric H$\\alpha$ disks tend to have the oldest values of $\\tq$ (i.e. NGC~4405, IC~3392, NGC~4580), but those values are all less than or equal to 500~Myr, a relatively short time compared to the orbital time in the cluster.  We do not know whether the star formation in these galaxies will fade or the star forming gas will be completely stripped on a subsequent core passage, but from the stellar populations of the outer disk, it appears that their current state is not long-lived.    \\subsection{Metallicities of Galaxy Populations}  It is generally understood that galaxies create metals through stellar evolution processes. This means that, broadly, there is an age-metallicity relationship: in a given galaxy, the stars that form later tend to be more metal rich than the stars that formed first. While this is broadly true, the properties of any one region of a galaxy can be largely determined by the {\\it local} star formation history.  In order to avoid adding free parameters to the SB99 models, we summed stellar models at a constant metallicity; all stars in a given model (regardless of the age) have a fixed metallicity. This is, of course, a gross over-simplification of the stellar population of the galaxy, but due to the relative insensitivity of the Balmer indices to metallicity (particularly at the young ages that are most relevant for our study), we have chosen {\\it not} to model any sort of enrichment history. As a result, the metallicities implied by the [MgFe]$^\\prime$ index are the luminosity-weighted average metallicity for the entire galaxy and not the metallicity of the current generation of stars. It should be noted that the luminosity-weighted average metallicity is slightly sub-solar for most of our sample of spiral galaxies. This is expected, as galaxies with solar-metallicity {\\it current} generation stars likely have a luminosity-weighted mean luminosity that is slightly sub-solar; in galaxies with realistic enrichment histories, the older stellar generations have fewer metals than the current stellar generation.     We have presented observations of several gas-stripped spiral galaxies in the Virgo cluster. These are galaxies that have apparently been partially stripped of their gas by interactions with the ICM, but all still have ongoing star formation in their centers. We present optical spectroscopy that demonstrates that the outer disks were stripped of their star-forming gas within the last 500 Myr.  We find that, while some galaxies are consistent with being stripped in the cluster core, others appear to have been stripped outside the core. There is evidence that some outer cluster galaxies have undergone gravitational interactions, but others appear consistent with ram pressure stripping. These galaxies provide evidence that the ICM is not static and smooth. In at least one case (NGC~4522), there is independent evidence that the ICM pressure is stronger than simple estimates would predict. If the effects of ram pressure are important well outside the cluster core, the HI deficiencies observed at large radii in clusters \\citep{solanes01} can be explained by ISM-ICM interactions. These interpretations imply that the ``reach'' of the ICM extends well beyond the cluster core and that ISM-ICM interactions may have a greater impact on morphologically-driven galaxy evolution than simple ICM models would suggest.  Furthermore, for three galaxies that have been modeled by gas dynamics simulations, we find agreement between our determined quenching times and those simulations. These results imply that our galaxies experienced essentially constant star formation prior to being quenched and that there was, at most, a modest starburst at the time of quenching. It also suggests that star formation is effectively halted by these strong ICM interactions and molecular gas is effectively stripped quickly by ICM pressure.  Taken together, our observations show that gas stripping is an important process in environmentally-driven galaxy evolution. As ISM-ICM interactions are capable of rapidly halting star formation, it is reasonable to suppose that a significant passive galaxy population may evolve from spiral galaxies in clusters. While there may be multiple paths to S0 galaxies, it appears that, in Virgo, we are observing a process that can efficiently halt star formation in cluster spiral galaxies.  \\vspace{3mm}  We gratefully acknowledge the advice and assistance of Jim Rose, which were critical at an early stage of this project. Discussions with Jacqueline van Gorkom, Aeree Chung, Bernd Vollmer, and Bob Zinn were also crucial in advancing the science case of our paper. Additionally, we are grateful to David Schiminovich for providing assistance and support with the GALEX data, as well as comments on the project and paper. We gratefully acknowledge NASA's support for construction, operation, and science analysis for the GALEX mission, developed in cooperation with CNES of France and the Korean Ministry of Science and Technology. Finally, we thank the anonymous referee for comments that improved this paper. This research is supported by NSF Grant AST-0071251."
    },
    "0807/0807.1742_arXiv.txt": {
        "abstract": "We present a light and fast, public available, ray-tracer {\\tt Splotch} software tool which supports the effective visualization of cosmological simulations data. We describe the algorithm it relies on, which is designed in order to deal with point-like data, optimizing the ray-tracing calculation by ordering the particles as a function of their ``depth'' defined as a function of one of the coordinates or other associated parameter. Realistic three-dimensional impressions are reached through a composition of the final color in each pixel properly calculating emission and absorption of individual volume elements. We describe several scientific as well as public applications realized with {\\tt Splotch}. We emphasize how different datasets and configurations lead to remarkable different results in terms of the images and animations. A few of these results are available online. ",
        "introduction": "According to the standard theory of the {\\it Big Bang}, our universe came into existence as a singularity around 13.7 billion years ago. After its initial appearance, it inflated from a very dense, hot and homogeneous state to the expanded, cooled and very structured state of our current universe. During its evolution, planets, stars, galaxies and all the different objects and structures that we observe have formed, as the result of a fine interplay of gravitational forces with a number of different physical processes, like fluid dynamics effects (shock waves, turbulence), magnetic fields, star formation and associated feedback (e.g.\\ supernova explosions). Cosmologists adopt numerical simulations as an effective instrument to describe, investigate and understand the evolution of such a complicated system within the framework of the expanding space-time of the universe. There are various numerical methods -- grid or particle based -- to perform such simulations. For a recent review on such numerical simulation methods within the cosmological context see \\citet{2008SSRv..134..229D}.  The resulting data is huge and complex and requires suitable and effective tools to be inspected and explored. Visualisation represents the most immediate and intuitive way to analyse and explore data and to understand results. Larger and larger data sets require an enormous computational effort to be systematically studied and more and more complex and expensive algorithms to identify patterns and features. Visualisation allows easily to focus on sub-samples and features of interest and to detect correlations and special characteristics of data. This is not a substitute for systematic analysis, but it is a fundamental support to accelerate and simplify the cognitive process.  Visualisation also is an effective instrument to introduce common people, especially young people, to even the most complicated and innovative aspects and concepts of science. The opportunity to use 3D digital visualisation and Virtual Reality facilities, further enhance the impact for the communication and divulgation and push toward a new approach to the cultural and scientific heritage.  The availability of tools to visualise scientific data in a comprehensible, self-describing, rich and appealing way is therefore crucial for researchers and for common people, both for creating knowledge and for disseminating it.   Such tools must be able to deal with a huge data volume, possibly leading to meaningful results in a reasonable time. Furthermore, they must meet the requirements of the particular discipline which produced the data and fit the specific characteristics and peculiarities of the same data. In particular, most of the state-of-the-art cosmological simulations describe the different matter components in the universe (stars, dark matter, gas \\dots) as fluid elements (either point-like or on a regular / irregular grid), which must be properly rendered. For data based on regular grids many visualization packages are available. However, simulations based on a particle-like description of the fluid elements are much more difficult to visualize. There are standard packages like {\\tt TIPSY}\\footnote{http://www-hpcc.astro.washington.edu/tools/tipsy/tipsy.html} or {\\tt VisIVO}\\footnote{http://visivo.cineca.it/} for displaying such data, however they are not designed to lead to ray-tracing like images, which might give more realistic impressions of 3D structures ore even artistic like impressions for public outreach.  In this paper we present {\\tt Splotch}, a public ray-tracing software. {\\tt Splotch} is specifically designed to render in a fast and effective way the different families of point-like data, results of a cosmological, but, more generally, astrophysical simulation. {\\tt Splotch} is based on a lightweight and fast algorithm, which will be described in Section 2. Section 3 will present some examples where {\\tt Splotch} was used to produce images for scientific publications and animations presented at various conferences. In Section 4 we describe a {\\it 4D Universe} project where {\\tt Splotch} was used to produce a public outreach movie which will be shown to the public in the Virtual Reality facility of the brand new Turin Planetarium. We will describe methods used to generate the data used for the {\\tt Splotch} rendering and the technique adopted to create the stereographic version of the movie.  \\setcounter{footnote}{0} ",
        "conclusions": "Visualisation is one of the most effective techniques to explore and present scientific data, understanding at a glance its basic features and properties. Proper graphical solutions must be developed and exploited in order to fully achieve such objective. Such solutions can change according to the target users, depending, for example, whether professionals or common people are addressed. Moreover, different application fields can have specific requirements and different data sets can have peculiarities which need to be treated with a suitable approach.  Furthermore, the software must be able to face the major challenge represented by present time scientific datasets: the huge data volume. Visualisation tools have to be able to manipulate Gbytes of data at once, producing the result in an acceptable time. The algorithm must be tuned on the data in order to get the best possible performance. 64bit architectures, large memories and multi-processor systems have to be supported and exploited in order to handle such extraordinary data processing efforts.  In this paper, we have shown how the {\\tt Splotch} software can be adapted to a variety of different applications in astronomy, ranging from professional usage to divulgation and outreach contents production. For point-like data sets, {\\tt Splotch} meets all the previous requirements, adding the further advantage of being open source, flexible and easily extensible. It can be used both as a stand-alone application and as a library function, callable from inside the simulation code.  {\\tt Splotch} was successfully used to produce a stereographic movie for the Turin Planetarium. The movie has the ambitious goal of describing the birth and the evolution of the universe, exploiting actual cosmological numerical simulations, observations and models. Therefore it helps to give the general public an exceptional insight in our understanding of the world on its largest scales.  As a side product of the increased complexity of the physical systems captured by state of the art, cosmological simulations, visualisations of them are extremely attractive for the public and start to reach the state where they are of comparable beauty than real observations, which traditionally reflected the extremely attractive nature of astronomy."
    },
    "0807/0807.2741_arXiv.txt": {
        "abstract": "The detection of transiting extrasolar planets requires high-photometric quality and long-duration photometric stellar time-series. In this paper, we investigate the advantages provided by the Antarctic observing platform \\mbox{Dome C} for planet transit detections during its long winter period, which allows for relatively long, uninterrupted time-series. Our calculations include limiting effects due to the Sun and Moon, cloud coverage and the effect of reduced photometric quality for high extinction of target fields. We compare the potential for long time-series from \\mbox{Dome C} with a single site in Chile, a three-site low-latitude network as well as combinations of \\mbox{Dome C} with Chile and the network, respectively. \\mbox{Dome C} is one of the prime astronomical sites on Earth for obtaining uninterrupted long-duration observations in terms of prospects for a high observational duty cycle. The duty cycle of a project can, however, be significantly improved by integrating \\mbox{Dome C} into a network of sites. \\vspace{1cm} ",
        "introduction": "Extrasolar planets are currently detected via indirect methods. The radial velocity technique has proved most successful in recent years. However, photometric measurements of planets transiting in front of their central star through the line-of-sight to Earth have already provided more than 30 planet detections in the past years\\footnote{See www.exoplanet.eu}. Many more transiting planet detections are expected in the very near future by the space missions CoRoT \\citep{corot} and Kepler \\citep{kepler}. Combining both methods allows the derivation of planetary radii, true masses and thus, the mean density of extrasolar planets. Transiting planets are therefore of particular importance to investigate basic planetary parameters.  The detection of extrasolar planets via the transit method requires long duration, high photometric quality lightcurves of a large number of stars. The major difficulties for ground-based transit searches are interruptions of the time series, e.g. due to the diurnal cycle or weather conditions, and the Earth's atmosphere limiting the photometric signal-to-noise-ratio \\citep{transit_method}. Located on the Antarctic Plateau at $75^\\circ$ southern latitude, \\mbox{Dome C} provides the potential for long, uninterrupted high-quality photometric time series observations from ground during its winter period.  The maximum number of astronomical nights at \\mbox{Dome C} in comparison to Mauna Kea has already been calculated by \\citet{kenyon}, including the influence of the Sun, the Moon and cloud coverage. While \\mbox{Dome C} has less astronomical nighttime, the percentage of cloud free observing hours is similar to Mauna Kea. Kenyon \\& Storey do not discuss optimizations for long time series observations of target fields. Their paper, however, presents a detailed discussion of the expected sky brightness at \\mbox{Dome C}. They estimate that atmospheric scattering should be exceptionally low, and even the sky brightness due to the Moon should be reduced. \\mbox{Dome C} has reduced zodiacal light, no light pollution and only minor disturbances due to aurora. Furthermore, scintillation noise as a major limiting factor of photometric quality is expected to be about 3.6 times lower at Dome C compared to Chilean sites for long integrations \\citep{kenyon2}, which is particularily important for the search of small planets. \\mbox{Dome C} may therefore be the best site on Earth for stable high-accuracy photometry.  Alternatively, networks distributed favorably in longitude can also provide high duty cycles. A well known example is the Global Oscillation Network Group (GONG) for asteroseismological studies which include six network sites (four in the northern and two in the southern hemisphere). An estimate of the expected duty cycle for spectroscopic asteroseismology from \\mbox{Dome C} in comparison to GONG has been presented by \\cite{mosser}. They conclude that \\mbox{Dome C} provides a better or similar performance than GONG for observing periods up to 100 days, attaining a duty cycle of up to 87\\%. The network performs better for observing periods of half a year or more, but with a duty cycle of only about 77\\%.  In this paper, we include the factors already investigated by Mosser \\& Aristidi, like the Sun and cloud coverage. In addition we include the Moon and the varying extinction of the target fields during observations. Furthermore, we discuss not only a comparison of \\mbox{Dome C} to low-latitude networks, but also a combination of \\mbox{Dome C} with other sites. A comparison between \\mbox{Dome C} and the CoRoT mission has already been presented \\citep{ARENA2}. The resulting potentials for long-duration time-series photometric observations are then discussed in the context of maximizing the probability to detect extrasolar photometric transits. ",
        "conclusions": "We have investigated the prospects for long-duration uninterrupted photometric time series from \\mbox{Dome C} and other low-latitude sites. Contrary to previous studies we \\textit{simultaneously} took into account the limitations of the Sun and Moon, clouds and target field extinction variations.  Due to the long winter, \\mbox{Dome C} has an outstanding performance in observing a target field over a long time period from a single site. Furthermore, the effective observing time period can be extended if \\mbox{Dome C} is combined with other sites which can provide data during the summer period. For example, the combination of \\mbox{Dome C} with a site in Chile extends the observational coverage significantly. Combining \\mbox{Dome C} with a low-latitude network of sites at Chile, Namibia and Australia provides prospects for observational duty cycles that can be improved upon only by observations from space.  To predict quantitative transit detection rates, the photometric quality of \\mbox{Dome C} must be investigated directly. However, for this we have to await the results of photometric site testing projects like the Antarctica Search for Transiting Extrasolar Planets (ASTEP, \\cite{astep}) for a detailed study. Nevertheless, the extensively long observational time baseline for photometric variability studies clearly demonstrates the extremely encouraging prospects of this site in Antarctica."
    },
    "0807/0807.0869_arXiv.txt": {
        "abstract": " ",
        "introduction": "~~~~~The recent  galactic redshift surveys showing a very inhomogeneous picture for the distribution of galaxies (de Lapparent, Geller \\& Huchra 1986; Saunders et al. 1991) have stimulated the trend of study of the galactic clustering problem which assumes that the large scale structure of the universe can be described as being a self--similar fractal system. Under this philosophy there have been studies in which the distribution of galaxies is assumed to form a multifractal system (Jones et al. 1988; Mart\\'{\\i}nez et al. 1990), but the first single fractal Newtonian model advanced as a description of the large scale structure is due to Pietronero (1987).  ~~~~~This paper is the second of a series where we develop a relativistic fractal cosmology obtained by the relativistic generalization of Pietronero's (1987) model. In paper I (Ribeiro 1992) we have argued that the all sky redshift surveys mentioned above present observations consistent both with the old Charlier's hypothesis of hierarchical clustering and with fractals, where the latter is basically a more precise conceptualization of the scaling idea implicit in the hierarchical clustering hypothesis. Also in paper I we have used Pietronero's basic hypothesis to propose similar ones in a relativistic framework and to develop observational relations compatible with fractals in Tolman's spacetime.  ~~~~~The purpose of this paper is to apply those observational relations to the specific and analytically manageable special case of the Einstein--de Sitter model, which is obtained by a particular choice of the three arbitrary functions of Tolman's spacetime. Our basic goal here is to answer the question of whether or not the Einstein--de Sitter model is compatible with a fractal description of galactic clustering as developed in paper I, and we intend to do so by studying the consequences of paper I's observational relations in a Einstein--de Sitter universe. We believe that this problem must be addressed because Einstein--de Sitter is spatially homogeneous and it is not at all clear what would be the behaviour of its fractal observational relations as we go down the geodesic, through hypersurfaces of constant $t$ each having a different value for the proper density. The calculations shown next also serve as an illustration of the theory developed in paper I.  ~~~~~The plan of the paper is as follows. In \\S 2 we describe very briefly the Tolman observational relations presented in paper I and discuss how the Tolman metric relates to its Friedmann counterpart. In \\S 3 we apply these observational relations to the Einstein--de Sitter case and analyse the results. There is a conclusion in \\S 4. ",
        "conclusions": ""
    },
    "0807/0807.2577_arXiv.txt": {
        "abstract": "We present {\\sc Rico}, a code designed to compute the ionization fraction of the Universe during the epoch of hydrogen and helium recombination with an unprecedented combination of speed and accuracy. This is accomplished by training the machine learning code {\\sc Pico} on the calculations of a multi-level cosmological recombination code which self-consistently includes several physical processes that were neglected previously. After training, {\\sc Rico} is used to fit the free electron fraction as a function of the cosmological parameters. While, for example at low redshifts ($z\\lesssim 900$), much of the net change in the ionization fraction can be captured by lowering the hydrogen fudge factor in {\\sc Recfast} by about $3\\%$, {\\sc Rico} provides a means of effectively using the accurate ionization history of the full recombination code in the standard cosmological parameter estimation framework without the need to add new or refined fudge factors or functions to a simple recombination model. Within the new approach presented here it is easy to update {\\sc Rico} whenever a more accurate full recombination code becomes available. Once trained, {\\sc Rico} computes the cosmological ionization history with negligible fitting error in $\\sim10$ milliseconds, a speed-up of at least $10^6$ over the full recombination code that was used here. Also {\\sc Rico} is able to reproduce the ionization history of the full code to a level well below $0.1\\%$, thereby ensuring that the theoretical power spectra of CMB fluctuations can be computed to sufficient accuracy and speed for analysis from upcoming CMB experiments like Planck. Furthermore it will enable cross-checking different recombination codes across cosmological parameter space, a comparison that will be very important in order to assure the accurate interpretation of future cosmic microwave background data. ",
        "introduction": "\\label{intro} Planck, the third generation of satellite missions to study the Cosmic Microwave Background (CMB), is scheduled for launch later this year. It will make accurate measurements of the temperature fluctuations in the CMB out to $\\ell\\sim2500$ as well as provide a detailed picture of the CMB polarization to $\\ell\\sim1500$ \\citep{Planck2006}.  In order to maximize the knowledge gained about the underlying cosmological parameters from this high resolution experiment, it is important that our theoretical calculations are equally precise.  Currently, uncertainty in the ionization history of the Universe remains one of the major factors limiting the accuracy of power spectrum calculations \\citep[e.g. see][]{Hu1995, Seljak2003}.  This fact has recently motived several studies on high precision computations of the cosmological hydrogen and helium recombination history (see Sect.~\\ref{sec:Rec_JC} for a detailed overview). The largest corrections thus far have been due to the non-equilibrium distribution of electrons within the hydrogen sub-states, the absorption of helium photons by neutral hydrogen and the $2^3 {\\rm P}_1-1^1 {\\rm S}_0$ triplet-singlet transition in helium. Each of the additional physical processes that have been discussed in the literature lead to $\\gtrsim 0.1\\%$ level corrections of the ionization history. While the individual changes partially cancel each other, the corresponding overall theoretical uncertainty in the CMB temperature and polarization power spectra, in particular at large $\\ell$, still exceeds the level of $0.1\\%$. It is also clear that some additional steps must still be taken, particularly in connection with the radiative transfer in hydrogen and the proper inclusion of two-photon processes (see Sect.~\\ref{sec:additions} for more details). These issues are still largely open questions and the subject of intensive ongoing study.  However, the simultaneous and self-consistent inclusion of all these processes makes the computations very difficult and time-consuming. It involves the simultaneous evolution of hundreds or thousands of differential equations for the occupation numbers of the individual levels of the hydrogen and helium atoms.  As an accurate calculation of the ionization history of a single cosmological model using the current version of our full multi-level recombination code takes many hours or even days of computational time, standard parameter estimation methods require some approximation of the ionization history that is fast to evaluate.  Currently the standard method used to evaluate the ionization fraction is the {\\sc Recfast} code \\citep{Seager1999}, which models the hydrogen and helium atoms as effective $3$-level systems.  In order to approximate the computations done using their multi-level recombination code \\citep{Seager2000}, the authors introduced a fudge factor, $\\fH$, to artificially speed up hydrogen recombination by about $14\\%$.  More recently, \\citet{Wong2008} updated {\\sc Recfast} to include a second fudge factor, $b_{\\rm He}$, to modify the recombination of helium.  This change, which allows speeding up \\ions{He}{i} recombination, was motivated by the results of \\citet{Switzer2008a}, \\citet{Kholupenko2007} and \\citet{Rubino-Mart'in2007}. It is intended to include the effect of the hydrogen continuum opacity during neutral helium recombination (see Sect.~\\ref{subsec:hopacity}).   Here we provide a different approach to computing a fast and accurate approximation of the recombination history. Our method uses a regression code based on {\\sc Pico} \\citep{Fendt2007,Fendt2007a}, which we will henceforth refer to as {\\sc Rico}, to model the ionization history, $X_{\\rm e}(z)=N_{\\rm e}(z)/N_{\\rm H}(z)$, as a function of the cosmological parameters. Here $N_{\\rm e}$ denotes the number density of free electrons, and $N_{\\rm H}$ is the number density of hydrogen nuclei in our Universe. Since $X_{\\rm e}$ is a smooth function, the polynomial approximation used by {\\sc Rico} is extremely accurate, and one retains nearly all of the information contained in the multi-level code used for training, the {\\it training code}.  As the accuracy of the training code is the major determinant of the accuracy of {\\sc Rico}'s approximation, in Sect.~\\ref{sec:Rec_JC} we review the physical processes included in our training code as well as an assessment of the processes that must still be considered and are currently under investigation. In Sect.~\\ref{sec:recfast} we compare the ionization history and corresponding power spectra computed with the multi-level recombination code used for the training of {\\sc Rico} to the different versions of {\\sc Recfast}.  Also we include a short discussion of how one could modify the approach employed by {\\sc Recfast} to capture the new physics included in our full code.  In Sect.~\\ref{sec:Rico} we show that the simple regression scheme used by {\\sc Rico} can accurately model the current full multi-level recombination calculation, giving power spectra that are sufficiently accurate for even a cosmic variance limited experiment.  Although the training code discussed here does take into account most of the important corrections discussed in the literature so far, it does not yet solve the problem completely. Still more physical processes should be included and the results validated by independent codes. As {\\sc Rico} can be trained equally well on {\\it any} recombination code it can facilitate crosschecks between those codes as well as a study of how approximations in these codes propagate to cosmological parameter constraints.  Our target is to ensure that the ionization history can ultimately be calculated to sufficient precision to avoid biasing parameter estimation from the next generation of CMB experiments. By making it easy to propagate advances in the calculation of the ionization history through to predictions of the CMB power spectra with {\\sc Rico}, future development can focus on the physics of recombination and to a lesser degree on the computational efficiency of the training code.  {\\sc Rico} and its future updates will be made available at \\verb+http://cosmos.astro.uiuc.edu/rico+. ",
        "conclusions": "We have presented a new code designed to compute the ionization history of the Universe.  The code includes previously neglected physics that leads to changes in the ionization fraction at the $2-3\\%$ level in some redshift regions.  This change leads to a correction to the CMB power spectra of more than $1\\%$ at $\\ell\\sim3000$.  As uncertainty in the ionization fraction is the main contributor to error in the theoretical CMB anisotropy power spectra, this code represents a significant step in our ability to compute the power spectra to a significant precision for upcoming experiments.  While it is possible to attempt to capture the changes to the ionization fraction by modification of the hydrogen fudge factor in {\\sc Recfast}, some residual error remains near the peak of the visibility function. The problem of correctly introducing fudge factors can be avoided entirely by using a regression code based on a training set of cosmological parameters and their corresponding ionization histories.  A simple polynomial fit is sufficient to compute ionization fraction at the level of $\\sim0.01\\%$ over the volume of parameter space relevant to current experimental data.  As it is straightforward to train {\\sc Rico} on multiple recombination codes, our algorithm enables the propagation of approximations made to the ionization history to biases on the cosmological parameters when analyzing data.  Since implementing the many new physical processes already discussed in the literature in an accurate and robust manner is very challenging, we hope that this ability facilitates a model by model cross validation among the recombination codes that have been developed. Given the complexity of the physical processes involved in the computation of the cosmological recombination history, such a comparison will become very important in order to ensure that the final result is correct. Our goal is to ensure that the ionization history can ultimately be calculated to sufficient precision to avoid biasing parameter estimation from the next generation of CMB experiments. By making it easy to propagate advances in the calculation of the ionization history through to predictions of the CMB power spectra with {\\sc Rico}, future development can focus on the physics of recombination and to a lesser degree on the computational efficiency of the physical recombination code.  {\\sc Rico} and its future updates will be made available at \\verb+http://cosmos.astro.uiuc.edu/rico+."
    },
    "0807/0807.3956_arXiv.txt": {
        "abstract": "Subtraction of astrophysical foreground contamination from ``dirty'' sky maps produced by simulated measurements of the Murchison Widefield Array (MWA) has been performed by fitting a $3^\\text{rd}$-order polynomial along the spectral dimension of each pixel in the data cubes.  The simulations are the first to include the unavoidable instrumental effects of the frequency-dependent primary antenna beams and synthesized array beams.  They recover the one-dimensional spherically-binned input redshifted 21~cm power spectrum to within $\\sim1$\\% over the scales probed most sensitively by the MWA ($0.01 \\lesssim k \\lesssim 1$ Mpc$^{-1}$) and demonstrate that realistic instrumental effects will not mask the EoR signal.  We find that the weighting function used to produce the dirty sky maps from the gridded visibility measurements is important to the success of the technique.  Uniform weighting of the visibility measurements produces the best results, whereas natural weighting significantly worsens the foreground subtraction by coupling structure in the density of the visibility measurements to spectral structure in the dirty sky map data cube.  The extremely dense $uv$-coverage of the MWA was found to be advantageous for this technique and produced very good results on scales corresponding to $|u| \\lesssim 500\\lambda$ in the $uv$-plane without any selective editing of the $uv$-coverage. ",
        "introduction": "Astrophysical foreground contaminants are five orders of magnitude brighter than the $\\sim10$ mK redshifted 21~cm emission expected from neutral hydrogen in the intergalactic medium (IGM) during the epoch of reionization (EoR). These foregrounds will severely complicate planned experiments and the interpretation of their results.  Galactic synchrotron radiation dominates the sky at radio frequencies near 150 MHz ($z \\approx 8$), accounting for $\\gtrsim70\\%$ of the 200 to 10,000~K total brightness temperature \\citep{1999A&A...345..380S}. Extragalactic continuum point sources are also especially strong and numerous, presenting a sea of confused point sources and  comprising the bulk of the remaining $\\sim30\\%$ of the sky brightness temperature. Galactic radio recombination lines (RRLs) and free-free emission from electrons in both the Galaxy and the IGM will additionally complicate the planned measurements.  Initial analyses indicated that these foregrounds were an insurmountable obstacle \\citep{2002ApJ...564..576D, 2003MNRAS.346..871O} because their angular variance dominates the expected fluctuations in the redshifted 21~cm background, but subsequent studies have suggested that multi-frequency observations and the application of appropriate statistical techniques should provide methods to separate the foregrounds from the redshifted 21~cm signal by exploiting the large coherence of the foregrounds with frequency \\citep{2004MNRAS.355.1053D, 2004ApJ...608..622Z, 2004ApJ...615....7M, 2004NewAR..48.1039F, 2004ApJ...608..611G, 2005ApJ...625..575S, 2006ApJ...650..529W, 2006ApJ...653..815M, 2008MNRAS.391..383G}. However, these studies have relied on a number simplifying assumptions, and in particular have neglected complications due to the frequency-dependent instrumental response. The point spread function and instrumental field-of-view are both frequency dependent, and can mix the angular structure that concerned \\citet{2002ApJ...564..576D} and \\citet{2003MNRAS.346..871O} into the frequency direction, masking the EoR signal. This transfer of power from the angular to frequency dimensions has been dubbed ``mode-mixing'', and is a significant concern for redshifted 21~cm measurements.  In this paper we explicitly model the mode-mixing effect for confusion level contaminants and explore ways of minimizing this power transfer into the frequency domain. With these new techniques we show that, for the Murchison Widefield Array (MWA), contamination due to mode-mixing can be reduced well below that of the expected 21~cm signal strength during the reionization epoch for a reasonable foreground model.  We begin in $\\S$2 by defining our instrument and sky models and establishing the mathematical foundation of the subtraction technique.  In $\\S$3 we analyze the ability of the method to produce a model of the confusion-level foreground contamination and, in $\\S$4, we show that the input 21~cm power spectrum in our simulations can be recovered to within $\\sim1$\\% (excluding thermal noise uncertainty) over most scales in the planned MWA measurements.  We conclude in $\\S$5 with a brief discussion and an analytic approximation to provide context to the results of the simulation.  A detailed study of the dependence of this foreground subtraction technique on the properties of the instrument and astrophysical sky has been conducted in parallel by \\cite{Adrian2008} for cases building on our fiducial MWA measurement framework.  In addition, \\citet{2008MNRAS.389.1319J} have recently developed a sophisticated astrophysical foreground model and applied it to analysis of subtraction techniques in the context of the expected antenna configuration of the Low Frequency Array (LOFAR).  Their results indicate that LOFAR should be able to constrain the reionization history over a wide range of redshifts by measuring the variance in image maps from individual spectral channels between 100 and 200~MHz.  Our present focus is on foreground subtraction in the context of three-dimensional power spectrum measurements including the full frequency-dependent instrumental properties of the MWA, but many of the results of all three investigations are relevant to both measurements and instruments, as well as other active redshifted 21~cm experiments including the Precision Array to Probe the Epoch of Reionization (PAPER), the reionization project with the Giant Metre-wave Radio Telescope (GMRT), and the future Square Kilometer Array (SKA). ",
        "conclusions": "In the preceding sections, we have shown that the effects of the frequency-dependent instrumental response of the MWA do not hinder the simple polynomial-fit diffuse foreground subtraction scheme to recover the redshifted 21~cm signal.  This is particularly encouraging because the fit has been performed in dirty sky maps, which are very sensitive to instrumental effects.  Thus, we have established, in some sense, a worst-case scenario that appears manageable.  Dense $uv$-coverage is a new advancement for radio arrays.  In the traditional paradigm, the $uv$-plane is almost always sparsely sampled (as with the VLA and other existing arrays), and there are essentially no regions where the coverage can be said to be ``complete'', as is the case for the MWA within $u\\lesssim500\\lambda$.  The reason that this is generally not a problem for telescopes like the VLA is that, in most cases, these instruments are not attempting to make observations at or below the confusion limit imposed by faint extragalactic continuum sources and diffuse Galactic structure. So, although the VLA has sparse coverage in the $uv$-plane, the sky that it is observing is also essentially sparse and deconvolution algorithms are able to reconstruct a reasonable model of the point sources in the field of view.  It is useful to construct an analytic approximation for the simulations discussed in the previous sections in order to foster a more intuitive understanding of the interaction between the array properties and the success of the foreground subtraction. \\citet{1974ApJ...188..279C}, \\citet{vla-sci-146}, and \\citet{2003A&A...409..787A} provide a foundation for such a model, and Sault and Wayth [2006, 2007, MWA project documents] have adapted the results of these efforts for analyzing confusion noise in measurements of the MWA.  As we outlined in Section~\\ref{sec_foregrounds_method}, a dirty sky map produced by an interferometer is, neglecting calibration errors, the result of multiplying the true sky by the primary antenna tile beam and then convolving by the synthesized array beam.  If the sky consists only of the faint extragalactic continuum sources with Poisson statistics, then the variance in the intensity of the dirty sky map, $\\sigma^2_D$, will be related to the variance in the true sky by \\begin{equation} \\label{eqn_confusion_analytic_full} \\sigma^2_D(\\theta_x, \\theta_y) = \\sigma^2_T \\int \\int B^2(\\theta_x-\\theta_x', \\theta_y - \\theta_y')~P^2(\\theta_x, \\theta_y) ~ d\\theta_x ~ d\\theta_y \\end{equation} where $\\sigma^2_T$ is the variance in the intensity of the true sky, $B$ describes the response of the synthesized array beam as a function of angle, and $P$ describes the primary antenna tile beam. This equation can be simplified by assigning simple models for the response profiles of the synthesized array beam and the primary antenna tile beam. We take the beams to be described by tophat functions, such that the response is defined to be one within a region of diameter $\\Theta_B$ (for the synthesized array beam; $\\Theta_P$, for the primary antenna tile beam). Outside this region, the response is taken to be $B_{rms}\\ll 1$ for the synthesized beam, and zero for the primary antenna tile beam. Thus, we include an allowance for the side lobes of the synthesized beam, but not for the primary beam. With these simplifications, the integrals in Equation~\\ref{eqn_confusion_analytic_full} can be performed and the expression reduces to \\begin{equation} \\sigma^2_D \\approx \\sigma^2_S \\left(1 + B^2_{rms} \\Omega_P /\\Omega_B\\right ) \\end{equation} where $\\sigma^2_S$ is the variance due to the faint sources in our model and accounts for the extra variance due to the weak angular clustering, $\\Omega_B\\approx\\Theta^2_B$ is the solid angle of the synthesized array beam, and $\\Omega_P\\approx\\Theta^2_P$ is the solid angle of the primary antenna tile beam.  Inspecting this simplified form of Equation~\\ref{eqn_confusion_analytic_full}, we see that when the synthesized array beam has no side lobes, the variance in the dirty sky map is equal to that of the true sky map (after gridding). This is equivalent to complete coverage of the $uv$-plane and corresponds generally to the cases considered by \\citet{2004MNRAS.355.1053D, 2004ApJ...608..622Z, 2004NewAR..48.1039F, 2004ApJ...608..611G, 2005ApJ...625..575S, 2006ApJ...650..529W, 2006ApJ...653..815M, 2008MNRAS.391..383G}.  But when side lobes are included in the model of the synthesized array beam, the variance in the dirty map is increased above that of the true sky map.  For sparse coverage in the $uv$-plane, $B_{rms}^2$ will be given approximately by the inverse of the number of independent measurements in the $uv$-plane. In the limiting case of a single, very short integration, (such as a ``snapshot'' observation), the number of independent measurements is equal to the number of baselines. For the MWA, with its $\\sim$125,000 instantaneous baselines, $B^2_{rms}\\approx10^{-5}$ in the worst-case, and is even better (lower) for long integrations when earth-rotation synthesis increases the number of independent measurements.  Using $\\Theta_B\\approx5'$ and $\\Theta_P\\approx30^\\circ$, we can estimate for the MWA that \\begin{equation} B^2_{rms} \\Omega_P /\\Omega_B \\approx 1. \\end{equation} Thus, the variance due to the side lobes in the synthesized array beam is comparable in magnitude to the inherent variance in the intensity of the true sky.  A large field of view acts to oppose the advantages of dense coverage of the $uv$-plane in this instance by introducing more variance into the dirty sky map. In this regime, in order to produce an estimate of the sky without the effects of the side lobes of the synthesized array beam, the full inversion of $\\mathbf{A}$ is required since it cannot be approximated as a sparse matrix.  Nevertheless, side lobe-induced noise in the dirty sky maps does not necessarily mean that the foreground subtraction will be adversely affected. If the side lobe patterns were the same for each image plane in the data cube, then the spectrum along each pixel would simply receive a faint mirror contribution from the smooth spectra along all the other pixels.  It is the variation with frequency of the side lobe-induced noise that can cause trouble in the planned foreground subtraction technique.  As we found in Section~\\ref{sec_foregrounds_subtraction}, however, the density of the MWA $uv$-coverage effectively eliminates the frequency-dependence of the sidelobe structure of the array beam over the large range of scales where the $uv$-plane is completely sampled at all frequencies.  Although it does not appear to be necessary for the MWA, less densely sampled arrays could control the spectral dependence of their array beam sidelobes by masking out regions of the $uv$-plane that are not sampled at all frequencies and only processing regions that do have coverage in each spectral channel (as considered in \\citet{2008MNRAS.389.1319J} for the more sparsely sample LOFAR).  Despite the increased susceptibility of long baselines to introducing undesirable fluctuations into the spectral domain, it is important to point out that they serve an important purpose.  Long baselines are still expected to be necessary in order to peel away the bright sources in the field of view since the peeling process requires a precise knowledge of the position of each bright source, which is greatly improved with long baselines.  It has been shown previously that the ideal baseline distribution for minimizing thermal uncertainty in measurements of the redshifted 21~cm 1D power spectrum is extremely condensed \\citep{2006ApJ...638...20B, 2007arXiv0711.4373L}.  In this work, we have further demonstrated that a compact array like the MWA also produces the dense $uv$-coverage necessary to efficiently subtract faint and diffuse foregrounds from dirty maps."
    },
    "0807/0807.4571_arXiv.txt": {
        "abstract": "Until recently, many of the dozens of quantitative predictions of the ambipolar-diffusion theory of gravitational fragmentation (or core formation) of molecular clouds have been confirmed by observations and, just as importantly, no prediction has been contradicted by any observation. A recent paper, however, claims that measurements of the variation of the mass-to-flux ratio from envelopes to cores in four clouds {\\it de}creases, in direct contrast to a prediction of the theory but in agreement with turbulent fragmentation (in the absence of gravity) and, therefore, the ambipolar-diffusion theory is invalid (Crutcher \\etal 2008). The paper treats magnetic-field nondetections as if they were detections. We show that the analysis of the data is fundamentally flawed and, moreover, the comparison with the theoretical prediction ignores major geometrical effects, suggested by the data themselves if taken at face value. The magnetic fluxes of the envelopes are also miscalculated. We carry out a proper error analysis and treatment of the nondetections and we show that the claimed measurement of the variation of the mass-to-flux ratio from envelopes to cores is not valid, no contradiction with the ambipolar-diffusion theory can be concluded, and no theory can be tested on the basis of these data. ",
        "introduction": "The ratio of the mass and magnetic flux of interstellar molecular clouds has received well-deserved observational attention in recent years (e.g., Crutcher 1999; Heiles \\& Cructher 2005). For a cloud as a whole, the mass-to-flux ratio is important input to the ambipolar-diffusion theory of fragmentation (or core formation) in molecular clouds. For fragments (or cores) within a cloud, the ambipolar-diffusion theory makes specific predictions of the mass-to-flux ratio and its evolution in time. Until now, observations have been in excellent quantitative agreement with the theoretical predictions (e.g., see Crutcher \\etal 1994; Crutcher 1999 and correction by Shu \\etal 1999, pp. 196 - 198; Ciolek \\& Basu 2000). Recently, however, Crutcher \\etal (2008) (hereinafter CHT) have looked at the variation of the mass-to-flux ratio from the envelopes to the cores of four molecular clouds and concluded that there is a fundamental discrepancy between the observations and the predictions of the ambipolar-diffusion theory.  In what follows, we examine critically the claimed discrepancy. Before we undertake a direct re-examination of the data and the conclusions drawn from them, it is important that we summarize the key elements of the ambipolar-diffusion theory of protostar formation. This is especially necessary because, often, what is a prediction of the theory is quoted as if it were a requirement or assumption, and vice versa, i.e., input of the theory taken directly from observations is often quoted as if it were a prediction. The ambipolar-diffusion timescale is also often a frequent victim of abuse in the astronomical literature. In this context, it is of utmost importance to clarify the difference in the way the process of ambipolar diffusion affects the ordered (or mean) component of the magnetic field in a molecular cloud and the magnetohydrodynamic (MHD) waves (or turbulence) that are superimposed on the mean field. It is also useful, for putting the current controversy in its proper context and perspective, to summarize how theory and observations of interstellar magnetic fields have interacted over the years and what the outcome of that interaction has been.  We summarize in \\S \\ref{theory} the main elements of the ambipolar-diffusion theory in a context relevant to the present controversy and to maintaining proper perspective. In \\S \\ref{zeeman} we first summarize the observational procedure followed by CHT and state their main conclusion, and then we state three major flaws in their analysis as well as three additional such flaws in their comparison of the results with the ambipolar-diffusion theory.  In \\S \\ref{A}, we explain the proper procedure for combining multiple observations of the envelope magnetic field in each cloud to derive an estimate of the mean and an associated uncertainty. An explicit assumption of the analysis by CHT is that line-of-sight effects (as encoded in $\\cos \\theta$) are the same for the core and the envelope. An implicit assumption of the same analysis is that the magnetic field does not vary intrinsically from one envelope location to another. None of these assumptions is justified by the observations. Relaxing them in a manner suggested by the data themselves alters the manner in which errors should be treated and increases significantly the uncertainties of all results.  In \\S \\ref{B}, we discuss the difference between a detection and an upper limit of the mean envelope magnetic field. Most of the individual observations yielded nondetections, a fact recognized by the authors. The authors state that for this experiment detections are not necessary. No upper limits are calculated, which would be appropriate for nondetections; instead, nondetections and detections are treated in the same way with respect to error propagation and handling, as well as for quoting results. We remedy this problem and explicitly calculate upper limits where appropriate, using a full Monte Carlo analysis of the error propagation to properly account for the (large) errors propagating through a nonlinear equation (see eq. [\\ref{ther}] below).  In \\S \\ref{C}, we elaborate certain theoretical issues (points 3a, b, c of \\S \\ref{zeeman}) which should be paid due consideration if a meaningful comparison between theory and observations is to be achieved at least as far as the variation of the mass-to-flux ratio from the envelope to the core of a molecular cloud is concerned. ",
        "conclusions": "\\label{disc}  We have examined critically the analysis and conclusions of recent OH Zeeman observations of four molecular cloud envelopes and cores by Crutcher \\etal (2008), and we have shown that: \\begin{enumerate} \\item The error analysis of measurements of the envelope magnetic field is flawed, and the uncertainties on the observed mean value of the field are systematically underestimated in every cloud. \\item Although the mean value of the envelope magnetic field is not statistically significantly different from zero in any cloud, the authors treat these values as if they were detections, and they do not give appropriate upper limits. \\item If the error analysis is performed correctly, allowing possible spatial variations of the line-of-sight magnetic field (suggested by the data), and upper limits are appropriately calculated for the mean magnetic field in the envelopes, these upper limits are very weak: the mean magnetic field in the envelopes is constrained to be smaller than $20-30{\\rm \\, \\mu G}$ at the $2\\sigma$ significance level. \\item If corresponding $2\\sigma$ upper limits are calculated for the value of the ratio $R$, which is given by equation (7), they are also weak: $R$ is constrained to be smaller than a few. For no cloud is the upper limit smaller than $1$. $R$ is never statistically significantly different from zero, and has not been ``measured'' for any cloud. \\item The magnetic fluxes of the clouds were calculated incorrectly, by averaging {\\it algebraically} the (presumed) values of the line-of-sight magnetic field $B_{\\rm los}$ over a cloud's area projected on the plane of the sky, despite claimed reversals in the field direction not along one and the same line-of-sight, but from one line-of-sight to another. \\item The possible gravitational breakup (of the mass {\\it along} the flux tubes threading the cores) into thermally supercritical but magnetically subcritical fragments, an effect of the ambipolar-diffusion theory studied early on (Mouschovias 1991c), was ignored by the authors. Although it is not clear whether this effect has occurred in the four clouds studied, it is nevertheless important to consider it properly before bold conclusions are drawn on the validitiy of the ambipolar- diffusion theory. \\item The ambipolar-diffusion theory can explain values of $R$ smaller than 1, especially if three-dimensional geometrical effects are properly accounted for. Although such effects are indeed missing from the published (axisymmetric) calculations against which CHT compare their observations, this is hardly new information. Focusing on effects that are known {\\em a priori} to be missing from specific theoretical calculations is not a proper way to test the underlying theory or to make observational progress toward distinguishing between alternative theories. Such observations, when reliable, would provide useful {\\it input} to the ambipolar-diffusion theory. \\end{enumerate}   \\appendix"
    },
    "0807/0807.1354_arXiv.txt": {
        "abstract": "We present a combined analysis of the low-mass initial mass function (IMF) for seven  star-forming regions. We first demonstrate that the ratios of stars to brown dwarfs are consistent with a single underlying IMF. By assuming that the underlying IMF is the same for all seven clusters and by combining the ratio of stars to brown dwarfs from each cluster we constrain the shape of the brown dwarf IMF and find it to be   consistent with a lognormal IMF . This provides the strongest constraint yet that the substellar  IMF turns over ($\\frac{dN}{dM}\\propto M^{-\\alpha}$, $\\alpha < 0$). ",
        "introduction": "We have compiled the ratio of stars to brown dwarfs in nearby, well-studied young embedded clusters and the Pleiades. The regions included in this study are described briefly below, where the ratio of stars (0.08--1.0 $M_\\odot$) to brown dwarfs (0.03--0.08 $M_\\odot$) is calculated. For all the regions, we consider the {\\it system} IMF, uncorrected for multiplicity within 200 AU. The sample is focused on  embedded clusters, in which  spectroscopy has been used to determine the age of the cluster, field star contamination has been taken into account, and an extinction-limited sample has been defined. Furthermore, we have included the Pleiades, because it is one of the best-studied open clusters and bacause its  substellar IMF has been estimated. The break point at 0.08 $M_\\odot$ has been adopted in accordance with the break point for the \\citet{kroupa} IMF, similar to the  characteristic mass in the  \\citet{chabrier} single object IMF\\@. Only a few of the clusters adopted here have the IMF derived in an extinction-limited sample reaching 0.02 $M_\\odot$ and we have opted for 0.03 $M_\\odot$ as a lower mass limit to obtain a larger  sample of clusters.  { \\it Taurus.} \\citet{luhman04} imaged a 4 deg$^2$ region of Taurus that focused on the denser filaments, to identify cluster candidates. Candidates were confirmed as cluster members,by use of follow-up intermediate-resolutionoptical spectroscopy, on the basis of their effective temperature, luminosity, and spectral features. In total, 112 objects were confirmed members with  derived masses between 0.03 and 1.0 $M_\\odot$ and  extinctions A$_\\mathrm{V} \\le 4$ mag. Some  96 objects were stars and 16 were brown dwarfs. Thus, the ratio of stars to brown dwarfs   in Taurus was found to be  $R=96/16=6.0^{+2.6}_{-2.0}$ where the errors are estimated using the method of \\citet{gehrels86}.  {\\it IC 348.} \\citet{luhman03b} imaged a 42\\arcmin$\\times$28\\arcmin\\ region of the IC 348 cluster to identify cluster candidates. By the use of intermediate-resolution spectroscopy, most of the candidates were confirmed as cluster members, on the basis of  their effective temperature, luminosity, and spectral features, which indicated that the objects were  young. In total, \\citet{luhman03b} found 168 cluster members with  masses between 0.03 and 1.0 $M_\\odot$ and  extinctions A$_\\mathrm{V} \\le 4$ mag. The ratio of stars to brown dwarfs was found to be $R=8.3^{+3.3}_{-2.6}$.  {\\it Mon R2.} \\citet{andersen} imaged the central 1\\arcmin$\\times$1\\arcmin\\ of the embedded cluster associated with Mon R2 by utilizing the Near-Infrared Camera andd Multi-Object Spectrometer on board the {\\it Hubble Space Telescope} ({\\it HST}). An extinction-limited sample A$_\\mathrm{V}\\le 10$ mag was defined and a total of 19 objects were detected with masses between 0.03 and 1 $M_\\odot$. The ratio of stars to brown dwarfs was found to be $R=8.5^{+13.6}_{-5.8}$.  {\\it Chameleon 1.} \\citet{luhman07b} obtained an extinction-limited sample  in Chameleon 1 that was complete down to 0.01 $M_\\odot$ for A$_\\mathrm{V}\\le 5$ mag, by use of   observations of a 0.22\\arcdeg $\\times$0.28\\arcdeg\\ region with the Advanced Camera for Surveyes on board {\\it HST} and a subsequent spectroscopic follow-up of cluster member candidates. The sub-sample from 0.03 to $1$ $M_\\odot$ includes  24 objects and the ratio $R$ was found to be $R=4.0^{+3.7}_{-2.1}$.  {\\it Pleiades. } The Pleiades is one of the best-studied open clusters, and numerous derivations of the IMF have been published. Here we focus on the  survey by \\citet{moraux} who covered a 6.4 deg$^2$ region of the Pleiades. The survey had a saturation limit of 0.48 $M_\\odot$. For higher masses, the survey was combined   with a mass function built using the  \\citet{prosser98} database. The Pleiades suffer relatively low ($A_\\mathrm{V} < 1$ mag), mostly uniform, extinction, with negligible impact on the completeness of this sample, so we did not apply a reddening criterion. The ratio of stars to brown dwarfs was found to be $R=4.9^{+1.5}_{-1.2}$.  {\\it The Orion Nebular Cluster.} The ONC has been the subject of extensive studies \\citep{hillenbrand,hillenbrandcarpenter,luhman00,muench}. We take the adopted ratio of stars to substellar objects from the study of \\citet{slesnick}. The total sample, covering the central 5.1\\arcmin$\\times$5.1\\arcmin, contains approximately 200 objects with masses between 0.02 and 0.6 $M_{\\odot}$ and A$_\\mathrm{V} \\le 15$ mag. Using their Figure 14, and extrapolating the slope from 0.08--0.6 to 1.0 $M_{\\odot}$ (one additional bin in their plot), we arrive at a ratio of stars to substellar objects of $R=3.3^{+0.8}_{-0.7}$.    {\\it NGC 2024.} The ratio of stars to brown dwarfs in NGC 2024 was found by \\citet{levine} from their photometric and spectroscopic  study, covering the central 10\\arcmin$\\times$10\\arcmin\\@. They assigned masses to the photometric objects on the basis of the mass distribution in each magnitude bin, determined from the spectroscopic sample as in \\citet{slesnick}.  The result was that a total of 148 objects in their survey area has  masses between 0.02 and 1 $M_{\\odot}$ and extinctions A$_\\mathrm{V}\\le$ 15 mag. Based on their Figure 9, we find  that there are 27 objects between 0.03 and 0.08 $M_{\\odot}$ resulting in a ratio of stars to substellar objects of $R= 3.8^{+2.1}_{-1.5}$.  Table~\\ref{results} shows the ratio of stars to brown dwarfs for nearby embedded clusters and the Pleiades, as described above, and the distribution of ratios is shown in Fig,~\\ref{figure1}. The weighted mean of the ratios is found to be 4.3, and the standard deviation of the weighted mean   is 1.6. All of the measurements presented are consistent with the weighted mean within 2$\\sigma$. There is thus  little evidence for variation in the low-mass IMF between the different regions and we have adopted the hypothesis that the IMF is  universal. Under this assumption, the complete set of IMF determinations can be combined to place  constraints that are stronger than for each of the individual measurements. ",
        "conclusions": "The results on the IMF presented here are based on the system IMF, including binaries unresolved  within 200 AU. As such, they may be difficult to compare directly with the locally derived (within 20 pc) field IMF discussed in \\citet{allen} that suffers from a much smaller fraction of unresolved binaries. Yet the overall binary  frequency for ultra-cool dwarfs (M6 and later) appears to be low \\citep[$\\sim$20\\%,][]{burgasser}, and furthermore the {\\it relative} number of companions with separations $>$15 AU and mass ratios q $>$ 0.4 may be extremely low around very cool stars $\\sim$ 1\\%; \\citet{allen_new}.  Indeed, if the companion mass ratio distribution follows the Chabrier IMF at wide separations, then one could expect fewer very low mass companions as one surveys progressively lower mass primaries \\citep[e.g.][]{siegler}, consistent with the observations by \\citet{mccarthy}. If the  IMF follows a Chabrier IMF in  the brown dwarf regime below 0.03 $M_\\odot$ (say, down to the opacity limit for fragmentation of $\\sim$ 0.001-0.004 $M_\\odot$; \\citet{whitworth}, then the number of stars below 1 $M_\\odot$ will outnumber brown dwarfs 4.7 to 1.  The sense of our results, that the mass function is falling in the BD regime, is consistent with various ideas put forward to explain the shape of the IMF (\\citet{bonnell07} and references therein).  Building on the ideas of \\citet{larson05}, \\citet{bonnellbate} produced an IMF that is only weakly dependent on the Jeans mass  through dynamical interactions in the cluster. However, \\citet{allen_new} show that the turbulent fragmentation models by  \\citet{batebonnell} predict too few low-mass binary systems. \\citet{goodwin}, on the other hand, suggest that the IMF should peak at higher masses in regions with low turbulence (e.g. Taurus) which would result in a higher ratio of stars to brown dwarfs. The lack of a strong variations in the ratio of stars to brown dwarfs is a problem for the turbulence models in general; for example, magnetic turbulence models predict strong variations in the low-mass IMF as a function of Mach number and density \\citep{padoan}. If the preliminary results indicated here are borne out through further observations, then models that depend only weakly on initial conditions would be required (e.g. \\citet{adams96}; \\citet{hennebelle}).  Possible IMF variations at least within 1 kpc are smaller than can be detected by comparing the currently observed  clusters. Thus, there are two challenges in detecting IMF variations: (1) One needs clusters with a well-sampled population to minimize the inherently stochastic nature of  populating an IMF, and  (2)  a larger set of clusters is needed to detect even small IMF variations with initial conditions. Although it appears that the variations in the IMF down to 30 $M_\\mathrm{Jup}$ are modest, we still expect that variations will be seen at the lowest masses  where the opacity limit for fragmentation can be reached \\citep{lowlynden-bell} and the metallicity of the star forming region could be imprinted in the lower mass limit."
    },
    "0807/0807.0164_arXiv.txt": {
        "abstract": "We evaluate the performance of the Multi-conjugate Adaptive optics Demonstrator (MAD) from \\h\\/ and \\ks\\/ imaging of 30~Doradus in the Large Magellanic Cloud.  Maps of the full-width half maximum (FWHM) of point sources in the \\h\\/ and \\ks\\/ images are presented, together with maps of the Strehl ratio achieved in the \\ks-band observations.  Each of the three natural guide stars was at the edge of the MAD field-of-view, and the observations were obtained at relatively large airmass (1.4-1.6).  Even so, the Strehl ratio achieved in the second pointing (best-placed compared to the reference stars) ranged from 15\\% to an impressive 30\\%.  Preliminary photometric calibration of the first pointing indicates 5$\\sigma$ sensitivities of \\ks$\\sim$21.75 and \\h$\\sim$22.25 (from 22 and 12 min exposures, respectively). ",
        "introduction": "Correction of atmospheric turbulence with adaptive optics (AO) is an increasingly important method to improve the performance of large telescopes.  AO-corrected instruments such as NAOS-CONICA\\cite{naco} now routinely deliver excellent image quality in the near-infrared when compared to natural seeing.  However, sky coverage has been limited by the requirement for a bright natural guide star (NGS) with which to monitor atmospheric variations.  The introduction of laser guide stars (LGS) at each of the major 8-m class observatories has opened-up the night sky for AO-assisted imaging and spectroscopy, but the field-of-view (FoV) over which the point spread function (PSF) can be improved remains limited to $\\sim$10-20$''$, except in very favourable conditions.  The concept of multi-conjugate adaptive optics (MCAO) overcomes this limitation, delivering a larger corrected FoV via tomographic reconstruction of multiple turbulent layers using multiple (natural or artificial) guide stars\\cite{mcao}.  As part of the technology development plan towards the European Extremely Large Telescope (E-ELT), the Multi-conjugate Adaptive optics Demonstrator (MAD)\\cite{marchetti2,marchetti3,marchetti} was developed by the European Southern Observatory (ESO) as a visiting instrument for the Very Large Telescope (VLT) in Chile.  Three Shack-Hartmann wavefront sensors are used to observe three NGS across a 2$'$ circular field, thereby allowing mapping of the turbulence.  The turbulence is then corrected using two deformable mirrors (operating at $\\sim$400 Hz), one conjugated to the ground-layer (i.e. 0\\,km), the second conjugated to 8.5\\,km above the telescope.  The high-resolution, near-IR camera used with MAD is the CAmera for Multi Conjugate Adaptive Optics (CAMCAO), designed and built by a consortium led by the University of Lisbon \\cite{camcao1,camcao2}. CAMCAO operates over the {\\em J}, \\h, and \\ks\\/ bands, with critical (2 pixel) sampling of the diffraction-limited PSF at 2.2$\\mu$m.  The detector is a HAWAII2 2k$\\times$2k HgCdTe array, with a pixel scale of 0.028$''$/pixel, giving a FoV of 57$''$x57$''$.  A useful feature is that the camera can be moved within the 2$'$ field without requiring positional offsets of the telescope, meaning that the AO loop can remain closed.  MAD was installed at the Nasmyth visitor focus of UT3 Melipal in early 2007 for technical evaluation.  The performance of the instrument was impressive\\cite{marchetti}, with the observations from the first run already beginning to yield science results\\cite{mad1,mad2}.  Here we present preliminary analysis of MAD \\h- and \\ks-band imaging of the central region of 30~Doradus, the most massive star-formation region in the Large Magellanic Cloud (LMC). ",
        "conclusions": "The observations presented here provide a further demonstration of the excellent performance of VLT-MAD. The image quality (as measured by the FWHM of the PSF) across Fields 1 and 2 was uniform to within 0\\farcs03. Such stability across `wide' fields (compared to classical AO observations) is very encouraging in advance of future multi-laser, MCAO systems on 8-m class telescopes, and the developement of `wide-field', AO-corrected, imagers and spectrometers for the E-ELT.  MAD has also provided us with a truly unique view of the stellar populations in 30~Doradus, in terms of both spatial resolution and spatial coverage compared to previous studies.  Work is now underway on the photometric calibrations and the resulting catalogue will be used to investigate the IMF and star-formation history of this important cluster."
    },
    "0807/0807.0508_arXiv.txt": {
        "abstract": "Oscillations occur in stars of most masses and essentially all stages of evolution.  Asteroseismology is the study of the frequencies and other properties of stellar oscillations, from which we can extract fundamental parameters such as density, mass, radius, age and rotation period.  We present an overview of asteroseismic analysis methods, focusing on how this technique may be used as a tool to measure stellar properties relevant to planet transit studies. We also discuss details of the Kepler Asteroseismic Investigation -- the use of asteroseismology on the Kepler mission in order to measure basic stellar parameters.  We estimate that applying asteroseismology to stars observed by Kepler will allow the determination of stellar mean densities to an accuracy of 1\\%, radii to 2--3\\%, masses to 5\\%, and ages to 5--10\\% of the main-sequence lifetime. For rotating stars, the angle of inclination can also be determined. ",
        "introduction": "Asteroseismology -- the study of stellar oscillations -- is a relatively new and growing research field in astrophysics.  The analysis of frequencies and other properties of stellar oscillations allows us to constrain fundamental parameters of stars such as density, mass, radius, age, rotation period and chemical composition.  Oscillations are found in stars of most masses and essentially all stages of evolution.  The amplitudes and phases are controlled by the energetics and dynamics of the near-surface layers.  The frequencies are determined by the internal sound-speed and density structure of the star, as well as rotation and (in some cases) magnetic fields.  Observationally, the frequencies can be determined with exceedingly high accuracy compared to any other quantity relevant to the internal properties of the stars. Analysis of the observed frequencies, including comparison with computed stellar models, allows determination of the properties of the stellar interiors and tests of the physics used in the model computation, applied under extreme conditions that cannot be matched in terrestrial laboratories.  Rotation induces fine structure in the frequency spectrum, in the form of rotational splitting.  The observed frequencies are determined by averages over the stellar interior of the rotation rate, which in general varies with position within the star. By comparing with independent determinations of the surface rotation rate, or from rotational splittings for a sufficient broad variety of modes, information about this variation can be obtained.  Some of the basic parameters that may be obtained from asteroseismology are crucial for understanding the fundamental properties of exoplanet systems. Hence, asteroseismology will potentially be an excellent tool for characterizing planet transit systems and, specifically, it will be able to measure stellar radii to a relative accuracy of 2--3\\%.  For a recent review on the present state of asteroseismology, see \\cite[Aerts \\etal\\ (2008)]{Aerts08} and \\cite[Bedding \\& Kjeldsen (2008)]{Bedding08}. In Figure~\\ref{fig:fourstars} we show examples of power spectra for four different stars.  In each case, we can clearly see the excess energy arising from stellar oscillations. The detailed properties of those oscillations form the observational basis for the asteroseismic analysis.   \\begin{figure}[b] \\begin{center} \\includegraphics[width=5.4in]{hkjeldsen02.ps} \\caption{\\label{fig:fourstars} Power spectra of time series observed in radial velocity for the four stars $\\xi$ Hydrae (G7\\,III; \\cite[Frandsen \\etal\\ 2002]{Frandsen02}), Procyon (F5\\,IV; \\cite[Arentoft \\etal\\ 2008]{Arentoft08}), $\\alpha$~Centauri~A (G2\\,V; \\cite[Bedding \\etal\\ 2004]{Bedding04}) and $\\alpha$~Centauri~B (K2\\,V; \\cite[Kjeldsen \\etal\\ 2005]{Kjeldsen05}).  The oscillation periods for the four stars are: around 3--4 hours for $\\xi$ Hydrae, 20--25 minutes for Procyon, 7 minutes for $\\alpha$~Cen~A and 4 minutes for $\\alpha$~Cen~B.  These differences, and those in the detailed structure of the power spectra, reflect the differences in stellar properties (radius, mass, surface temperature and age). Note that the vertical scales are normalized -- the actual amplitudes decrease from top to bottom in the figure.  } \\label{fig1} \\end{center} \\end{figure}    \\begin{figure}[b] \\begin{center} \\includegraphics[width=5.2in]{hkjeldsen03.ps} \\caption{\\label{fig:golf} The power spectrum of oscillations in the Sun. The data are full-disk radial velocity measurements obtained over 30 days using the GOLF instrument on board the SoHO spacecraft (\\cite[Ulrich \\etal\\ 2000]{Ulrich00}, \\cite[Garc{\\'{\\i}}a \\etal\\ 2005]{Garcia05}). The inset shows the details of the p-mode structure, and we indicate the so-called large and small separations, which contain information on the basic stellar properties.} \\label{fig2} \\end{center} \\end{figure}  \\begin{figure}[b] \\begin{center} \\includegraphics[width=5.2in]{hkjeldsen04.ps} \\caption{\\label{fig:CDdiagram} The so-called asteroseismic HR-diagram (\\cite[Christensen-Dalsgaard 1993)]{Christensen-Dalsgaard93}, where the axes are the large and small frequency separations ($\\Dnu{}$ and $\\dnu{02}$).  The solid lines are evolutionary tracks (for fixed mass) and the dashed lines show constant core-hydrogen content, which is related to the age.  Based on simulations we have estimated the uncertainties for the measurements of mass, radius and stellar age using the seismic HR-diagram to be 5\\% for the mass, 2--3\\% for the radius and 5--10\\% for the age (relative to the total time on the main sequence). The mean densities can be measured to 1\\%.  The positions of the Sun and $\\alpha$ Centauri A and B are indicated, and it can be seen that $\\alpha$~Cen~A is more evolved than the Sun while $\\alpha$~Cen~B is less evolved, in agreement with traditional modeling.} \\label{fig3} \\end{center} \\end{figure}   \\begin{figure}[b] \\begin{center} \\includegraphics[width=5.2in]{hkjeldsen05.ps} \\caption{\\label{fig:rotation1}A close-up of the power spectrum of solar oscillations, from full-disk radial velocity measurements.  The data are based on 805 days of observing the Sun using the GOLF instrument on board the SoHO spacecraft (\\cite[Ulrich \\etal\\ 2000]{Ulrich00}, \\cite[Garc{\\'{\\i}}a \\etal\\ 2005]{Garcia05}).  We show a small region of the power spectrum just below 2\\,mHz. In this region one can see the power from modes with $l=2$, $n=12$ and $l=0$, $n=13$.  The $l=2$ multiplet has 5 components ($m=0, \\pm1, \\pm2$) but, due to the inclination of the rotation axis of the Sun, only modes with $m=-2$, 0 and 2 are visible. The strength of the individual $m$-components can be used to measure the inclination axis of the stellar rotation.  Modes with $l=0$ have only $m=0$ and are therefore not split by rotation.  The upper curve is a smoothed version of the power spectrum, shifted vertically for clarity.} \\end{center} \\end{figure}   \\begin{figure}[b] \\begin{center} \\includegraphics[width=5.2in]{hkjeldsen06.ps} \\caption{\\label{fig:rotation2} Same as Figure~\\ref{fig:rotation1}, but around a frequency of 3 mHz.  We see the multiplet with $n=20$ and $l=2$, and the singlet with $n=21$ and $l=0$.  } \\end{center} \\end{figure}   \\begin{figure}[b] \\begin{center} \\includegraphics[width=5.2in]{hkjeldsen07.ps} \\caption{\\label{fig:rotation3} Same as Figure~\\ref{fig:rotation1}, but at a frequency just below 2.3\\,mHz.  In this region one can see the power from the triplet with $l=1$, $n=15$, $m=0, \\pm1$.  Due to the inclination of the rotation axis of the Sun, only modes with $m=-1$ and 1 are visible.  } \\end{center} \\end{figure}   \\begin{figure}[b] \\begin{center} \\includegraphics[width=4.8in]{hkjeldsen01.ps} \\caption{Using asteroseismology we are able to measure detailed basic properties of stars that have a transiting planet. Asteroseismology is expected to become a crucial tool for the Kepler mission when measuring stellar parameters.} \\end{center} \\end{figure} ",
        "conclusions": ""
    },
    "0807/0807.0200_arXiv.txt": {
        "abstract": "We study the properties of strange quark matter in equilibrium with normal nuclear matter. Instead of using the conventional bag model in quark sector, we achieve the confinement by a density-dependent quark mass derived from in-medium chiral condensates. In nuclear matter, we adopt the equation of state from the Brueckner-Bethe-Goldstone approach with three-body forces. It is found that the mixed phase can occur, for a reasonable confinement parameter, near the normal nuclear saturation density, and goes over into pure quark matter at about 5 times the saturation. The onset of mixed and quark phases is compatible with the observed class of low-mass neutron stars, but it hinders the occurrence of kaon condensation. ",
        "introduction": "By far the study of neutron stars has been mainly focused on the relationship between the equation of state (EoS) of nuclear matter and the observed maximum mass. The connection has been achieved by solving the hydrostatic equilibrium equations based on general relativity. The first generation observed masses exhibited an average value around 1.5 solar masses. This value requires a soft EoS that can be easily obtained by introducing new degrees of freedom like hyperons, kaons or quarks accompanied or not by a phase transition. Sometimes the softening was so large that the neutron star is predicted to collapse into a black hole, as for the SN1987A \\cite{BB}. In the new generation of observations the masses are distributed within a large range, up to 2 solar masses, that requires a stiff EoS, i.e., hadronic matter without new degrees of freedom. Since it is hard to imagine pure hadronic matter to sustain the high pressure predicted in the inner core, new scenarios have to be advanced to explain the coexistence, in the phenomenology of neutron stars, of low and high mass spectra.  Recently \\cite{HAE} it has been argued that the two observed classes of neutron stars might correspond to two different evolutionary scenarios of neutron stars. In one case, the hot and dense remnant of the supernova explosion rapidly evolves into a hybrid star, where the transition to a quark phase softens the nuclear matter so that $M\\approx 1.5 M_\\odot$; in the other case a slow evolution could lead the neutron star to a large mass via a mass accreting from the coupling with a white dwarf. From this point of view the destiny of the remnant is strongly affected by the initial conditions, i.e. density, temperature, leptonization degree etc. For instance, if the mass of the remnant is below the mass threshold for quark nucleation the transition to the quark phase is forbidden \\cite{bombaci}. If the mass is slowly accreting the transition is allowed. The role of temperature or other parameters defining the initial state of a new born neutron star has not yet been studied.  To investigate the possible phase transition to quark matter in neutron stars, we need also to know the EoS of quark matter. Although we have in hand the fundamental theory of strong interactions, i.e., Quantum Chromodynamics (QCD), we still do not know the true ground state. It is now generally expected that quark matter is in the color-flavor locked phase (CFL) \\cite{Alford1999} at extremely high densities when the finite current mass of strange quarks becomes unimportant. In the density range from nuclear saturation to CFL, there may exist a rich and varied landscape of phases, e.g., the 2SC, g2SC, gCFL etc. Presently, however, these phases suffer from the so-called chromomagnetic instability problem for both the two- \\cite{HuangM2004} and three-flavor \\cite{Casalbuoni2005,Alford2005,Fukushima2005} cases. On the other hand, experiments show that quarks become asymptotically free rather slowly \\cite{LiTD2005}. Therefore, in the present study we are dealing with the ordinary strange quark matter (SQM) \\cite{WittenPRD,Jaffe}.  The special problem in studying the EoS of ordinary quark matter is to treat quark confinement in a proper way. In the conventional standard approach, an extra constant term, the famous bag constant $B$, is added to the energy density of the system, which provides a negative pressure to confine quarks within a finite volume, usually called a `bag'. The quark mass is infinitely large outside the bag, and a finite constant within the bag. A vast quantity of investigation have been performed within the framework of the bag model \\cite{WeinerIJMPE15}.  As is well known, however, particle masses vary with environment. Such masses are usually called effective masses. Effective masses of hadrons and quarks have been extensively discussed, e.g., within the Nambu-Jona-Lasinio model \\cite{Buballa99plb} and within a quasi-particle model \\cite{SchertlerNPA616}. In principle, the density dependence of quark masses should be connected to the in-medium chiral condensates \\cite{LombardoPRC72,PenggxNPA747}.  Taking advantage of the density dependence, one can describe quark confinement without using the bag constant. Instead, the quark confinement is achieved by the density dependence of the quark masses derived from in-medium chiral condensates \\cite{PenggxPRC61,PenggxPRC62,WenxjPRC72}. The two most important aspects in this model are the quark mass scaling~\\cite{PenggxPRC61,WenxjPRC72} and the thermodynamic treatment~\\cite{PenggxPRC62,WenxjPRC72}. Both aspects will be reviewed in this paper.  In the present contribution, the transition from hadron phase (HP) to strange quark phase (SQP) in the inner core of a neutron star is investigated within the fully consistent nuclear and quark models. In the hadron sector we adopt the equation of state from Brueckner-Bethe-Goldstone (BBG) approach with three-body forces (TBF) \\cite{grange:1989,Zuowei02prc,shq1998prl}. This theory, being a completely microscopic approach, can easily incorporate degrees of freedom such as nucleon resonances [$\\Delta$(1232) or $N^*$(1440)], which are expected to appear at higher hadron densities. It is found that the mixed hadron-quark phase can occur, for reasonable values of the confinement parameter, a little above the normal saturation density, and can undergo the transition  to pure quark matter at about 5-6 times the saturation. This result is quite different from the previous results from Nambu-Jona-Lasinio (NJL) model in which the mixed quark phase can not appear at neutron-star densities~\\cite{Schertler99prc,njl}. Afterwards, the influence of the mixed and quark phases on the structure of compact stars is discussed by solving the Tolman-Oppenheimer-Volkov (TOV) equation and extracting the mass-radius plots for neutron stars. Finally, it is shown that the transition to the deconfined phase turns out to be incompatible with the onset of kaon condensation. ",
        "conclusions": "In this paper the neutron star inner structure has been studied. The lack of strong observational constraints demands for sophisticated models of the NS composition and interaction mechanisms. In this study we only included baryons and quarks in equilibrium with leptons and kaons. Since in the NS interior high baryon densities are reached, N$\\bar{\\mathrm{N}}$ and nucleonic excitations are expected to play a  major role. Their effects can be incorporated in a three body force. The baryon EoS in weak coupling equilibrium with electrons was derived within the BBG theory suitably extended so to include the three body force. In the quark sector, we adopted the semiphenomenological CDDM quark model, which exhibits a confinement mechanism alternative to the crude MIT bag model. Furthermore,  in contrast with the extension of the MIT model, where the density dependence is introduced artificially \\cite{ddbag}, CDDM quark model shapes its density dependence in agreement with chiral requirements \\cite{LombardoPRC72,PenggxNPA747}.  Hadrons and quarks in $\\beta$-equilibrium were considered at zero and finite temperature. Since typical temperatures of protoneutron stars are as high as $40$ MeV and beyond, rigorous thermodynamic potentials have to be derived. This was possible with CDDM quark model, due to its simplicity, whereas some approximations are needed to calculate the hadron phase in the Brueckner theory. The latter is still a main drawback of the finite temperature microscopic theory of strongly interacting Fermi systems.  The transition from the low density hadron phase to high density quark phase in beta equilibrium was studied under the Glendenning hypothesis of total charge neutrality. The Gibbs construction enabled to follow the evolution of the mixed hadron-to-quark phase, varying the temperature $T$ and the confinement parameter $D$. The EoS, in terms of pressure, energy density and chemical composition, was constructed as a function of $T$ and $D$. Moreover the phase diagram $T$-$D$ was also depicted. The transition density from hadron to mixed phase is strongly dependent on the confinement parameter and, for D small enough, it becomes lower than the nuclear saturation density. However, at that point nuclear matter is so strongly neutron rich that it hard to imagine that it can be observed in terrestrial laboratory experiments. At high temperature the transition density is comparatively smaller.  The TOV equations have been solved in the above discussed model of neutron stars. The transition to quark matter produces a strong reduction of the maximum mass, from $M=2.3M_\\odot$ to $1.5M_{\\odot}$ for the lowest value of the confinement parameter $D$, but the corresponding radius is not changed, being in both cases about $10$ km. The other $M$-$R$ configurations ly in a range $R\\simeq 10\\sim 13$ km depending on the value of $D$. With increasing $D$, the maximum mass increases from $1.5$ to $1.7$ solar masses. Other quark models have been adopted to describe the NS transition to a deconfined phase (see, for instance, Refs.~\\cite{Schertler99prc,njl,nico}). Except for the Nambu-Jona Lasinio model which exhibits instability of neutron stars \\cite{njl}, the other models point to a softening of the EoS of nuclear matter. The MIT model with density-dependent bag constant predicts a beta-stable quark phase quite similar to our model \\cite{nico}, but the comparison on the NS structure turns out to be difficult, since in that calculation the transition to quark phase is built on top of hyperonized nuclear matter.  The effect of temperature on the maximum mass is negligible up to $T=20$ MeV, but as expected the heated system can support much less gravitational mass in the same volume for all $M$-$R$ configurations. Thus, an evolutionary study of isolated neutron stars would show a strong compression from the newborn phase to the long era phase as fast as the star cools down from $40$ MeV to approximately zero.  Beyond the confinement parameter, the transition to quark phase is interrelated to other properties, in particular the possible onset of kaon condensation. In our analysis we concluded that the quark phase could not be compatible with the kaon condensation in neutron stars, or at least could hinder it very strongly.  The model developed in this paper misses some important aspects. The first one  is the effect of neutrinos. Neutrinos play the major role in new born neutron stars, when they are still trapped in the interior. In this case the neutronization process is still hindered and the EoS is that of symmetric nuclear matter. Moreover, the onset of kaon condensation is shifted to higher densities \\cite{kao2} and probably kaons have no more any chance to compete with quarks.  The second one is the inclusion of hyperons and their competition with other mechanisms such as kaon condensation. In Refs.~\\cite{nico} the hybrid stars are studied by including hyperons in the hadron phase, which makes the hadron EoS very soft and brings the maximum mass to $M=1.25M_\\odot$. The introduction of the quark phase rises up again the maximum mass to $1.5M_\\odot$, a value which is consistent with our prediction without hyperons. On the other hand, the appearance of hyperons much depends on the threshold of kaon production.  % Therefore the interplay between hyperons and kaons turns out to be quite important and deserves additional investigation, as soon as more reliable empirical inputs will be available, especially on the hyperon-nucleon and hyperon-hyperon interaction."
    },
    "0807/0807.0941_arXiv.txt": {
        "abstract": "We explore the stellar population of M31 in a Spitzer Space Telescope survey utilizing IRAC and MIPS observations. Red supergiants are the brightest objects seen in the infrared; they are a prominent evolutionary phase. Due to their circumstellar envelopes, many of these radiate the bulk of their luminosity at IRAC wavelengths and do not stand out in the near infrared or optically. Going fainter, we see large numbers of  luminous asymptotic giant branch stars, many of which are known long period variables. Relative to M33 the AGB carbon star population of M31 appears sparse, but this needs to be spectroscopically confirmed. ",
        "introduction": "The disk and bulge of M31 were first resolved into stars by \\cite{hub36} and \\cite{baa44}, respectively. Since then, many optical surveys have been carried out, including those of \\cite{hs53}, \\cite{bs63}, \\cite{hod76}, \\cite{ber88}, \\cite{mag92}, \\cite{moc01} and \\cite{mas06}. More recent papers that included spectra and optical and near-infrared photometry of individual stars are \\cite{h79}, \\cite{h88}, \\cite{hmf90}, \\cite{h94} and \\cite{mas98}. Near-infrared point-source photometry of M31 has been hampered by the lack of large-format detectors; for example, the recent work of \\cite{ols06} covers only a small portion of M31. The 2MASS Large Galaxy Atlas \\citep{jar03} presented integrated photometry of M31, but no galaxy-wide assessment of the M31 point source population has been published.  Spitzer Space Telescope characterization of the point source population of other Local Group galaxies includes the Large Magellanic Cloud \\citep{blu06}, M33 \\citep{ver07,mcq07}, and NGC 6822 \\citep{can07}. Other Local Group dwarfs studied include IC1613 and the WLM galaxy \\citep{jaa07,jab07,jac06}. The Andromeda galaxy, M31, presents a unique opportunity to study stellar evolution at a chemical composition similar to that of the Milky Way. The red giant branch was resolved in V \\& I by \\cite{mk86},  and the brightest evolved stars have been studied by \\cite{mas98}. However, an infrared survey of point sources in M31 has awaited the arrival of the Spitzer Space Telescope \\citep{wer04,geh07} and 2MASS \\citep{jar03}.  In the mid-infrared, integrated measurements of M31 have been made by IRAS \\citep{hab84,xuh96}, COBE/DIRBE \\citep{ode98},  MSX \\citep{kra02} and {\\it Spitzer} \\citep{gor06,bar06}. With the aim of reaching stars with the most significant mass loss rates this paper presents a point source survey with the IRAC \\citep{faz04} and MIPS \\citep{rie04}  cameras. We also utilize the 2MASS survey in order to include shorter wavelengths. We present extensive tables, giving positions of the sources we identify, to enable follow-up of these observations. ",
        "conclusions": "Our Spitzer Space Telescope survey reveals a rich population of evolved stars in M31. Although there is strong foreground contamination, supergiants are present as bright as $1 - 2.5 \\times 10^6$~$L_\\odot$. These evolve into stars with extensive circumstellar envelopes like $\\eta$~Carinae and VY CMa in the Milky Way. A number of the brightest supergiants at 3.6\\mic are so red that they will be missed in optical surveys. We also resolve stars below the tip of the AGB to $M_{\\rm bol} = -5$~mag. Many of these are LPVs, which also have dusty envelopes. Photometric candidates for AGB carbon stars are much fewer in M31 than M33, but spectroscopy is required to quantify this."
    },
    "0807/0807.3345_arXiv.txt": {
        "abstract": "A specialized data mining algorithm has been developed using wide-field photometry catalogues, enabling systematic and efficient searches for resolved, extremely low surface brightness satellite galaxies in the halo of the Milky Way (MW). Tested and calibrated with the Sloan Digital Sky Survey Data Release 6 (SDSS-DR6) we recover all fifteen MW satellites recently detected in SDSS, six known MW/Local Group dSphs in the SDSS footprint, and 19 previously known globular and open clusters. In addition, 30 point source overdensities have been found that correspond to no cataloged objects. The detection efficiencies of the algorithm have been carefully quantified by simulating more than three million model satellites embedded in star fields typical of those observed in SDSS, covering a wide range of parameters including galaxy distance, scale-length, luminosity, and Galactic latitude.  We present several parameterizations of these detection limits to facilitate comparison between the observed Milky Way satellite population and predictions. We find that all known satellites would be detected with $>90\\%$ efficiency over all latitudes spanned by DR6 and that the MW satellite census within DR6 is complete to a magnitude limit of $M_V\\approx-6.5$ and a distance of 300 kpc. Assuming all existing MW satellites contain an appreciable old stellar population and have sizes and luminosities comparable to currently known companions, we predict lower and upper limit totals of 52 and 340 Milky Way dwarf satellites within $\\sim260$ kpc if they are uniformly distributed across the sky. This result implies that many MW satellites still remain undetected. Identifying and studying these elusive satellites in future survey data will be fundamental to test the dark matter distribution on kpc scales. ",
        "introduction": "The dwarf galaxy population of the Milky Way (MW) provides invaluable insight into galaxy formation and evolution.  Their resolved stars reveal their formation histories and enable precise measurements of their structural parameters, ages and metallicities.  These histories of individual, nearby systems provide a unique approach to studying the Universe across the cosmic ages. Dwarf galaxies are also the most numerous type of galaxy in the Universe and are thought to be the building blocks of larger galaxies. Owing to their low masses, their properties may be strongly influenced by ionizing radiation in the early Universe and by the energy released by supernovae.  The impacts of both of these are weak links in our understanding of structure formation. Finding and studying nearby dwarfs of the lowest masses and luminosities is thus an essential component to understanding galaxy formation on all scales.  The Milky Way dwarf galaxy population is also at present the most direct tracer of the abundance, mass spectrum, characteristic size, and spatial distribution of dark matter on sub-galactic scales. Standard $\\Lambda$CDM simulations of MW-size dark matter haloes predict many more dark matter sub-haloes than are observed as dwarf galaxies \\citep{Klypin99,Moore99}. The recent ``Via Lactea'' simulation contains 2,000 dark matter sub-halos within 289 kpc of the simulated primary galaxy \\citep{vialactea} which have no observed optically luminous counterparts. This discrepancy gives rise to the questions how and in what mass regime do the baryons disappear from dark matter clumps? Studies of the spatial distributions of Milky Way and M31 dwarf galaxy companions have also highlighted possible discrepancies between $\\Lambda$CDM theory and observations \\citep{Kroupa05,Kang05,Metz07}.  The most obvious reason for these apparent discrepancies in the number and spatial distributions of dwarf galaxies is substantial incompleteness as the Milky Way halo has not yet been uniformly searched for dwarf galaxy companions to low enough luminosities and star densities, in particular close to the Galactic plane where foreground contamination is severe.  For example, \\cite{Willman04} compared the spatial distribution of MW satellites to that of M31's population, as well as that of a simulated dark matter halo and concluded that some dwarfs may have been missed at low Galactic latitudes, and that the total number of MW satellites with properties similar to the known objects could be as many as triple the known population.  The viability of this solution, at least in part, has been underscored by the recent discoveries of fourteen new Galactic companions from the photometric data of the Sloan Digital Sky Survey (SDSS). These objects all appear to be dominated by old ($\\gtrsim10$ Gyr) stellar populations, with the exception of Leo T \\citep{leotpop}. Nine of these companions were immediately identified as dwarf spheroidal (dSph) galaxies: Ursa Major, Canes Venatici, Bo\\\"{o}tes, Ursa Major II, Canes Venatici II, Hercules, Leo IV, Coma Berenices, Leo T and Leo V \\citep{uma,cvn,bootes,uma2,quintet,leot,leov}. Spectroscopic follow-up has confirmed that they all are highly dark matter dominated dwarf galaxies \\citep{simon07,spectro}.  Willman 1, Segue 1 and Bo\\\"otes II \\citep{willman1,quintet,boo2} occupy a region of size-luminosity space not previously known to be inhabited by old stellar populations, at the intersection of MW dSphs and globular clusters. Spectroscopic studies \\citep{spectro} showed that Willman 1 may be resident inside a dark matter subhalo with a mass-to-light ratio of $\\sim470$. If these ambiguous objects are gravitationally bound, then tidal arguments also favor them being dark matter dominated \\citep{boo2mmt}. The remaining two objects discovered in SDSS, Koposov 1 and 2, \\citep{k12} are extremely faint Galactic globular clusters.  Numerous authors have shown that the predictions of $\\Lambda$CDM simulations can be reconciled with the small number of observed MW dwarf galaxies if simple models for baryonic physical processes are taken into account when interpreting the results of numerical simulations (e.g. \\citealt{bullock01,tumultuous,simon07}).  For example, \\citet{strigari07} show that the central masses ($M_{0.6 kpc}$) of the Milky Way dwarf galaxies are well-constrained by the data and that their mass function closely matches the $M_{0.6 kpc}$ mass function of both the earliest forming sub-halos and the most massive accreted sub-halos in the Via Lactea simulation.  A well-defined, deep survey of Milky Way dwarf galaxies over a large fraction of the sky is critical to assess any of the above scenarios. The dwarf galaxies detected (or not) by such a survey will provide one of the best ways to rigorously test the $\\Lambda$CDM paradigm by comparing a variety of metrics (distribution, mass, scale sizes, and number) of the Milky Way dwarfs with the predictions of $\\Lambda$CDM + galaxy formation models. Willman et al (2002) and \\cite{koplf} have previously conducted automated searches for Milky Way dwarfs in the SDSS and their corresponding detection limits.  The original Willman et al survey was only performed over a couple of thousand square degrees of sky.  The \\cite{koplf} survey was performed with a more sensitive algorithm (critical, because they found many new satellites to be on the edge of detectability), but few galaxies were used to accurately quantify their detection limits.  In this paper, we present critical improvements to the present characterization of the detectability of Milky Way dwarf galaxies over the $\\sim$9,500 square degrees of the SDSS in Data Release 6 (DR6).  We also present an improved detection algorithm over previous searches.  We aim to construct the most complete, well-defined census of MW satellites by embarking on a Milky Way all sky satellite hunt.  This search will ultimately combine SDSS Data Release 6 (\\citealt{dr6}), The 2-Micron All Sky Survey (2MASS; \\citealt{2mass}) and the upcoming Southern Sky Survey \\citep{SkyMapper}.  The Southern Sky Survey will cover the entire 20,000 deg$^2$ below $\\delta<0^\\circ$ using the new Australian National University (ANU) SkyMapper Telescope equipped with a 5.7 sq degree wide- field camera that is currently under construction with survey operation expected to commence early 2009. ",
        "conclusions": "The dwarf galaxy satellites of the Milky Way provide excellent opportunities to further our understanding of galaxy formation and near-field cosmology. They can be resolved into individual stars allowing detailed studies of their structure, kinematics and composition. They have also been cause for concern regarding their interpretation as the luminous components of dark matter substructure; it has been argued that the number and spatial distribution of these satellites are inconsistent with $\\Lambda$CDM structure formation scenarios. The commencement of the Sloan Digital Sky Survey triggered a cascade of discoveries, with fourteen new satellites discovered. The limited spatial coverage and photometric depth of SDSS suggests that many, if not most, MW satellites are still yet to be discovered. The coming years are likely to bring the MW satellite census towards completeness as new survey telescopes such as SkyMapper, Pan-STARRS and LSST come online.  We present here the method used to search the Sloan Digital Sky Survey Data Release 6 for ultra-faint Milky Way satellite galaxies. By screening for stars consistent with an old population at a fixed distance, we enhance the signal of a dSph over the Milky Way foreground. Smoothing with a kernel corresponding to the expected surface density profile further elevates the dSph above the foreground, and our comprehensive thresholds account for varying stellar density and more diffuse objects.  Applying our algorithm to SDSS DR6, we recover the ``classical'' and recently discovered dSphs, as well as 17 globular clusters and two open clusters. We also have 30 unidentified detections, some of which may be new satellites. The discovery of Leo V demonstrates the difficulty in following-up dwarf candidates; while we detect Leo V, there are several unknown detections of greater significance that may prove to be something. However observing these weakest candidates is a rather hit and miss affair, as pointed out by \\cite{leov}.  To compare the known dwarf galaxy population of the Milky Way with predictions, its essential to have a very well-defined dwarf selection function.  To do this, we thoroughly model the detection efficiency of systems covering a wide range in parameter space by simulating more than 3,000,000 galaxies. We fit various functions to the resulting detection efficiency contours to semi-analytically describe efficiency as a function of magnitude, size, distance and Galactic latitude. Using the results of our detailed investigation of dwarf detectability, we show that:  \\begin{itemize} \\item{Assuming a Via Lactea subhalo radial distribution and that $R_{complete}=R_{90}$, there should be $\\sim13$ satellites with DR6, 12 of which are known. If we include Segue 1, Willman 1, and Bo\\\"otes II in this calculation, this estimate jumps to $\\sim56$, only 15 of which are known.}  \\item{Dwarf detectability shows a smooth transition from 100\\% to 0\\% over size and distance. For example, the distance at which a CVn II-like object is detected with 90\\% efficiency is 200 kpc, compared to 316 kpc for 10\\% efficiency.}  \\item{Galactic latitude does not significantly impact the detection of satellites over the DR6 footprint, and surveys of similar quality should still detect dwarfs as low as $b\\approx15^{\\circ}$. All of the satellites discovered in SDSS would have been detected at any latitude.}  \\item{The census of MW satellites brighter than $M_V=-6.5$ should be complete out to 300 kpc, and all objects brighter than $M_V=-5$ would be detected with at least 50\\% efficiency out to this distance.}  \\item{Given the present data, the spatial anisotropy of the MW satellites within DR6 cannot be confirmed or ruled out.} \\end{itemize}  We provide several different parameterizations of our detection limits to facilitate comparisons between the known Milky Way dwarf galaxy population and predictions.  We provide software that returns the detection efficiency of a dwarf galaxy as a function of its luminosity, scale size, distance, and latitude.  There are two different codes provided for this; one is based on an analytic description of our detection limits and the other provides a direct interpolation from our 3,825,000 simulated galaxies.  We also provide a linear fit as a function of $M_V$ of the distance out to which dwarfs are detected with each of 50\\%, 90\\%, and 99\\% efficiency. These fits assume an underlying dwarf galaxy population with combinations of sizes and luminosities similar to those known.  2009 will bring about the beginning of the Southern Sky Survey, and with it a way to uniformly search a further 20,000 square degrees of sky for new MW dwarfs. We will apply our algorithm to the incoming data and produce the most complete and well characterized census of the MW neighbourhood possible to date. We may also implement improvements to the algorithm to optimize for stellar streams or young stellar populations. \\cite{lgtrio} have shown that M31 satellites occupy as yet unexplored size-luminosity space around the Milky Way. Surveys beyond SDSS and SkyMapper, such as Pan-STARRS and LSST, will be needed to carefully search for such systems.  Even with our carefully characterized detection limits the true number of MW satellites remains highly uncertain."
    },
    "0807/0807.3459_arXiv.txt": {
        "abstract": "Fermi acceleration can take place at ultra-relativistic shock waves if the upstream or downstream magnetic field has been remodeled so that most of the magnetic power lies on short spatial scales. The relevant conditions under which Fermi acceleration become efficient in the presence of both a coherent and a short scale turbulent magnetic field are addressed. Within the MHD approximation, this paper then studies the amplification of a pre-existing magnetic field through the streaming of cosmic rays upstream of a relativistic shock wave. The magnetic field is assumed to be perpendicular in the shock front frame, as generally expected in the limit of large shock Lorentz factor. In the MHD regime, compressive instabilities seeded by the net cosmic-ray charge in the shock precursor (as seen in the shock front frame) develop on the shortest spatial scales but saturate at a moderate level $\\delta B/B \\sim 1$, which is not sufficient for Fermi acceleration. As we argue, it is possible that other instabilities outside the MHD range provide enough amplification to allow successful Fermi acceleration. ",
        "introduction": "The physics of collisionless shock waves is of paramount importance in modern astrophysics, as it governs a variety of phenomena, most notably the emission of high energy radiation. In particular, the afterglow radiation of gamma-ray bursts is generally attributed to the synchrotron emission of electrons that have been accelerated around the relativistic external shock wave, with shock Lorentz factor $\\Gamma_{\\rm sh}\\,\\sim\\,100$ (Paczy\\'nski \\& Rhoads 1993, Katz 1994, M\\'esz\\'aros \\& Rees 1997, Sari \\& Piran 1997, Vietri 1997, Waxman 1997; see Piran 2005 for a review). In this framework, it can be shown that the magnetic field, downstream of the shock wave, must have been amplified by orders of magnitude beyond the shock compression of a typical interstellar magnetic field (see for instance Waxman 1997; see also Piran 2005 and references therein). Recent work by Li \\& Waxman (2006) further shows that even the upstream magnetic field must have been amplified by at least one or two orders of magnitude in order to explain observed $X-$ray afterglows.  These results should be put in perspective with recent studies on the process of Fermi acceleration in the test particle limit around relativistic shock waves. In particular, Niemiec \\& Ostrowski (2006) have shown that Fermi acceleration becomes inoperative in a turbulent magnetic field with power spectra of the Kolmogorov or scale invariant type. Through analytical calculations, Lemoine, Pelletier \\& Revenu (2006) have shown that relativistic Fermi acceleration is inefficient if the magnetic field power is distributed on scales larger than the typical Larmor radius of the accelerated population, in agreement with the above numerical result. Therefore the interpretation of the gamma-ray burst afterglow emission from the relativistic external shock wave requires the magnetic field to have been amplified on short spatial scales.  It has been suggested that the relativistic two stream Weibel instability could amplify the downstream magnetic field to the level required by afterglow modeling of gamma-ray bursts (Gruzinov \\& Waxman 1999; Medvedev \\& Loeb 1999). Subsequent studies have however argued that this instability should saturate early on (Wiersma \\& Achterberg 2004; Lyubarsky \\& Eichler 2006). The question of the long term evolution of the downstream magnetic field (on timescales $\\gg\\,\\omega_{p}^{-1}$, with $\\omega_{p}$ the plasma frequency) also remains open. Ongoing particle-in-cell simulations should eventually shed light on this issue (Silva {\\it et al.} 2003; Frederiksen {\\it et al.} 2004; Medvedev {\\it et al.} 2005; Kato 2007; Chang, Spitkovsky \\& Arons 2008; Keshet {\\it et al.} 2008; Spitkovsky 2008).  Regarding the amplification of the upstream magnetic field, the present situation is reminiscent of results obtained for non-relativistic supernovae remnant shock waves, where the interstellar magnetic field has apparently been amplified by one or two orders of magnitude (see V\\\"olk, Berezhko \\& Ksenofontov 2005, Parizot {\\it et al.} 2006).  In this case, the leading candidate for the instability is the streaming instability, seeded by the cosmic ray precursor in the upstream plasma (see Bell 2004, 2005; Pelletier, Lemoine \\& Marcowith 2006, Marcowith, Lemoine \\& Pelletier 2006; Reville {\\it et al.} 2007; Niemiec et al. 2008; Amato \\& Blasi 2008; Zirakashvili, Ptuskin \\& V\\\"olk 2008; Reville et al. 2008). The generalization of this instability to the relativistic regime has been studied on phenomenological grounds by Milosavljevi\\'c \\& Nakar (2006) who have concluded that it should be able to account for the degree of amplification inferred in gamma-ray bursts. More recently, Reville, Kirk \\& Duffy (2006) have derived in detail the dispersion relation for a parallel shock wave and the saturation due to thermal effects.  In the present work, we propose to explore in detail the generalization of this type of instability to the ultra-relativistic regime $\\Gamma_{\\rm sh}\\,\\gg\\,1$. One crucial difference with the previous works on the streaming instability is that we consider the most natural case of superluminal shock waves (with a magnetic field perpendicular to the shock normal in the shock front frame). This case is more generic than the parallel configuration studied previously because the transverse component of the magnetic field is boosted by the shock Lorentz factor when going to the shock frame. Another important difference is that we bring to light a new type of instability, of a compressive nature. Finally, in contrast with most particle-in-cell simulations performed to date, our study focusses on magnetized shock waves, for which there exists a coherent upstream magnetic field (whose dynamical influence on the shock jump conditions can be neglected however). Nevertheless there exist pioneering PIC simulations in the moderately relativistic regime, which include both a mean field and a significant mass ratio between electrons and ions, see Hededal \\& Nishikawa (2005), Dieckmann, Shukla, \\& Drury (2008).  We adopt a simplified description in which the cosmic-ray distribution is modeled as a step function out to some distance $\\ell_{\\rm cr}$ and we neglect the cosmic-ray response to the disturbance. This latter assumption is justified by the fact that the instability is maximal on the shortest spatial scales, orders of magnitude below the typical Larmor radius of accelerated particles.  The paper is organized as follows. In Section~\\ref{sec:general}, we introduce the main scales of the problem, most notably the diffusion scale of the cosmic rays; we then calculate the level of amplification that is necessary to make Fermi acceleration operative. Section~\\ref{sec:instab} is devoted to the investigation of the instabilities under the condition that some cosmic-rays have undergone a first Fermi cycle. Section~\\ref{sec:conc} summarizes out results and provides some outlook.  Details of the calculations are provided in Appendix~\\ref{sec:appperp}. ",
        "conclusions": "\\label{sec:conc} In this study, we have examined the amplification of a pre-existing magnetic field upstream of an ultra-relativistic shock wave. We have assumed that the magnetic field is fully perpendicular to the shock normal in the shock front frame, as generally expected in the ultra-relativistic limit $\\Gamma_{\\rm sh}\\,\\gg\\,1$. In the shock frame, the cosmic rays do not induce any current at zeroth order, only a net charge density, which is partly screened by the inflowing plasma. This charge distribution then triggers an instability on very short spatial scales, with a growth rate increasing with the wavenumber. Cosmic rays do not respond to this amplification as it takes on scales much shorter than the typical Larmor radius.  It is important to stress that these instabilities are of a different nature than the resonant or non-resonant instability studied by Bell (2004) in the case of a non-relativistic parallel shock wave, most notably because they are essentially {\\it compressive}. Since density depletions are naturally limited, our linear analysis indicates that these instabilities saturate at a moderate level of amplification, $\\delta B/B_0 \\, \\sim\\, 1$. Furthermore, there is no unstable incompressible mode within the present MHD approximation which could provide a higher level of amplification. Therefore, we conclude that, within the framework of ideal MHD, instabilities at ultra-relativistic magnetized shock waves appear limited in their efficiency.  Such instabilities cannot therefore account for the degree of amplification that has been inferred from the modeling of the afterglow radiation of gamma-ray bursts. We have also argued in Section~\\ref{sec:fermi} that the ratio $\\delta B/B_0$ sets the dynamic range of energies over which powerlaw spectra can be produced through Fermi acceleration. The present instabilities thus cannot allow successful Fermi acceleration. Nevertheless, it would certainly be useful to pursue the present investigation with dedicated numerical simlations, which would allow one to go beyond the present linear approximation and study whether non-linear effects can push the small scale magnetic field to higher values.  If the standard interpretation of gamma-ray burst afterglows as the synchrotron light of electrons accelerated at the forward shock (but see Uhm \\& Beloborodov 2006, Genet, Daigne \\& Mochkovitch 2007 for recent alternatives involving the reverse shock) holds, one is led to conclude that some other instability is able to amplify the upstream magnetic field on short spatial scales, or that some other form of acceleration mechanism is operating (see for instance Hoshino 2008 for an alternative involving radiative pressure effects). The present work indicates that such cosmic ray induced instabilities would involve non-MHD effects. The short spatial scales $\\lambda_{\\rm c}\\,\\ll\\,r_{\\rm L}/\\Gamma_{\\rm sh}$ required are in conflict with the usual synchrotron resonance between cosmic rays and MHD waves; this remark is actually one of the motivation of the present work in a full MHD framework. However, one cannot rule out that other types of resonance could become relevant and yet involve incompressible modes, such as Alfv\\'en waves. One promising instability, currently under study, involves a Cerenkov resonance ($\\omega = k_xc$) between plasma waves and the cosmic ray beam, generating a modified two stream instability. The growth rates appear quite promising and the details will be given in a forthcoming article.   \\appendix"
    },
    "0807/0807.4615_arXiv.txt": {
        "abstract": "We report the discovery of a faint ($\\rm L_x \\sim 4 \\pm 1.5 \\times 10^{37}$ \\ergs, 0.5-2 keV), out-flowing gaseous hot interstellar medium (ISM) in NGC 3379.  This represents the lowest X-ray luminosity ever measured from a hot phase of the ISM in a nearby early type galaxy.  The discovery of the hot ISM in a very deep {\\it Chandra} observation was possible thanks to its  unique spectral and spatial signatures, which  distinguish it from the integrated stellar X-ray emission, responsible for most of the unresolved emission in the {\\it Chandra} data.  This hot component is found in a region of $\\sim 800$ pc in radius at the center of the galaxy and has a total mass M$\\sim 3 \\pm 1 \\times 10^5$ M$_{\\sun}$.  Independent theoretical prediction of the characteristics of an ISM in this galaxy, based on the intrinsic properties of NGC~3379, reproduce well the observed luminosity, temperature, and radial distribution and mass of the hot gas, and indicate that the gas is in an outflowing phase, predicted by models but not observed in any system so far. ",
        "introduction": "The nature of the X-ray emission of elliptical galaxies has been debated since their discovery as X-ray sources with the {\\it Einstein Observatory}, which led to the early suggestion of ubiquitous hot X-ray emitting halos in hydrostatic equilibrium  (Forman et al. 1985). These halos would originate from the stellar ejecta resulting from their normal evolution, accumulated during the lifetime of the galaxies, and retained by their deep dark matter potentials. However, from the analysis of these early observations and of ROSAT and ASCA data, a more complex picture of the X-ray emission emerged, suggesting that the non-nuclear X-ray emission of these galaxies consist of varying amounts of hot ISM, and of a baseline component related to the stellar population, either from the finale stages of stellar evolution (i.e. low-mass X-ray binaries, LMXBs, Trinchieri \\& Fabbiano 1985; Canizares et al. 1987; Kim, Fabbiano \\& Trinchieri 1992; Fabbiano et al. 2006; Matsushita et al. 1994) or possibly from coronal stellar emission (Pellegrini \\& Fabbiano 1994). Although controversial at the time, this picture has important implications for the evolutions of the stars and of the gaseous components of ellipticals, suggesting halo retention in the most massive systems, and partial or full winds in most galaxies (Ciotti et al, 1991; David et al. 1991; Pellegrini \\& Ciotti 1998).  With {\\it Chandra} for the first time the X-ray emission of nearby ellipticals can be imaged, and the existence of the different emission components can be pursued by direct observational means. These observations have led to the detection and study of LMXB populations (see review in Fabbiano 2006), and the removal of their contamination from the diffuse emission component (e.g. NGC~1316, Kim \\& Fabbiano 2003). In some galaxies this diffuse emission is due to gaseous halos; however, in other galaxies, where the X-ray luminosity function can be probed to very low luminosities and the residual emission is itself of low luminosity, unresolved LMXBs and possibly stellar emission may account for a large amount or even the bulk of the X-ray luminosity (e.g.NGC821, Pellegrini et al 2007a,b; M32, Revnintsev et al. 2007).  We can now attempt to constrain  the presence and properties of vestigial hot halos in these gas-poor galaxies. Since ellipticals are expected to host super-massive nuclear black holes (e.g., Richstone et al. 1998), which are however radiatively quiescent in most cases (Fabian \\& Canizares 1988, Pellegrini 2005, Ho 2008), these measurements provide useful constraints for models of halo evolution and nuclear feedback (e.g., Springel, Di Matteo \\& Hernquist 2005; see Pellegrini et al 2007b).  The nearby, gas-poor, early-type galaxy NGC3379 provides an excellent target for these investigations. This is a very well studied prototypical elliptical, with $\\rm L_B = 1.5 \\times 10^{10}$ $\\rm L_{B \\sun}$. A series of very deep  exposures were obtained for this galaxy by our team (PI: Fabbiano), as part of a very large {\\it Chandra} program aimed at unraveling and studying the different X-ray emission components of ellipticals. Together with  pre-existent archival data (David et al 2005), these observations resulted in an integrated exposure of 337 ks with ACIS S-3 (see Brassington et al. 2008 for details). In this paper we present our analysis of the diffuse/unresolved X-ray emission component of this galaxy. This analysis leads us to conclude that in addition to the unresolved emission of faint LMXBs and other stellar sources (Revnivtsev et al. 2008), a hot ISM is likely to be present in the central 800 pc of NGC3379.  We discuss our data analysis in Section 2, and our interpretation of these results in Section 3. Our conclusions are summarized in Section 4. We assume a distance of 10.6 Mpc (Tonry et al. 2001), which gives a scale of 51 pc/\\arcsec. ",
        "conclusions": "With the very deep {\\it Chandra} observation of NGC 3379 we have scored two firsts:  we have the first evidence of a very small amount of hot ISM in a nearby galaxy, down to the level of $\\rm L_x \\sim 4 \\times 10^{37}$ \\ergs\\ (0.5-2.0 keV), and at the same time we have the first evidence that this gas is in an outflow phase, which has been predicted by models but as of yet has not been observed. As a third and equally important result, we find that detailed modelling based on the optical data (stellar mass, total mass and supernova rate) reproduces remarkably well the observed total luminosity, total mass  and spatial distribution of the hot phase of the gas.  After detection and removal of point sources down to a luminosity of $\\rm L_x \\sim 10^{36}$ \\ergs\\ we have been able to study the unresolved emission in NGC 3379. The detailed spectral analysis has led us to the evidence of three separate components: a very soft (plasma at kT$\\sim$0.3 keV) and a hard (Bremsstrahlung at kT=7 keV) components can be attributed to the integrated emission of unresolved sources, closely linked to the stellar population (binary systems and coronally active stars, as discussed by Revnivtsev et al. 2008).  However we also  detect a hot ISM in this galaxy, within the innermost $\\sim$ 800 pc, with a luminosity $\\rm L_x \\sim 4 \\pm 1.5 \\times 10^{37}$ \\ergs\\ (0.5-2.0 keV). This is lowest detection of a hot gas in a nearby early-type galaxy.  The  total gas mass involved is also very small, of a few $10^5$ M$_{\\sun}$.  With the aid of hydrodynamical simulations and a galaxy model tailored specifically on NGC3379, we have reproduced both  the luminosity and the average radial distribution of this hot gas component. For a SNIa's explosion rate consistent with current estimates, the gas is outflowing,  even for the \"maximum\" dark matter model, where the amount of dark mass is the maximum allowed for by optical studies."
    },
    "0807/0807.3279_arXiv.txt": {
        "abstract": " ",
        "introduction": "The evidence for the existence of non-baryonic dark matter (DM) now comes from many sources and ranges over many distance scales \\cite{review}, from the rotation curves of galaxies, the dynamics of clusters, lensing data and the characteristics of large-scale structure, to the features of the cosmic microwave background (CMB) fluctuation spectrum and the success of big bang nucleosynthesis (BBN). All of these pieces of astronomical data point to a similar cosmological density of dark matter, several times that of visible baryonic matter. However, this data only probes the gravitational interaction of dark matter and, while it presents us with one of the most compelling arguments for physics beyond the Standard Model, gaining insight into its non-gravitational interactions remains a primary experimental focus, both through underground detectors, particle colliders, and the observation of photon and neutrino fluxes from overdense regions in the galaxy and beyond.  This lack of any direct information on how dark matter may couple to the Standard Model (SM) means we are forced to rely on various theoretical expectations. In particular, the many successes of standard cosmology motivate a simple thermal mechanism for populating the universe with dark matter with a well-defined freeze-out abundance as the universe expands. This in turn requires a specific annihilation cross-section which is necessarily non-gravitational in origin. The fact that a weakly interacting particle with a weak-scale mass has an annihilation rate in the right ballpark \\cite{lw,review}, combined with our expectations for new physics at the electroweak scale, has rightly led to the prevailing WIMP paradigm for cold dark matter. Nonetheless, persistent problems in understanding the small-scale gravitational clustering properties in cold dark matter simulations and on galactic scales has motivated variants of this picture where the dark matter may be somewhat lighter, with masses down to the keV range. Masses in this range imply a super-weak interaction strength between dark matter and the SM sector, indeed many orders of mangitude below weak-scale cross sections. This follows from the necessity to have early thermal decoupling of the DM sector, prior to the electroweak epoch at $T\\sim 100$ GeV, in order to satisfy the conflicting requirements of not having too much energy density in dark matter, and the strong lower bounds on $m_{DM}$ coming from the analysis of structure formation. If early decoupling can be achieved, then masses in the keV range may withstand these combined constraints and at the same time provide a rather attractive mechanism for ensuring the correct dark matter energy density.   An important feature of keV-scale dark matter is that, unlike the majority of electroweak scale WIMP models, it need not be stable against decays to light SM degrees of freedom, {\\em e.g.} photons and neutrinos. Given the super-weak strength of its interaction such decays may be strongly inhibited, but nevertheless keV dark matter can be emitted or absorbed in astrophysical environments and in terrestrial experiments, and thus is possibly subject to additional constraints. The main question to be studied in this paper  is  whether the application of cosmological and astrophysical constraints leaves any  realistic chance for the direct detection of such keV-scale candidates. The best studied models in this class are the sterile neutrino and gravitino, examples of  fermionic superWIMPs \\cite{sterile,swimp}, where it is well-known that there are no chances for direct detection. In contrast, the present paper studies bosonic superWIMP models, which need not entail such a pessimistic conclusion. In particular, we show that the potential sensitivity of direct detection experiments to bosonic keV dark matter can indeed be competitive with astrophysical bounds as well as with the lifetime and cosmic gamma-background constraints.  The majority of underground direct-detection experiments are specifically tuned to the nuclear recoil energy regime and all $\\gamma$- or $\\beta$-like events are typically considered as background. Currently, the experiments with the best limits are now sensitive to picobarn-scale cross-sections for WIMP-nucleus elastic scattering, when the WIMP mass is on the order of the electroweak scale. To illustrate our main point, we  present a simple estimate that compares the rate of absorption of  10~keV-mass axion-like dark matter ($a$-particles),  coupled to electrons with strength $m_e/f_a =m_e(10^{10}~{\\rm GeV})^{-1}$, with the recoil signal from a $1$ TeV WIMP that has an elastic scattering cross section on the  nucleus  of $\\sigma_{el} \\sim 10^{-36}$cm$^2$: \\be \\fr{\\sigma_{abs} v_{\\rm DM}n_a}{\\sigma_{el} v_{\\rm DM}n_{\\rm WIMP}}\\sim \\fr{m_{\\rm WIMP} }{m_a } \\times  \\fr{c}{v_{\\rm DM}} \\times \\fr{(10^{10}~{\\rm GeV})^{-2}}{10^{36}~{\\rm cm}^2} \\sim 10^{8}\\times 10^{3}\\times 10^{-11} \\sim O(1). \\label{estimate} \\ee For this rather crude estimate we simply took $\\sigma_{abs} v_{\\rm DM} \\sim f_a^{-2}$, and $m_an_a\\sim m_{\\rm WIMP}n_{\\rm WIMP}\\sim \\rho_{DM}$. Although the superWIMP absorption cross section is  orders of magnitude below the weak cross section, it is compensated by the tremendous gain in the local number density of dark matter particles, and by the fact that the inelastic cross section $\\sigma_{abs}$ scales inversely with the dark matter velocity. Although rather imprecise, the estimate (\\ref{estimate}) illustrates that the absorption of 10 keV dark matter can indeed produce signals that are well within modern detection capabilities. Moreover, the scale $f_a \\sim 10^{10}$ GeV is at or above the limits imposed by even the most stringent constraints on star cooling. We believe that this is an important point, and the sensitivity to new and viable dark matter scenarios  can (and should) be explored by direct detection experiments. For comparison, fermionic candidates in the superWIMP class, a well-studied example of which is sterile neutrino dark matter, do not fall into the class covered by the estimate (\\ref{estimate}) for many different reasons; the primary one being that sterile neutrinos are not fully absorbed but rather converted to active neutrinos that carry away most of the rest energy.    Thus far, the issue of direct detection of keV dark matter has only been discussed by the DAMA collaboration \\cite{DAMA1}\\footnote{As this paper was being prepared for publication, another experimental collaboration, CoGeNT \\cite{cogent}, reported their results on the absorption of keV-scale superWIMPs.} in connection with an annual modulation signal initially reported by DAMA/NaI \\cite{DAMA3} and recently confirmed by DAMA/Libra \\cite{DAMA2}. If indeed the annual modulation signal of DAMA were to be attributed to the absorption of keV dark matter,\\footnote{See \\cite{dama_wimp,cogent} for recent analyses of WIMP models in relation to the DAMA results.} which would cause ionization but no significant recoil, it could provide a plausible explanation for why other front-running experiments such as CDMS and Xenon \\cite{cdms,xe} see no signal. This was the main point of Ref.~\\cite{DAMA1}, and in this paper we provide a re-analysis",
        "conclusions": "With vast resources now justifiably being devoted toward the direct detection of dark matter, it appears all but clear that the scientific scope of these searches should be diversified. While a characteristic elastic WIMP-nucleus recoil may remain the main benchmark scenario for these experiments, the detection possibilities for other generic classes of dark matter  should certainly be exploited. In this paper, we have shown that bosonic superWIMPs represent a legitimate and feasible target.  The models analyzed in this paper can naturally produce the required relic abundance of dark matter, once the couplings $\\alpha'$ or $(m_e/f_a)^2$ are in the superweak $\\sim O(10^{-20})$ range. The feeble nature of this coupling to the SM is partially overcome by a factor of $(c/v_{DM}) \\times (m_{\\rm WIMP}/m_{a(V)})$ which enhances superWIMP absorption relative to WIMP scattering and renders direct detection feasible. We have analyzed models that do not require any special fine-tunings and enjoy protection for the small mass scales due to the symmetries of their interactions: gauge symmetry for the model of secluded U(1)$'$, and shift symmetry for the pseudoscalar model.  The result of our analysis has revealed that for the pseudoscalar model, direct detection sensitivity may compete with the red giant, solar and SN constraints. However, we found that the model is most strongly constrained by limits on monochromatic X-ray lines in our galaxy, that rule out much of the interesting regions of the parameter space. In vector models of dark matter, the production of X-rays is strongly inhibited by gauge invariance, and for the model studied here direct searches for ionization are apparently capable of probing the most interesting range of the coupling--mass parameter space, namely that consistent with the observed dark matter energy density. Interestingly, part of this range is not strongly constrained by indirect astrophysical probes, away from the low mass region where stellar energy loss constraints become important and the higher mass range where the $\\gamma$-background is too large. Finally, contrary to some existing claims in the literature,  we have found that the absorption of superWIMPs does not lead to an annual modulation of the ionization signal at a level that would be of experimental interest.   \\subsection*{Acknowledgements}  M.P. would like to thank Georg Raffelt for correspondence and for numerous helpful comments following the submission of v1 of this paper to the arXiv. We would also like to thank Marieke Postma and Javier Redondo for pointing out an error in the stellar bound shown on an earlier version of Fig.~3. The work of A.R. and M.P. is supported in part by NSERC, Canada, and research at the Perimeter Institute is supported in part by the Government of Canada through NSERC and by the Province of Ontario through MEDT. The work of M.V. is supported in part by the DOE grant DE-FG02-94ER40823."
    },
    "0807/0807.1083_arXiv.txt": {
        "abstract": "{We present preliminary results on a study based on contemporaneous photometric and spectroscopic observations of the young K0-1V star SAO 51891. We find that SAO\\,51891, a possible member of the Local Association, shows emission cores in the \\ion{Ca}{ii} H\\&K~and fillings in the H$\\alpha$ and \\ion{Ca}{ii} Infra-Red Triplet (IRT) lines. Moreover, we detect absorption lines of \\ion{He}{i}-D3 and \\ion{Li}{i} and measure a $v\\sin i$ of 19 km s$^{-1}$. A clear rotational modulation of both the light and the photospheric temperature, due to photospheric spots, has been detected. The net H$\\alpha$ chromospheric emission does not show any detectable variation, while the \\ion{Ca}{ii} IRT emission displays a fair modulation. ",
        "introduction": "Stars just arrived on the Zero Age Main Sequence (ZAMS) or on their way to reach it are in an important evolutionary phase because they start to spin up getting free from their circumstellar disks which can begin to condensate giving rise to proto-planetary systems. At the same time, they start to loose angular momentum via magnetic braking. SAO\\,51891 is in this evolutionary stage; it is indeed a young star counterpart of an EUV source with an IR excess attributed to dust around the star (\\citealt{Najita2005}).   The spectra were acquired in August 2006 with FOCES@CAHA at a spectral resolution $R\\approx40000$ in the wavelength range 3720-8850 \\AA~with a signal-to-noise ratio higher than 200. The contemporaneous photometry was performed at OACt in the $BV$ Johnson bands. ",
        "conclusions": "From the study of photospheric and chromospheric inhomogeneities based on spectroscopic and photometric monitoring of SAO 51891, we find a clear light and $T_{\\rm eff}$ rotational modulation due to spots and a modulation of the total IRT \\ion{Ca}{ii} emission due to plages. This chromospheric diagnostics seems to indicate a possible shift between spots and plages. Thus, as a follow/up of our previous works (\\citealt{Biazzo07b,Frasca05,Frasca08}), we aim to develop a spot/plage model for reproducing the observed behaviors at photospheric and chromospheric levels and for deriving spot/plage parameters. SAO\\,51891, and other weak-line T Tauri and PTT, and ZAMS stars already observed by us with FOCES@CAHA and SARG@TNG, are important to explore the correlations between global stellar parameters (e.g., $\\log g$, $T_{\\rm eff}$) and spot characteristics (e.g., filling factor and temperature) in stars with different evolutionary stage and activity level."
    },
    "0807/0807.0390_arXiv.txt": {
        "abstract": "Photospheric granulation may excite transverse kink pulses in anchored vertical magnetic flux tubes. The pulses propagate upwards along the tubes with the kink speed, while oscillating wakes are formed behind the wave front in a stratified atmosphere. The wakes oscillate at the kink cut-off frequency of stratified medium and gradually decay in time. When two or more consecutive kink pulses with different polarizations propagate in the same thin tube, then the wakes corresponding to different pulses may superimpose. The superposition sets up helical motions of magnetic flux tubes in the photosphere/chromosphere as seen by recent Hinode movies. The energy carried by the pulses is enough to heat the solar chrmosphere/corona and accelerate the solar wind. ",
        "introduction": "Recent high resolution movies obtained by Hinode spacecraft clearly show continuous helical motions of spicule axes \\citep{dep}. This phenomenon is known for a long time (Beckers 1972), but no satisfactory explanation has been done yet. \\citet{dep} suggested that these motions are caused by Alfv\\'en waves excited in the photosphere by granular motions or acoustic oscillations. Photospheric magnetic field is concentrated in thin flux tubes and therefore may support the propagation of kink and torsional Alfv\\'en waves (the tubes also support the propagation of sausage waves, but here we consider only transverse waves). The kink wave is a tube wave i.e. the whole tube oscillates with one frequency even if the Alfv\\'en speed changes across the tube, while the frequency of torsional waves is different at different surfaces \\citep{van}. Therefore, the excitation of torsional waves in photospheric magnetic tubes is complicated. The oscillations of spicule axis have also been explained in terms of kink waves \\citep{kuk,zaq}, but the chromospheric magnetic field has rather expanded structure. Therefore, \\citet{dep} suggested through Monte-carlo simulation that spicules are not wave guides for kink waves. The problem is currently under debate and more observations are needed to understand the real process.  On the other hand, granular buffeting on an anchored magnetic tube may easily excite a transverse kink pulse. The pulse propagates through the stratified photosphere with the kink speed, but the oscillating wake is formed behind the wave front. The wake oscillates at cut-off frequency of kink waves and decays as time progresses \\citep{rob1,rob,rae,spruit,has,mus}. However, the magnetic tube undergoes continuous buffeting of granular cells from different sides. Therefore, another pulse polarized in a different plane may quickly follow. The second pulse again propagates with the kink speed and form another wake oscillating with the same cut-off frequency. Therefore, the superposition of the two oscillations may set up the {\\it helical motion} of the tube axis in the photosphere. The helical motion will have the photospheric kink cut-off period $\\sim$ 7-8 min.  However, when the pulse penetrates into the chromosphere, then two possible scenarios can be developed. If chromospheric magnetic field has a tube structure, then the pulse continues to propagate as the kink one. But if the magnetic field is not concentrated in tubes, then it will be transformed into the Alfv\\'enic pulse. Yet the pulse will have the same main properties in both cases: it will propagate at either the kink or Alfv\\'en speed and the wake oscillating at the chromospheric cut-off frequency will be formed behind the pulse. The superposition of the wakes corresponding to different pulses may set up helical motions of magnetic field lines just as in the photospheric case. These helical motions may be responsible for the oscillations of spicule axes as seen by Hinode movies.  Here we study the phenomenon using the Klein-Gordon equation for wave propagation in the stratified atmosphere. ",
        "conclusions": "We suggest that the propagation of consecutive transverse pulses in the stratified atmosphere, which are excited by photospheric granular buffeting, may set up the helical motions of magnetic flux tubes through the superposition of oscillating wakes formed behind the wave fronts. This scenario may explain the continuous motions of spicule axes seen in recent Hinode movies. The pulses carry almost whole energy of initial perturbations, while the energy in wake oscillations is much smaller. Therefore, the energy carried into corona by transverse pulses can be much higher than it is estimated by observed oscillations. More observations and numerical/analytical works need to look further into this problem."
    },
    "0807/0807.3736_arXiv.txt": {
        "abstract": " ",
        "introduction": "In the framework of supersymmetry with conserved $R$ parity, the lightest supersymmetric particle (LSP) can have neither electromagnetic nor strong interactions: otherwise it would have bound to conventional matter and been detected in searches for anomalous heavy nuclei \\cite{isotopes}. Within the minimal supersymmetric extension of the Standard Model (MSSM), weakly-interacting candidates for the LSP are the lightest sneutrino $\\tilde\\nu$, the lightest neutralino $\\chi$, and the gravitino ${\\tilde G}$. The (left-`handed') sneutrino LSP hypothesis is excluded by a combination of neutrino counting at LEP and direct dark matter searches \\cite{FOS94}. Accordingly, general attention is focused on the neutralino (NDM) \\cite{EHNOS} and gravitino dark matter (GDM) \\cite{gdm,FengGDM} possibilities, and in this paper we assume the latter.  The next question is the possible nature of the next-to-lightest supersymmetric particle (NLSP) in a GDM model. Two natural possibilities are the other candidates for the LSP, namely the sneutrino and the neutralino, but the NLSP could equally well be charged and even coloured. Indeed, the lighter stau slepton is a natural candidate for the NLSP~\\cite{FengGDM3} within the constrained MSSM (CMSSM) with gravity-mediated supersymmetry breaking, in which the soft supersymmetry-breaking scalar masses $m_0$, trilinear parameters $A_0$ and gaugino masses $m_{1/2}$ are each assumed to be universal at a GUT input scale \\cite{eoss,cmssmwmap}. The lighter stau is also a natural possibility within minimal supergravity (mSUGRA), in which the gravitino mass is fixed: $m_{\\tilde G} = m_0$, and there is an additional relation between trilinear and bilinear soft supersymmetry-breaking parameters \\cite{vcmssm}. Another possibility for the NLSP within a scenario with non-universal Higgs masses (NUHM) \\cite{nonu,ourNUHM1,ourNUHM2,baerNUHM,eosk4} is the lighter stop~\\cite{deos}.  In this paper we study the other possibility for the NLSP within the NUHM with gravity-mediated supersymmetry breaking, namely the lightest sneutrino, assuming that the gravitino provides the cold dark matter~\\footnote{The possibility of a sneutrino NLSP has also been studied within gaugino-mediated models of supersymmetry breaking~\\cite{BCKS}.}. Like other NLSP candidates in gravity-mediated scenarios, the sneutrino NLSP within the NUHM is expected to be very long-lived. The dominant decays of the other NUHM NLSP candidates produce particles with copious interactions such as charged particles and photons, which are subject to strong cosmological limits \\cite{decay1}. These limits are strong enough to exclude effectively all of the parameter space where the lightest neutralino is the NLSP \\cite{CEFOS}~\\footnote{If the gravitino mass is much lighter than the neutralino mass, these limits might still be satisfied.}. However, the dominant decay mode of the sneutrino is ${\\tilde \\nu} \\to {\\tilde G} \\nu$, and the cosmological limits on neutrino injection are much weaker than those on the injection of photons and charged particles \\cite{KKKM-1,KKKM-2}. Therefore the cosmological limits on NUHM ${\\tilde \\nu}$ NLSP scenarios are relatively weak, leaving considerable scope for a sneutrino NLSP. On the other hand, the sneutrino must appear at the end of the decay chain of every MSSM sparticle produced at a collider, and the particles produced in supersymmetric decay cascades provide distinctive experimental signatures for a sneutrino NLSP \\cite{CK}. In particular, the charged leptons produced in association with a ${\\tilde \\nu}$ NLSP provide tools for diagnosing its flavour.  The structure of this paper is as follows. In Section~2 we discuss sneutrino properties within the NUHM, including its mass and lifetime. In Section~3 we discuss the relic abundance of sneutrinos after freeze-out from a primordial plasma in thermal equilibrium. In Section~4 we analyze the NUHM parameter space and identify regions where the NLSP may be either the ${\\tilde \\nu_{e,\\mu}}$ or the ${\\tilde \\nu_\\tau}$. In Section~5 we discuss the cosmological constraints on ${\\tilde \\nu}$ NLSP scenarios, and show that they are not severe. In Section~6 we discuss some signatures of metastable sneutrinos in different NUHM scenarios, in particular those with different lepton flavours accompanying the ${\\tilde \\nu}$ NLSP. Our conclusions are summarized in Section 7. Calculations of the three-body decay ${\\tilde \\nu} \\to {\\tilde G} \\nu \\gamma$ are described in an Appendix. ",
        "conclusions": "We have analyzed in this paper the possibility of a sneutrino NLSP in NUHM models with a gravitino LSP. This possibility does not exist in the CMSSM, but is quite generic in the NUHM, as we have illustrated with various specific examples. The sneutrino might well be the ${\\tilde \\nu_\\tau}$, but the ${\\tilde \\nu_\\mu}$ and ${\\tilde \\nu_e}$ are also possible candidates for the NLSP. A sneutrino NLSP would be metastable and subject to cosmological constraints on late-decaying particles, but we have shown that these are not difficult to respect. There are various different possible scenarios for the spectrum of sparticles heavier than the sneutrino, which would have distinctive signatures at colliders. In addition to events  with missing energy carried away by the invisible $\\snu$, there would also be events with accompanying charged leptons. The flavours of such leptons would help identify the flavour of the sneutrino NLSP.  As particle physics embarks on the study of the TeV scale with the LHC, much unknown physics will surely be revealed. Supersymmetry is occasionally regarded as a `known unknown' in the sense that, whereas we do not know whether it exists,we think we know what it would look like if it does exist. This paper reminds us that supersymmetry should rather be regarded as an `unknown unknown', in the sense that not only do we not know whether it exists, but we also do not know what it would look like. In the conventional `known unknown' scenario, the LSP is the lightest neutralino and supersymmetry would produce missing-energy events. The latter would also be the signature of a scenario with a gravitino LSP with a neutralino NLSP, at least in gravity-mediated scenarios. However, once the `Pandora's box' of a gravitino LSP has been opened, many other NLSP candidates fly out. In addition to the relatively familiar case of the lighter stau and the more radical case of the lighter stop, there are other possibilities including the sneutrino NLSP scenarios discussed here. All of these scenarios have distinctive features, as illustrated here, so the LHC and subsequent experiments have good prospects for detecting and distinguishing between the various `unknown unknowns'. No `unknown unknown' stone should be left unturned in the search for supersymmetry."
    },
    "0807/0807.1733_arXiv.txt": {
        "abstract": "The nonlinear perturbation theory of gravitational instability is extended to include effects of both biasing and redshift-space distortions, which are inevitable in predicting observable quantities in galaxy surveys. Weakly nonlinear effects in galaxy clustering on large scales recently attracted a great interest, since the precise determination of scales of baryon acoustic oscillations is crucial to investigate the nature of dark energy by galaxy surveys. We find that a local Lagrangian bias and redshift-space distortions are naturally incorporated in our formalism of perturbation theory with a resummation technique via the Lagrangian picture. Our formalism is applicable to any biasing scheme which is local in Lagrangian space, including the halo bias as a special case. Weakly nonlinear effects on halo clustering in redshift space are analytically given. We assume only a fundamental idea of the halo model: haloes form according to the extended Press--Schechter theory, and the spatial distributions are locally biased in Lagrangian space. There is no need for assuming the spherical collapse model to follow the dynamical evolution, which is additionally assumed in standard halo prescriptions. One-loop corrections to the power spectrum and correlation function of haloes in redshift space are explicitly derived and presented. Instead of relying on expensive numerical simulations, our approach provides an analytic way of investigating the weakly nonlinear effects, simultaneously including the nonlinear biasing and nonlinear redshift-space distortions. Nonlinearity introduces a weak scale dependence in the halo bias. The scale dependence is a smooth function in Fourier space, and the bias does not critically change the feature of baryon acoustic oscillations in the power spectrum. The same feature in the correlation function is less affected by nonlinear effects of biasing. ",
        "introduction": "Introduction }  The nonlinear perturbation theory of gravitational instability has recently attracted renewed interest. As precision measurements of the large-scale structure of the universe become possible, theoretically accurate modeling is essential to interpret the observational data. It is recognized that the linear perturbation theory is not satisfactory for this purpose. The importance of nonlinear perturbation theory resides in the era of precision cosmology.  A strong motivation for developing the perturbation theory is to model the baryon acoustic oscillations (BAOs) imprinted in the large-scale structure. Acoustic waves which propagate in the baryon-photon plasma of the early universe freeze out at the recombination epoch, and the sound horizon at that epoch is imprinted in spatial fluctuations of photons and baryons \\cite{BAOtheory,CMBobs}. The BAOs provide a standard ruler \\cite{BAOSR} to geometrically investigate the expansion history of the Universe \\cite{AP79}. The dark energy component is efficiently constrained by galaxy surveys of an intermediate- to high-redshift universe \\cite{GeomTest,MSDE}. Using the BAOs as a standard ruler, large galaxy surveys are expected to provide a robust, promising way of constraining the nature of dark energy \\cite{BAODE,Mat04}. Recent observations of BAOs in modern galaxy surveys \\cite{BAOLSS} prove the method works well.  Even though the BAO scale is quite large $\\sim 100\\,h^{-1}{\\rm Mpc}$, detailed structure of BAOs in galaxy clustering is affected by gravitationally nonlinear evolution after the recombination epoch. The BAO signature in the power spectrum or in the correlation function is deformed by nonlinearity in a lower redshift universe where realistic galaxy surveys are possible \\cite{BAONL,BAONLRD,Ang05,ESW07,SBA08}.  In galaxy redshift surveys, the clustering of dark matter is not directly observable. There are two major sources in the difference between the clustering pattern of dark matter and that of galaxies: galaxy biasing and redshift-space distortions. The spatial pattern of galaxy distribution is not necessarily the same as that of dark matter, and the galaxies are biased tracers of mass \\cite{Bias, BBKS}. The redshift of a galaxy does not purely reflect the Hubble flow, and the Doppler shift by a peculiar velocity is inevitably added. Thus the spatial pattern of clustering of galaxies is distorted in redshift space \\cite{RDist,Kaiser87}. In the linear regime, the power spectrum of biased objects $P_{\\rm obj}(k)$ is usually assumed to be proportional to that of mass $P_{\\rm m}(k)$ in real space: $P_{\\rm obj}(k) = b^2 P_{\\rm m}(k)$, and there is a theoretical reason, which is known as ``the local bias theorem'' \\cite{LocBias,SW98}. Assuming linear dynamics, a linear bias, and a linear velocity field, the power spectrum of biased objects in redshift space is given by \\cite{Kaiser87} \\begin{equation} P^{\\rm (s)}_{\\rm g}(\\bm{k}) = b^2 \\left(1 + \\beta \\mu^2 \\right)^2 P_{\\rm m}(k), \\label{eq:0-1} \\end{equation} where $b$ is the linear bias factor between galaxies and mass, $\\beta \\simeq \\Omega_{\\rm m}^{0.6}/b$, and $P_{\\rm m}(k)$ is the power spectrum of mass in real space.  The linear formula of Eq.~(\\ref{eq:0-1}) applies only when the redshift is large enough, or when the wave number $k$ is small enough. However, the linear formula in redshift space is not sufficiently accurate for typical redshift surveys \\cite{Sco04}. Extending the linear formula of Eq.~(\\ref{eq:0-1}) to include nonlinear effects is far from trivial. A straightforward application of the nonlinear perturbation theory \\cite{SPT,SPTdiag,MSS,BCGS02} with a local biasing scheme \\cite{FG93} has been studied so far \\cite{HMV98,Taruya00,McD06,HaloPT}. In the local biasing scheme, the number density of biased objects is a local function of the smoothed density field of mass in Eulerian space. This scheme seems to work as long as the tree-level perturbation theory is adopted \\cite{SCF99}. However, when loop corrections are considered, there appears a conceptual problem in the local biasing scheme. It turns out that the one-loop corrections, even on large scales, strongly depend on the artificial smoothing scale \\cite{HMV98}. In Ref.~\\cite{McD06}, a way to remove the dependence on the smoothing scale by renormalizing bias parameters is proposed. While this could be a solution to phenomenologically represent the power spectrum of biased objects \\cite{JK08}, it is not guaranteed that the procedure actually reflects the physical nature of biasing, since the biasing is more or less a nonlocal process.  It is true that the biasing is difficult to be included exactly in an analytic framework, since the galaxy formation is a highly nonlinear process. The next best thing is to find a good analytic model. In this respect, a plausible model is provided by a halo approach \\cite{MW96,MJW97,SL99,ST99,Sel00,PS00,MF00,SSHJ01,CS02}, which is based on the extended Press and Schechter (PS) model \\cite{PS74,ExtPS}. In this approach, the galaxy biasing is described through two steps. In the first step, the formation and clustering of dark matter haloes are analytically modeled by using the extended PS model. The spherical collapse model is combined to include the gravitational evolution of halo positions. In the second step, the distribution of dark matter or galaxies within haloes is empirically modeled with a number of assumptions \\cite{SB91,Sel00,PS00,MF00,SSHJ01,CS02}, which include complicated physics such as gas cooling, star formation, feedback effects from supernovae, and so forth \\cite{SemiAna}.  While the second step involves many parameters which should be fitted by observations or simulations, the first step is theoretically less uncertain because of its purely gravitational nature. On large scales, the second step is not important and the power spectrum of nonlinear objects is approximately the same as that of haloes. Therefore, determining the power spectrum of haloes is an important step toward understanding the nonlinear structure formation and biasing. In the usual halo approach, the halo clustering is modeled by linear dynamics and linear bias factors \\cite{MW96,Sel00,PS00,MF00,SSHJ01}. Recently, an attempt to incorporate the halo bias with nonlinear perturbation theory in a framework of local biasing scheme has been made \\cite{HaloPT}. There still remains conceptual arbitrariness in the latter formulation, mainly because the halo bias is intrinsically nonlocal in Eulerian space and does not fit well into the local Eulerian biasing scheme. From the construction, the halo bias is local in Lagrangian space \\cite{MW96}, and is therefore nonlocal in Eulerian space, because positions of mass and biased objects are displaced by dynamical evolution.  In this paper, a natural approach to incorporate the halo bias and redshift-space distortions with nonlinear perturbation theory is newly developed. To account for the locality of halo bias in Lagrangian space, our formulation is based on the Lagrangian perturbation theory (LPT) \\cite{LPT}, instead of the standard Eulerian perturbation theory (EPT). In the original halo approach, the spherical collapse model is adopted to take into account the dynamical evolution in Eulerian space. However, the dynamical evolution of haloes is more naturally described in the Lagrangian picture \\cite{LHalo}. It is not straightforward to obtain an analytically useful formalism which combines the LPT and the halo approach. A main obstacle is that the observable quantities reside in Eulerian space, while the calculations in LPT directly give Lagrangian quantities. In paper I \\cite{Mat08}, a new approach is developed to overcome this point, partially expanding the Lagrangian variables in Eulerian space. The resulting expression contains an infinite series of perturbations in terms of the EPT, and this approach offers a simplified technique of resumming cosmological perturbations, such as done in renormalized perturbation theory and its variants \\cite{RPTvar, NewPT}. While our technique is not suitable enough to describe the fully nonlinear regime which other renormalization techniques aim at, we have shown in paper I that our technique is accurate enough in the quasilinear regime, and most importantly, predictions in redshift space are straightforward. To date, our technique offers the only way of obtaining resummed power spectra in redshift space. In this paper, we show that the local Lagrangian bias, including the halo bias, is also straightforward to be incorporated in the approach of paper I, on top of the redshift-space distortions. The local Lagrangian biasing scheme is not equivalent to the local Eulerian biasing scheme, and the conceptual problem about the strong dependence on smoothing scales which appears in the local Eulerian biasing scheme described above is not present in our Lagrangian approach.  The rest of the paper is organized as follows. In Sec.~\\ref{sec:LLBias}, our basic formalism is described. This formalism is applicable not only to the halo bias but also to a local Lagrangian bias in general. Linear and one-loop results are presented. In this general theory, we have parameters which are related to the Lagrangian biasing scheme. In Sec.~\\ref{sec:HaloLPT}, those Lagrangian bias parameters are calculated from the fundamental concept of the halo model. Effects of halo bias and redshift-space distortions on BAO scales are demonstrated both in the power spectrum and in the correlation function. Our conclusions are summarized in Sec.~\\ref{sec:concl}. In the Appendix, details of one-loop calculations in our framework are outlined. ",
        "conclusions": "}  In this paper, we show that the nonlinear perturbation theory via the Lagrangian picture is naturally incorporated with the local Lagrangian bias on top of the redshift-space distortions. Explicit results of one-loop power spectrum with a halo bias both in real space and in redshift space are given. Our approach does not suffer from a conceptual problem which appears in the EPT with local Eulerian bias. The halo bias, as a local Lagrangian bias, is properly treated in the Lagrangian picture of perturbations. Simultaneous inclusion of redshift-space distortions is also natural in the Lagrangian picture. Our approach involves a partial resummation of Eulerian perturbations, and provides a better description in the quasilinear regime than the standard EPT.  In our general framework of local Lagrangian bias, the bias is characterized by local Lagrangian bias parameters $\\langle F^{(n)} \\rangle$. We do not need a Taylor expansion of the bias function as frequently adopted in local Eulerian biasing scheme. Only the first two of the local Lagrangian bias parameters are needed in our one-loop calculations. The results in Sec.~\\ref{sec:LLBias} are applicable for all biasing schemes as long as the bias is local in Lagrangian space. The local Lagrangian bias parameters are considered as free parameters if the bias model is not specified. Because of the nonlocal nature of gravitational evolution, the local Lagrangian biasing scheme is not equivalent to the local Eulerian biasing scheme.  There is a successful model of Lagrangian bias, i.e., the halo bias. The general halo approach consists of several pieces of assumptions. We adopt only the most fundamental concept of the halo approach, the halo bias in Lagrangian space. The clustering of galaxies or dark matter is dominated by halo clustering on large scales, where the perturbation theory is safely applied. As a result, the local Lagrangian bias parameters are calculated for halo bias without ambiguity. The resulting power spectrum does not have any free parameters once the mass or mass range of haloes is specified.  In usual linear analysis, the halo bias is independent on scales. However, nonlinear effects introduce the scale dependence into the halo bias. Such scale dependence could affect the determination of the BAO scales unless the effect is properly quantified. Such effects are quantitatively calculated without numerical simulations. We find that the scale dependence of bias is a smoothly varying function. Therefore, the BAO scales are shifted by change of the broadband shape in the power spectrum. The correlation function is also affected by scale dependence of nonlinear halo bias. Compared to the power spectrum, the shape of the BAO peak in the correlation function seems to be less affected. Those observations are consistent with the recent results of numerical simulations \\cite{Ang05,SBA08}.  Our formalism is compatible with any biasing scheme which is local in Lagrangian space. For example, the peak bias is approximately considered as a local in Lagrangian space in a limit of the peak-background split \\cite{BBKS,MJW97}, when the constraints imposed by the spatial derivatives to define peaks can be neglected. However, the last approximation may not be appropriate for accurately predicting the BAO signature \\cite{Des08}. The exact biasing mechanism in the real world is definitely not local both in Eulerian space and in Lagrangian space. The success of the halo approach indicates the local Lagrangian bias is a good approximation at least on large scales. However, extending the model of local Lagrangian bias to a nonlocal one is an option to make the theory more accurate and general. It is straightforward to extend our formalism in Sec.~\\ref{sec:LLBias} to include a nonlocal bias.  Although our formalism contains a partial resummation of higher-order Eulerian perturbations, Lagrangian perturbations are truncated at the one-loop level. Recent developments of the renormalized perturbation theory and its variants \\cite{RPTvar} show that it is possible to reorganize and resum higher-order perturbations using the concept of propagators. It would be interesting if one could use the concept of propagators in Lagrangian space \\cite{BV08} to further renormalize the present formalism and to describe the deeply nonlinear regime, $k > k_{\\rm NL}/2$, which is not accessible by the present formalism.  We consider only the halo bias in this paper. The resulting power spectrum corresponds to that of halo centers, or the two-halo term of the halo model on large scales. In the halo approach, the nonlinear power spectrum is given by a superposition of the one-halo term and the two-halo term, which are given by convolutions with a model of density profile of galaxies or dark matter. A model of density profile is dominantly relevant to clustering on small scales. In a context of the halo approach, our formalism improves the description of the two-halo term, in which only the linear dynamics is usually included because of simplicity. It is possible to include a model of density profile of halo model in redshift space \\cite{HaloRed}.  As pointed out in paper I, our approach does not have much power on small scales in the power spectrum, $k > k_{\\rm NL}/2$. The main source of the powerless is the exponential damping factor, which originates from random motion of the displacement field. A fully nonlinear description of this factor may dramatically improve the applicability of the present formalism for dark matter clustering. However, in the presence of bias, the nonlinear regime is dominated by the scale dependence of bias, which may be more appropriately described by halo approach with the one-halo term, including a model of density profile or halo occupation dynamics."
    },
    "0807/0807.1505_arXiv.txt": {
        "abstract": "We analysed two-dimensional maps of 48 early-type galaxies obtained with the SAURON and OASIS integral-field spectrographs using kinemetry, a generalisation of surface photometry to the higher order moments of the line-of-sight velocity distribution (LOSVD). The maps analysed include: reconstructed image, mean velocity, velocity dispersion, $h_3$ and $h_4$ Gauss-Hermite moments. Kinemetry is a good method to recognise structures otherwise missed by using surface photometry, such as embedded disks and kinematic sub-components. In the SAURON sample, we find that 31\\% of early-type galaxies are single component systems. 91\\% of the multi-components systems have two kinematic subcomponents, the rest having three. In addition, 29\\% of galaxies have kinematically decoupled components, nuclear components with significant kinematic twists. We differentiate between slow and fast rotators using velocity maps only and find that fast rotating galaxies contain disks with a large range in mass fractions to the main body. Specifically, we find that the velocity maps of fast rotators closely resemble those of inclined disks, except in the transition regions between kinematic subcomponents. This deviation is measured with the kinemetric $k_5/k_1$ ratio, which is large and noisy in slow rotators and about 2\\% in fast rotators. In terms of E/S0 classification, this means that 74\\% of Es and 92\\% of S0s have components with disk-like kinematics. We suggest that differences in $k_5/k_1$ values for the fast and slow rotators arise from their different intrinsic structure which is reflected on the velocity maps. For the majority of fast rotators, the kinematic axial ratios are equal to or less than their photometric axial ratios, contrary to what is predicted with isotropic Jeans models viewed at different inclinations. The position angles of fast rotators are constant, while they vary abruptly in slow rotators.  Velocity dispersion maps of face-on galaxies have shapes similar to the distribution of light. Velocity dispersion maps of the edge-on fast rotators and all slow rotators show differences which can only be partially explained with isotropic models and, in the case of fast rotators, often require additional cold components. We constructed local (bin-by-bin) $h_3 - V/\\sigma$ and $h_4 - V/\\sigma$ diagrams from SAURON observations. We confirm the classical anti-correlation of $h_3$ and $V/\\sigma$, but we also find that $h_3$ is almost zero in some objects or even weakly correlated with $V/\\sigma$. The distribution of $h_4$ for fast and slow rotators is mildly positive on average. In general, fast rotators contain flattened components characterised by a disk-like rotation. The difference between slow and fast rotators is traceable throughout all moments of the LOSVD, with evidence for different intrinsic shapes and orbital contents and, hence, likely different evolutionary paths. ",
        "introduction": "\\label{s:intro}   The classification of galaxies both acknowledges the complexity of these celestial objects and attempts to understand their formation and evolution. The Hubble Classification of galaxies \\citep{1936RNeb..........H, 1961hag..book.....S, 1975gaun.book.....S,1994cag..book.....S} recognises the dichotomy between, broadly speaking, disk and ellliptical galaxies, for historical reasons now often called late- and early-types. The classification works well on the late-type galaxies in particular, dividing the class into a number of subgroups which correlate with properties such as bulge-to-disk ratio, morphology of spiral arms, gas and dust content, to name a few, but it fails to bring a physical insight to our understanding of early-types \\citep{1987IAUS..127..367T}, where the classification is based on apparent shape and thus dependant on viewing angles.  In an effort to eliminate the unsatisfactory situation \\citet{1996ApJ...464L.119K} proposed a revision of the Hubble classification. It was based on two discoveries in the 1970s and 1980s, both enabled by an improvement in the technical capabilities of astronomical instruments. A series of papers \\citep{1975ApJ...200..439B, 1977ApJ...218L..43I} showed that bright elliptical galaxies do not rotate as fast as they should, if they were oblate isotropic systems supported by rotation \\citep{1978MNRAS.183..501B}, whereas less bright and generally smaller systems, including also bulges of spirals, generally agree with such predictions \\citep{1983ApJ...266...41D, 1982modg.proc..113K, 1982ApJ...256..460K}. A complementary discovery \\citep{1988A&A...193L...7B, 1988A&AS...74..385B,1989A&A...217...35B} that the fast rotating galaxies are more likely to have disky isophotes, while the slow rotating galaxies have boxy isophotes, linked again the kinematics and shape of galaxies. \\citet{1996ApJ...464L.119K} changed the uniformity of early-type galaxies to a dichotomy (disky vs boxy, fast vs slow, brighter vs less bright) and linked the whole Hubble sequence from right to left, from Sc--Sb--Sa to S0--E types, where 'rotation decreases in dynamical importance compared to random motions'.  This important step forward introduced a readily measurable parameter related to some physical properties. However, the higher order variations in the isophotal shape (diskiness/boxiness) are not measurable at all inclinations regardless of the prominence of the disks \\citep{1990ApJ...362...52R}, and finally, they are used to infer the dynamical state of the galaxy. This might be a decent approximation, especially if one assumes that all fast rotating galaxies comprise of spheroidal slow rotating components and disks seen at different inclination \\citep{1999ApJ...513L..25R}, but the edge-on observations of spheroidal components in spiral galaxies showed that bulges are rotating fast as well \\citep{1982ApJ...256..460K}. To complicate things further the updated classification continues to distinguish S0s from Es, keeping a viewing angle dependent definition of S0s \\citep{1990ApJ...348...57V}.  The choice of using the fourth order Fourier term in the isophotometric analysis for classification is natural, because {\\it (i)} it is much easier to take images of galaxies than to measure their kinematics, and {\\it (ii)} until recently it was not realistic to spectroscopically map their two-dimensional structure.  The advent of panoramic integral-field units, such as SAURON \\citep{2001MNRAS.326...23B}, is changing the technical possibilities and the field itself; it is now possible to systematically map kinematics of nearby galaxies up to their effective radii. We have observed 72 nearby E, S0 and Sa galaxies as part of the SAURON survey \\citep[][hereafter Paper II]{2002MNRAS.329..513D}. Focusing here on a subsample of 48 early-type galaxies (E, S0), these observations clearly show the previously hinted rich variety of kinematic substructures such as: kinematically decoupled cores, kinematic twists, counter-rotating structures and central disks \\citep[][hereafter Paper III]{2004MNRAS.352..721E}.  Analysing the global properties of the SAURON velocity and velocity dispersion maps, \\citet[][hereafter Paper IX]{2007MNRAS.379..401E} were able to separate the early-type galaxies into two physically distinct classes of {\\it slow} and {\\it fast} rotators, according to their specific (projected) angular momentum measured within one effective radius, $\\lambda_R$. This finding augments the view that led to the revision of the classification, but the SAURON observations provide the crucial quantitive data. Moreover, the results of Paper IX suggest a way to dramatically improve on the Hubble classification and substitute S0s and (misclassified) disky ellipticals with one class of fast rotators.  \\citet[][hereafter Paper X]{2007MNRAS.379..418C} addressed again the issue of orbital anisotropy of early-type galaxies. They constructed the ($V/\\sigma$, $\\epsilon$) diagram \\citep{1978MNRAS.183..501B} using an updated formalism \\citep{2005MNRAS.363..937B}, and compared it with the results from general axisymmetric dynamical models for a subsample of these galaxies \\citep[][hereafter Paper IV]{2006MNRAS.366.1126C}. They found that slow and fast rotators are clearly separated on the ($V/\\sigma$, $\\epsilon$) diagram (unlike Es and S0s), such that slow rotators are round, moderately isotropic and are likely to be somewhat triaxial, while fast rotators appear flattened, span a larger range of anisotropies, but are mostly oblate axisymmetric objects. This finding is in a partial agreement with previous studies which either found round early-type galaxies radially anisotropic \\citep{1991MNRAS.253..710V}, moderately radially anisotropic \\citep{2001AJ....121.1936G} or only weakly anisotropic with a range of anisotropies for flattened systems \\citep{2003ApJ...583...92G}. The results of Paper X, however, clearly show that intrinsically flatter galaxies tend to be more anisotropic in the meridional plane. The models also indicate that the fast rotators are often two component systems, having also a flat and rotating, kinematically distinct, disk-like component.  Dynamical models are often time consuming and difficult to construct. Ultimately, one would like to be able to classify galaxies by their observable properties only. Is it possible to learn about the intrinsic shapes of the early-type galaxies from observations only? Surface photometry, being but the zeroth moment of the ultimate observable quantity for distant galaxies, the line-of-sight velocity distribution (LOSVD), cannot give the final answer. It is necessary to look at the other moments of the LOSVD: mean velocity, velocity dispersion and higher-order moments, commonly parameterised by Gauss--Hermite coefficients, $h_3$ and $h_4$ \\citep{1993ApJ...407..525V, 1993MNRAS.265..213G}, which measure asymmetric and symmetric deviation of the LOSVD from a Gaussian, respectively.  Indeed, in the last dozen years several studies investigated higher moments of the LOSVD of early-type galaxies observing them along one or multiple slits \\citep[e.g.][]{1994MNRAS.269..785B,1994MNRAS.268..521V, 1994MNRAS.270..271V,2000A&AS..144...53K, 2000A&AS..145...71K,2001MNRAS.326..473H,2002A&A...395..753W, 2006MNRAS.371..633H,2008ApJS..175..462C}. These studies deepened the dichotomy among early-type galaxies showing that fast rotating galaxies with disky isophotes also exhibit an anti-correlation between $h_3$ and $V/\\sigma$. This is consistent with these galaxies being made of two components: a bulge and a disk. The symmetric deviations, on the other hand are usually smaller than asymmetric ones, and somewhat positive in general. In addition, the observed higher order moments of the LOSVD can be used to constrain the possible merger scenarios of early-type galaxies and their formation in general \\citep[e.g.][]{1991A&A...249L...9B, 2000MNRAS.316..315B,2006MNRAS.372L..78G,2006MNRAS.372..839N,2007MNRAS.376..997J}. However, observations along one or two slits are often not able to describe the kinematical richness of early-type galaxies.  In this work we use \\kinemetry~\\citep{2006MNRAS.366..787K}, a generalisation of surface photometry to all moments of the LOSVD, to study SAURON maps of 48 early-type galaxies. The purpose of this work is to investigate observational clues from resolved two-dimensional kinematics for the origin of the differences between the slow and fast rotators.  In Section~\\ref{s:dat} we briefly remind the reader of the SAURON observations and data reduction. Section~\\ref{s:met} describes the methods and definitions used in this work. The main results are presented in Section~\\ref{s:res}. In Section~\\ref{s:disc} we offer an interpretation of the results and we summarise the conclusions in Section~\\ref{s:conc}. In Appendix~\\ref{s:psf} we discuss the influence of seeing on the two-dimensional kinematics and in Appendix~\\ref{s:profiles} we present the radial profiles of the kinemetric coefficients used in this study. ",
        "conclusions": "\\label{s:conc}   Using kinemetry we analysed two-dimensional maps of 48 early-type galaxies observed with SAURON and velocity maps of a sub-sample observed with OASIS. The analysed maps are: reconstructed image, mean velocity, velocity dispersion, $h_3$ and $h_4$ Gauss-Hermite moments. The reconstructed images and the maps of the mean velocity were analysed along the best fit ellipses. Maps of the higher moments of the LOSVD were analysed either along the isophotes (velocity dispersion, $h_4$) or best-fitting ellipses from velocity maps ($h_3$). We presented the profiles of kinemetric coefficients for velocity maps, being the dominant kinematic moment of the LOSVD and having the highest signal-to-nose ratio.  Kinemetry and its kinemetric coefficients allow us {\\it (i)} to differentiate between slow and fast rotators from velocity maps only and {\\it (ii)} to indicate that fast rotating galaxies contain disks with a larger range in the mass fractions to the main body. The results of this work can be summarised as follows:  \\begin{itemize}  \\item Following the kinemetric analysis of the velocity maps it is possible to distinguish between galaxies with {\\it Single} and {\\it Multiple Components}. Components can be described as having a {\\it Disk-like rotation}, {\\it Low-level velocity}, {\\it Kinematic misalignment}, {\\it Kinematic twist} or being {\\it Kinematically Decoupled Component}. The sorting of galaxies in these groups is based on kinemetric coefficients: \\PAk, \\qk, $k_1$ and $k_5/k_1$.  \\item The majority of early-type galaxies are MC systems (69\\%)  \\item The total fraction of galaxies with a DR component (including all that have $k_5/k_1<0.02$) is 81\\%. In terms of S0/E classification, 92\\% of S0s and 74\\% of Es have components with disk-like kinematics.  \\item KDCs are found in 29\\% of galaxies. These KDC are of different sizes, some are not resolved at SAURON resolution and appear as CLVs, but are clearly detectable in the OASIS observations at higher spatial resolution.  \\item Early-type galaxies with constant \\PAk ($\\Delta PA <10\\degr$ outside any central component) are fast rotators. All fast-rotators have a DR component.  \\item Most fast rotators have $\\langle$\\qk$\\rangle \\lesssim \\langle$\\qp$\\rangle$. Their kinematics often show a structure flatter than the distribution of light. This means that images alone are not sufficient to recognise all fast rotators (e.g. NGC524 or N3379 would be missed)  \\item In face-on galaxies, isophotes are coincident with iso-$\\sigma$ contours. In edge-on galaxies, however, there are differences which can be detected with kinemetry. Specifically, the edge-on fast rotators contain dynamically cold components which changes the shape of iso-$\\sigma$ contours.  \\item Slow rotating galaxies have low $V/\\sigma$ and a range in $h_3$ amplitudes. Fast rotating galaxies have large spread in $h_3$ which anti-correlates with $V/\\sigma$. There is, however, a significant number of fast rotating galaxies which have $h_3$ radial profiles that change from strongly negative to near to zero $h_3$ values. Similar behaviour was found in bar galaxies or in remnants of colissionless mergers.  \\item Allowing for large uncertainties in the determination of $h_4$, our data show a trend where slow and fast rotators have similar $h_4$ values, with a weak tendency for an increased spread of $h_4$ in fast rotators, also dependant on $V/\\sigma$. These trends can be explained through a combination of dry and wet mergers of both equal and unequal in mass progenitors.  \\item Dissimilarities between slow and fast rotators originate in their different internal structures. Slow rotators are mildly triaxial objects supporting a variety of orbital families. Fast rotators are axisymmetric spheroids with embedded flattened components of different mass fraction ranging from completely disk dominated systems to small central disks visible only in kinematics.  \\end{itemize}       \\vspace{+1cm} \\noindent{\\bf Acknowledgements}\\\\  DK acknowledges the support from Queen's College Oxford and hospitality of Centre for Astrophysics Research at the University of Hartfordshire.  The SAURON project is made possible through grants 614.13.003, 781.74.203, 614.000.301 and 614.031.015 from NWO and financial contributions from the Institut National des Sciences de l'Univers, the Universit\\'e Lyon~I, the Universities of Durham, Leiden, and Oxford, the Programme National Galaxies, the British Council, PPARC grant `Observational Astrophysics at Oxford 2002--2006' and support from Christ Church Oxford, and the Netherlands Research School for Astronomy NOVA.  RLD is grateful for the award of a PPARC Senior Fellowship (PPA/Y/S/1999/00854) and postdoctoral support through PPARC grant PPA/G/S/2000/00729. The PPARC Visitors grant (PPA/V/S/2002/00553) to Oxford also supported this work.  GvdV acknowledges support provided by NASA through grant NNG04GL47G and through Hubble Fellowship grant HST-HF-01202.01-A awarded by the Space Telescope Science Institute, which is operated by the Association of Universities for Research in Astronomy, Inc., for NASA, under contract NAS 5-26555. JFB acknowledges support from the Euro3D Research Training Network, funded by the EC under contract HPRN-CT-2002-00305. This paper is based on observations obtained at the William Herschel Telescope, operated by the Isaac Newton Group in the Spanish Observatorio del Roque de los Muchachos of the Instituto de Astrof\\'{\\i}sica de Canarias. Based on observations obtained at the Canada-France-Hawaii Telescope (CFHT) which is operated by the National Research Council of Canada, the Institut National des Sciences de l'Univers of the Centre National de la Recherche Scientifique of France, and the University of Hawaii. This project made use of the HyperLeda and NED databases. Part of this work is based on data obtained from the ESO/ST-ECF Science Archive Facility."
    },
    "0807/0807.4266_arXiv.txt": {
        "abstract": "Recently two groups used first-principles computer simulations to model Jupiter's interior.  While both studies relied on the same simulation technique, density functional molecular dynamics, the groups derived very different conclusions. In particular estimates for the size of Jupiter's core and the metallicity of its hydrogen-helium mantle differed substantially. In this paper, we discuss the differences of the approaches and give an explanation for the differing conclusions. ",
        "introduction": "In this section we discuss the differences in the first-principles simulations performed by the two groups. Although there are differences in the level of accuracy, they are unlikely to be the main reason for the differences in the Jupiter models. Both groups used density functional molecular dynamics (DFT-MD) simulations and should derive identical equations of state for hydrogen, helium, and their mixtures. However, computational details are important for accuracy of the derived EOS.  The NHKFRB group used exclusively VASP code while MHVTB used CPMD as well as the VASP code. MHVTB verified that both codes yield identical results when run accurately.  The MHVTB group performed simulations of hy\\-dro\\-gen-helium mixtures using a mixing ratio that closely resembles Jupiter's composition within the limitations of a finite simulation cell. The NHKFRB instead relied on the ideal mixing approximation using EOSs of pure hydrogen and helium. This approximation neglects all interactions between hydrogen and helium. In Refs.~\\citep{Vo07,Vo07b}, it was shown that this interaction leads to significant corrections in the derived pressures and energy but, more importantly, the presence of helium increases that stability of the hydrogen molecules for given $T$ and $P$. These interaction effects also have an impact on the derived adiabat that will be discussed in the next section. However, linear mixing approximation does not affect the predicted core mass very much. In fact, the MHVTB group used their own set of simulations for pure hydrogen and helium~\\citep{Mi08} in a separate Jupiter model and derived a core mass that differed by only 2 Earth masses.  MHVTB use this independent calculation to estimate the uncertainty of the predictions.  The NHKFRB group performed DFT-MD simulations with 500-2000 time steps. MHVTB performed simulations for 2 picoseconds with 5000 time steps. Longer simulations lead to more accurate averages for thermodynamic variables such as pressure and internal energy.  Concerning the size of the simulation, the NHKFRB paper states that simulations with between 32 and 162 atoms for hydrogen, helium, and water were performed. A conservative estimate of the precision of the EOS 5\\% is given. Conference presentations by this group showed results with 64 hydrogen atoms and 32 helium atoms (64 electrons). The MHVTB group used simulations with 128 electrons throughout, which are more accurate.  The NHKFRB used $(\\frac{1}{4},\\frac{1}{4},\\frac{1}{4})$ Baldereschi point to sample the Brillouin zone for all simulations. The MHVTB group analyzed the $k$-point dependence using simulation with up to 4x4x4 k-points and then performed simulation with 2x2x2 in the metallic regime and using the $\\Gamma$ point at lower densities. It was shown that $\\Gamma$ point only simulation overestimate the pressure near Jupiter's core-mantle boundary by 1.6\\%. ",
        "conclusions": ""
    },
    "0807/0807.1443_arXiv.txt": {
        "abstract": "A fundamental presupposition of modern cosmology is the Copernican Principle; that we are not in a central, or otherwise special region of the Universe. Studies of Type Ia supernovae, together with the Copernican Principle, have led to the inference that the Universe is accelerating in its expansion.  The usual explanation for this is that there must exist a `Dark Energy', to drive the acceleration. Alternatively, it could be the case that the Copernican Principle is invalid, and that the data has been interpreted within an inappropriate theoretical frame-work.  If we were to live in a special place in the Universe, near the centre of a void where the local matter density is low, then the supernovae observations could be accounted for without the addition of dark energy. We show that the local redshift dependence of the luminosity distance  can be used as a clear discriminant between these two paradigms. Future surveys of Type Ia supernovae that focus on a redshift range of $\\sim 0.1-0.4$ will be ideally suited to test this hypothesis, and hence to observationally determine the validity of the Copernican Principle on new scales, as well as probing the degree to which dark energy must be considered a necessary ingredient in the Universe. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0807/0807.2056_arXiv.txt": {
        "abstract": "We describe Spitzer IRAC and MIPS observations of the nearby 25 Myr-old open cluster NGC 2232. Combining these data with ROSAT All-Sky Survey observations, proper motions, and optical photometry/spectroscopy, we construct a list of highly probable cluster members. We identify 1 A-type star, HD 45435, with definite excess emission at 4.5--24 $\\mu$m indicative of debris from terrestrial planet formation.   We also identify 2--4 late-type stars with possible 8 $\\mu$m excesses, and 8 early-type stars with definite 24 $\\mu$m excesses. Constraints on the dust luminosity and temperature suggest that the detected excesses are produced by debris disks. From our sample of B and A stars, stellar rotation appears correlated with 24 $\\mu m$ excess, a result expected if massive primordial disks evolve into massive debris disks. To explore the evolution of the frequency and magnitude of debris around A-type stars, we combine our results with data for other young clusters.  The frequency of debris disks around A-type stars appears to increase from $\\sim$ 25\\% at 5 Myr to $\\sim$ 50--60\\% at 20--25 Myr. Older A-type stars have smaller debris disk frequencies: $\\sim$ 20\\% at 50--100 Myr. For these ages, the typical level of debris emission rises from 5--20 Myr and then declines. Because 24 $\\mu$m dust emission probes icy planet formation around A-type stars, our results suggest that the frequency of icy planet formation is $\\eta_{i}$ $\\gtrsim$ 0.5--0.6. Thus, most A-type stars ($\\approx$ 1.5--3 M$_{\\odot}$) produce icy planets. ",
        "introduction": "Recent studies of young stellar clusters from the \\textit{Spitzer Space Telescope} \\citep{We04} provide powerful insights into the emergence of mature planetary systems from primordial disks surrounding young stars.  Spitzer observations indicate that the initial stages of primordial disk evolution proceed rapidly.  Within $\\sim$ 1-3 Myr, many optically-thick primordial disks undergo dust settling, form inner regions cleared of dust, and undergo grain growth to sizes larger than the sub-micron sizes characteristic of the interstellar medium \\citep[e.g.][]{Cal05, La06, Sa06, Bou08}. As these processes occur, the level of inner disk emission drops and the frequency of primordial disks declines (e.g. \\citealt{He07}, see also \\citealt{Hi05}).  Primordial disks around early type stars disappear faster than around later-type stars \\citep{Ca06, He07, Ck08}.  By $\\sim$ 10--15 Myr, primordial disks are exceedingly rare. The disk population is then dominated by optically thin, gas-poor \\textit{debris disks} \\citep{He06, Cu07a, Ck08}.  Because debris dust is removed on timescales much shorter than the age of the star, the dust must be replenished to sustain disk emission. Collisions between $\\gtrsim$ km-sized planetesimals excited by large-scale planet formation \\citep[e.g.][]{Bp93, Kb04} can provide this replenishment. Though debris disks around early-type stars can appear as soon as $\\sim$ 3 Myr \\citep{Ck08}, the mid-IR debris emission does not peak until $\\approx$ 10--20 Myr \\citep{Cu08a}.  This rise in debris emission is consistent with planetesimal collisions during the growth of icy protoplanets in the outer disk \\citep{Kb08}. After protoplanets reach their maximum size at $\\approx$ 20--30 Myr, the level of debris emission decays as $\\sim$ t$^{-1}$ \\citep{Ri05, Su06}. This decay is consistent with steady-state collision models \\citep[e.g.][]{Wy07}.  The epoch of maximum debris emission ($\\sim$ 10--30 Myr) is also important for terrestrial planet formation.  Planet formation models and radiometric dating suggest that terrestrial planets reach their final mass by $\\sim$ 10--30 Myr \\citep{Ws93, Yi02, Kb06}.  Terrestrial planet formation produces debris emission observable at $\\approx$ 5--10 $\\mu m$. Mid-IR Spitzer observations reveal terrestrial zone dust emission around many 13 Myr old stars in the massive Double Cluster, h and $\\chi$ Persei, and several stars in other $\\lesssim$ 40 Myr old clusters/associations including Sco-Cen, NGC 2547, and the $\\beta$ Pic Moving Group \\citep{Zs04, Ch05, Ch06, Cu07b, Rh07a, Go07, Cu08a, Li08}\\footnote{At least two other, much older stars harbor warm debris disks that may be indicative of stochastic collisions in the terrestrial zone: BD +20$^{o}$ 307 and HD 23514 \\citep{Sz05, Rh07b}}. Observations of many warm, IRAC-excess stars in h and $\\chi$ Persei point to a spectral-type/stellar-mass dependent evolution of terrestrial zone dust, suggesting that the terrestrial planet formation process runs to completion faster for high-mass stars than for intermediate-mass stars \\citep{Cu07a, Cu08b}.  Reconstructing the time history of terrestrial and icy planet formation at the 10--40 Myr epoch requires Spitzer IRAC and MIPS observations of many clusters besides h and $\\chi$ Persei. While observations of stars in 10--20 Myr old Sco-Cen and the $\\beta$ Pic moving group provide important studies of cold dust probed by MIPS \\citep{Ch05, Re08}, strong constraints on warm dust probed by $\\le$ 10--15 $\\mu m$ broadband photometry of cluster members are limited to IRAS, MSX, and ground-based campaigns with much poorer sensitivity.  The FEPS (\\textit{Formation and Evolution of Planetary Systems}; \\citealt{Me06}) Legacy Program observed $\\sim$ 35 10--20 Myr-old Sco-Cen sources in IRAC with a range of spectral types ($\\sim$ F8/G0 to K3).  However, this sample size is too small to provide statistically robust constraints given the low overall frequency of warm dust emission \\citep{Ma04, Cu07a, Go07}.  Though h and $\\chi$ Persei is massive enough to investigate the frequency of warm dust from terrestrial planet formation in one environment, observations of other clusters are needed to constrain the time evolution of warm dust in many environments.  The Spitzer survey of NGC 2547 \\citep{Go07} potentially fill this void in sampling, yielding a large population of cluster stars ($\\sim$ 400-500) and revealing warm dust around at least $\\sim$ 6 FGKM stars.  However, recent work has revised the age of NGC 2547 upward from 25 Myr to $\\sim$ 35--40 Myr \\citep{Nj06, Mn08}.  Spitzer studies of other $\\sim$ 15--35 Myr-old clusters would then clearly aid in constraining the evolution of the observable signatures of terrestrial planet formation.  Understanding the evolution of cold dust emission from debris disks also requires additional observations of 15--35 Myr-old clusters.  While 24 $\\mu m$ debris disk emission around early-type stars peaks at $\\approx$ 10--20 Myr, the evolution of debris emission at $\\approx$ 20--30 Myr is not well constrained. Few clusters in this age range have been observed with Spitzer.  The most populous cluster previously assigned to this age range, NGC 2547, is likely older.  Observing 20--30 Myr old clusters may also reveal a delay in the peak debris emission for later type, lower-mass stars relative to early-type stars as implied by planet formation models \\citep[e.g.][]{Kb08}. Observations of young stars at 70 $\\mu m$ may provide tentative evidence for this delay \\citep{Hi08}.  Furthermore, the disk environment around late type, low-mass stars may be qualitatively different due to the importance of corpuscular wind drag in removing debris dust \\citep{Pl05}. Recent Spitzer studies argue that detectable debris disks around M stars may be rarer than around intermediate/high-mass stars \\citep{Ga07, Pl08}. To further investigate this possibility, MIPS observations of M dwarf stars provide new constraints on planet formation processes.  In this paper, we investigate debris disk evolution and planet formation in the young open cluster NGC 2232 located in Gould's Belt close to the Orion Nebula Cluster.  Until very recently, the most complete studies of this cluster date from over 30 years ago \\citep{Cl72, Lm74}, which identify NGC 2232 as a group of proper motion stars with B-F spectral types and an age of $\\sim$ 20-25 Myr.  Based on pre-main sequence evolutionary models,  \\citet{Ly06} estimate an age of 25--30 Myr and find a 25 Myr nuclear age.  Both \\citet{Cl72} and \\citet{Ly06} show that the cluster is nearby ($\\sim$ 320-360 pc).  At this distance, IRAC observations can probe well into the M dwarf regime; MIPS observations can detect the photospheres of early and intermediate-mass stars.  The cluster has an extremely low mean reddening (E(B-V) $\\sim$ 0.07; \\citealt{Ly06}). Thus, NGC 2232's distance, age, and reddening combine to make investigating debris disk evolution feasible in a way that complements previous studies of h and $\\chi$ Persei and NGC 2547.  Our investigation is organized as follows. After describing the observations, image processing, and photometry of Spitzer IRAC/MIPS sources, we match sources with 2MASS/optical VI data in \\S 2.  Analysis of these sources shows that many may have excess emission from circumstellar dust. To identify cluster members, we rely on ROSAT x-ray observations, proper motion data, optical/IR color-magnitude diagrams and spectroscopy in \\S 3. In \\S 4 and 5, we analyze Spitzer data of x-ray selected cluster members to identify stars with evidence for circumstellar dust, identify stars with active terrestrial planet formation, and explore any statistical trends in the disk population with the stellar properties and in the level of debris emission with time.  NGC 2232 harbors a substantial population of stars with 24 $\\mu m$ excess with a range of spectral types from late B to late K spectral type.  At least one star (HD 45435) shows evidence for warm 8 $\\mu m$ excess.  Constraining the disk luminosity and dust temperature of this source and comparing its SED to debris disk models shows that its warm dust emission is plausibly produced by terrestrial planet formation.  This survey yields robust constraints on the evolution of cold debris disks. We identify a correlation between stellar rotation and 24 $\\mu m$ excess for high-mass stars that may reveal an evolutionary link between massive primordial disks and massive debris disks.  When combined with Spitzer data of other young clusters, our survey shows that the level of debris emission around A stars in the cluster is consistent with a peak in the 24 $\\mu m$ debris emission at $\\sim$ 10--20 Myr, confirming the results of recent surveys \\citep{Cu08a}. The frequency of 24 $\\mu m$ debris emission \\textit{increases} from $\\sim$ 5 Myr and peaks between 10 Myr and 30 Myr.  Finally, this work has important implications for the frequency of icy planets around high-mass, early-type stars. The fraction of A stars ($\\approx$ 1.5--3 M$_{\\odot}$) with 24 $\\mu m$ excess emission is high ($\\ge$ 50\\%).  Because this emission likely comes from disk regions beyond the ice line ($\\approx$ 170 K), these results imply that icy planets around A stars are common. ",
        "conclusions": "We have described the first search for debris disks in the 25 Myr-old nearby open cluster, NGC 2232. Using ROSAT x-ray observations and optical/IR colors and spectroscopy, we identified probable cluster members and cross correlated this list with Spitzer IRAC and MIPS data. By comparing these data with results for other 5-40 Myr old clusters, we reconstruct the time evolution of 8 $\\mu m$ and 24 $\\mu m$ emission from debris disks and probe the connection between debris emission and stellar rotation.  Our analysis identifies one A-type cluster member, HD 45435, with strong warm dust emission and 1--3 late-type stars with weakly-emitting warm dust. HD 45435 has dust with a temperature of T $\\sim$ 350--800 K and a luminosity $L_d/L_{\\star} \\sim 5 \\times 10^{-3}$. For the late-type stars, the lack of 24 $\\mu$m excesses yields good limits for the dust luminosity, $L_d/L_{\\star} \\sim 1-10 \\times 10^{-4}$, but poor limits on the dust temperature, T $\\lesssim$ 200-300 K. The vast majority of cluster stars lack evidence for marginal warm dust emission.  Combined with results for h and $\\chi$ Per and NGC 2547, our analysis suggests that warm dust emission from debris disks around 10--30 Myr old stars is rare, with a frequency $\\lesssim$ 10\\% (see also \\citealt{Ma04, Cu07a, Go07}).  In contrast, cold dust emission around 10--30 Myr old A-type stars is common. Roughly half of all A stars with ages between 10 Myr and 30 Myr have excess emission at 24 $\\mu$m. Upper limits to the excess emission at 8 $\\mu$m yield good constraints on the dust temperature, $T \\sim$ $\\lesssim$ 200-300 K, and luminosity, $L_d/L_{\\star} \\lesssim$ 10$^{-3}$.  These results place interesting limits on the frequency of planet around young stars. For most A-type stars with 24 $\\mu$m excesses, dust emission probes active planet formation beyond the ice line, according to debris disk evolutionary models \\citep[e.g.][]{Kb08}.  Thus, at least 50\\% of A-type ($\\sim$ 1.5--3 $M_{\\odot}$) stars likely have icy planets. Warm dust emission probes terrestrial planet formation inside the ice line.  Thus, $\\lesssim$ 10\\% of all stars show evidence for rocky planet formation. Despite fairly large samples, it is not yet clear whether terrestrial planet formation (a) is rare, (b) produces only weak emission concealed by intrinsic photometric scatter, or (c) concludes much more rapidly than icy planet formation, as predicted by planet formation models \\citep{Kb04}.  For well-sampled clusters with ages 5--50 Myr, the frequency and magnitude of cold dust emission among early-type stars depends on stellar age. The frequency of 24 $\\mu$m emission increases from 5--10 Myr, peaks at 10--30 Myr, and then declines. The typical amount of cold dust emission correlates with the frequency of dust emission. Both of these trends are consistent with predictions from debris disk models \\citep{Kb04}.  Comparing the rotation rates of B and A stars with their 24 $\\mu m$ excesses shows that slower rotators have stronger excesses. Because primordial disk emission is also stronger for slower rotators \\citep{Cb07}, there may be an evolutionary link between massive primordial disks and massive debris disks.  Despite our conclusion that icy planets are common around higher-mass stars, we cannot infer the masses of planets from the luminosity of the debris disk emission.  The objects responsible for debris-producing planetesimal collisions have radii $\\gtrsim$ 500 km and masses $\\ge$ 10$^{-4}$ M$_{\\oplus}$. While the debris luminosity increases as planets grow to hundreds of kilometers in size during runaway growth, the luminosity is independent of the planet size during oligarchy\\footnote{Even though planetesimals collide at higher velocities around more massive planets, the scale height of the planetesimals also increases.  Thus, the collision rate, which controls the debris luminosity, does not increase dramatically (S. Kenyon, unpublished; \\citealt{Kb08})}. The observed L$_{d}$/L$_{\\star}$ depends on the initial disk mass -- and thus the total mass in larger objects -- more than the mass of the largest object \\citep{Kb08}.  Radial velocity surveys also constrain the frequency of planets around high-mass stars\\footnote{However, constraining the frequency of \\textit{icy} planets around high-mass stars requires an observing baseline that is much longer than that which is currently available.}. Current surveys show that the frequency of planets is higher around $\\ge$ 1.5 M$_{\\odot}$ stars than around solar and subsolar-mass stars \\citep{jjohn07}, consistent with theoretical expectations \\citep{gk08}. From the standpoint of debris disk studies, however, it is not clear whether a) high-mass stars should make more icy planets than solar-mass stars or b) high-mass stars produce more debris than they do around solar-mass stars simply because their disks are likely more massive.  This work highlights two problems that impede our understanding of terrestrial planet formation from surveys of young, open clusters with Spitzer observations.  The lack of massive young clusters prevents robust constraints on terrestrial planet formation around high and intermediate-mass stars. The low luminosity of warm debris around low-mass stars limits our ability to construct large samples of stars with warm dust. High-precision photometry of nearby clusters like NGC 2232 may probe terrestrial planet formation around low-mass stars. However, NGC 2232 and most other nearby clusters are not massive enough to provide constraints on terrestrial planet formation around intermediate and high-mass stars. The only cluster studied so far that is clearly massive enough to provide robust constraints is h and $\\chi$ Persei \\citep{Cu07a, Cu07b}.  To investigate the evolution of warm debris emission around high and intermediate-mass stars more conclusively, observations of clusters more populous than NGC 2232 are necessary.  Spitzer Cycle 4 and 5 observations of h and $\\chi$ Persei and other 8--20 Myr old clusters will fill this role.  Preliminary analysis of MIPS h and $\\chi$ Per data taken during March 2008 show a huge increase in the number of cluster stars detected, including several terrestrial zone debris disk candidates from \\citet{Cu07b} that previously had MIPS upper limits.  Observations for h and $\\chi$ Persei will be bracketed by observations of massive 8 and 20 Myr old clusters, NGC 6871 and NGC 1960, to provide better observational constraints on terrestrial planet formation during the critical 10--30 Myr epoch.  This work justifies the need for follow-up observations of NGC 2232.  New optical photometry, spectroscopy, and x-ray surveys will provide a far better constraint on cluster membership and member properties, which will improve the utility of NGC 2232 as a laboratory for studying debris disk evolution and planet formation. \\citet{Ly06} showed that optical colors provide a useful diagnostic of cluster membership; a deep spectroscopic survey to identify late-type members is also feasible with multi-object spectrographs like Hectospec on the MMT and Hydra at Kitt Peak. Identifying late-type members with a deep Chandra survey of NGC 2232 would better constrain the evolution of x-ray activity from 10 Myr to 50 Myr in addition to identifying late-type, chromospherically active cluster stars.  All these observations should provide a much more complete census of cluster members, which in turn will yield a better analysis of the disk population for late-type stars.  Finally, the proximity of NGC 2232 to the Sun allows for follow-up ground-based/space-based observations of disks, some of which are not possible with distant, populous clusters such as h and $\\chi$ Persei, NGC 6871, and NGC 1960.  The hot dust emission from ID 18566 may be resolvable by ground-based interferometers such as the Large Binocular Telescope Interferometer. The ESO space telescope \\textit{Herschel} should be able to survey NGC 2232 for evidence of cold debris disks emitting longwards of $\\sim$ 40 $\\mu m$.  With a significant disk population established, NGC 2232 should be an excellent target for mid-IR observations with the \\textit{James Webb Space Telescope}.  Although NGC 2232 has been virtually ignored by the star/planet formation community for the past 30 years, this work shows it to be a potentially important laboratory for understanding debris disk evolution and planet formation which deserves further study."
    },
    "0807/0807.0924_arXiv.txt": {
        "abstract": "We study the luminosity function, the peculiar velocities, and the sizes of voids in the Local Volume (LV) in observational samples of galaxies which contain galaxies down to $M_B=-10$ and to $M_B=-12$ within the distance $4- 8$~Mpc.  When we compare the results with the predictions of the standard cosmological LCDM model, we find that the theory faces a sever problem: it predicts a factor of ten more dwarf haloes as compared with the observed number of dwarf galaxies.  In the LV we identify voids with sizes ranging from 1 to 4.5\\,Mpc and compare the observational distribution of void sizes with the voids in very high resolution simulations of the LCDM model with WMAP1 and WMAP3 parameters. The theoretical void function matches the observations remarkably well {\\it only} if we use haloes with circular velocities $V_c$ larger than $40-45$\\,km/s ($M_{\\rm vir}=(1-2)\\times 10^{10}\\msun$) for models with with $\\sigma_8 = 0.9$ and $V_c>35$\\,km/s ($M_{\\rm vir}=(6-8)\\times 10^{9}\\msun$) for $\\sigma_8 = 0.75$.  We exclude the possibility that in the LCDM model haloes with circular velocities $< 35$\\,km/s can host galaxies as bright as $M_B=-12$: there are too many small haloes in the LCDM model resulting in voids  being too small as compared with the observations.  The problem is that many of the observed dwarf galaxies have HI rotational velocities below 25\\,km/s that strictly contradicts the LCDM predictions. Thus, the LCDM model faces the same overabundance problem, which it had with the number of satellites in the LG.  We also estimate the rms deviations from the Hubble flow $\\sigma_H$ for galaxies at different distances from the Local Group and find that in most of our model LV-candidates the rms peculiar velocities are consistent with observational values: $\\sigma_H=50$\\,km/s for distances less than 3\\,Mpc and $\\sigma_H=80$\\,km/s for distances less than 8\\,Mpc. At distances 4 (8) Mpc, the observed overdensities of galaxies are 3.5-5.5 (1.3-1.6) -- significantly larger than typically assumed. ",
        "introduction": "\\label{sec:intro} Voids in the distribution of galaxies are an important ingradient of the distribution of light and mass in the Universe.  They constitute a natural outcome of structure formation via gravitational instability, and, thus, can be used to constrain theories of galaxy formation. Emptiness of voids -- the number of small galaxies in the voids -- is an interesting question for both the observations and the theory to tackle \\citep{Peebles2001,Gottloeber2003,Patiri2006,Tinker08}. Cosmological simulations predict \\citep[e.g.,][]{Gottloeber2003} that many small DM haloes should reside in voids. There seems to be no disagreement between the $\\Lambda$CDM theory and the observations regarding the giant voids defined by $M_*$ galaxies or by $\\sim 10^{12}M_{\\odot}$ haloes \\citep{Patiri2006}. The situation is less clear on smaller scales.  In the region of $\\sim 10$\\,Mpc around the Milky Way, where observations go to remarkably low luminosities, small voids look very empty: dwarf galaxies do not show a tendency to fill the voids and voids are still relatively large. The theory predicts that many dwarf dark matter haloes should be in the voids, which puts it on a collision course with observations. Yet, below some mass the haloes are expected to stop producing galaxies inside them. There are different arguments for that: stellar feedback \\citep{Dekel1986} or photoionization \\citep{Bullock00} may play a significant role in quenching star formation in too small haloes. For example, \\citet{Loeb08} made a simple estimation of the limiting circular velocity below which haloes have essentially no gas infall due to increase of Jeans mass caused by UV background at the epoch of reionization: $V_{lim} = 34 \\times (T_{IGM}/1.5\\times10^4K)^{1/2}$\\,km/s, where $T_{IGM}$ is the temperature of intergalactic medium gas ionized by stars. \\citet{Hoeft06} studied formation of dwarf DM haloes in cosmological void regions using high-resolution hydrodynamic simulations and assuming that cosmological UV-background photo-evaporates baryons out of haloes of dwarf galaxies, and thereby limits their cooling and star formation rate.  \\citet{Hoeft06} give characteristic mass $M_c = 6 \\times 10^9 h^{-1}M_\\odot$ below which haloes start to fail accreting gas.  Theoretical estimates for the least massive luminous halo are still uncertain.  It is difficult to get a definite answer because the physics of  dwarf galaxies at high redshifts is quite complicated. Star formation histories of some isolated irregular dwarfs indicate that starbursts may produce enough power to throw gas away, but not sufficient for galaxy to get rid of it -- gas again returns to potential well of the DM halo, hosting the galaxy \\citep{Quillen08}.  Satellites of the Local Group shed light on the problem from a different direction.  The $\\Lambda$CDM model predicts that thousands of dwarf DM haloes should exist in the Local Group \\citep{Klypin99, Moore99, Madau08}, while only $\\sim 50$ are observed. Recent discoveries of very low luminosity dwarfs \\citep{Simon07} and careful analysis of incompleteness effects in SDSS \\citep{Strigari07,Simon07,Tollerud08} bring the theory and observations a bit closer, but the mismatch still seems to be present.  The currently favored  explanation of the overabundance of the dark matter subhaloes \\citep{Bullock00,Kravtsov2004} assumes that dwarf haloes above $V_c \\approx 30-50$\\,km/s were forming stars before they fall into the Milky Way or M31 and that smaller haloes never formed any substantial amount of stars. Once the haloes above the limit fall into the halo of the Milky Way or M31, they get severely stripped and may substantially reduce their circular velocity producing galaxies such as Draco or Fornax with the rms line-of-sight velocities onlya few km/s. The largest subhaloes retain their gas and continue to form stars, while smaller ones may lose the gas and become dwarf spheroidals. Haloes below the limit never had substantial star formation. They are truly dark. This scenario implies that circular velocity before the infall $V_c\\approx 30-50$\\,km/s is the limit for star formation in haloes. \\citet{Moore06} give additional arguments in favor of this scenario.  If this picture is correct, it can be tested using the abundance and the distribution of dwarf galaxies outside of the Local Group. Because dwarfs can only be detected at small distances, useful observational samples are limited to distances less than 10\\,Mpc.  While the Local Volume sample of galaxies is not as deep as the sample of satellite galaxies in the Local Group, we argue that the Local Volume dwarfs provide a unique opportunity to study the smallest and the darkest galaxies. The problem with using the LG satellites is related with the fact that the satellites have been tidally stripped by the Milky Way or by M31. Indeed, theoretical estimates indicate that a very substantial mass loss occurs even at very small distances from the centres of the dwarfs \\citep[e.g.,][]{Madau08}. Thus, we really do not know whether we deal with a truly low mass and low circular velocity satellite or with a satellite, which was much more massive in the past and was later severely stripped by its parent galaxy. The Local Volume dwarfs represent a more pristine sample in this respect.  The problem of the emptiness of voids was recently re-visited by \\citet{Tinker08} with the conclusion that voids are not a problem for the $\\Lambda$CDM model.  Most of the observational results used in \\citet{Tinker08} are based on the void probability function in the SDSS DR6 sample and on the nearest neighbor statistics in the ORS catalog as analyzed by \\citet{Peebles2001}.  The void probability function was estimated for relatively bright galaxies with $M_r<-17$. The nearest neighbor statistics goes slightly deeper: the ORS sample is formally complete to $m_B=14.5$, which gives $M_B\\approx -15$ at the distance of 8\\,Mpc. Our sample is 3 magnitudes deeper, which is essential for the ``void phenomenon''. We also note that the quality of our sample is much better than that of the ORS sample. For example, we use real distances to galaxies, not redshifts. Special effort was made to ensure that the sample is complete for LSB galaxies and for galaxies in the zone of avoidance.  \\citet{Tikhonov2006} presented results on statistics of nearby voids in the Local Volume. Here we continue the analysis using an updated list of galaxies (Karachentsev, private communication).  We characterize the spatial distribution of galaxies in the LV mostly by studying the distribution of sizes of  quasi-spherical regions - voids. The voids may still contain gas and small dark matter haloes. For our purpose void statistic is reasonably robust since 3d void maps are not very sensitive to a total number of objects in a sample. The distribution of void sizes have the advantage that they are sensitive to appearance of galaxies in very low density environments: an important property for studying the smallest existing galaxies.  We compare the spectrum of void sizes in the LV with the distribution of voids in high-resolution cosmological simulations. The simulations give us detailed information (positions, velocities, masses, circular velocities, and so on) for dark matter haloes and their satellites. However, they do not provide luminosities of galaxies. Theoretical predictions of the luminosity of a galaxy hosted by a halo with given mass, circular velocity and merging history are quite uncertain and cannot be used for our analysis. Instead, we ask a more simple question: what luminosity a halo or subhalo with given circular velocity {\\it should have} in order to reproduce the observed spectrum of void sizes. When doing this, we assume that haloes with larger circular velocities should host more luminous galaxies. We will later see that matching of the void spectrum in simulations and with the observations puts significant constraint on relation of the halo circular velocity and the luminosity of a galaxy hosted by the halo. If we take too large circular velocity, there are too few galaxies and sizes of voids become too large. Instead, if very small haloes host galaxies, the number of large voids declines well below what is observed in the Local Volume.  In addition to the analysis of distribution of voids in the Local Volume we also re-visit the problem of the deviations from the Hubble flow. The flow of field galaxies in LV appeared to be rather ``cold'': deviations from the Hubble velocity are rather small.  For example,  using the Tully-Fisher distances and applying error correction via quadrature subtraction \\citet{Shlegel94} have found for 15 galaxies within 500~km/s and outside Local Group the rms peculiar velocity $\\sim$60~km/s. They noted that such values are very rare in the CDM models, but are not uncommon in MDM models that include massive neutrinos.  Comparable values have been derived by \\citet{KM96} for galaxies within 7~Mpc from the Local Group.  \\citet{Maccio05} based on their galaxy sample obtained $\\sigma_H \\sim 52$\\,km/s within 3\\,Mpc.  \\citet{KM01} found evidence of anisotropy of Hubble flow in LV. There is a dipole component, which is due to the LG motion relative to galaxies in the LV. This is typically described as apex motion of the LV.  There is also a quadrupole component of the deviations from the Hubble flow interpreted as anisotropic expansion of Local Volume. \\citet{KM01} removed the apex motion and the quadrupole component from estimates of $\\sigma_H$.  They also removed galaxies with large velocity deviations, which were considered due to infall of galaxies onto large groups inside the LV. Galaxies inside groups were also not considered. Using only the galaxies with accurate measurements of distances  and assuming that errors in distances increase the rms  peculiar velocities, \\citet{KM01} estimate $\\sigma_H$ in distance range 1-3 Mpc from the centre of LG may be as low as $\\sim 30$ km/s. \\citet{Kar03} found $\\sigma_H \\sim 40$\\,km/s inside a sphere of 5~Mpc using distances from the luminosity of the tip of the red giant branch stars of about 20 dwarf galaxies. Their procedure of $\\sigma_H$ estimation include determination of the ``local'' Hubble constant, which is slightly different from universal expansion rate: the true Hubble constant.  All these effects and corrections are valid, but they have a tendency to systematically underestimate the deviations from the Hubble flow. Some of the corrections can be mimiced in simulations, but it is difficult to do all of them. At the same time, one should not use these corrections: any deviation from the global Hubble flow must be included in the estimates. We make only one exclusion. We still consider the apex motion because it is related with the selection of the reference frame in which the whole sample does not have a net velocity.  There were several attempts to study $\\sigma_H$ in simulations. \\citet{Gov97} emphasized  that the dispersion of random motions of field galaxies and the centres of groups allow to discriminate between models with different values of matter density $\\Omega_M$. \\citet{Klypin03} obtained in their constrained ``Local Supercluster'' simulation the peculiar line-of-sight velocity dispersion within $7h^{-1}$\\,Mpc of the model LG  $\\sigma_H\\sim 60$\\,km/sec comparable to the observed velocity dispersion of nearby galaxies. They emphasize that there is no need in exotic explanations of the ``coldness'' of Hubble flow. \\citet{Maccio05} have found that in their simulations the Hubble flow is significantly colder around model LG-candidates selected in a $\\Lambda$CDM cosmology than around LG-candidates in open or critical models. Their estimation of $\\sigma_H$ was much simpler then \\citet{KM01} approach -- they calculate rms around mean value of peculiar velocity with local Hubble constant (best fit to data).  A key question for our kind of investigation is to define what an ``LV-candidate'' is.  In other words, it is important to define our local environment. Usually the conditions for selection of LG-candidates are rather simple and include distances of the LG-candidate to the nearest Virgo-type cluster \\citep[e.g.,][]{Hoffman08} and the overdensity of a sphere of 8~Mpc radius \\citep{Maccio05}.  \\citet{Shlegel94} used the overdensity 0.25. \\citet{Maccio05} used $0.1 < \\delta\\rho / \\rho < 0.6$. \\citet{Karachentsev2004} noted that inside the radius of 8\\,Mpc centred on the Local Group the luminosity density in the B-band is 1.8 -- 2 times larger than the average luminosity density. ",
        "conclusions": "in the hierarchical scenario any void is filled with small haloes. The only question is what is the mass of the haloes. For the Local Volume with the completeness limit $M_B=-12$ this appears to be $V_c>35-45$~km/s.  If this is true, haloes with $V_c\\approx 20$\\,km/s should not host galaxies. {\\it The problem is that in reality they do}: in the Local Volume many luminous galaxies with these absolute magnitudes  rotate with velocities $V_{\\rm rot}\\approx 20$\\,km/s or smaller.  In order to demonstrate this, we select all isolated galaxies with $-13.2<M_B<-11.8$. The isolation criteria are very strict because we would like to be sure that the rotational velocities of galaxies were not reduced by stripping or by any other interaction with large neighboring galaxies. We select the galaxies, which are more than 1\\,Mpc away from any galaxy brighter than $M_B<-19$ and do not have companions within 200~kpc, which are brighter than the galaxy itself. Note that the galaxies are so small that the expected virial radius is smaller than 100\\,kpc. All the galaxies are dwarf irregulars, and for half of them there are measurements of the 21\\,cm HI line width.  Using the HI full-width-half-maximum $W_{50}$ measurements, we estimate the rotational velocity of a galaxy: $V_{\\rm rot} = W_{50}/2\\sqrt{1-(b/a)^2}$, were $b/a$ is the axial ratio. A large fraction of the HI line width is likely produced by random 8--10~km/s velocities. We do not subtract those  because we compare the results with the circular velocities of dark matter haloes. Table~4 presents the results for galaxies with detected HI emission and with $V_{\\rm rot} < 25$~km/s.  Galaxies in the Table~4 have rotational velocities well below those required by the LCDM model.  Half of the galaxies rotate slower than 20~km/s. We also studied galaxies, which are not so isolated and galaxies, which are slightly brighter than those presented in the Table~4. Results are qualitatively the same.  The disagreement between the theory and observations also shows up in the Tully-Fisher (TF) relation for galaxies in the Local Volume. In order to construct the relation, we use all galaxies in the \\citet{Karachentsev2004} sample with measured HI line width $W_{50}$, which have morphological type $ > 0$ (i.e. spirals and irregulars). The rotational velocities are corrected for inclinations. In order to allow the comparison with  theory, we do not make any corrections for internal gas motions. Open circles with error bars in Figure~\\ref{fig:TF} show the TF relation of galaxies in the Local Volume. Two lines in the plot also show the TF relation from \\citet{TF2000} and \\citet{GHASP}, which we extrapolate down to small magnitudes (GHASP results extend down to $M_B\\approx -16$). The observational TF estimates agree reasonably well for bright galaxies with $M_B<-16$. At smaller luminosities the LV results go slightly above the extrapolations from brighter samples.  When assigning luminosities to dark matter haloes we follow the prescription of \\citet{Conroy2006}. Specifically, we rank by luminosity all the galaxies in the LV sample and we rank by circular velocity all the haloes in our LV-candidate samples in simulations.  We then take the luminosity of the brightest galaxy and assign it to the halo with the largest circular velocity. Then we take the second brightest galaxy and give its luminosity to the second halo and so on. According to \\citet{Conroy2006}, this prescription reproduces clustering properties of galaxies in the SDSS sample. We add a small ($20\\%$) correction to the circular velocity of haloes to accommodate the effect of adiabatic contraction due to infall of baryons. Figure~\\ref{fig:TF} shows that the LCDM model gives a good match to observations at the bright end of the luminosity function ($M_B<-17$). The model with lower normalization produces a better fit, but even the high normalization model cannot be excluded: a more accurate treatment of the adiabatic infall may slightly improve the situation. At low luminosities the theory and observations gradually diverge, and at $M_B=-12$ the differences are quite substantial: a factor of two in circular velocities. This is the same problem, which we found using the spectrum of voids: haloes with $V_c\\approx 35-45$~km/s should have luminosities $M_B=-12$ in order to match the observational data.   We would like to emphasize that {\\it the disagreement with the theory is staggering.} The observed spectrum of void sizes disagrees at many sigma level from the theoretical void spectrum if haloes with $V_c>20$\\,km/s host galaxies brighter than $M_B=-12$. We can look at the situation from a different angle. In the LCDM model with $\\sigma_8=0.9$ there are $\\sim 320$ haloes with $V_c>45$\\,km/s -- the same number as the number of galaxies in the Local Volume with the $M_B=-12$ limit. In the same volume in the LCDM model there are $\\sim 3500$ haloes with $V_c>20$\\,km/s. If all these haloes host galaxies brighter than $M_B=-12$, the theory predicts a factor of ten more haloes as compared with the observations.  The problem has the same roots as the overabundance of substructure in the Local Group: the LCDM model predicts too many dwarf dark matter (sub)haloes as compared with the observed dwarf galaxies \\citep{Klypin99, Moore99, Madau08}. We suggest that the solution of the problem of the overabundance of the dwarfs in the Local Volume may be similar to current explanations of the substructure problem in the LG:  \\begin{itemize} \\item{}The observational sample is not complete: there are ten times more dwarf galaxies down to limiting magnitude $M_B=-12$ than listed in the \\citet{Karachentsev2004} sample. The ``missed'' dwarfs are unlikely to be dwarf irregulars because they would have HI emission and would have been detected by blind HI surveys such as HIPASS \\citep{HIPASS}. Dwarf spheroidal galaxies is a possibility. They do not have gas and cannot be detected in HI. They have very low surface brightness, which makes it difficult to detect them on photographic plates. So, it is likely that many of the galaxies were missed. Still, we do not know whether a large population of dSph galaxies exists in the LV.  If this is so, we will have another problem: how to form thousands of dwarf spheroidals in very low density environments without any tidal stripping or interaction with massive parent galaxy. The slope of the luminosity function also will be much steeper: $\\alpha \\approx 2-2.5$. \\item{} The observed galaxies with $V_{\\rm rot}\\approx 20$\\,km/s are hosted by significantly more massive haloes. The overabundance problem would be solved, if the circular velocity of a dark matter halo is $V_c \\approx 2V_{\\rm rot}$. This is somewhat similar to the solution of the overabundance problem in the LG \\citep[e.g.,][]{Penarubbia08} \\item{} Most of the dwarf haloes with $V_c<35$\\,km/s in local voids failed to form stars because they collapsed after the epoch of reionization \\citep{Bullock00}. \\end{itemize}  We also estimate the rms deviations from the Hubble flow $\\sigma_H$ for galaxies at different distances from the Local Group and find that in most of our model LV-candidates the rms peculiar velocities are consistent with observational values: $\\sigma_H=50$\\,km/s for distances less than 3\\,Mpc and $\\sigma_H=80$\\,km/s for distances less than 8\\,Mpc. At the distances 4 (8) Mpc the observed overdensities of galaxies are 3.5-5.5 (1.3-1.6) -- significantly larger than typically assumed."
    },
    "0807/0807.1113_arXiv.txt": {
        "abstract": "Observations of the Cosmic Microwave Background (CMB) provide increasingly accurate information about the structure of the Universe at the recombination epoch.  Most of this information is encoded in the angular power spectrum of the CMB. The aim of this work is to propose a versatile and powerful method for spectral estimation on the sphere which can easily deal with non-stationarity, foregrounds and multiple experiments with various specifications. In this paper, we use needlets (wavelets) on the sphere to construct natural and efficient spectral estimators for partially observed and beamed CMB with non stationary noise.  In the case of a single experiment, we compare this method with Pseudo-$C_\\ell$ methods. The performance of the needlet spectral estimators (NSE) compares very favorably to the best Pseudo--$C_\\ell$ estimators, over the whole multipole range.  On simulations with a simple model (CMB + uncorrelated noise with known variance per pixel + mask), they perform uniformly better.  Their    distinctive ability to aggregate many different experiments, to control the propagation of errors and to produce a single wide-band error bars is highlighted. The needlet spectral estimator is a powerful, tunable tool which is very well suited to angular power spectrum estimation of spherical data such as incomplete and noisy CMB maps. ",
        "introduction": "\\label{sec:framework}  \\subsection{Observation model} \\label{sec:observation-model}    Let $T$ denote the temperature anisotropy of the CMB emission. For the sake of simplicity, we consider the following observation model \\begin{multline} \\label{eq:model} X(\\xi_k) = W(\\xi_k) \\left( (B * T)(\\xi_k) + \\sigma(\\xi_k) Z_k \\right)   , \\\\   \\; k = 1,\\dots,\\npix \\end{multline} where $(\\xi_k)$ is a collection of pixels on the sphere, $W$ denotes a (0-1)-mask or any weight function $0\\leq W \\leq 1$, $B$ denotes the instrumental beam. An additive instrumental noise is modelled by the term $\\sigma(\\xi_k)Z_k$  with the assumption that $(Z_k)$ is an independent standard Gaussian sequence.  Further, we assume that $\\sigma$, $W$ and $B$ are known deterministic functions and that $B$ is axisymmetric. Typically, the variance map $\\sigma^2$ writes $\\sigma^2(\\xi_k) = \\sigma_0^2/N_{\\text{obs}}(\\xi_k)$, where $N_{\\text{obs}}$ is referred to as the \\emph{hit map}, that is, $N_{\\text{obs}}(\\xi_k)$ is the number of times a pixel in direction $\\xi_k$ is seen by the instrument.  We assume that the observations have been cleaned from foreground emissions or that those emissions are present but negligible outside the masked region. When observations from several experiments are jointly considered, the model becomes \\begin{multline} \\label{eq:model2} X_e(\\xi_k) = W_e(\\xi_k) \\left( B_e * T(\\xi_k) + \\sigma_{e}(\\xi_k)Z_{k,e}) \\right) \\\\ k = 1,\\dots,\\npix \\; , \\; e = 1,\\dots, E \\end{multline} where $e$ indexes the experiment.  The CMB sky temperature $T$ is the same for all experiments but the instrumental characteristics (beam, coverage) differ (see for example Table~\\ref{tab:fusionparams}), and the respective noises can usually be considered as independent.  \\subsection{Definition and implementation of a needlet analysis}  We recall here the construction and practical computation of the needlet coefficients.  Details can be found in Guilloux {\\it et al.} (2007\\nocite{guilloux:fay:cardoso:2008}; see also Narcowich {\\it et al.}, 2006\\nocite{narcowich:petrushev:ward:2006}).  Needlets are based in a decomposition of the spectral domain in bands or `scales' which are traditionally index by an integer $j$. Let $b_\\ell\\ji$ be a collection of window functions in the multipole domain, with maximal frequencies $\\ell\\ji_{\\max}$ (see Figure~\\ref{fig:windowfunctions} below).  Consider some pixelization points $\\xi\\ji_k, k=1,\\dots,\\npixj$, associated with positive weights $\\lambda\\ji_k, k = 1,\\dots,\\npixj$ which enable exact discrete integration (quadrature) for spherical harmonics up to degree $2\\lmax\\ji$, that is, equality \\begin{equation*} \\int_{\\Sset} Y\\lm(\\xi) \\dd \\xi = \\sum_{k=1}^{\\npixj} \\lambda_k\\ji Y\\lm(\\xi\\ji_k) \\end{equation*} holds for any $\\ell, m$ such that $\\ell \\leq 2\\ell\\ji_{\\max}$, $|m| \\leq \\ell$. Needlets are the axisymmetric functions defined by \\begin{equation} \\label{eq:defneedlet} \\psi\\ji_k(\\xi)=\\sqrt{\\lambda\\ji_{k}}\\sum_{\\ell=0}^{\\lmax\\ji} b\\ji_\\ell L_\\ell (\\xi \\cdot \\xi\\ji_k ), \\end{equation} where $L_\\ell$ denote the Legendre polynomial of order $\\ell$ normalized according to the condition $L_\\ell(1)=\\frac{2\\ell+1}{4\\pi}$. For proper choices of window functions $\\{b_\\ell\\ji\\}_j$, the family $\\left\\{\\psi_k\\ji\\right\\}_{k,j}$ is a \\emph{frame} on the Hilbert space of square-integrable functions on the sphere $\\Sset$. In a $B$-adic scheme, it is even a \\emph{tight frame} \\citep{narcowich:petrushev:ward:2006}. Though redundant, tight frames are complete sets which have many properties reminiscent of orthonormal bases (see \\eg \\cite{daubechies:1992}, chap.3).  For any field $X$ on the sphere, the coefficients $\\gamma\\ji_k := (\\lambda_k\\ji)^{-1/2}\\langle X , \\psi\\ji_k \\rangle$ are easily computed in the spherical harmonic domain as made explicit by the following diagram \\begin{equation} \\label{eq:diagram} \\begin{array}{rcc} \\left\\{X(\\xi_k)\\right\\}_{k=1,\\dots,\\npix} & \\stackrel{\\textrm{SHT}}{\\longrightarrow} & a\\lm \\\\ & & \\Downarrow \\times \\\\ \\left\\{\\gamma_k\\ji)\\right\\}_{k=1,\\dots,\\npixj} & \\stackrel{\\; \\; \\textrm{SHT}^{-1}}{\\Longleftarrow} & b\\ji_\\ell a\\lm \\end{array} \\end{equation} Double arrows denotes as many operations (\\textit{e.g.} spherical transforms) as bands.  The initial resolution must be fine enough to allow an exact computation of the Spherical Harmonics Transform up to degree $\\lmax\\ji$. If, say, the HEALPix pixelization is used, the $\\nside{}$ parameter of the original map determines the highest available multipole moments and, in turn, the highest available band $j$.   \\subsection{Distribution of the needlet coefficients}  A square-integrable random process $X$ on the sphere is said to be centered and stationary (or isotropic) if $\\esp(X(\\xi)) = 0$, $\\esp(X(\\xi)^2)<\\infty$ and $\\esp(X(\\xi)X(\\xi')) = (4\\pi)^{-1} \\sum_\\ell C_\\ell L_\\ell(\\xi \\cdot \\xi')$, with $C_\\ell$ referred to as the angular power spectrum of  $X$. The next proposition summarizes the first and second order statistical properties of the needlet coefficients of such a process.  They are the building blocks for any subsequent spectral analysis using needlets. \\begin{prop} \\label{prop:coeff-need-gauss} Suppose that $X$ is a stationary and centered random field with power spectrum $C_\\ell$.  Then the needlet coefficients are centered random variables and, for any 4-tuple $(j,j',k,k')$ \\begin{gather} \\cov[\\gamma_k\\ji, \\gamma_{k'}^{(j')}] = \\sum_{\\ell\\geq0} b_\\ell\\ji b_\\ell^{(j')} C_\\ell  L_\\ell\\left(\\cos\\theta\\right)\\label{eq:covgammajk} \\end{gather} where $\\theta=\\theta(j,k,j',k')$ is the angular distance between $\\xi_k\\ji$ and $\\xi_{k'}^{(j')}$. In particular \\begin{equation} \\label{eq:varneedcoeff} \\var[\\gamma_k\\ji] = C\\ji \\ . \\end{equation} where \\begin{equation} \\label{eq:defcj} C\\ji := (4\\pi)^{-1}\\sum_{\\ell\\geq0} \\left(b_\\ell\\ji\\right)^2 (2\\ell+1) C_\\ell. \\end{equation} \\end{prop} In other words, the variance of the coefficients $\\gamma_k\\ji$ is the power spectrum of $X$ properly integrated over the $j$-th band. \\begin{rem} It also follows from~(\\ref{eq:covgammajk}) that if the bands $j$ and $j'$ are non-overlapping (this is the case for any non-consecutive filters of Figure~\\ref{fig:windowfunctions}), all the pairs of needlet coefficients $\\gamma_k\\ji$ and $\\gamma_{k'}^{(j')}$ are uncorrelated and then independent if the field is moreover Gaussian. \\end{rem}  Suppose now that \\begin{equation*} X(\\xi_k)=\\sigma(\\xi_k)Z_k,\\ k=1,\\cdots, \\npix, \\end{equation*} is a collection of independent random variables with zero mean and variance $\\sigma^2(\\xi_k)$, where $\\sigma$ is a band-limited function. This is a convenient and widely used model for residual instrumental noise (uncorrelated, but non-stationary). Needlet coefficients are centered and, moreover, if the quadrature weights are approximately uniform ($\\lambda_k \\simeq 4\\pi/N_{\\textrm{pix}}$, as is the case of HEALPix) and $\\sigma$ is sufficiently smooth, then \\begin{equation*} \\cov[\\gamma_k\\ji,\\gamma_{k'}^{(j')}] \\simeq \\int_\\Sset \\sigma^2(\\xi)\\psi_k\\ji(\\xi)\\psi_{k'}^{(j')}(\\xi)\\dd\\xi. \\end{equation*} We denote \\begin{equation} n_k\\ji(\\sigma) := \\left( \\int_\\Sset\\sigma^2(\\xi)|\\psi_k\\ji(\\xi)|^2\\dd\\xi \\right)^{1/2}. \\end{equation} the standard deviation of the needlet coefficient of scale $j$ centered on $\\xi_k$. When the noise is homogeneous ($\\sigma$ is constant), it reduces to $\\frac{\\sigma^2}{\\npix}\\sum_{\\ell\\geq0} (2\\ell+1)\\left(b_\\ell\\ji\\right)^2$.    \\subsection{Mask and beam effects}\\label{sec:appr-deal-with}  As already noticed in the Introduction, missing or masked data makes the angular power spectrum estimation a non trivial task. Simple operations in Fourier space such as debeaming become tricky.  Needlets are also affected by both the mask and the beam.  The effect on needlets of beam and mask can be approximated as described below. These approximations, which lead to simple implementations, are validated in numerical simulations in relatively realistic conditions in Section~\\ref{sec:monte-carlo-studies}.   \\paragraph{Mask.}  Recalling that the needlets are spatially localized, the needlet coefficients are expected to be insensitive to the application of a mask on the data if they are computed far away from its edges.  Numerical and theoretical studies of this property can be found in \\cite{baldi:etal:2008a} and \\cite{guilloux:fay:cardoso:2008}.  In practice, we choose to quantify the effect of the mask on a single coefficient $\\gamma_k\\ji$ by the loss induced on the $\\mathbb L^2$-norm of the needlet $\\psi_{k}\\ji$, \\ie a purely geometrical criterion.  More specifically, needlet coefficients at scale $j$ are deemed reliable (at level $t\\ji$) if they belong to the set \\begin{equation} \\label{eq:Wt} \\mathcal{K}\\ji_{t\\ji} := \\left\\{ k=1,\\dots,\\npixj\\; :\\; \\frac{\\|W\\psi\\ji_k\\|^2_2}{\\|\\psi\\ji_k\\|^2_2} \\geq t\\ji \\right\\} \\end{equation} Parameter $t\\ji$ is typically set to $0.99$ or $0.95$ for all bands.  Note that $t\\ji \\mapsto \\mathcal{K}\\ji_t$ is decreasing, $\\mathcal{K}\\ji_0 = K\\ji$ and $\\mathcal{K}\\ji_{1^+} = \\emptyset$. % In practice, this set is computed by thresholding the map obtained by the convolution of the mask with the axisymmetric kernel $\\xi \\mapsto \\left(\\sum_\\ell b_\\ell\\ji L_\\ell(\\cos \\theta)\\right)^2$. This operation is easy to implement in the multipole domain.   \\paragraph{Beam.}  Consider now the effect of the instrumental beam. Its transfer function $B_\\ell$ is assumed smooth enough that it can be approximated in the band $j$ by its mean value $B\\ji$ in this band, defined by \\begin{equation} \\label{eq:defBj} (B\\ji)^2 := (4\\pi)^{-1}\\sum_{\\ell \\geq 0} (2\\ell + 1) (b_\\ell\\ji)^2  B_\\ell^2 . \\end{equation} In the following, the beam effect for spectral estimation is taken into account in each band.  Indeed, with Definition~(\\ref{eq:defBj}), with no noise, no mask and a smooth beam, Eq.~(\\ref{eq:varneedcoeff}) translates to \\begin{equation} \\label{eq:varneedcoeffdebeam} \\var\\left[ \\frac{\\gamma_k\\ji}{B\\ji}\\right] \\simeq  C\\ji \\; , \\; k=1,\\dots,\\npixj. \\end{equation} In other words, thanks to the relative narrowness of the bands and to the smoothness of the beam and CMB spectrum, the attenuation induced by the beam can be approximated as acting uniformly in each band and not on individual multipoles. Numerically, with typical beam values from WMAP or ACBAR experiments (see Table~\\ref{tab:fusionparams}), the relative difference (statistical bias) between the goal quantity $C\\ji$ and the estimated one $\\var(\\gamma_k\\ji)/(B\\ji)^2$ remains under 1\\% for bands below $j=27$ ($\\lmax = 875$) for WMAP-W, and below $j=39, \\lmax=2000$ for ACBAR. ",
        "conclusions": "We have presented some potentialities of the needlets on the sphere for the angular power spectrum estimation. This tool is versatile and allows to treat consistently the estimation from a single map or from multiple maps. There remains many ways of improving or modify the method described in Section~\\ref{sec:needl-spectr-estim}.  In the future, it is likely that again complementary data sets will co-exist. This is the case, in particular, for polarisation, for which Planck will measure the large scale CMB power on large scales with moderate sensitivity, while ground-based experiments will measure very accurately polarisation on smaller scales. Extensions to polarisation of the approach presented hers will likely be important for the best exploitation of such observations."
    },
    "0807/0807.4393_arXiv.txt": {
        "abstract": "{In star-forming regions like Taurus-Auriga, it has been found that most young stars are born as multiples, which theories for star formation should definitely take into account. The R CrA star-forming region has a small dark cloud with quite a number of protostars, T~Tauri stars, and some Herbig Ae/Be stars, plus a number of weak-line T Tauri stars around the cloud found by ROSAT follow-up observations.} {We would like to detect multiples among the young stars in and around the R CrA cloud in order to investigate multiplicity in this region.} {We performed interferometric and imaging observations with the speckle camera SHARP~I at the ESO 3.5\\,m NTT and adaptive optics observation with ADONIS at the ESO 3.6\\,m telescope, all in the near-infrared bands JHK obtained in the years 1995, 2000, and 2001.} {We found 13 new binaries among the young stars in CrA between 0.13~arcsec (the diffraction limit) and 6~arcsec (set as an upper separation limit to avoid contamination by chance alignments). While most multiples in CrA are binaries, there are also one quadruple (TY CrA), and one triple (HR 7170) which may form a quintuple together with the binary HR 7169.  One of the newly detected companions with a large magnitude difference found near the M3-5 type T~Tauri star [MR 81] H$\\alpha$ 17 could be a brown dwarf or an infrared companion with an edge-on disk.  Among seven Herbig Ae/Be stars in CrA, six are multiple.} {The multiplicity frequency in CrA is as high as in similar star forming regions.  By comparing with the period distribution of main-sequence stars and extrapolating to separations not probed in this survey, we conclude that the companion-star frequency is $(95\\pm23)\\,\\%$; i.e.\\ the average number of companions per primary is 0.95.} ",
        "introduction": "The star-forming region {\\em Corona Australis} (southern crown), abbreviated {\\em CrA} or {\\em R CrA}, is today known as one of the nearest regions of ongoing and/or recent intermediate- and low-mass star formation. The dark cloud around R CrA has an extinction of up to $A_{V} \\sim 45$ mag.  The age ranges between $\\le 1$ Myrs for the protostars to $9 \\pm 4$ Myrs for some T Tauri stars \\citep{james2006}. See \\citet{NeuhReview2008} for a recent review.  \\citet{cas98} determined the distance towards the eclipsing double-lined spectroscopic binary (SB) TY CrA to be $129 \\pm 11$~pc from their orbit solution.  Because this is consistent with other earlier estimates \\citep[see discussion in][]{NeuhReview2008} and, at the same time, the best current estimate, we use $\\sim 130$ pc as the distance for the young stars in CrA.  \\citet{NeuhReview2008} give a list of all known optically visible members compiled from various sources. Only a few multiples have been published among the CrA members so far; see Table~1 for a listing. Hence, while six out of seven Herbig Ae/Be stars are known or suspected to be multiple, only five T~Tauri stars (TTS) were known to be multiple, a low number compared to other star-forming regions.  In addition, the brown dwarf member Denis1859 is a known binary \\citep{bouy2004}.   \\begin{table*}[ht] \\caption[]{Previously known multiples among optically visible young stars in and around CrA (c,d)} \\label{PrevBin} \\input 9897tab1.tex  Remarks: (a) Always $\\le 1$ here by definition. (b) For consistency, we give here the reciprocal value of the flux ratio given by Ghez et al. (c) RXJ1846.7-3636 was listed as $\\sim 8''$ binary in (11), i.e. with separation above the 6'' upper limit used here, hence regarded as two separate stars here. (d) The very low-mass star or brown dwarf CrA 444 may also be a binary \\citep{lop2005}.  References: (1) \\citet{joy44}, (2) \\citet{Ghez97}, (3) \\citet{Baier85}, (4) \\citet{reip93}, (5) \\citet{Bailey98}, (6) \\citet{chelli95}, (7) \\citet{Prato03}, (8) \\citet{chau03}, (9) \\citet{Ageorges97}, (10) \\citet{Walter97}, (11) \\citet{Neuh2000}, (12) \\citet{bouy2004}, (13) \\citet{gap34}, (14) \\citet{cas98}, (15) \\citet{proust81} (16) \\citet{CCDM} \\end{table*}  Therefore, we performed this new homogeneous multiplicity survey of most CrA members with the infrared speckle camera SHARP~I and the infrared adaptive optics (AO) instrument ADONIS. We explain our observations and the data reduction in Sect. 2, list all results in Sect. 3, compare our observations of multiples with previous observations in Sect. 4, and conclude in Sect. 5. ",
        "conclusions": "We found 13 new binaries among young stars in the R CrA association by simple speckle imaging, speckle interferometry, and AO imaging, all with 3.5--3.6 meter telescopes.  The companion found next to [MR 81] H$\\alpha$ 17 should be mentioned explicitely (Fig.\\ \\ref{VisBinFig}): It has a brightness ratio of only $0.046 \\pm 0.004$ between the faint NE component and the bright SW component (Table~\\ref{BinTab}). The unresolved system (i.e.\\ the bright component) has a spectral type of M3-5 \\citep{pat98}. Hence, the fainter object ($\\rm K\\simeq13.1\\,mag$) may be a substellar companion or a so-called infrared companion (e.g.\\ with edge-on circumstellar disk) or a background object. We continue to observe it in more detail (AO at larger telescopes and spectroscopy).  All other companions found are almost certainly normal TTS.   Among eight known cTTS, there are only two binaries ($25\\pm18\\,\\%$), but among 28 known wTTS, there are 13 binaries found ($46\\pm13\\,\\%$). No mixed systems were found, but for most companions and even some primaries, its not known whether they are wTTS or cTTS.  The system [MR 81] H$\\alpha$ 17 SW-NE may be mixed.  The multiplicity in CrA is high and not significantly different from other similar star-forming regions like Sco-Cen or Tau-Aur. The multiplicity among new ROSAT TTS around the dark cloud is high, just as it is in the case Tau-Aur. There is no significant difference between on-cloud TTS and off-cloud TTS in CrA, but a tendency to more binaries among the off-cloud TTS.  The multiplicity among the seven Herbig Ae/Be stars is high: one quadruple, one triple, and four binaries (from SBs to 6 arcsec separation). If going out to 13 arcsec separation, we would even have one quintuple.  However, the sample of known Herbig Ae/Be stars in CrA is too small to compare to other regions. All triples and quadruples known in CrA are hierarchical.  For a comparison of the distribution of separations (or orbital periods) found here in CrA with those of nearby solar-type, main-sequence stars (DM91), see Fig.~\\ref{MSfig}.  We can extrapolate to the total companion-star frequency (at any separation) in CrA being $\\sim95\\,\\%$, also similar to Sco-Cen and Tau-Aur."
    },
    "0807/0807.4949_arXiv.txt": {
        "abstract": "We have  selected a  sample of X-ray-emitting active  galactic nuclei (AGNs)   in   low-mass   host   galaxies   ($\\sim$5$\\times10^{9}$- 2$\\times10^{10}$ M$_{\\odot}$)  out to $z$$\\sim$1. By  comparing to AGNs in more massive  hosts, we have found that  the AGN spatial number density and the fraction of galaxies  hosting AGNs depends strongly on the host mass, with the AGN host mass function peaking at intermediate mass and with  the AGN  fraction increasing  with host  mass.  AGNs  in low-mass hosts  show strong cosmic  evolution in  comoving number  density, the fraction of such galaxies hosting active nuclei and the comoving X-ray energy density.   The integrated X-ray luminosity function  is used to estimate the  amount of the accreted  black hole mass in  these AGNs and places a strong lower limit of  12\\% to the fraction of local low-mass galaxies hosting black  holes, although a more likely  value is probably much  higher ($>$50\\%)  once the  heavily obscured objects  missed in current X-ray surveys are accounted for. ",
        "introduction": "Active  galactic nuclei  (AGNs),  the manifestation  of accretion  onto massive  black holes  (MBHs),  have   been  recognized  as  a  critical ingredient in galaxy formation  and evolution. The demography of local galaxies suggests  that most --  perhaps all -- massive  galaxies host MBHs  at their centers  and that  MBH masses  are correlated  with the galaxy  bulge properties  \\citep{Kormendy95,  Magorrian98, Gebhardt00, Ferrarese00, Haring04},  implying the  coevolution of  the galaxy and MBH.   The good match between  the local BH mass  density and the mass density of AGN relics further  suggests that all massive galaxies have experienced an AGN  phase during their  evolution \\citep[e.g.][]{Aller02, Shankar04,  Marconi04}.   Energy  feedback  from  AGNs  to  their  host galaxies  is invoked to  explain different  aspects of  massive galaxy evolution.  AGNs may serve as  the heating sources for cooling flows in clusters \\citep[e.g.][]{McNamara07}.  They may suppress star formation in their host  galaxies and cause them to migrate  from the blue cloud to  the  red-sequence  in  the color-magnitude  plot  \\citep{Croton06, Nandra07,  Georgakakis08}  and they  may  account for  ``down-sizing'' galaxy  evolution  \\citep[e.g.][]{Cowie96}.    As  the  most  abundant population in the universe,  low-mass (defined as stellar mass $M_{*}$ $<$ 2$\\times$10$^{10}$ M$_{\\odot}$ through this paper) galaxies act as the  building  blocks  of  massive  galaxies.   However,  the  current understanding of  MBHs' role in low-mass galaxy  evolution is limited, leaving  us with  some basic  questions: Is  the existence  of  BHs in low-mass galaxies  as common  as it is in massive  galaxies?  Are  there two types of low-mass galaxy populations  (i.e.  ones with BHs vs. ones without  BHs)?  How many  low-mass galaxies  experience an  AGN phase? All of these questions are  related to a more basic  astrophysical problem: what is the  black hole occupation  function  (BHOF; the  fraction of  galaxies hosting either active or quiet BHs) in low-mass galaxies?   Searching for  BHs in low-mass systems offers an unique opportunity to extend the  MBH-$\\sigma$ correlation  for massive galaxies,  which not only tests  the universality of  the relation but  also is required  to understand the  origin of  the relation.   Although the MBH-$\\sigma$ relationship  has been confirmed for  low-mass systems in the   local  universe   \\citep{Barth05},  some   works   suggest  that the MBH-$\\sigma$ relationship  may be  replaced at low  mass by  a similar relation    involving   compact   stellar    nuclei   \\citep{Wehner06, Ferrarese06, Rossa06}.  A  better understanding  of primordial  MBH seed  growth in  the early universe may also benefit from  the study of BHs in low-mass galaxies, as  such work  could reveal  aspects of  the accretion  mode in  a low gravitational  potential and  low metallicity  environment and  of the relative  importance  of mass  growth  through  accretion and  merging processes. For example, in  low-mass galaxies, the impulsive kick from anisotropic gravitational emission during  BH-BH mergers is thought to be  strong enough  to eject  central BHs  \\citep{Favata04, Merritt04}. Understanding AGN activity in  low-mass galaxies can further constrain MBH  seed formation  theories.  In  a cold  dark matter  universe, the primordial  MBH  seeds  form  as  remnants  of  Population  III  stars \\citep[e.g.][]{Madau01} or through the direct collapse of pre-galactic gas disks  \\citep{Lodato06}.   The  efficiencies  in  different  formation scenarios predict a wide range  in the local BHOF in low-mass galaxies which can vary from zero to unity, while the prediction of the BHOF in massive galaxies  is invariably  unity as constrained  by observations \\citep{Volonteri08}.    The current study of BHs in low mass galaxies is mostly limited to low redshift ($z$  $<$ 0.3).  Despite  a variety of investigations,  it is still  unclear  how many  local  low-mass  galaxies  harbor BHs.   The dynamical searches  for MBHs have confirmed  the existence of  a BH in M32 \\citep{Verolme02}, but not in M33 \\citep{Merritt01, Gebhardt01} or NGC205 \\citep{Valluri05}.  Searching  for actively accreting BHs (i.e. AGNs)  has provided  a stronger  lowerlimit and  more complete  view of local BHOFs as a function  of the system mass through optical emission lines  \\citep[e.g.][]{Filippenko89,  Maiolino95,  Ho95,  Kauffmann03a, Greene04,     Greene07c,     Decarli07, Dong07, Shields08}, mid-infrared line selections \\citep{Satyapal07, Satyapal08}     and     X-ray     emission \\citep[e.g.][]{Gallo07, Ghosh08}.  It has been shown that the fraction of  galaxies hosting active  BHs depends  on both  BH mass  and galaxy stellar mass, peaking at an intermediate mass range and falling toward higher and lower mass \\citep{Kauffmann03a, Heckman04, Greene07b}.   As the local  BH masses are most likely  accumulated at high redshift, the  search for  AGN activity  in low-mass  galaxies at  high redshift should  provide independent  and  possibly better  constraints on  the local BHOF. Such study is also significant for understanding the evolution of  nuclear activity  in low-mass  galaxies and  the role  of  AGNs in low-mass galaxy evolution.  With the  advent of deep {\\it Chandra} and {\\it XMM-Newton}  X-ray surveys, we  have searched for  X-ray emitting AGNs in low-mass host galaxies out to $z$ $\\sim$ 1.  In this paper, we first describe  the identification of  AGNs in low-mass  galaxies (see \\S~\\ref{DATA} and  \\S~\\ref{DAGN_DEF}).  We then  present the 1/$V_{\\rm max}$ method  to correct for  incompleteness in \\S~\\ref{Method_NumDen} and study  their spatial number density and  X-ray luminosity function in  \\S~\\ref{RESULT}.  In \\S~\\ref{Discussion},  we discuss  the local BHOF in  low-mass galaxies constrained  by our study  of high redshift AGNs  in   such  galaxies.   Our  conclusions   are  presented  in \\S~\\ref{Conclusions}.  Throughout  this paper, ``low-mass''  refers to normal galaxies  or AGN  host galaxies with  stellar mass  $M_{*}$ $<$ 2$\\times$10$^{10}$  M$_{\\odot}$ and  ``massive'' indicates  those with stellar mass  $M_{*}$ $>$ 2$\\times$10$^{10}$ M$_{\\odot}$.   We adopt a cosmology with $H_{0}$=70 km s$^{-1}$ Mpc$^{-1}$, $\\Omega_{\\rm m}$=0.3 and  $\\Omega_{\\Lambda}$=0.7.  All  magnitudes  are defined  in the  AB system. ",
        "conclusions": "We have studied a sample of  X-ray AGNs in low-mass host galaxies with stellar  mass   of  5$\\times$10$^{9}$  M$_{\\odot}$   $<$  $M_{*}$  $<$ 2$\\times$10$^{10}$   M$_{\\odot}$  out  to   $z$  $\\sim$1.    Our  main conclusions are:    (1) By  including  AGNs  in   more  massive  host  galaxies,  we  have constructed the  stellar mass function of AGN  host galaxies extending down to the  low-mass regime in two redshift intervals  of 0.1 $<$ $z$ $<$ 0.4 and 0.4  $<$ $z$ $<$ 0.7.  We have found  the AGN host stellar mass  function peaks  at an  intermediate  mass range,  with the  peak shifting toward higher mass at higher redshift.   (2) By comparing to AGNs in more massive hosts, we have found that the fraction  of galaxies hosting  active nuclei  depends strongly  on the host mass, increasing from  low-mass galaxies continuously to the most massive ones.  The fraction of low-mass galaxies hosting active nuclei suggests that  a large fraction  of such galaxies  at 0 $<$ $z$  $<$ 1 harbor black holes  at their centers as long as  the low-mass host AGN lifetime  with $L_{X}$  $>$ 10$^{42}$  erg s$^{-1}$  is not long ($<$ 0.06 Gyr).  (3) The X-ray luminosity functions of AGNs in low-mass hosts have been constructed in two  redshift intervals of 0.1 $<$ $z$  $<$ 0.4 and 0.4 $<$ $z$ $<$ 0.7.  AGNs  in low-mass hosts contribute $\\sim$10\\% of the X-ray energy  density of all the AGNs  in 0 $<$ $z$  $<$ 1, indicating that such  AGNs make an energetically significant  contribution to the cosmic X-ray background.  (4) AGNs  in low-mass hosts  show strong  redshift evolution  in their comoving  number density, the  fraction of  such galaxies  with active nuclei and the comoving  radiation energy density.  By integrating the X-ray luminosity function  of these AGNs over the  redshift range of 0 $<$  $z$  $<$  1,  the  accreted  black hole  mass  in  galaxies  with 5$\\times$10$^{9}$  M$_{\\odot}$   $<$  $M_{*}$  $<$  2$\\times$10$^{10}$ M$_{\\odot}$ is  (3.9$\\pm$0.9)$\\times$ 10$^{3}$ M$_{\\odot}$ Mpc$^{-3}$. This number  gives a  strong lower  limit of 12\\%  to the  fraction of local  low-mass galaxies  harboring  black holes,  which  may be  much higher ( $>$ 50\\%) if the dusty torus of the AGNs in low-mass galaxies has  similar structure  to or  is  more opaque  than that  of AGNs  in massive host galaxies."
    },
    "0807/0807.3044_arXiv.txt": {
        "abstract": "We present observations of the nuclear star cluster in the nearby edge-on spiral galaxy NGC~4244 using the Gemini Near-Infrared Integral Field Spectrograph (NIFS) with laser guide star adaptive optics.  From a previous study of edge-on galaxies, this nuclear star cluster was found to be one of a sample of clusters that appear flattened along the plane of their host galaxies disks.  Such clusters show evidence for multiple morphological components, with younger/bluer disk components and older/redder spheroidal components.  Our new observations of NGC~4244 show clear rotation of 30~km\\,s$^{-1}$ within the central 10~pc (0.5$\\arcsec$) of the cluster.  The central velocity dispersion is found to be 28$\\pm$2~km\\,s$^{-1}$.  The multiple stellar populations inferred from the optical colors and spectra are seen as variations in the CO line strength in the NIFS spectra.  The rotation is clearly detected even in the older, more spheroidal stellar component.  We discuss evidence for similar structures and kinematics in the nuclear star clusters of other galaxies including M33 and the Milky Way.  Our observations support two possible formation mechanisms: (1) episodic accretion of gas from the disk directly onto the nuclear star cluster, or (2) episodic accretion of young star clusters formed in the central part of the galaxy due to dynamical friction. ",
        "introduction": "Prominent nuclear star clusters (NSCs) are found in $\\sim$75\\% of late-type spiral galaxies \\citep{boker02} and early-type galaxies \\citep[][]{cote06}.  These NSCs are intrinsically very luminous, with typical $M_I \\sim -12$, and sizes similar to globular clusters \\citep[$r_{eff} \\sim 5$pc;][]{boker04a}.  Recently, multiple studies \\citep{ferrarese06,wehner06,rossa06} have shown that NSC masses are coupled to the mass of their host galaxy, following a similar $M$-$\\sigma$ relation as massive black holes (MBHs) and extending this relation to lower mass galaxies.  This result suggests that the formation of both types of objects are linked to the formation of their host galaxies.  Some galaxies host both NSCs and MBHs, but the relationship between these two types of objects is unknown \\citep{seth08}.  Unlike black holes, the assembly history of NSCs can, in principle, be inferred from studies of their stellar populations and kinematics.  Thus NSCs provide important tools for studying the history of galaxy centers and their connection to the galaxies as a whole.  In spiral galaxies, integrated optical spectra of NSCs show that they nearly always have multiple stellar populations \\citep{walcher06,rossa06}. More specifically, NSCs typically have a young stellar population ($<$100~Myr) but are dominated in mass by stars older than $>$1~Gyr.  Further support to this picture comes from the observations of several distinct populations in the color-magnitude diagram of the nucleus of the Sgr dwarf spheroidal \\citep[M54;][]{monaco05,siegel07}.  In \\citet{seth06} (Paper~I), we found from a study of edge-on galaxies that in some NSCs these multiple stellar populations appear to be associated with different morphological components.  These clusters are flattened along the major-axis of the galaxy and have color maps and morphologies that suggest they consist of young disks and older more spheroidal components.  We present here follow-up observations of the NSC in NGC~4244, the nearest galaxy in the Paper~I sample \\citep[D=4.3~Mpc, 1$\\arcsec$=21~pc;][]{seth05a}.  In Paper~I, we presented integrated optical spectra of this cluster that indicate the presence of multiple stellar populations.  Although the detailed star formation history is not well-constrained, populations of 50-100~Myr combined with older populations are required to provide a good fit to the spectra. Redshifted emission identified with an HII region offset by 19~pc along the major axis of the cluster gives a lower-limit to the mass of $\\sim$2.5$\\times$10$^6$~M$_\\odot$.  This mass measurement is consistent with the masses derived from the stellar population fitting of the spectrum as well.  In \\S2 we describe the observations and data reduction.  Our results on the kinematics, morphology and stellar populations of the cluster are presented in \\S3.  We discuss the results in the context of other galaxies and formation models in \\S4. ",
        "conclusions": "In NGC~4244, we find a flattened nuclear star cluster with multiple components, a younger disk structure and an older spheroidal structure (Paper~I).  From our NIFS data, we have shown that the cluster is rotating at $\\pm$30~km\\,s$^{-1}$ within the central 10~pc of the cluster. This rotation is not confined just to the younger disk component, but also appears to be present in the older, more spheroidal component.  There are two basic models for NSC formation: (1) globular clusters which collect at the center due to dynamical friction \\citep{tremaine75,lotz01}, or (2) episodic accretion of gas onto the center followed by {\\it in situ} star formation \\citep{milosavljevic04,bekki07}.  A third alternative (3), combining these two scenarios, has recently been proposed by Milosavljevic \\& Agarwal (2008, {\\it private communication}) in which NSCs form from young star clusters in the inner part of the galaxy disk which are massive enough to accrete onto the nucleus.  These mechanisms are not exclusive; NSCs could form from a combination of these effects, and different types of galaxies may have different dominant mechanisms. However, our observations for NGC~4244 strongly conflict with the first scenario.  If a significant fraction of the NSC were formed from globular clusters distributed in a spheroidal distribution no rotation would be expected.  Thus the primary formation mechanism must be through episodic accretion of material from the disk of the galaxy, either gas or young star clusters.  Distinguishing between gas and star cluster accretion may be possible through comparison of our kinematic data with simulations.  Direct evidence of gas accretion into the central $\\sim$10 parsecs is seen in a couple nearby nuclear star clusters using CO observations \\citep{schinnerer03,schinnerer06}.   There is substantial evidence that the structure we see in NGC~4244 is not unique, and in fact may be typical for NSCs in spiral galaxies. In Paper~I, 5/9 NSCs in edge-on spirals were significantly flattened along the major axis of the galaxy disk, with three showing clear composite morphologies.  NGC~4244 is the nearest of these galaxies and thus presents the clearest view of these structures.  There is also evidence for similar structures in earlier type galaxies; \\citet{balcells07} find a number of disky ``nuclear extended components'' with scale lengths as small as 5 pc in a sample of early-type spirals and elliptical galaxies.  Previous studies of nuclei in large samples of spirals \\citep{boker02,carollo02} have focused on face-on galaxies making the structures in the disk plane difficult to detect.  And very few galaxies have the pc-scale quality of kinematic data that we have here.  However, where this data exists there is evidence for nuclear structures similar to those presented here.  The Milky Way hosts a nuclear star cluster centered on Sgr~A$^*$, first noted by \\citet{becklin68}.  The luminosity profile of the inner $\\sim$100~pc of the Milky Way is distinct from the underlying bulge and disk profiles \\citep[e.g.][]{serabyn96}. This component is clearly flattened along the galactic plane \\citep[axial ratio of 0.45;][]{catchpole90}, but less so in the central 30 pc, which led \\citet{launhardt02} to propose the existence of two components, a 3$\\times$10$^7$ Msol NSC dominating the central 30~pc (with $r_{eff} \\sim 10$~pc) and a more massive nuclear stellar disk extending out to 100-200~pc.  Numerous young OH/IR stars are found, primarily within the central 40~pc and these have a flattened distribution with rotation of $>$100~km\\,s$^{-1}$ \\citep{lindqvist92}. Based on integrated K-band spectra, \\citet{mcginn89} find that the central dispersion is $\\sim$125~km\\,s$^{-1}$ and drops quickly with increasing radius, while the rotation increases outwards equalling the rotation at a radius of $\\sim$3~pc, the edge of their study.  Compared to the NGC~4244 NSC, the Milky Way appears quite similar, but scaled up by $\\sim$10$\\times$ in mass and $\\sim$3$\\times$ in size, rotation, and dispersion as would be expected for a more massive galaxy.  The very center ($<$1~pc) of the MW NSC has been intensively studied and shows complex dynamics with striking differences between older stars and very young ($<$10 Myr) stars. The latter show ordered rotation (and counter rotation!) and provide evidence for processes 2 \\citep{genzel03} and 3 \\citep[e.g.][]{maillard04,lu05} described above.  There is also evidence for flattening and rotation in the M33 NSC. This cluster is elongated along the major axis of the galaxy with an axial ratio of $\\sim$0.85 \\citep{lauer98,matthews01}.  At a radius of $\\sim$1$\\arcsec$ (4~pc), the rotation is $\\sim$8~km\\,s$^{-1}$ and dispersion is 27~km\\,s$^{-1}$ \\citep{gebhardt01}.  Even considering that the orientation is likely not edge-on, it appears that the rotation does not dominate in the M33 NSC to the same extent as in NGC~4244.  In conclusion, we have shown above that the nuclear star cluster in NGC~4244 is rotating.  The rotation is seen in both the young disk and older spheroid components, suggesting that accretion of gas or star clusters from the disk is the dominant mechanism in the formation of this NSC.  While there is some evidence that other NSCs in spiral galaxies have similar structures (and thus formation histories), a better understanding of NSC formation across Hubble type will require a detailed study resolving the stellar populations and kinematics of numerous nearby NSCs.  Acknowledgments: The authors thank Tracy Beck, Andre Wong, and Richard McDermid for their help with the observation planning and analysis. We thank Tim Davidge for giving us his spectral templates.  This paper was much improved by the suggestions of the referee, Torsten B\\\"oker. We also acknowledge helpful conversations with Thomas Puzia, Kevin Covey, Dawn Erb, and Andr\\'es Jord\\'an.  Anil Seth gratefully acknowledges the support of the CfA Postdoctoral Fellowship.  Based on observations obtained at the Gemini Observatory, which is operated by the Association of Universities for Research in Astronomy (AURA), Inc., under a cooperative agreement with the National Science Foundation (NSF) on behalf of the Gemini partnership: the NSF (United States), the Science and Technology Facilities Council (United Kingdom), the National Research Council (Canada), CONICYT (Chile), the Australian Research Council (Australia), the Minist\\'erio da Ci\\'encia e Tecnologia (Brazil), and SECYT (Argentina).  {\\it Facilities:} \\facility{Gemini:Gillett (NIFS/ALTAIR)}, \\facility{HST (ACS/WFC)}"
    },
    "0807/0807.4452_arXiv.txt": {
        "abstract": "{It has been demonstrated that active galactic nuclei are powered by gas accretion onto supermassive black holes located at their centres. The paradigm that the nuclei of inactive galaxies are also occupied by black holes was predicted long ago by theory. In the last decade, this conjecture was confirmed by the discovery of giant-amplitude, non-recurrent X-ray flares from such inactive galaxies and explained in terms of outburst radiation from stars tidally disrupted by a dormant supermassive black hole at the nuclei of those galaxies. } {Due to the scarcity of detected tidal disruption events, the confirmation and follow-up of each new candidate is needed to strengthen the theory through observational data, as well as to shed new light on the characteristics of this type of events. } {Two tidal disruption candidates have been detected with \\emph{XMM--Newton} during slew observations. Optical and X-ray follow-up, post-outburst observations were performed on these highly variable objects in order to further study their classification and temporal evolution.} {We show that the detected low-state X-ray emission for these two candidates has properties such that it must still be related to the flare. The X-ray luminosity of the objects decreases according to theoretical predictions for tidal disruption events. At present, optical spectra of the sources do not present any evident signature of the disruption event. In addition, the tidal disruption rate as derived from the \\emph{XMM--Newton} slew survey has been computed and agrees with previous studies.} {} ",
        "introduction": " ",
        "conclusions": ""
    },
    "0807/0807.1870_arXiv.txt": {
        "abstract": "{We study the chaotic behavior of the  synchronization phase transition in  the Kuramoto model. We discuss the  relationship with analogous features found  in the Hamiltonian Mean Field (HMF) model. Our numerical results support the connection between the two models, which can be considered as limiting cases (dissipative and conservative, respectively) of a more general dynamical system of damped-driven coupled pendula. We also show that, in the Kuramoto model, the shape of the phase transition and the largest Lyapunov exponent behavior are strongly dependent on the distribution of the natural frequencies. }   \\begin{document} ",
        "introduction": "Long-range interacting systems have been intensively studied in the last years and new methodologies have been developed in the attempt to understanding their intriguing features. One of the most promising directions is the combination of statistical mechanics tools and methods adopted in dynamical systems \\cite{long-range}. In particular, phase transitions have been extensively explored in both conservative and dissipative long-range systems. The Hamiltonian Mean Field (HMF) model \\cite{hmf} and the Kuramoto model \\cite{kuramoto,strogatz,acebron} represent two paradigmatic toy models, the former conservative and  the latter dissipative, for many real systems with long-range forces and have several applications. Both  models share the same order parameter and display a spontaneous phase transition from an homogeneous/incoherent phase to a magnetized/synchronized one. \\\\ In \\cite{kura-hmf} we already observed that HMF and Kuramoto models can be considered as limiting cases (respectively conservative and overdamped) of a more general model of driven-damped coupled inertial oscillators. In this paper we present new numerical results which support a common scenario for the two models. More precisely, first we discuss the well known equilibrium features of the second order phase transition in the HMF model, then we study the stationary asymptotic behavior of the Kuramoto model as a function of the coupling strength. On one hand, through new numerical simulations of large size systems, we confirm that, as also pointed out by other authors \\cite{pazo} , the shape of the dynamical phase transition in the Kuramoto model changes from a continuous to an abrupt one, depending on the distribution of the natural frequencies of the oscillators. On the other hand, and this is our main result, we clearly show  that, as  for the  HMF model, the largest Lyapunov exponent (LLE) of  the Kuramoto model exhibits a peak just around the critical value, confirming the generality of this microscopic signature for a phase transition. Chaotic behavior in the Kuramoto model was discussed previously in ref.\\cite{maistrenko}. However those authors compute the entire spectrum of the Lyapunov exponents  only for small sizes and only for one kind of natural frequencies distribution, without discussing the strong dependence of the chaotic behavior on the shape of that distribution and its persistence in the thermodynamical limit. Finally, by tuning the width of the natural frequencies distribution, we show how the phase transition changes continously from 2nd-order-like behavior towards  a 1st-order-like one and we draw a complete synchronization phase diagram. As far as we know,  these results are reported for the first time and we think that they could provide  with  new insights for the study of dynamical phase transitions in systems displaying collective synchronization. ",
        "conclusions": "We presented new numerical evidence of the presence of  chaotic behavior in the Kuramoto model for very large system sizes, discussing the analogies with the Hamiltonian Mean Field (HMF) model. We studied the phase transition features and the LLE behavior for both  models. The latter can be also regarded as the dissipative and the conservative version of a more general model of coupled driven/damped pendula. Our simulations confirm that two different kinds of dynamical phase transitions occur in the Kuramoto model, depending on the distribution of the natural frequencies adopted as driving term. A uniform $g(\\omega)$ gives rise to a sharp 1st-order-like transition, where both homogeneous and synchronized stationary states coexist. Instead, a Gaussian $g(\\omega)$ yields a continuous 2nd-order-like transition, very similar to the true thermodynamical phase transition observed in the HMF model. On the other hand, the presence in the Kuramoto model of a peak observed in the LLE correspondingly to the critical region and regardless of the kind of distribution $g(\\omega)$ reflects the fact that in this region the competition between locked and drifting oscillators activates a microscopic chaotic dynamics which is a good dynamical indicator of the phase transition. Again, such a chaotic behavior shows many analogies with the one observed in the HMF model, which exhibits as well a peak just before the critical point, where there are large fluctuations in the main thermodynamical quantities characterizing the macroscopic phase transition."
    },
    "0807/0807.1794_arXiv.txt": {
        "abstract": "The peak amplitude of linear polarization detected recently from an extrasolar hot giant planet HD 189733b, is a few times of $10^{-4}$, more than an order of magnitude higher than all theoretical predictions. Rayleigh scattering off $H_2$ and $He$ may although give rise to a planet-star flux ratio of the order of $10^{-4}$ in the blue band, it cannot account for the high polarization unless the planet has an unusually extended atmosphere. Therefore, it is suggested that the high polarization should be attributed to the presence of a thin cloud of sub-micron size dust grains in the upper visible atmosphere which supports the observational finding of an almost feature-less transmission spectrum in the optical with no indication of the expected alkaline absorption features. It is found that the polarimetry observation allows for a small eccentricity of the orbit that is predicted from the time delay of the secondary eclipse of the planet. The estimated longitude of the ascending node is $16^o\\pm6$ which interestingly coincides with the observationally inferred location of the peak hemisphere-integrated brightness. ",
        "introduction": "Polarization has always been an efficient tool to probe the physical properties in the environment of various astrophysical objects. \\cite{sen01} predicted detectable amount of linear polarization from L dwarfs because of the presence of condensates in the visible atmosphere. Subsequently, linear polarization attributed to dust scattering \\citep{sen03,sen05} has been detected in several brown dwarfs \\citep{menard,osorio} whose atmospheres resemble those of irradiated extrasolar planets, the 'Hot Jupiters'.  The use of polarimetry in detecting and understanding the physical properties of extra-solar planets, especially the close-in-planets or the so called roaster such as the first discovered extra-solar planet 51 Peg b \\citep{mayor95}, is emphasized by \\cite{seager00,seager03,sen06}.  Recently, \\cite{berd} have reported detection of linear polarization in the blue band from one of the well-studied 'Hot Jupiter' HD 189733b \\citep{bouchy}. If the observation is correct, then it should have several important implications to our understanding of the physical processes of the planetary atmosphere. The observed peak amplitude of polarization is about $10^{-4}$ implying the B-band flux ratio of the planet and the star to be more than one order of magnitude higher than that predicted by current theroetical models \\citep{showman,burr07,barman}. In order to model and interpreate the observed polarization, \\cite{berd} have assumed the planet as a Lambert sphere of perfectly reflecting surface. A Lambert sphere follows Lambert's law of diffuse reflection. On this law the diffusely reflected light is isotropic in the outward hemisphere and is natural, independently of the state of polariation and the angle of incident light (Chandrasekhar 1960, p. 147).  That a Lambert sphere of perfectly reflecting surface yields zero polarization is also shown by \\cite{stam06}. Consequently, the model fit and the inference by \\cite{berd} that the observed polarization implies that the planet has an abnormally large scattering radius of 1.5-1.7 $R_J$ is unphysical.  In the present letter, I show that single scattering of photons by sub-micron size grains in a thin layer of silicate cloud in the upper atmosphere of HD 189733b may explain the observed polarization. I discuss the polarization model in the next section. The cloud model adopted is discussed in section~3. The results are discussed in the fourth section and specific conclusions are made in the last section. ",
        "conclusions": "As mentioned in section~1, \\cite{berd} obtained a fit of the observed polarimetric data by assumeing a perfectly reflecting Lambert sphere with the radius as large as 1.5-1.7 $R_J$. While a Lambert sphere does not yield any polarization, the argument in favour of an unusually large Rayleigh scattering exosphere is somewhat prematured at this stage. Taking the scattering radius to be the same as the optical radius measured through transit photometry, the geometric albedo as implied by the polarization is larger than 2/3 which is the geometric albedo of a Lambert sphere of perfectly reflecting surface. This may not be impossible as some Solar system objects have geometric albedo exceeding unity due to strong back-scattering. However, \\cite{berd} interpreated the polarization by considering smaller albedo but an abnormally large planetary radius. Recent observation of the exoplanet TrES-3 by \\cite{winn08} does not support such interpretation. These authors find upper limits on the planet's geometric albedo in the $i,z$, and $R$ bands as 0.30, 0.62 and 1.07 respectively. Thus they rule out the presence of highly reflective clouds in the atmosphere of TrES-3. It is worth mentioning that the geometric albedo and the degree of polarization are two distinct physical quantities. The geometric albedo is calculated from the ratio of the incident starlight to the sum of the isotropic and anisotropic components of the emergent radiation from the planet. On the other hand, the degree of polarization from an unresolved planet is the ratio of the emergent anisotropic radiation from the planet to the star light. While albedo is always non-zero, polarization could be zero if there is no anisotropy in the emergent planetary radiation. Therefore, the photometric study by \\cite{winn08} implies that the degree of polarization of TrES-3 should be very small although it can have high albedo.  However, the conclusion for TrES-3 may not be applicable to other planets, and in particular to HD 189733b. In fact, \\cite{pont} have reported an almost featureless transmission spectrum between 550 and 1050 nm with no indication of the expected alkaline absorption features which suggests the presence of a haze of sub-micron particles in the upper atmosphere of HD 189733b.  \\cite{burr08}  have found that a high star-to-planet flux ratio in the blue is possible due to Rayleigh scattering off $H_2$ and He. This although explains a high geometric albedo implied by the observed polarization, it cannot explain the peak amplitude of the polarization itself unless the scattering radius is much larger than that infered from transit observation.  Absorption in the stellar Lyman $\\alpha$ line observed during the transit of the extrasolar planet HD 209458b by \\cite{vidal03} revealed high velocity atomic hydrogen at great distances from the planet. This is interpreated by \\cite{vidal04} as hydrogen atoms escaping from an extended atmosphere of the planet that is possibly undergoing hydrodynamic blow-off. This interpretation however, fails to get theoretical support \\citep{hub07} and lead to controversy \\citep{ben07}. Recently, \\cite{holm08} have provided a viable alternative interpretation of the measured transit-associated Lyman alpha absorption as the interaction between the exosphere of HD 209458b and the stellar wind. These authors suggest a slow and hot stellar wind near the planet at the time of observation. This interpretation is consistent with the energetic neutral atoms around Solar System planets that is observed to form from charge exchange between solar wind protons and neutral hydrogen from the planetary exospheres. Under such situation and in the absence of any signature obtained by any other method, it would highly be premature to consider an extended exosphere of HD 189733b in order to explain the observed polarimetric data.  The re-binned polarization data shows negative value of Q at all phase angles. Because of the large error-bars in U, I choose the value of U corresponding to the maximum value of Q and find that $U/Q=0.63\\pm0.3$. Hence, for a small value of U in the scattering reference plane, the longitude of the ascending node $\\Omega$ can be estimated to be $16^o\\pm6$. Coincidentally, the observation of HD189733b over half an orbital period indicates \\citep{knut} the peak hemisphere-integrated brightness to occur $16^o\\pm6$ before opposition. Similar type of observation along with the polarimetry data for other planets may decide if there is any correlation. Taking the inclination angle $i=85^o.76$  as determined by transit method \\citep{winn07}, I obtain the best fit by setting $\\Omega=20^o$ and the modal grain diameter $d_0=0.8 \\mu m$. The theoretical model alongwith the observed data is presented in figure~1. Since the inclination angle of the planet is nearly $90^o$, polarization is zero at the transit $(\\Lambda=180^o)$ and at the secondary eclipse $(\\Lambda=0^o)$. These are the positions when the planet's night and day sides are turned towards the observer. The polarization peaks near $\\Lambda=90^o$ because polarization of light that is singly scattered is the largest for a scattering angle $\\pi-\\Lambda=90^o$ and in a sufficiently thin medium, single scattering is favoured over multiple scattering. In a non-circular orbit, the peak polarization shifts towards the longitude of pericentre $\\omega$ \\citep{sen06}. The observed data for Q allows an eccentricity as high as e=0.06 so that $e\\cos\\omega=0.001$ for $\\omega=89^o$ which is consistent with the time delay of the secondary eclipse \\citep{knut}. However, more stringent limit is needed for $e\\sin\\omega$ in order to constrain the eccentricity \\citep{winn07}.  Contrary to the claim \\citep{berd} that the orbital inclination angle greater than $90^o$ can be detected  through polarization, I do not find any change in the polarization profile if  the inclination is taken to be $90^o+i$.  At small values of x and the oblateness, the higher harmonic coefficients of the Stokes parameters are expected to be negligible as compared to the second one for Mie scattering. The ratio  of the second harmonic coefficients is equal to $(1+\\cos 2i)/2\\cos i$ and hence Fourier transformation of more accurately measured polarization data may provide the inclination angle of any extrasolar planet, transits as well as non-transits.  Unlike the case of Brown Dwarfs which needs nonspherical photosphere to make incomplete cancellation of polarization when integrated over the disk, a perfectly spherical planet can yield non-zero polarization because the illumination of the planetary surface by the star-light does not cover the entire disk except when the planetary phase angle is $0^o$ or $180^o$.  The rotation induced oblateness or tidal distrotion of the planetary disk introduces additional asymmetry which increases the amount of polarization. For a non-spherical planetary disk, the shape of the planet could also affect the time variation of the polarization. However, as estimated in section~\\ref{obl}, the rotation induced oblateness of the tidally locked planet HD 189733b is too small to yield any significant effect on the polarization profile as comapred to the case of a perfectly spherical geometry. In fact, the estimated oblateness of the planet increases the polarization by about twice its value calculated by considering spherical geometry. Therefore, the observed high peak amplitude of polarization cannot be achieved by considering only non-spherical photosphere. Consequently, the presence of dust in the photosphere of the planet becomes essential in order to explain the observed polarization.  Finally, it remains to be verified if the inclusion of forsterite dominated thin cloud can give rise to a flux ratio of about $10^{-4}$ in the blue band."
    },
    "0807/0807.4387_arXiv.txt": {
        "abstract": "{} {We determine carbon, oxygen, and their isotopic abundances  based on CO and OH spectra in 23 red giant stars, and identify possible origin of difficulty in abundance  analysis of cool luminous stars. } {We apply the line-by-line analysis based on the classical micro-turbulent model and  1 D model photospheres. } {We found purely empirically that there is a critical value of log\\,${W/\\nu} \\approx -4.75$ ($W$ is the equivalent width and $\\nu$ the wavenumber) above which the observed lines do not follow the classical line formation theory based on the micro-turbulent model and that the classical abundance analysis can be applied only to the weak lines of log\\,${W/\\nu} \\la -4.75$. The carbon, oxygen, and their isotopic abundances in 23 K - M giant stars obtained from  such weak lines are roughly consistent with the recent evolutionary models, although the $^{12}$C/$^{13}$C puzzle remains unsolved. On the other hand, it is already known that the very strong lines of log ${W/\\nu} > -4.4$ are contaminated by the contribution from extra molecular layers. The less strong but saturated lines of $ -4.75 < {\\rm log}\\,{W/\\nu} \\la -4.4 $ (intermediate-strength lines) cannot be understood at all with the  abundance and turbulent velocity based on the weak lines. From the behavior shown by these lines and considering other observations such as the detection of H$_2$O lines, not only in the late M giants but also in the early M and K giants, we found that the intermediate-strength lines are  essentially the same as the strong lines in that they also include contamination from the extra molecular layers.  Thus the behavior of the intermediate-strength lines, including those with LEP as high as 2\\,eV, appears to be nothing but a manifestation of the warm molecular envelope or MOLsphere. } {The infrared spectra of K - M giant stars are a hybrid of at least two components originating in the photosphere and MOLsphere. Only the weak lines, mostly  high excitation, are relatively free from the effect of MOLsphere and can be used to extract photospheric abundances. The strong lines and the intermediate-strength lines which actually dominate the observed infrared spectra, are badly disturbed by contamination from the MOLsphere. For this reason, they provide little information about the photosphere,  but  instead  can be new proves of the warm molecular envelope for which little is known yet. In the interpretation and analysis of the infrared spectra of K - M giant stars, it is essential to keep their hybrid nature in mind. } ",
        "introduction": "The basic principle  of stellar abundance determination was established already about half a century ago (e.g. Uns\\\"old 1955), but it is only recently that high quality spectroscopic data are made available for accurate abundance analyses. Actual abundance determinations have largely been done using the classical methods based on  grids of model photospheres (e.g. Gustafsson et al. 1975; Tsuji 1978; Kurucz 1979; Plez et al. 1992). In many applications, local thermodynamical equilibrium (LTE) was assumed to hold (e.g. Uns\\\"old 1955)  and favored for its simplicity in  abundance determinations. The next major step was to treat line formation without the restriction of the LTE assumption (e.g. Mihalas 1978). More recently, the classical model photospheres describing the structure of the photosphere only in the radial direction (1 D) are  being challenged by the new models considering the hydrodynamics of convection in 3 dimension (3 D model) (e.g. Asplund 2005). In the regime of the coolest stars such as red giant and supergiant stars, however, the situation may be a bit different and there may be another problem.  In cool giant stars representing the evolution from the red giant branch (RGB) to the asymptotic giant branch (AGB), CNO and their isotopic abundances experience drastic changes  and accurate abundance determination should be vital not only to testing the stellar evolutionary models but also to understanding the chemistry of interstellar matter which is enriched by the mass-loss from red giant stars. Despite their importance, few abundance analyses based on high resolution spectra have been done for cool giant stars, compared to other type of stars. It should be remembered that high quality infrared spectra needed for this purpose have been available, thanks to the Fourier transform spectroscopy (FTS) pioneered by P. \\& J. Connes (e.g. Connes 1970) and developed further in the 1970's (e.g. Maillard 1974; Hall et al. 1979; Ridgway \\& Brault 1984), even before high efficient detectors such as CCD innovated optical spectroscopy. Thus the problem here was not in the quality of the observed data but should be due to some inherent difficulties in the spectral analysis of very cool stars. So far, some efforts to determine the CNO and their isotopic abundances were made for oxygen-rich case (e.g. Harris \\& Lambert 1984; Smith \\& Lambert 1985, 1986, 1990; Harris et al. 1985;  Tsuji 1986, 1991; Aoki \\& Tsuji 1997) as well as for carbon-rich case (e.g. Lambert et al. 1984; Harris et  al.1987; Ohnaka \\& Tsuji 1996, 1999; Abia \\& Isern 1997; Ohnaka et al. 2000; Wahlin et al. 2006; Abia et al. 2008). Although many interesting results were shown by these works, we were not convinced entirely that the stellar abundance analysis could be done consistently for cool giant stars.  For example,  in our previous analysis of the CO first overtone bands in M giant stars, we found that most lines could be interpreted consistently with the classical model in terms of the micro-turbulent velocity and carbon abundance (Tsuji 1986, hereafter Paper I). The very strong lines of low excitation could not be fitted into this scheme but it was shown that these strong lines were contaminated by the contributions of outer molecular layers (Tsuji 1988, hereafter Paper II). We then analyzed the CO second overtone bands of the same sample of M giant stars and found that the resulting carbon abundances appear to be larger by about a factor of 2 compared with those based on the CO first overtone bands (Tsuji 1991, hereafter Paper III). The origin of these discrepant results remains unsolved yet. The difficulty in abundance analysis of very cool stars was also noticed by Smith \\& Lambert (1990), who found that the strong OH 1-0 and 2-0 bands gave unreasonably large oxygen abundances compared with those based on the weaker bands and that strong CO lines cannot be used for abundance analysis as well.  The carbon and oxygen abundances in other stars are by no means well established yet. Even for the sun, the classical values widely adopted (e.g. Anders \\& Grevesse 1989) had to be challenged by a new approach based on the 3 D models, and a reduction of carbon and oxygen abundances by more than 50\\% from the classical values was suggested (Allende Pietro et al. 2002). The classical values, however, were defended by a detailed analysis of CO lines with semi-empirical model solar photosphere by Ayres et al. (2006). These results remind us that the accurate modeling of the photospheric structure is a critical issue in abundance determinations. It is unknown if the difficulties in abundance analysis in M giant stars outlined above are due to a limitation of the classical 1 D model photosphere for cool giant stars or if they are related to other problems. For example, spectra of red giant stars revealed such an unexpected feature as H$_2$O lines, not only in  the early M giants and K5 giant $\\alpha$ Tau (Tsuji 2001) but also in the earlier K giant star $\\alpha$ Boo (Ryde et al. 2002). These features cannot be fitted into our present understanding of the photospheres of red giant stars, and suggest a new problem  in the photosphere and/or in the atmosphere extending beyond the photosphere. For  this reason, we analyze the nature of the difficulty in interpreting the spectra of cool giant stars in some detail in this paper. ",
        "conclusions": "We initiated this study with a hope to clarify the nature of the difficulty encountered in our previous analysis of CO (Papers I \\& III ), since we felt that something is still not well understood in abundance analysis of cool giant stars. At first, we had to recognize that the line spectra of cool giant stars are in fact highly complicated and may consist of at least three different groups of lines that may have different origins: The weak lines of log $W/\\nu \\la -4.75$, the intermediate-strength lines of $ -4.75 < {\\rm log} W/\\nu \\la -4.4$, and the strong lines of log $W/\\nu > -4.4$. Of these three groups, only the weak lines could be understood well on the basis of the classical micro-turbulent model of line formation. On the other hand, it has been known that the  strong lines could not be understood as being of the photospheric origin alone but should include an extra component originating in the outer molecular layers (Paper II).  The most difficult case was the intermediate-strength lines. However, we are finally led to conclude that the intermediate-strength lines  are essentially the same as the strong lines. This means that not only the low excitation strong lines but also medium strong lines with LEP as high as 2\\,eV (Fig.\\,12)  are disturbed by the contamination originating in the outer molecular envelope, which we refer to as MOLsphere for simplicity. This in turn means that most of the dominant molecular lines observed in the infrared spectra of K - M giant stars should be badly disturbed by the contamination from the extra component beyond the photosphere. At the same time, this finding provides further support for the presence of the MOLsphere, on the basis of the high resolution spectra of dozens of red giant stars. Thus, the difficulty of the abundance analysis we have encountered is nothing but a warning that the spectra of cool luminous stars cannot be analyzed by the conventional methods using the classical model photospheres. Instead, the influence of the extra constituent of the atmosphere, which we tentatively referred to as MOLsphere, should be quite essential not only in the strong low excitation lines but also in many weaker lines that dominate the infrared spectra.  Now we had to know that the largest fault in our previous analysis (Paper I) was to have assumed that the critical value discriminating the weak and the intermediate-strength lines is log $W/\\nu \\approx -4.5$ rather than -4.75. For this reason, many lines which are disturbed by the contamination from the warm molecular envelope were included in the abundance analysis of Paper I. It should certainly be difficult to attempt abundance analysis on such lines including the contamination of the non-photospheric origin by the use of a conventional model photosphere, and  the reason why abundance analysis of cool giant stars often encountered difficulties appeared to be quite simple, at least partly. This result also implies that it should be difficult in principle to apply low resolution infrared spectra to abundance analyses of cool luminous stars, and the hybrid nature of the infrared spectra reinforces the reason why high resolution spectroscopy should be indispensable in abundance analyses.  On the other hand, it is fortunate that the weak lines are relatively free from disturbance by extra component of non-photospheric origin and we are convinced that the abundance analysis of cool giant stars can be possible with these lines. With a limited number of the weak lines, we could determine the C, O, and their isotopic abundances with reasonable accuracy. The resulting CNO and their isotopic abundances are roughly consistent with the predictions of the first and second dredges-up during the red giant phase. We noticed, however, that there are some deviations from the evolutionary models especially in \\element[][12]{C}, \\element[][13]{C}, and \\element[][14]{N} abundances in some stars, and it seems that some additional episodic events may be needed to explain them. Since the CNO isotopes in the universe, especially \\element[][13]{C}, \\element[][14]{N}, and \\element[][17]{O},  are largely supplied by red giant stars, further refinements of both abundance analyses and evolutionary modelings should be pursued.  In conclusion, it turns out that high resolution infrared spectra include wealth of information not only on  photospheric abundances but also on the structure of the entire atmosphere extending beyond photosphere. For this reason, the infrared spectra must be of hybrid nature with contributions from multiple components. In the interpretation and analysis of the infrared spectra of cool luminous stars, it is essential to keep their hybrid nature in mind."
    },
    "0807/0807.3334_arXiv.txt": {
        "abstract": "\\small{ A sudden small change in the second derivative of the inflaton potential can result in a universal local feature in the spectrum of primordial perturbations generated during inflation. The exact solution describing this feature \\cite{minu} is characterized by a step in the spectral index modulated by characteristic oscillations and results in a large running of the spectral index localized over a few e-folds of scale. In this paper we confront this step-like feature with the 5 year WMAP results and demonstrate that it provides a better fit to this data than a featureless initial spectrum. If such a feature exists at all, then it should lie at sufficiently large scales $k_0 \\lesssim 0.003 {\\rm Mpc}^{-1}$ corresponding to $l \\lesssim 40$. The sign of the effect is shown to correspond to the negative running of $n_s$ localized near this scale. This feature could arise as a result of a `mini-waterfall'-type fast second order phase transition experienced by an auxiliary heavy field during inflation, in a model similar to hybrid inflation (though for a different choice of parameters). If this is the case, then the auxiliary field should be positively coupled to the inflaton. } ",
        "introduction": "\\label{sec:intro} The great precision of current cosmological data and the enormous volume of data expected in coming years leads to the hope that cosmological parameters will soon be determined to great accuracy. Since cosmological parameters are intimately linked to an underlying theoretical model, the increasing depth and sophistication of cosmological data sets, especially those associated with the Cosmic Microwave Background (CMB), will, in all likelihood, lead to a more refined and deeper understanding of such important issues as the form of the initial perturbation spectrum, the energy scale of inflation, etc.  Indeed, the cosmological world view which has emerged during the past two decades has been quite intriguing! New cosmological data has, while providing support to earlier constructs such as inflation, also paved the way for radically new propositions such as dark energy.     An important issue concerns the form of the primordial spectral index $n_s(k) \\equiv d\\ln P(k)/d\\ln k$. The large quantity of CMB data which has become available over the past few years indicates that the departure of the spectral index from exact scale invariance is likely to be small, $|n_s(k)-1|\\ll 1$, which is in good agreement with predictions of the simplest inflationary scenarios (see e.g. \\cite{LSV07} for a recent model-independent analysis of observational data assuming local analytic behaviour of the inflaton potential $V({\\phi})$ in the observable window).  An associated question concerns the value of the running of the spectral index $\\alpha(k)\\equiv d\\,n_s(k)/ d\\ln k$. Most inflationary models with smooth inflaton potentials do predict a non-vanishing value of $\\alpha(k)$, usually $|\\alpha(k)|\\sim |n_s(k)-1|^2\\ll 1$. However, recent WMAP$^1$ {\\footnotetext[1] {Wilkinson Microwave Anisotropy Probe (WMAP)}} results \\cite{wmap,wmap5,komatsu} may be suggesting a somewhat larger value $|\\alpha(k)|\\sim |n_s(k)-1|$. Taken together with other `anomalies' such as the `Archeops feature' at $l\\sim 40$ \\cite{archeops,wmap,wmap5}, these recent data may be providing a subtle hint that inflationary models are slightly more complex than the simplest single-field models suggested during the early 1980's.    Note that the only way for a large value of the running $\\alpha(k)$ to be accommodated within the inflationary paradigm is for it to be localized (\\ie have the form of a {\\em feature}), since, otherwise, the number of inflationary e-foldings turn out to be too small \\cite{EP06}. Generation of features in the primordial inflationary spectrum has been the subject of considerable discussion. A partial list of references may be found in \\cite{feature,S92,sarkar,G05,YY06}. Several model independent attempts to search for such features in the primordial spectrum directly from the WMAP data have led to some positive, though inconclusive, results \\cite{directpps}. In this paper we focus on the exact solution for the primordial scalar perturbation spectrum generated during inflation found in \\cite{minu}, in which the effective mass of the inflaton, $V''({\\phi})$, experiences a sudden small change. This model assumes a much smoother potential than those which had previously been considered, which dealt either with a sudden step in the potential \\cite{S92,sarkar, G05}, or in its first derivative \\cite{S92,G05}. As a result, the corresponding feature in the perturbation spectrum is the weakest (in particular, superimposed oscillations are very small). Still it is observable, and can even produce the dominant contribution to the running of the slope of the power spectrum over the observable range of scales. We shall show that this new universal (i.e. not dependent on details of the change) feature in the inflationary perturbation spectrum is in excellent agreement with recent WMAP5 results.   A microscopic model producing such feature in the inflaton potential may be chosen similar to that used in the hybrid inflationary scenario \\cite{hybrid}, but with different numerical values of parameters in order that a fast (quasi-equilibrium) second order phase transition experienced by another heavy scalar field should occur {\\em during} inflation, and not at its end. Also this second field should be sufficiently weakly coupled to the inflaton to produce a small change, $|\\Delta m| \\ll H$, in the inflaton mass where $H$ is the Hubble parameter, $H\\equiv \\dot a/a$, at the moment of the transition. Thus, such a transition may be called a `mini-waterfall' in analogy with the  `waterfall' transition used to end inflation in the hybrid inflationary scenario. The sign of this coupling may be arbitrary, in principle, but we shall show that the (more natural) positive coupling is favoured by the WMAP5 data since it corresponds to the negative running of the power spectrum. ",
        "conclusions": "In this paper we confront the WMAP 5-year data with an exact solution \\cite{minu} for the primordial power spectrum of perturbations generated when the inflaton effective potential $V(\\phi)$ has a sudden small change in its second derivative with respect to $\\phi$, i.e. in $m_{\\phi}^2$. The spectrum possesses a local universal feature having the form of a step in the primordial spectral index $n_s$, modulated by comparatively weak oscillations (by universal is meant that the form of the spectral feature does not depend upon the structure of the discontinuity in $V(\\phi)$). It results in a large (but local) running of $n_s$.    The simplest microscopic realization of such behaviour of $V(\\phi)$ is provided by a model having an auxiliary heavy scalar field which experiences a rapid second order phase transition (a mini-waterfall) during inflation in the observably accessible range of scales. Coupling this field to the inflaton leads to the desired type of local discontinuity in $V(\\phi)$. The model is similar to that used in the hybrid inflationary scenario, but in contrast to the latter, its parameters are chosen in such a way that: (i) the transition occurs during inflation and not at its end, (ii) the change in the inflaton mass is small compared to the Hubble parameter at this moment. That is why we call it a `mini-waterfall', in contrast to a `waterfall' which provides an end to inflation in the hybrid scenario. Should such a mini-waterfall be detected through the corresponding feature in the primordial power spectrum of scalar perturbations, it would provide direct experimental evidence for the naturalness of a similar (though larger) waterfall in the hybrid scenario (the latter not being directly observable since it lies at a very small comoving scale).    We find that the best $\\chi^2_{eff}$ for this model shows an improvement by $3.052$ over the best fit obtained assuming a featureless power law for the primordial spectrum. (This improvement comes at the cost of introducing two additional parameters.) It is shown that such a feature in the primordial spectrum, if exists at all, should lie on large scales $k_0 \\lesssim 0.003$ Mpc$^{-1}$. Anyway, this feature is not excluded by the present observational data. Better data expected from future CMB experiments will help to settle the question about its existence.    An interesting problem not considered in this paper is the amount of non-Gaussianity in the statistics of primordial perturbations. Here, strictly speaking, one needs to distinguish two different, though related cases: the single-field inflationary model with an effective inflaton potential having the studied type of local non-analytic behaviour, and the two-field model (\\ref{eq:hybrid}). These models produce the same results for the power spectrum in the leading approximation, but may become non-equivalent at the level of deviations from Gaussian behaviour. Since the second field in the latter model is in the fast-rolling regime, one may expect larger amount of non-Gaussianity for it. We hope to return to this question elsewhere."
    },
    "0807/0807.4208_arXiv.txt": {
        "abstract": "We present observations of  the Sunyaev-Zel'dovich effect (SZE) in the Bullet cluster (1E 0657--56) using the APEX-SZ instrument at 150$\\,$GHz with a resolution of $1\\arcmin$. The main results are maps of the SZE in this massive, merging galaxy cluster. The cluster is detected with $23\\,\\sigma$ significance within the central 1\\arcmin\\ radius of the source position. The SZE map has a broadly similar morphology to that in existing X-ray maps of this system, and we find no evidence for significant contamination of the SZE emission by radio or IR sources. In order to make simple quantitative comparisons with cluster gas models derived from X-ray observations, we fit our data to an isothermal elliptical $\\beta$ model, despite the inadequacy of such a model for this complex merging system. With an X-ray derived prior on the power-law index, $\\beta = 1.04^{+0.16}_{-0.10}$, we find a core radius $r_c =142 \\pm 18 \\arcsec$, an axial ratio of $0.889 \\pm 0.072$, and a central temperature decrement of $-771 \\pm 71\\,\\mu{\\rm K_{CMB}}$, including a $\\pm 5.5\\%$ flux calibration uncertainty. Combining the APEX-SZ map with a map of projected electron surface density from Chandra X-ray observations, we determine the mass-weighted temperature of the cluster gas to be $T_{mg}=10.8 \\pm 0.9\\,$keV, significantly lower than some previously reported X-ray spectroscopic temperatures. Under the assumption of an isothermal cluster gas distribution in hydrostatic equilibrium, we compute the gas mass fraction for prolate and oblate spheroidal geometries and find it to be consistent with previous results from X-ray and weak lensing observations. This work is the first result from the APEX-SZ experiment, and represents the first reported scientific result from observations with a large array of multiplexed superconducting transition-edge sensor bolometers. ",
        "introduction": "\\label{SEC:introduction}  Clusters of galaxies are a unique probe of the growth and dynamics of structure in the Universe. In particular, active mergers of sub-clusters provide a window to the processes by which massive clusters are assembled. In these systems, the galaxies and associated dark matter are essentially collisionless. In contrast, the ionized intracluster gas, typically at temperatures of $T \\sim 10^8$~K, is strongly interacting and experiences complex dynamics. In extreme cases, the normally associated dark matter and intracluster gas can be significantly separated.  The Bullet cluster (1E 0657--56) at $z=0.296$, is a massive cluster consisting of two sub-clusters in the process of merging. The smaller sub-cluster or ``bullet\" has passed through the larger main cluster. X-ray observations infer a bow shock velocity of $\\sim$4700 km/s~\\markcite{markevitch2006}({Markevitch} 2006), while simulations of the collision yield a substantially lower speed for the sub-cluster \\markcite{springel2007}({Springel} \\& {Farrar} 2007). This collision is perpendicular to the line of sight, providing an ideal system for studying interacting sub-clusters~\\markcite{clowe2006}({Clowe} {et~al.} 2006).  The mass surface density of the Bullet cluster has been measured using weak and strong gravitational lensing of light from background galaxies. There are significant angular offsets between the peaks of the X-ray surface brightness, which trace the baryonic gas through thermal bremsstrahlung emission, and the peaks of the lensing surface density, which are associated with the majority of the mass. The combined weak and strong lensing analyses of \\markcite{bradac2006}{Brada{\\v c}} {et~al.} (2006) show that the main cluster and sub-cluster are separated from their associated X-ray peaks at $10\\,\\sigma$ and $6\\,\\sigma$ significance respectively. This result has been recognized as providing direct evidence for the presence of collisionless dark matter in this system~\\markcite{clowe2006}({Clowe} {et~al.} 2006).  The Sunyaev-Zel'dovich effect (SZE) provides an independent probe of the intracluster gas. In the SZE, a small fraction ($\\sim1\\%$) of cosmic microwave background (CMB) photons undergo inverse Compton scattering from intracluster electrons~\\markcite{sunyaev1970,birkinshaw1999}({Sunyaev} \\& {Zel'dovich} 1970; {Birkinshaw} 1999). This process distorts the Planck blackbody spectrum of the CMB and produces a signal proportional to the gas pressure integrated along the line of sight. At $150\\,$GHz, the SZE produces a temperature decrement with respect to the unperturbed CMB intensity. Early detections of the SZE in the Bullet cluster include work by \\markcite{andreani1999}{Andreani} {et~al.} (1999) and \\markcite{gomez2003}{Gomez} {et~al.} (2004).  Unlike the X-ray surface brightness, the peak SZE surface brightness for a given cluster is independent of redshift. Therefore, the SZE has the potential to be an effective probe of intracluster gas out to the redshifts at which clusters are assembled. SZE measurements of galaxy clusters provide complementary constraints on cluster properties typically derived from X-ray measurements such as central electron density, core radius of the intracluster gas, cluster gas mass, and fraction of the total cluster mass in gas. Since the SZE and X-ray signals are proportional to the line-of-sight integral of the electron density and electron density squared respectively, SZE results will be less sensitive to clumping of the intracluster gas. For all comparisons between SZ and X-ray data, we assume a $\\Lambda$CDM cosmology, with $h = 0.7$, $\\Omega_{\\rm{m}} = 0.27$, and $\\Omega_{\\Lambda} = 0.73$.  In this paper, we present a 1\\arcmin\\ resolution SZE image of the Bullet cluster at $150\\,$GHz made with the APEX-SZ instrument. It is the first reported scientific result from observations with a large array of multiplexed superconducting transition-edge sensor bolometers. In \\S\\,\\ref{SEC:observations}, we discuss the instrument and observations. Calibration is discussed in \\S\\,\\ref{SEC:calibration}. In \\S\\,\\ref{SEC:datareduction}, we describe the data reduction procedure, and in \\S\\,\\ref{SEC:results}, we present the results of fits to the SZE surface brightness with cluster models, including mass-weighted electron temperature and gas mass fraction calculations. We summarize the conclusions and discuss future work in \\S\\,\\ref{SEC:conclusions}. ",
        "conclusions": "\\label{SEC:conclusions} Measurements of the SZE provide a robust and independent probe of the intracluster gas properties in galaxy clusters. The APEX-SZ $150\\,$GHz observations detect the Bullet system with $23\\,\\sigma$ significance within the central 1\\arcmin\\ radius of the SZE centroid position. We do not expect to see a resolved signal from the Bullet sub-cluster in the $150\\,$GHz 85\\arcsec\\ FWHM resolution SZE maps, and no obvious feature, such as a secondary peak, is present. We expect no significant contamination of the observed SZE decrement due to radio sources, and there is no evidence for significant contamination by a known bright lensed dusty galaxy.  We process an elliptical $\\beta$ model through the observation transfer function and fit it to the measured temperature decrement map. We also measure the cluster mass-weighted electron temperature and gas mass fraction with the SZE data. Combining the APEX-SZ map with a map of projected electron surface density from Chandra X-ray observations, we determine the mass-weighted temperature of the cluster gas to be $T_{mg}=10.8 \\pm 0.9\\,$keV. This value is consistent with the lowest X-ray spectroscopic temperatures reported for this cluster and should be less sensitive to the details of the cluster thermal structure. The derived baryon fraction is also found to be in reasonable agreement with previous X-ray and weak lensing determinations.  Throughout this work, we make an assumption of isothermal cluster gas. Clearly, incorporating thermal structure, measured by X-ray observations, in the analysis of the SZE data would improve the determination of the gas distribution and gas mass fraction. Ultimately, a more sophisticated analysis could be implemented that combines X-ray, SZE, and weak lensing data and relaxes assumptions of hydrostatic equilibrium between the gas and dark matter components of the cluster. This is particularly important for a detailed understanding of actively merging systems such as the Bullet cluster."
    },
    "0807/0807.1331_arXiv.txt": {
        "abstract": "We study the scaling relations between global properties of dwarf galaxies in the Local Group.  In addition to quantifying the correlations between pairs of variables, we explore the ``shape'' of the distribution of galaxies in log parameter space using standardised Principal Component Analysis (PCA), the analysis is performed first in the 3-D structural parameter space of stellar mass $\\Ms$, internal velocity $V$ and characteristic radius $\\Rs$ (or surface brightness $\\mus$).  It is then extended to a 4-D space that includes a stellar-population parameter such as metallicity $Z$ or star formation rate $\\sfr$.  We find that the Local-Group dwarfs basically define a one-parameter ``Fundamental Line''(FL), primarily driven by stellar mass, $\\Ms$.  A more detailed inspection reveals differences between the star-formation properties of dwarf irregulars (dI's) and dwarf ellipticals (dE's), beyond the tendency of the latter to be more massive.  In particular, the metallicities of dI's are typically lower by a factor of 3 at a given $\\Ms$ and they grow faster with increasing $\\Ms$, showing a tighter FL in the 4-D space for the dE's.  The structural scaling relations of dI's resemble those of the more massive spirals, but the dI's have lower star-formation rates for a given $\\Ms$ which also grow faster with increasing $\\Ms$.  On the other hand, the FL of the dE's departs from the fundamental plane of bigger ellipticals.  While the one-parameter nature of the FL and the associated slopes of the scaling relations are consistent with the general predictions of supernova feedback \\citep{dek03}, the differences between the FL's of the dE's and the dI's remain a challenge and should serve as a guide for the secondary physical processes responsible for these two types. ",
        "introduction": "\\label{intro}  \\subsection{Scaling Relations, Fundamental Distributions and Galaxy Formation in Giant Galaxies}  In the standard $\\Lambda$CDM picture for galaxy formation, galaxies form from the gas that cools within a dark matter halo into a disc of stars whose specific angular momentum is conserved \\citep{fal80}.  Since the gas initially shares the spatial distribution and dynamical properties of the dark matter halo, the structural properties of the resulting gaseous disc, specifically the mass of the disc of stars $\\Ms$, the size of the disc $\\Rs$ and its rotation velocity $\\Vrot$, are expected to correlate strongly with those of the halo.  The halo itself is expected to be characterised by the virial theorem, \\ie $M_{\\rm vir} \\prop V_{\\rm vir}^3 \\prop R_{\\rm vir}^3$.  The actual scalings of the stellar disc will be affected by the physics of the baryon gas and its interactions with the dark matter halo \\citep{dut07}.  Thus the observed scaling relations between these structural parameters in galaxies help to shed light on the physical processes governing their formation.  Several recent studies of scaling relations of galaxies have contributed to significant improvements in our appreciation of galaxy formation physics. For example, the scaling relation between $\\Ms$ and the stellar surface density $\\mus$ (related to $\\Rs$) for 123,000 galaxies of all types in the Sloan Digital Sky Survey (SDSS) shows a transition at $\\sim3 \\times 10^{10 } \\Msun$ \\citep{kau03a}. Above the transition, $\\mus$ only weakly depends on $\\Ms$ while below it, $\\mus$ scales as $\\Ms^{0.6}$, consistent with simple theoretical predictions of supernova feedback (SNF) \\citep{dek86,dek03}.  The transition itself can also be explained by considering the energy constraints of gas ejection by supernovae \\citep{dek86}.  The interpretation of the scaling relations between $\\Vrot$ and luminosity $L$, and $\\Rs$ and $L$ for $\\sim1300$ local spiral galaxies by \\cite{cou07} also suggests that the interaction of the disc and halo during galaxy formation is a fundamental constraint to $\\Lambda$CDM models \\citep{dut07}.  The study of scaling relations for elliptical galaxies has revealed the existence of a \"fundamental plane'' (FP) \\citep{djo87,dre87,ber03}. These galaxies lie on a plane in the parameter space spanned by the logarithms of $L$, $R$, and the velocity dispersion $\\sigma$. The projections of this plane onto the $\\sigma-L$ and $R-L$ planes are known as the Faber-Jackson \\citep{fab76} and the Kormendy \\citep{kor77} relations.  Recently, \\cite{zar06,zar06b,zar07} proposed that the slight tilt of the FP can be straightened out essentially by adding mass-to-light ratio as an extra dimension to the parameter space. \\cite{dek06} used hydrodynamical simulations to show how the properties of the fundamental plane of elliptical galaxies, including its tilt with respect to virial theorem predictions, can be explained by considering the dissipative properties of the major mergers that form them.  Besides the structural parameters ($\\Ms$, $\\mus$, $\\Rs$, $V$), the physics of the gaseous systems that later become galaxies will also affect the global metallicity $Z$, star formation history, and consequently the current star formation rate $\\sfr$ (SFR) of these galaxies.  Since $Z$ and $\\sfr$ are related to star formation we call these the ``SF'' quantities for brevity.  Studies that involve the SF parameters with stellar mass have also yielded insight into the gas and stellar population physics of galaxies.  For example, \\citet{tre04} show from SDSS measurements of galaxies with $10^{8.5} \\Msun < \\Ms < 10^{10.5} \\Msun$ that $Z$ decreases systematically with decreasing $\\Ms$, while higher mass systems only show a weak dependence on $\\Ms$.   These scalings are consistent with SNF model predictions \\citep{dek03,tas03}. As another example, \\cite{bri04} show trends of current star formation rate $\\sfr$ with $\\Ms$, and \\cite{kau04} show environmental dependence of this relation.  This poorly explored scaling relation may reveal insight into the effect of the disc self-gravity on gas processes that combine to form stars.  \\subsection{Dwarf Galaxies}  These investigations of the scaling relations and the FP have led to major advancements in our understanding of galaxy formation with respect to giant galaxies, and this motivates a similar investigation of dwarf galaxies.  Dwarf galaxies are the central players in one of the most puzzling problems in the $\\Lambda$CDM picture of galaxy formation, namely, the ``missing dwarf problem.''  This relates to the apparent discrepancy between the predicted distribution of halo masses at the low-mass end and the relatively few dwarf galaxies actually observed - see for example (\\citealp{kly99,moo99,sto02}).  Similarly, the luminosity function of galaxies below the Schechter's characteristic $L_*$ is observed to be flatter than predicted in $\\Lambda$CDM cosmology, implying that the ratio of stellar mass to total mass decreases with decreasing $M_{\\rm tot}$.  \\citet{dek03} used this systematic variation in $\\Ms/M_{\\rm tot}$ to predict the scaling relations of $\\mus \\prop \\Ms/R^2$ and $V$ with $\\Ms$ for dwarf galaxies in the SNF scenario.  (See also \\citealp{mal02} who show how SNF helps solve the angular momentum problem, and \\citealp{dekbir06} who detail the role of SNF in small galaxies).  A dependence of the ratio of $\\Ms/M$ on $M$ for low-mass galaxies, and thus a dependence of the overall star formation efficiency on $M$, should also have consequences on the scaling relations of the SF quantities with $\\Ms$.  The theoretical developments concerning the scaling relations between the properties of dwarf galaxies motivate an investigation of the fundamental distribution of the dwarf galaxies in parameter space. For example, \\cite{ben92} plot the galaxies in what they call ``$\\kappa$''-space. \\citet{pra02} (hereafter PB02) find a  ``fundamental line'' (FL) of dwarf galaxies in the space of $M_{\\rm tot}/L$, metallicity [Fe/H], and the central surface brightness $\\Sigma_V$, which they explain with a simple chemical enrichment model.  PB02 find that this FL is independent of Hubble type \\citep[see also][]{sim06}.   Similarly, \\cite{vad05b} find a fundamental plane of dI's in the parameter space of absolute $K$ magnitude $M_K$, central surface brightness $\\Sigma_K$ and \\HI line width $W$. The same authors find that blue compact dwarfs (BCD's) lie on the same fundamental plane as dI's \\citep{vad06}. \\cite{car92} and \\cite{zar06b} find that the dE's lie in the same ``fundamental manifold'', the latter putting them on the same manifold as larger elliptical galaxies.  However these fundamental distributions lie in parameter spaces whose axes are various combinations of structural and SF quantities.  In order to better understand the underlying physics behind these fundamental distributions, a disentangling of these axes into more ``fundamental'', or linearly independent quantities is in order.  The scaling relations between the structural and SF quantities in the dwarf galaxies were first presented as luminosity-dependent relations; \\eg $V$-$L$: \\citealp{fab76} and \\citealp{tul77}; $\\Sigma-L$: \\citealp{ben92}; [Fe/H]-$L$, [O/H]-$L$: \\citealp{zar94} and \\citealp{ric98}. However, for the purpose of constraining and comparing with physical models, the use of physical quantities is preferable.  We wish to improve on these investigations.  Our goals are two-fold: (i) present the scaling relations of dwarf galaxies in physical quantities; and (ii), find a fundamental distribution in a parameter space spanned by linearly independent physical quantities that can constrain physical models. We will show that in such a parameter space the dwarf galaxies lie in a linear distribution.  Such a distribution, or FL, would be described by a one-parameter family of equations which narrowly constrains the possible physical scenarios of their origin.  We will show this parameter space to include both structural and SF quantities.  The dwarf galaxies of the Local Group (LG), whose stellar masses range from as low as $10^{5.5} \\Msun$ to the SDSS transition scale of $10^{10.5} \\Msun$, provide an excellent opportunity to achieve our goals. In this paper, we compile data for LG dwarf galaxies and present scaling relations for the structural and SF properties of the dwarf galaxies in physical units.  We use a Principal Component Analysis (PCA) to search for  linear distributions in structural and SF space and quantify their tightness.  We identify the space, involving both structural and SF quantities, in which the LG dwarf galaxies form a fundamental line.  The outline of this paper is as follows: the data base is described in \\sec{data}, followed in \\sec{stellarmass} by our derivation of stellar mass and other global properties. In \\sec{tools} we describe the tools of our analysis, namely the fitting prescription and principal component analysis. In \\sec{stranalysis}, we present the scaling relations of LG dwarf galaxies in structural parameter space. In \\sec{speanalysis}, we augment the scaling relations with SF parameters and mass to derive the fundamental linear distributions of the dwarf galaxies in the parameters space of structural + SF parameters. Our results are discussed and summarised in \\sec{conc}. ",
        "conclusions": "Our analysis of structural and spectroscopic data for LG galaxies has provided important constraints for modelling the physical processes that govern the formation of dwarf galaxies.  In particular, successful models will need to reproduce the fundamental line in structural space +$Z$ (\\ie SNF space) obeyed by all dwarf galaxies. One such model is the simple supernova feedback prescription of \\cite{dek03} which quite accurately predicts the slopes of the scaling relations using $\\Ms$ as the one primary parameter determining the distribution of the data.  What supernova feedback does not deal with is the zero-points of the scaling relations, particularly the different zero-points between early and late-type dwarf galaxies. Future successful models must include secondary physical processes that distinguish between early and late-type dwarf galaxies, producing the separate higher-dimensional fundamental lines that they are observed to follow in SNF and $\\Ms,\\mus,V,Z,\\sfr$ spaces. The graphical display of PCA output as we have done in this paper can also help in visualising and understanding scatter in model dwarf galaxies.  Several unexplained challenges include: \\begin{itemize} \\item the different normalisations of the metallicity and surface density scaling relations for different Hubble type; \\item the presence of recent star formation episodes in late type galaxies, and hence the secondary fundamental line that dI's follow (\\ie in the parameter space of $\\Ms,\\mus,V,Z,\\sfr$); and \\item the dependence of the dwarf galaxy SFR scaling relation with environment, particularly the steepening of the relation in lower density environments, and the order of magnitude drop in SFR in lower density environments. \\end{itemize}"
    },
    "0807/0807.0930_arXiv.txt": {
        "abstract": "We investigate the physical processes occuring in the multiphase gas of a damped Ly$\\alpha$ system (DLA). We base our analysis on a high quality Keck HIRES spectrum  of the QSO J1211+0422 in which a DLA is detected at $z\\simeq 2.377$. There is little contamination of the high-ion (\\ovi, \\nv, \\civ, \\siiv) absorption, allowing us to  explore the properties of the highly ionized gas and its connection to other gas-phases.  The metallicity ($[{\\rm Z/H}] = -1.41 \\pm 0.08$), \\hi\\ column density ($\\log N($\\hi$)= 20.80 \\pm 0.10$), full-width velocity ($\\Delta v_{\\rm neut} \\simeq 70$ \\km) and relative abundances ($[{\\rm Si/Fe}] = +0.23 \\pm 0.05$ and $[{\\rm N/Si}] = -0.88 \\pm 0.07$)  of this DLA are not unusual. However, we derive the lowest \\cii*\\ cooling rate in a DLA, $l_c < 10^{-27.8}$ erg\\,s$^{-1}$ per H atom (3$\\sigma$). Using this stringent limit, we show that the neutral gas (confined at $|v| < +39$ \\km) must be warm and the star formation rate is $<7.1\\times 10^{-3}$ M$_\\odot$\\,yr$^{-1}$\\,kpc$^{-2}$. Surprisingly, the gas shows strong, complex absorption profiles from highly ionized gas whose kinematics appear connected to each other and the low ions. The total amount of highly and weakly ionized gas is very large with $N_{\\rm tot}($\\hii$)/N_{\\rm tot}($\\hi$) \\ga 1.5$. At $|v| \\ga +39$ \\km, the gas is fully and highly ionized (H$^+/$H\\,$\\sim 1$, $N($\\civ$)\\gg N($\\cii$)$, $N($\\siiv$)\\gg N($\\siii$)$). Based on ionization models, \\ovi\\ and \\nv\\ are generally difficult to produce by hard photons, while \\siiv\\ and \\civ\\ can be photoionized to a large extent. There is, however, no evidence of \\ovi-bearing gas at $T\\sim 10^6$~K associated with this DLA. In contrast, there is some evidence for narrow \\ovi, \\nv, and \\civ\\ components (unexplained by photoionization), implying too low temperatures ($T < 10^5$~K) for simple collisional ionization models to produce their observed column densities. Stellar feedback is a possible source for producing the high ions, but we cannot rule out accretion of non-pristine material onto the protogalaxy. ",
        "introduction": "Damped Ly$\\alpha$ absorbers (DLAs) are defined as absorbers with $\\log N($\\hi$) \\ge 20.3$  \\citep[][and references therein]{wolfe05}. As such these absorbers are often treated as dominantly neutral entities. This characteristic separates the DLAs from the other intervening absorbers seen in the QSO sightlines, the so-called Ly$\\alpha$ forest ($\\log N($\\hi$) \\le 17$) and Lyman limit systems  ($17 < \\log N($\\hi$) < 20.3$ ) where the neutral gas is a minor or non-existent phase. The presence of neutral, cold, and molecular gas is crucial to link the DLAs to star-forming galaxies  \\citep{wolfe03,wolfe05,howk05,ledoux03,noterdaeme08}, although the exact nature of the DLAs and their relation to present-day galaxies is still under debate and study \\citep{wolfe05,wolfe06,rauch08,meiksin08}.  Surveys of DLAs have also shown systematic absorption from weakly (e.g., \\aliii) and highly (\\siiv, \\civ, \\nv, \\ovi) ionized species in DLAs \\citep{lu96,wolfe00a,fox07a,fox07}, demonstrating a complex interstellar structure as observed in present-day galaxies. The detection of ionized gas, and more specifically highly ionized gas, is important because while high ions can be produced by hard photoionization, their ionization potentials are also high enough to be directly associated with galactic feedback following supernovae and stellar winds from massive stars \\citep[e.g.,][]{kawata07}.  Stellar feedback is important since it may produce galactic outflows that are the dominant mechanisms for enriching the intergalactic medium (IGM) with metals, as other dynamical processes, including ram-pressure stripping or tidal interactions between galaxies, are rather inefficient at polluting the IGM on large scales \\citep[e.g.,][]{aguirre01}. Energetic feedback from stellar winds and supernovae not only pollutes the IGM but is also believed to serve as a regulator of star formation, even cutting it off altogether by super heating or removing (almost) all the gas from a galaxy. Hence stellar feedback is an essential ingredient for cosmological models of the formation and evolution of galaxies \\citep[e.g.,][]{mashchenko08,dalla08}. For example, \\citet{white91} showed that without feedback, star formation is far too efficient compared to the observations. However, feedback is extremely difficult to model because it requires enough resolution to resolve individual supernova  and the ability to model the complex hydrodynamic processes in a highly multiphase interstellar gas \\citep[e.g.][]{dalla08}. Characterizing feedback signatures via observations of the low and high $z$ Universe is therefore critical.  At all redshifts, evidence for outflows resulting from feedback processes has been discovered.  High-$z$ Lyman break galaxies (LBGs) show signs of outflows \\citep{pettini01,shapley03}, analogous to the complex, multiphase outflows seen in low- and intermediate-$z$ starburst galaxies \\citep{heckman00,martin06,tremonti07,weiner08}. At low redshift, winds have been found in even some modestly star forming galaxies such as the Large Magellanic Cloud \\citep{lehner07} and the Milky Way \\citep{zech07,yao07,bland03}.  At $z\\ga 2$, two important DLA surveys undertaken by \\citet{fox07a,fox07} have revealed the statistical properties of the high ions associated with DLAs. They presented circumstantial evidence that \\civ\\ and \\ovi\\ absorption at high velocities relative to the DLAs is associated with protogalactic outflows. However, there is not yet a detailed analysis of the high ions (\\ovi, \\nv, \\civ, and \\siiv) in DLAs in order to test their origin(s) \\citep[although for detailed analysis of a DLA with \\civ\\ and \\siiv\\ absorption, see][]{prochaska02}. Part of the difficulty is that, while \\civ\\ and \\siiv\\ have generally little contamination, \\nv\\ absorption is weak and \\ovi\\ is often mixed up in the Ly$\\alpha$ forest. Among the 9 QSOs with intervening DLAs presented by \\citet{fox07}, there is not a single intervening DLA where these four ions and both transitions of each doublet are simultaneously detected, complicating the physical interpretation of their profiles. In particular, it is not clear if the apparent differences in breadth between the \\ovi\\ and \\civ\\  component-profiles are because \\ovi\\ probes truly hot gas or the \\ovi\\ profiles are significantly blended with intervening Ly$\\alpha$ absorbers.  In this work, we analyze a Keck/HIRES spectrum of the QSO J1211+0422 (SDSS\\,J121117.59+042222.3, $z_{\\rm em} = 2.5412$) in order to study the DLA  at  $z\\simeq 2.377$ observed along this sightline. This DLA, not surprisingly, shows high-ion absorption. Remarkably, however, the profiles of the high ions, as well as other tracers of neutral and weakly ionized gas, have little contamination from other intervening absorbers, allowing us to pursue detailed modeling of their profiles. Notably, both lines of the \\nv\\ doublet, while weak, are well detected. The \\ovi\\ doublet is partially tainted by Ly$\\alpha$ forest, but there is enough information in each transition to interpret the \\ovi. With our detailed modeling, the main issues we address in this work are the connection  of the highly ionized with the neutral and weakly ionized gas, the origin(s) of the highly ionized gas in DLAs at  $z\\sim 2.4$, and a search for possible feedback signatures.  The outline of this paper is as follows: in \\S\\ref{sec-obs} we describe the observation, give an overview of the velocity structure in the DLA, present the analysis techniques and results (kinematics, column densities, and $b$-values of the high ions). In \\S\\ref{sec-prop} we present and discuss the results for the low-velocity gas ($|v| < +39$ \\km), while in \\S\\ref{sec-ion} the higher-velocity ($|v|\\ge +39$ \\km), fully ionized gas is presented. In \\S\\ref{sec-disc} we discuss some of the properties of this DLA in the context of other studies. Our main results are summarized in \\S\\ref{sec-sum}. ",
        "conclusions": "\\label{sec-disc} Before summarizing our main findings, we argue that this DLA is typical in many regards. From a standpoint of the metallicity ($[{\\rm Z/H}] = -1.41 \\pm 0.08$), the main component of the DLA at $z=2.377$ toward J1211+0422 is actually quite ordinary \\citep[see, e.g.,][for a summary of the metallicity evolution of DLAs]{prochaska03}. The full velocity width of the neutral absorption of about 70 \\km\\ is also not far from the median value of 90 \\km\\ derived in a large sample of DLAs \\citep{prochaska97}. The $N($\\hi$)$ value lies near the median of the DLA \\hi\\ column density distribution. The $\\alpha$-element enhancement relative to Fe-peak elements and N is also not out of the ordinary \\citep{lu96,prochaska02,petitjean08}. But what may really distinguish this DLA from others is the absence of cold neutral gas signature. However, $32/76 = 42$\\% of the DLAs studied by \\citet{wolfe08} have no detectable \\cii*\\ absorption, possibly indicating the neutral gas is warm. None of the upper limits derived on $l_c$ are as stringent as in the present DLA and improving the signal in the spectra of these DLAs would help to understand if the population of DLAs with WNM is important.  We also saw that the high, intermediate, and low ions in this DLA seem to follow each other quite well kinematically. \\citet{wolfe00a} found that \\civ\\ and low ions rarely follow the same kinematic distribution, \\siiv\\ and \\civ\\ appear strongly correlated, and in many cases the \\aliii\\ profiles are better correlated with those of the low ions than with those of \\civ\\ and \\siiv\\ \\citep[see also][]{lu96}. However, even though the components of the high-ion profiles are generally disjoint from the low-ion profiles, \\citet{wolfe00a} noted that the high-ion and low-ion profiles are roughly centered on similar velocities and that the high-ion profiles have a full-width systematically similar to or larger than that low-ion profiles, properties similar to the present DLA \\citep[see also][]{fox07a,fox07}.  The ionized  DLA described by \\citet{prochaska02a}  shows some correlation between the velocity centroids of the low, intermediate, and high ions. The absence of clear kinematic alignment between the high, intermediate, and low-ion profiles in every DLA may, however, not be entirely surprising. In our own Galaxy, there is a significant variety in the type of associations between the various phases of the interstellar gas, which are generally attributed to sightline projection effects and the physical processes that give rise to the high-ions.  We also found that the amount of ionized gas associated to this DLA is very large. To gauge the importance of the highly ionized gas relative to the neutral gas we can estimate the amount of hydrogen in the highly ionized phase from $N($\\hii$) = N($\\nv$)/(f_{\\rm N\\,V} \\, ({\\rm N/H})_\\odot 10^{[{\\rm N/H}]})$, where $f_{\\rm N\\,V} = N($\\nv$)/N({\\rm N}) \\simeq 0.2$ is the ionization fraction of N at $T \\simeq 2 \\times 10^5$ K when using the NECI calculations from \\citet{gnat07}. At $T$ greater or lower than $2 \\times 10^5$ K, $f_{\\rm N\\,V}< 0.2$. We assume that the N abundance is the same in the neutral and ionized phase,  so that $[{\\rm N/H}] = [$\\nni/\\hi$] \\simeq -2.3$. For the LVC, $\\log N($\\hi$) = 20.8$, $\\log N($\\hii$) \\sim 20.5$ (highly ionized gas) and  $\\log N($\\hii$) \\sim 20.1$ (photoionized gas). For the NHVC, $\\log N($\\hii$) \\sim 20.5$ (highly ionized gas) and  $\\log N($\\hii$) \\sim 20.2$ (photoionized gas). Using the lower limits on $N($\\ovi$)$ (Tables~\\ref{t4a} and \\ref{t5}), we find consistent \\hii\\ amounts for the highly ionized gas with $\\log N($\\hii$) > 20.3$ in the LVC and NHVC (in that case  $[{\\rm O/H}] =  -1.4$ and $f_{\\rm O\\,VI}  \\la 0.2$). The $+39$ \\km\\ component and PHVC are negligible with respect to hydrogen amounts in the LVC and NHVC. Summing the various components, $N_{\\rm tot}($\\hii$)/N($\\hi$) \\ga 1.5$, where $N_{\\rm tot}($\\hii$)$ includes contribution from both the photoionized gas and highly ionized gas. As we discuss above there is some uncertainty in the processes that give rise to \\ovi\\ and \\nv. Nevertheless the ionization fraction of these species is unlikely to be much more than 0.2 based on CIE/NECI and photoionization models. The metallicity correction in the high-velocity components is uncertain because no direct metallicity estimates can be done. If the metallicity is not similar to the DLA itself, it is more likely to decrease in these components, although if there is enough pollution from other nearby protogalaxies an increase may be possible. The limit  $N_{\\rm tot}($\\hii$)/N($\\hi$) \\ga 1.5$ is larger than those typically derived by \\citet{fox07} for DLAs with similar total \\hi\\ column density, but is comparable to ionized-to-neutral ratios derived in DLAs with lower $N($\\hi$)$ and sub-DLAs \\citep{fox07,fox07c}. The  use of \\nv\\  may have allowed us to simply derive a more stringent lower limit.  Finally, we observe that \\citet{fox07} argue that the highly ionized gas may help to solve (in part) the missing metal ``problem\" at high redshift \\citep{pettini99} since if the \\ovi-bearing gas has $T\\sim 10^6$ K, then  $f_{\\rm O\\,VI} \\sim 0.003 $ (according to CIE models), which would imply  that the \\ovi\\ phase has enough metals to solve the issue. In our analysis, several components have $T \\la 3\\times 10^5$ K, consistent with $f_{\\rm O\\,VI} \\la 0.2$. However, there is also evidence for \\ovi-bearing gas with $T<9\\times 10^4$ K, which would also imply $f_{\\rm O\\,VI} \\ll 0.2$ if CIE/NECI applies. But CIE/NECI (and photoionization) provides a poor description of this gas. To quantify accurately the importance of the highly ionized phase in DLAs requires first a better understanding of the high-ion absorption in a large and (relatively) uncontaminated sample (with the purpose of obtaining a statistical sample where the distribution of $b$ in the {\\em individual}\\ components could be studied), the ionization fractions of the high ions (requiring models with more realistic physics), and metallicities in the ionized plasma (which is likely to rely on similar assumptions made in the present work)."
    },
    "0807/0807.0101_arXiv.txt": {
        "abstract": "We make a search for H$\\alpha$ emitting galaxies at $z \\approx 0.24$ in the Subaru Deep Field (SDF) using the archival data set obtained with the Subaru Telescope. We carefully select H$\\alpha$ emitters in the narrowband filter $NB816$, using $B$, $V$, $R_{\\rm c}$, $i^\\prime$, and $z^\\prime$ broad-band colors. We obtain a sample of 258 emitting galaxies with observed equivalent widths of (H$\\alpha$+[N{\\sc ii}]6548,6584) greater than 12 \\AA. We also analyze a sample of H$\\alpha$ emitters taken from the Sloan Digital Sky Survey (SDSS) to constrain the luminous end of H$\\alpha$ luminosity function. Using the same selection criteria as for the SDF, and after excluding AGNs, we obtain 317 H$\\alpha$ emitting star-forming galaxies. Combining these two samples of H$\\alpha$ emitters found in both SDF and SDSS, we derive a H$\\alpha$ luminosity function with best-fit Schechter function parameters of $\\alpha = -1.31^{+0.17}_{-0.17}$, $\\log \\phi^{*} = -2.46^{+0.34}_{-0.40} {\\rm Mpc^{-3}}$, $\\log L^{*} = 41.99^{+0.08}_{-0.07} {\\rm ergs \\; s^{-1}}$. An extinction-corrected H$\\alpha$ luminosity density is $4.45^{+2.96}_{-1.75} \\times 10^{39}\\;{\\rm ergs \\; s^{-1} \\; Mpc^{-3}}$. Using the Kennicutt relation between the H$\\alpha$ luminosity and star formation rate, the star formation rate density in the survey volume is estimated as $0.035^{+0.024}_{-0.014} ~ M_{\\odot} \\; {\\rm yr^{-1} \\; Mpc^{-3}}$. The angular two-point correlation function of H$\\alpha$ emitters over 875 $\\rm arcmin^2$ at $z = 0.24$ is well fitted by a power-law form with $w(\\theta) = 0.047^{+0.017}_{-0.013} \\theta^{-0.66 \\pm 0.08}$, corresponding to the correlation function of $\\xi(r) = (r/{\\rm 2.6^{+1.0}_{-0.8}Mpc})^{-1.66\\pm0.08}$. The small correlation length of H$\\alpha$ emitters may imply the weak clustering of active star-forming galaxies. ",
        "introduction": "Since the cosmic star formation history in galaxies from high-redshift to the present day is one of important measurements for the understanding of both formation and evolution of galaxies, a number of observational studies have been made for this decade (e.g., Madau et al. 1996; Madau, Pozzetti, \\& Dickinson 1998; Pettini et al 1998; Steidel et al. 1999; Barger et al. 2000; Giavalisco et al. 2004; Dickinson et al. 2004; Taniguchi et al. 2005; Bouwens \\& Illingworth 2006; see also Hopkins 2004; Hopkins \\& Beacom 2006). Although much interest has been addressed to the star formation in higher-redshift universe, large programs of spectroscopic surveys have provided important observations of galaxies located in lower-redshift universe; e.g.,  Canada-France Redshift Survey (CFRS: Lilly et al. 1995, 1996; Tresse \\& Maddox 1998; Tresse et al. 2002), the Sloan Digital Sky Survey (SDSS: York et al. 2000; Salim et al. 2007), the VIMOS VLT Deep Survey (VVDS: Le Fevre et al. 2005; Tresse et al. 2007). All these observations have provided us a rough trend of the cosmic star formation history from high-$z$ to the present day: the star-formation rate density (SFRD, $\\rho_{\\rm SFR}$) rises from $z \\simeq 0$ to $z \\simeq 1$, is nearly constant between $z \\sim 2$ and $z \\sim 5$ and decline at $z > 5$. Recent observations by the {\\it Galaxy Evolution Explorer} (GALEX) and the {\\it Spitzer Space Telescope} have confirmed that $\\rho_{\\rm SFR}$ increases from $z \\sim 0$ to $z \\sim 1$ (e.g., Schiminovich et al. 2005; Le Floc'h et al. 2005).  Remaining problems include the faint-end and luminous-end of the luminosity function and clustering properties of star-forming galaxies. For this purpose, deep imaging observations  with a narrowband filter are highly useful because such surveys could probe fainter star-forming galaxies by detecting their H$\\alpha$ emission. The H$\\alpha$ luminosity is directly connected to the ionizing photon production rate. In particular, wide-field imagers on 8m class telescopes, such as Suprime-Cam on the Subaru Telescope (Miyazaki et al. 2002) are most suitable because of their superior observational capability.  In this paper, we present our new search for faint H$\\alpha$ emitting galaxies at $z = 0.24$ in the Subaru Deep Field (SDF: Kashikawa et al. 2004). One of the uniqueness of the SDF is the field is its very wide field of view ($34^\\prime \\times 27^\\prime$). Another uniqueness is that deep photometry in narrowband filters has been carried out. Since the SDF is one of the deepest fields observed so far by ground-based telescopes, we can probe much fainter star-forming galaxies. Although the sky coverage of SDF is fairly wide, it is difficult to sample a statistically significant number of bright H$\\alpha$ emitters. This makes it difficult to obtain a reliable H$\\alpha$ luminosity function (H$\\alpha$LF) because of large ambiguity in the bright end of luminosity function. Therefore, in this paper, we also use the SDSS data to probe the bright-end of the luminosity function of H$\\alpha$ emitters in the same redshift range.  We note that Ly et al. (2007, hereafter L07) discussed low and intermediate redshift NB704, NB711, NB816, and NB921 emitters in the SDF using the official photometric catalogs. Since their selection criteria are different from those of this work, it is hard to compare with our previous work (Shioya et al. 2008, hereafter S08). We therefore reanalyze the public catalog of the SDF. Details of the differences between our sample and the sample of L07 are discussed in the following section.  Throughout this paper, magnitudes are given in the AB system. We adopt a flat universe with $\\Omega_{\\rm matter} = 0.3$, $\\Omega_{\\Lambda} = 0.7$, and $H_0 = 70 \\; {\\rm km \\; s^{-1} \\; Mpc^{-1}}$. ",
        "conclusions": "\\subsection{H$\\alpha$ Luminosity Function at $z = 0.24$}  In order to investigate the star formation activity in galaxies at $z \\approx 0.24$, we construct the H$\\alpha$LF by using the relation, $\\Phi(\\log L(\\textrm{H}\\alpha)) \\Delta\\log L(\\textrm{H}\\alpha) = \\sum_i 1/V^i$, where $V^i$ is the comoving volume, and the sum is over galaxies with H$\\alpha$ luminosity within an interval of $\\log L(\\textrm{H}\\alpha) \\pm0.1 \\Delta\\log L(\\textrm{H}\\alpha)$. It is unlikely that all the H$\\alpha$ luminosity is produced by star formation because active galactic nuclei (AGNs) can also contribute to the H$\\alpha$ luminosity. In the SDSS sample, we have distinguished the star-forming galaxies from AGNs by using the excitation diagram (Fig \\ref{BPTdiagram}). However, in our SDF sample, we cannot apply the same method because no spectroscopic information is available. We therefore adopt a statistical method following the manner of Pascual et al. (2001). The amount of AGNs depends on the selection criteria. For example, it is 8-17\\% in the CFRS low-$z$ sample (Tresse et al. 1996), 8\\% in the local UCM (Gallego et al. 1995), and 17 -- 28\\% in the 15R survey (Carter et al. 2001) for the number of galaxies. In the following analysis, we adopt that AGNs contribute to 8\\% of the number density in the SDF H$\\alpha$LF.  Combining the two H$\\alpha$ emitter samples obtained from SDF and SDSS, we make a fitting with the Schechter function (Schechter 1976) using STY method (Sandage, Tammann, \\& Yahil 1979) with error estimations adopting the method by Marchesini et al. (2007), and then obtain the H$\\alpha$LF parameters; $\\alpha = -1.31^{+0.17}_{-0.17}$, $\\log \\phi^{*} = -2.46^{+0.34}_{-0.40} {\\rm Mpc^{-3}}$, and $\\log L^{*} = 41.99^{+0.08}_{-0.07} {\\rm ergs \\; s^{-1}}$. Our result is shown in Figure \\ref{LF} together with H$\\alpha$LFs of previous studies at $z < 0.3$ (Tresse \\& Maddox (1998); F03; Hippelein et al. 2003; L07 and S08). We note that F03, L07, S08, and this work are based on the $\\mathit{NB816}$ imaging obtained with the Subaru Telescope.  First, we compare our H$\\alpha$LF with that derived by L07. Although their H$\\alpha$ luminosity was derived from the official catalog of the SDF as our work, their best-fit Schechter function parameters ($\\alpha = -1.71$, $\\log \\phi_*=-3.7$, $\\log L_*=42.2$) are quite different from those of our H$\\alpha$LF. Comparing our data points with those shown in Fig.10b of L07, we find that the slope of their LF at the faint end ($\\log L({\\rm H}\\alpha) \\sim 39.5$-- 41.0) is similar to ours, although their number density is lower than ours by a factor of two. This difference of the number density is explained by the different criteria of the sample selection: they selected emission line galaxies with larger $EW$s than ours. The origin of the significant difference of Schechter parameters between ours and L07's is considered that their fit is overweighted towards the faintest data points with the smallest error bars.  Next, we compare our H$\\alpha$LF with the other H$\\alpha$LFs. Although our H$\\alpha$LF is similar to those of Tresse \\& Maddox (1998) and Hippelein et al. (2003), the H$\\alpha$LF of F03 shows a steeper faint-end slope and a higher number density for the same luminosity than ours. L07 and S08 pointed out that the H$\\alpha$ emitter sample of F03 may include [O{\\sc iii}] or H$\\beta$ emitters because of the inappropriate selection criteria of H$\\alpha$ emitters. We consider this contaminant sources make the faint-end slope of H$\\alpha$ luminosity function of F03 steeper than ours.  \\subsection{H$\\alpha$ luminosity density and star formation rate density}  The H$\\alpha$ luminosity density is obtained by integrating the H$\\alpha$LF: \\begin{equation} L(\\textrm{H}\\alpha) = \\int_0^\\infty \\Phi(L) L\\textrm{d}L = \\Gamma (\\alpha + 2) \\phi^{*} L^{*}. \\end{equation} We then find a total H$\\alpha$ luminosity per unit comoving volume $4.45^{+2.96}_{-1.75} \\times 10^{39} {\\rm ergs \\; s^{-1}}$ at $z\\approx0.24$ from our best fit H$\\alpha$LF.  The star formation rate (SFR) can be estimated from the H$\\alpha$ luminosity using the relation, $SFR = 7.9 \\times 10^{-42} L(\\textrm{H}\\alpha) ~ M_{\\odot}$ yr$^{-1}$, where $L(\\textrm{H}\\alpha)$ is in units of ergs s$^{-1}$ (Kennicutt 1998). Thus, the H$\\alpha$ luminosity density can be translated into the SFR density (SFRD) of $\\rho_{\\rm{SFR}} \\simeq 0.035^{+0.024}_{-0.014} ~ M_{\\odot} \\; {\\rm yr^{-1} \\; Mpc^{-3}}$ . This value is higher by a factor of 3  than the local SFRD.  We compare our result with the previous investigations compiled by Hopkins (2004) in Figure \\ref{SFRD}. We also show the evolution of the SFRD derived from the observation of GALEX (Schiminovich et al. 2005) with mean attenuation of $A_{\\rm UV}^{\\rm meas} = 1.8$, evaluated from the relation between $A_{UV}$ and FUV slope for starburst galaxies (Meurer et al. 1999). Recent observations showed that the typical value of $A_{UV}$ is 1 mag for nearby galaxies detected in the UV by GALEX (Seibert et al. 2005; Buat et al. 2005; Tresse et al. 2007). If we adopt the attenuation of $A_{\\rm UV} = 1$, the SFRD derived from the FUV becomes smaller by factor 2.  Our SFRD estimated above seems roughly consistent with the previous evaluations, e.g., Tresse \\& Maddox (1998) and F03. We note that the SFRD of F03 was overestimated because of the contamination of [O{\\sc iii}] emitters. The SFRD of S08 is slightly smaller than ours. Since our selection criteria are similar to those of 08, the difference between S08 and our work seems to be real; e.g., the cosmic variance. The SFRD of L07 is smaller than other evaluations. The small SFRD of L07 is originated from the different selection criterion of H$\\alpha$ emitters as we mentioned in section 4.1.  \\subsection{Spatial Distribution and Angular Two-Point Correlation Function in the SDF}  We show the spatial distribution of our 258 H$\\alpha$ emitter candidates in Figure \\ref{XY}. There are some concentrations of H$\\alpha$ emitters in the corners. To discuss the clustering properties more quantitatively, we derive the angular two-point correlation function (ACF), $w(\\theta)$, using the estimator defined by Landy \\& Szalay (1993), \\begin{equation} w(\\theta) = \\frac{DD(\\theta)-2DR(\\theta)+RR(\\theta)}{RR(\\theta)}, \\label{two-point} \\end{equation} where $DD(\\theta)$, $DR(\\theta)$, and $RR(\\theta)$ are normalized numbers of galaxy-galaxy, galaxy-random, and random-random pairs, respectively. The random sample consists of 100,000 sources with the same geometrical constraints as the galaxy sample. The formal error in $w(\\theta)$ is described by \\begin{equation} \\sigma_w=\\sqrt{[1+w(\\theta)]/DD} \\end{equation} (Hewett 1982). We show the ACF for all our 258 H$\\alpha$ emitter candidates in Figure \\ref{ACF} together with that for 980 H$\\alpha$ emitters in the COSMOS field (S08). The FOV areas from which we derive $w(\\theta)$ are $37^\\prime \\times 27^\\prime$ for the SDF and $1.4^\\circ \\times 1.4^\\circ$ for the COSMOS field. The amplitude of our ACF is positive within the 0.06 degree ($\\sim 4$ arcmin) and it is larger than that of the COSMOS field. We consider that this clustering property reflects the concentration of H$\\alpha$ emitters in the corners. Our ACF is fit well by power law, $w(\\theta) = 0.047^{+0.017}_{-0.013} \\theta^{-0.66 \\pm 0.08}$.  It is useful to evaluate the correlation length $r_0$ of the two-point correlation function $\\xi(r) = (r/r_0)^{-\\gamma}$. A correlation length is derived from the ACF through Limber's equation (e.g., Peebles 1980). Assuming that the redshift distribution of H$\\alpha$ emitters is a top hat shape of $z=0.242 \\pm 0.009$, we obtain the correlation length of $r_0 = 2.6^{+1.0}_{-0.8}$ Mpc. Therefore, the two-point correlation function for all H$\\alpha$ emitters is written as $\\xi(r) = (r/{\\rm 2.6^{+1.0}_{-0.8}Mpc})^{-1.66 \\pm 0.08}$. This correlation length is larger and the slope of the correlation function is shallower then those evaluated for H$\\alpha$ emitter at $z \\approx 0.24$ in COSMOS; $\\xi(r) = (r/{\\rm 1.9Mpc})^{-1.88}$ (S08). We note that the correlation lengths referred in this section are evaluated for the cosmological parameters of $\\Omega_{\\rm matter}=0.3$, $\\Omega_\\Lambda=0.7$, and $H_0=70 {\\rm km \\; s^{-1} \\; Mpc^{-1}}$. This difference reflects the difference of the amplitude of the ACF within $\\sim 4$ arcmin. The correlation length is smaller than that of nearby $L_*$ galaxies ($\\sim 7$ Mpc; Loveday et al. 1995; Norberg et al. 2001; Zehavi et al. 2005) and $z \\sim 1$ galaxies ($\\sim$ 4--5 Mpc; Coil et al. 2004) It is known that the correlation length is smaller for fainter galaxies and is smaller for later-type galaxies in the nearby universe (Loveday et al. 1999; Norberg et al. 2001; Zehavi et al. 2005). Although the absolute $R$-band magnitudes of our sample ($M_R \\sim -18$) is faint, the correlation length is still smaller than the nearby galaxies with the same absolute magnitude (3.8 Mpc; Zehavi et al. 2005). The clustering strength of emission-line galaxies (active star-forming galaxies) may be weaker than that of non-active star-forming galaxies.  \\vspace{1pc} We would like to thank the Subaru Telescope staff for their invaluable help. We would like to thank Nobunari Kashikawa who is the PI of SDF, and the other SDF members. We also thank Chun Ly and Matt Malkan who gave us useful information of their independent study of H$\\alpha$ emitters in SDF prior to the publication. This work has been financially supported in part by grants of JSPS (Nos. 15340059 and 17253001)."
    },
    "0807/0807.2568_arXiv.txt": {
        "abstract": "We present very sensitive polarimetry of 55 Cnc and $\\tau$ Boo in an attempt to detect the partially polarised reflected light from the planets orbiting these two stars. 55 Cnc is orbited by a hot Neptune planet (55 Cnc e) at 0.038 AU, a hot Jupiter planet (55 Cnc b) at 0.11 AU, and at least 3 more distant planets. The polarisation of this system is very stable, showing no sign of the periodic variations that would be expected if a short period planet were detected. The measured standard deviation of the nightly averaged Stokes Q/I and U/I parameters is 2.2$\\times 10^{-6}$. We derive upper limits on the geometric albedo, $A_G$ and planetary radius using Monte Carlo multiple scattering simulations of a simple model atmosphere. We assume Rayleigh-like scattering and polarisation behaviour (scaled by the maximum polarisation, $p_m$ at 90$^{\\circ}$) and pressure insensitive extinction. Atmospheres in which multiple scattering plays only a small role have an almost linear relation between polarisation and $A_G$. In this case, the 4$\\sigma$ upper limit is $A_G<0.13(R/1.2 R_{Jup})^{-2}p_m^{-1}$ for 55 Cnc e. This is most easily explained if 55 Cnc e is relatively small, like GJ436b, and therefore not a pure H-He planet. The data do not provide a useful upper limit for 55 Cnc b. $\\tau$ Boo is orbited by an unusually massive hot Jupiter planet. The data show a standard deviation in the night to night average Stokes Q/I and U/I polarisation parameters of 5.1$\\times 10^{-6}$. The 4$\\sigma$ upper limit is $A_G<0.37(R/1.2 R_{Jup})^{-2}p_m^{-1}$ for $\\tau$ Boo b, adopting the fairly well established orbital inclination $i$$\\sim$40$^{\\circ}$. This extends the similar upper limits reported previously for this planet to longer wavelengths. The fact that the $\\tau$ Boo data show more scatter, despite the smaller photon noise for this bright star, may be due to the spot activity detected photometrically by the {\\it MOST} satellite. These results contrast markedly with the recent claim of a 3$\\sigma$ detection of a periodic polarisation signal from HD189733 with amplitude P$=2\\times 10^{-4}$, attributed to the planet HD189733~b. ",
        "introduction": "Since the seminal discovery of 51 Peg b (Mayor \\& Queloz 1995), $\\sim$300 extrasolar planets have been found (e.g. Butler et al.2006, Schneider 2008). Approximately one third of these orbit very close to their central star, with semi-major axes, $a<0.1$~AU (Schneider 2008). The number of short period planets is increasing rapidly with the growing success of transit based searches such as SuperWASP (Pollacco et al.2006) and HATNet (Bakos et al.2004), which are even more strongly biased towards such detections than the radial velocity method.  These close in planets offer the best opportunity to study the physical characteristics of the planets, as opposed to their orbital parameters. The combination of transit data and radial velocity data yields a precise planetary mass, size and density. Furthermore, recent studies with the {\\it Spitzer} Space Telescope (e.g. Knutson et al.2007; 2008; Burrows et al.2007; Tinetti et al.2007) and the {\\it Hubble} Space Telescope (e.g. Pont et al.2008; Swain et al.2008; Lecavelier des Etangs 2008) have obtained a great deal of information about the chemical composition and pressure temperature profiles of the atmospheres of the two brightest transiting ``hot Jupiter'' planets, HD189733~b and HD209458~b. These studies used infrared photometry and spectroscopy of radiation emitted by the planets and optical spectrophotometry of the transit. All required exceptionally high signal to noise data for the integrated light of the star and planet. These results build on the more limited data provided by earlier space based studies, e.g. Charbonneau et al.(2002); Harrington et al.(2006).  While space-based observations have now produced high quality data on the atmospheres of these two planets, ground based observations have yet to produce a confirmed direct detection of an extrasolar planet. Neither space based nor ground based efforts have yet detected an extrasolar planet in reflected light. The MOST satellite has provided strong upper limits on the reflected light from HD209458~b (Rowe et al.2007) and ground based searches have also produced meaningful upper limits for a few planets (e.g. Charbonneau et al.1999; Leigh et al.2003; Winn et al.2008; Snellen 2005; Snellen \\& Covino 2007). Recently, Berdyugina et al.(2008) have reported a possible detection of reflected light from HD189733~b at low signal to noise, using ground based polarimetry with a conventional instrument employing slow modulation. However, the reported fractional polarisation is so large ($P \\sim 2 \\times 10^{-4}$) that it exceeds that expected for a perfectly reflective planet (i.e. a Lambert sphere with Rayleigh scattering polarisation properties) with the size indicated by the transit data, see Seager, Whitney \\& Sasselov (2000).  \\begin{table*} \\begin{minipage}{100mm} \\caption{Orbital parameters} \\begin{tabular}{lccccc} \\hline Planet & $m.sin(i)$ (M$_{Jup}$) & a (AU) & Period (d) & e & i \\\\ \\hline 55 Cnc e &   0.034    \t      &  0.038   & 2.817  & 0.07$\\pm$0.06 & $\\sim55^{\\circ}$     \\\\ 55 Cnc b &   0.824    \t      &  0.115   & 14.652  & 0.014$\\pm$0.008 & $\\sim55^{\\circ}$  \\\\ $\\tau$ Boo b &  4.13   \t      &  0.0481   & 3.312  & 0.023$\\pm$0.015 & $\\sim40^{\\circ}$  \\\\ \\hline \\end{tabular} \\end{minipage} \\end{table*}  Polarimetry is nonetheless a promising method to detect extrasolar planets in reflected light. The basic principle is that reflected light is partially polarised (Seager et al.2000; Stam, Hovenier \\& Waters 2004; Stam et al.2006; Stam 2008) whereas the direct light emitted by stellar photospheres has negligible linear polarisation. Kemp et al.(1987) measured an upper limit on the integrated linear polarisation from the solar disc of $P<2\\times 10^{-7}$ in the V band. Polarimetry therefore circumvents the contrast problem associated with extrasolar planets and their central stars. It can be used to detect planets that are spatially unresolved from the central star by searching for changes in the polarisation that have the same period as the planet's orbit (Seager et al.2000). In this case the time varying polarisation due to the planet will usually be superimposed on a constant polarisation that is caused by dichroic extinction in the interstellar medium along the line of sight to Earth. The interstellar polarisation of stars in the local ionised bubble within $\\sim$100~pc of the sun is very small (typically P$\\sim$10$^{-5}$-10$^{-6}$, see Hough et al.2006) and is assumed to be constant during the period of an observing run, see $\\S$4.4. Alternatively, polarimetry can be used to assist in detecting planets that are spatially resolved from the central star by very high resolution imaging. This approach is be attempted with new instruments such as Hi-CIAO on Subaru, {\\sc NICI} and {\\sc GPI} on Gemini and {\\sc SPHERE} on the VLT.  In this paper we have adopted the first approach, using an aperture polarimeter to observe exoplanet systems in which the planet cannot be spatially resolved in the forseeable future. Hough et al.(2006) described the {\\sc Planetpol} instrument, which was specifically designed to achieve a sensitivity to fractional linear polarisations of order 10$^{-6}$ by using a fast modulator system, similar in principle to that described by Kemp et al.(1987) and references therein. Polarimetry with fast modulators (which induce a rapid periodic variation of the relative retardance of electromagnetic vibrations in orthogonal planes) permits this very high sensitivity by separating the polarised component of the incident radiation field from the unpolarised component, converting the former into a high frequency signal in the time domain. This means that very small fractional polarisations can be measured without the need for very precise flux measurements with the detector. {\\sc Planetpol} achieves photon noise limited performance for bright exoplanet systems so the main limitation on the polarisation sensitivity is the need to gather in excess of 10$^{12}$ photons in order to measure fractional polarisations as low as 10$^{-6}$.  The fast modulation is performed by Photoelastic Modulators (PEMs) operated at 20 kHz, which is fast enough to remove any effects of the Earth's atmosphere.  By contrast, conventional polarimeters for night time astronomy employ slow modulation, rotating a waveplate with a fixed retardance of 0.5 wavelengths to several different position angles. Measurement of fractional polarisations of order $10^{-6}$ with such a system would require a precision of order $10^{-6}$ in the measured flux on the detector, which would be very challenging. In practice such polarimeters rarely achieve a sensitivity better than $P=10^{-4}$.  While the recent space based measurements have revealed much about the atmospheres of two transiting planets the energy balance of these systems cannot yet be understood without measurement of their albedos. Most of the space based measurements have been obtained during transits, which only sample the upper layers of the atmosphere, not the principal reflecting layer. Even if space based spectroscopy is obtained outside transit events, it is difficult to infer the presence or absence of dust particles and their contribution to the radiative transfer without polarisation data. Successful polarimetric detection of exoplanets could address these questions and permit determination of the inclination of the planetary orbit (transiting systems are not required). For example, Hansen \\& Hovenier (1974) used ground-based aperture polarimetry to determine that sulphuric acid clouds, with $\\sim$1~$\\mu$m-sized droplets, dominate the reflection of sunlight from Venus.  Here we describe the results of {\\sc Planetpol} polarimetry of two exoplanet systems, 55 Cnc and $\\tau$ Boo. The orbital parameters of their potentially detectable inner planets are listed in Table 1, taken from Butler et al.(2006) and Fischer et al.(2008) (with the exception of the inclinations, see below). These systems were selected on the basis that their bright I band magnitudes (brighter than any of the known transiting planets) and the short orbital period of their closest planetary companions made them good prospects to produce a detectable polarised flux.  The 55 Cnc system (HR3522) is a wide stellar binary, consisting of a G8IV-V primary (Baliunas et al.1997) with Cousins I mag=5.1 and an M4 secondary (Patience et al.2002; Eggenberger et al.2004; Raghavan et al.2006). The primary is now believed to be orbited by 5 planets (Fischer et al.2008; Butler et al.1997, Marcy et al.2002; McArthur et al.2004) and the inclination of the planetary orbits has some constraints from high precision astrometry of reflex motion caused by the outermost planet (McArthur et al.2004). It has an R$\\prime$$_{HK}$ index of -4.97 (Soderblom 1985; Baliunas et al.1996) and a rotation period of 44 days (Soderblom 1985). The innermost planet, 55 Cnc e, is a ``hot Neptune'' and the next planet out, 55 Cnc b, is a hot Jupiter. Both planets could reasonably be expected to produce a detectable signal, although the relatively large orbit of 55 Cnc b (see Table 1) would require it to have a high albedo and a large radius, and the Neptune mass planet 55 Cnc e might also require a high albedo, depending on the size of the planet.  The $\\tau$ Boo system (HR5185) is also a wide stellar binary composed of an F7IV-V primary (Baliunas et al.1997) with Cousins I mag=4.0 and an M2V secondary (Raghavan et al.2006). The primary is orbited by one of the most massive known hot Jupiter planets (Butler et al.1997), which is commonly referred to as $\\tau$ Boo b. Baliunas et al.(1997; 1996) found the star to be photometrically stable at the millimagnitude level, with an R$\\prime$$_{HK}$ index of -4.73. However, recent space based photometry by MOST has found that the stellar flux is variable at the millimagnitude level (Walker et al.2008). The period is consistent with the 3.3 day planetary orbit, suggesting that there is magnetic activity on the stellar surface that is linked to the planet. The changes in flux appear to be due mainly to a single star spot. However, the amplitude of the changes varies on long timescales, in some years apparently corresponding to a bright spot rather than a dark spot. $\\tau$ Boo is believed to be a unique case where the massive planet has tidally spun up at least the outer layers of the star to have the same rotation period as the planet's orbit. The inclination of the stellar rotation axis can be determined from the measured photometric period and the $v sin(i)$ Doppler broadening of stellar lines as $i$ $\\approx$ 40$^{\\circ}$, e.g. Leigh et al.(2003). Recent mapping of the magnetic field using (circular) spectropolarimetry (Catala et al.2007) also finds a similar intermediate inclination. Assuming that the planetary orbit shares this inclination, this allows the mass of the planet to be determined as $\\sim$ 6-7~M$_{Jup}$. ",
        "conclusions": "The strong upper limit for 55 Cnc e, $A_G$$<$0.13$(R/1.2 R_{Jup})^{-2}$$p_m^{-1}$, is hard to explain in terms of a low albedo. At the time the observations were planned there were no available calculations or data concerning the likely size of very strongly irradiated Neptune mass planets. Subsequent calculations by Baraffe et al.(2006) suggest that while a wide range of sizes (0.4~$R_{Jup}$ to $\\sim 1~R_{Jup}$ are possible, it is likely that hot Neptune planets will usually have radii $\\le 0.6 R_{Jup}$. The larger radii can occur for hydrogen and helium dominated Neptune mass planets but these may evaporate on a relatively short timescale. Unfortunately the physics of the evaporation process is still not clearly established. The sole known transiting hot Neptune, GJ436b, has a radius between 0.38 and 0.44~R$_{Jup}$ (Gillon et al.2007a; 2007b; Deming et al.2007; Bean et al.2008; Torres et al.2007), though this may not be a useful comparison planet owing to the much weaker irradiation from its M2V star. Our non-detection of 55 Cne e is most easily explained by a relatively small planet. The data allow us to rule out a Jupiter sized planet with even a low geometric albedo and also rule out a 0.8$~R_{Jup}$ planet with a moderate albedo. However, we caution that it is possible for planets to have very low geometric albedos. For example the radiative transfer could be dominated by very small dust particles with a very small single scattering albedo. Also, if there is no high altitude dust layer, strong atomic and molecular absorption features such as the KI resonant line at 770 nm could reduce the geometric albedo in the {\\sc Planetpol} passband. Furthermore, Hood et al.(2008) have shown that 3D structure in planetary atmospheres can reduces the geometric albedo by a large fraction, relative to homogeneous atmospheres such as those modelled in this work. Hence a planetary radius as high as $R$=1.2~$R_{Jup}$ cannot be entirely ruled out.  We conclude from the upper limit for $\\tau$ Boo b that this is a fairly dark planet in our 590-920 nm passband. This extends the similar limits that were found by Charbonneau et al.(1999) and Leigh et al.(2003) at shorter wavelengths by searching for Doppler shifted absorption lines in the stellar spectrum. A non-detection at an upper limit of $A_G$$<$0.37$(R/1.2 R_{Jup})^{-2}$$p_m^{-1}$ is not particularly surprising in view of the possible absorption mechanisms described above. The most plausible example model of Seager et al.(2000), in which the scattering is dominated by a mixture of three types of 0.1~$\\mu$m grains, gave $A_G=0.175$. Nonetheless, the data indicate that $\\tau$~Boo b is less reflective than Jupiter and Saturn, which have $A_G$$\\approx$0.5 throughout most of the {\\sc Planetpol} passband (Karkoschka et al.1994). This is of course subject to the assumption that the scattering particles have Rayleigh-like polarisation behaviour, i.e. the scattering is dominated by molecules or by dust particles $\\la$0.1~$\\mu$m in radius. Both our data and the {\\it MOST} data indicate that this planet is not an ideal target for the highest precision studies.  Our multiple scattering model atmosphere calculations indicate that the large amplitude periodic polarisation signal from the HD189733 system that was reported by Berdyugina et al.(2008) cannot be explained in terms of reflected light from the planet HD189733b. If the observations are confirmed it would be important to consider the possible contribution of star spots to the polarisation of the system, given that HD189733 is an active star with much larger photometric variations than $\\tau$ Boo or 55 Cnc, e.g. Winn et al.(2007).  The measurements for nearby stars with low polarisation are stable on timescales of a week. It is not yet clear whether they are stable at the 10$^{-6}$ level on timescales of a year. Finally, we note that these results do not represent the full potential of the instrument. The new detectors offer a 40\\% improvement in sensitivity, so observations of stable bright hot Jupiter systems such as $\\upsilon$~And and 51 Peg would offer a good chance of a successful detection. The new class of very hot Jupiters recently detected by transit surveys around fainter stars may also offer good prospects for observation. Transiting planets, which have $i$$\\approx$90$^{\\circ}$, display the full range of phase angles (0-180$^{\\circ}$) during the course of an orbit, thereby offering the maximum possible amount of information. For the purpose of simple detection, there is less need to sample the full phase curve since the most likely timing of maximum polarisation is predictable if we know that $i$$\\approx$90$^{\\circ}$ and assume that the scattering phase function is not highly aniostropic. However, the expected amplitude of the periodic polarisation changes in Stokes Q/I and U/I is slightly smaller at $i$=90$^{\\circ}$ than for smaller inclinations, and we caution that  planets in such small orbits might induce variable polarisation on their central star due to star spots."
    },
    "0807/0807.4933_arXiv.txt": {
        "abstract": "We examine the properties of 18 long period ($80-210$ days) and very luminous (median absolute magnitude of M$_I=-7.86$ and M$_V=-6.97$) Cepheids to see if they can serve as an useful distance indicator. We find that these Ultra Long Period (ULP) Cepheids have a relatively shallow Period-Luminosity (PL) relation, so in fact they are more ``standard candle''-like than classical Cepheids. In the reddening-free Wesenheit index, the slope of the ULP PL relation is consistent with zero. The scatter of our sample about the W$_I$ PL relation is $0.23$ mag, approaching that of classical Cepheids and Type Ia Supernovae. We expect this scatter to decrease as bigger and more uniform samples of ULP Cepheids are obtained. We also measure a non-zero period derivative for one ULP Cepheid (SMC HV829) and use the result to probe evolutionary models and mass loss of massive stars. ULP Cepheids main advantage over classical Cepheids is that they are more luminous, and as such show great potential as stellar distance indicators to galaxies up to $100$ Mpc and beyond. ",
        "introduction": "\\label{sec:intro}  A reliable method of measuring the physical distance to astrophysical objects has always been sought after in observational astronomy \\citep[\\eg,][]{Bessel1839}. In the era of ``precision cosmology'', the need for accurate physical distance measurements has been amplified \\citep[\\eg,][]{Spergel03, Riess04, Tegmark04}. Accurate and precise distance indicators hold the key to pinning down the value of the Hubble constant $(H_0)$ and many other cosmological parameters \\citep[see discussion in, \\eg,][]{Macri06}. A number of methods have been employed to determine extragalactic distances, with varying degree of success \\citep[\\eg,][]{Freedman01}. The construction and reliability of the ``cosmological distance ladder'' depends crucially on Cepheid variables being able to provide precise and accurate distances to nearby ($d\\lesssim20\\;$Mpc) galaxies.  The quest for such distances has been an arduous journey for almost a hundred years, with many dramatic twists and turns \\citep[for a review of early years, see][for a recent review see, \\eg, \\citealt{Macri05}]{Baade1956}.  Cepheids offer several advantages as distance indicators. Massive stars $(\\ge5 M_{\\odot})$ make an excursion through the instability strip and most, if not all, of them become Cepheid variables. These variable stars are relatively bright ($M_V\\sim -4$ for a $P\\sim 10\\;$day Cepheid) and often have large brightness variations (amplitude $\\sim1$ mag) with a characteristic ``saw-tooth'' shaped light curve. Their intrinsic brightness, combined with their light curve shape and colors, make them easy to distinguish from other classes of variable stars. As a result, Cepheids have been detected and studied in a significant number of star-forming galaxies. The physical mechanisms underlying Cepheid pulsation are well understood, including the observed tight period-luminosity (PL) relationship \\citep[\\eg,][]{Chiosi93}. The small scatter in the PL relation allows distance measurements precise to $\\sim 5\\%$ \\citep[\\eg,][]{Macri06}. For these reasons, Cepheids are commonly used to calibrate other distance indicators, forming the base of the cosmological distance ladder.  Despite their many advantages as a distance indicator, Cepheid distances also have some shortcomings. Most Cepheids have an intrinsic brightness of $M_V \\ge -5$, so with the current instrumentation they can be only used to measure distances to $\\lesssim30\\;$Mpc (the largest Cepheid distance in \\citealt{Freedman01} is $\\sim 22\\;$Mpc). Observations of Cepheids in distant galaxies are also hindered by blending \\citep{Mochejska00}--- as young stars, Cepheids live in close proximity to the crowded star-forming regions of their host galaxies, and are thus likely to have another star of similar brightness on the scale of a typical instrumental point-spread-function (PSF). The effect of blending becomes worse as the square of the distance to the host galaxy \\citep{Stanek99}, again limiting the usefulness of Cepheids to measuring distances $\\lesssim30\\;$Mpc even with high resolution instruments such as the Hubble Space Telescopes ({\\em HST}). Ideally, we would like to find a distance indicator that shares the good properties of classical Cepheids, but is even more luminous, allowing us to observe it further away and be less susceptible to blending. In this paper we discuss such a possible distance indicator, namely Ultra Long Period (ULP) Cepheids.  We define ULP Cepheids as fundamental mode Cepheids with pulsation periods greater than 80 days. Several such Cepheids have been already discovered in the pioneering study of \\citet{Leavitt1908}. However, ULP Cepheids have traditionally been ignored for distance measurements as they are PL outliers.  Indeed, the observed PL relation flattens for Cepheids with periods greater than 100 days \\citep[\\eg,][]{Grieve85,Freedman92}. \\citet{Grieve85} suggests that long period Leavitt Variables could be used for distance measures--- unfortunately that idea has not permeated through the community.  We argue that the flattening of the PL at long periods actually improves the usefulness of ULP Cepheids as distance indicators because it makes them a good standard candle in the traditional sense. We note several additional advantages of ULP Cepheids over lower period Cepheids due to their increased luminosity. ULP Cepheids could be used as a stellar distance measure to the Hubble Flow (up to $\\sim150$ Mpc)--- several times the current observational limit. In Section~\\ref{sec:sample} we describe our sample compiled from the literature. The ULP Cepheid PL relation is discussed in Section~\\ref{sec:DM}. Section~\\ref{sec:massloss} demonstrates how ULP Cepheids may provide the additional benefit of testing massive stellar evolutionary models. We summarize our results in Section~\\ref{sec:conclusion}. ",
        "conclusions": "\\label{sec:conclusion}  We have presented, for the first time, a collection of Ultra Long Period ($P>80\\;$days) Cepheids from the literature and demonstrated their viability as a distance indicator. In the past, ULP Cepheids have been ignored as distance indicators, and in fact many stellar variability searches did not extend their cadence and search strategy to allow for their discovery.  In $V$, $I$, and especially reddening-free W$_I$ magnitude, ULP Cepheids have a relatively flat PL relation compared to the respective relations for classical Cepheids \\citep{Udalski99}.  The dispersion in both the classical and ULP Cepheid populations about their respective PL relationships becomes smaller as one moves from $V$ to $I$ to W$_I$, and the discrepancy between the slopes increases. Most notably, the slope of the ULP Cepheid W$_I$ PL relation is significantly shallower than the standard SMC PL relation (slope is $-0.05$ vs. $-3.28$). The reddening-free Wesenheit index produces the tightest ULP PL relation, with RMS residual of only $0.23$ mag (See Figure~\\ref{fig:PL}). Other papers \\citep[\\eg][]{Kanbur07} have found non-linearites in the PL relationship at lower periods. However, the change in slope seen here is much more dramatic and it is unlikely that it shares a physical connection with the PL changes at lower periods.   Our ULP PL scatter in W$_I$ is already less than that of the initial peak brightness vs. absolute magnitude relation for Type Ia Supernova ($\\sim 0.3$ mag; \\cite{Phillips93}).  A huge amount of effort has been necessary to increase the number of observed Type Ia SNe, refine the calibration technique \\citep[\\eg,][]{Prieto06}, and to obtain the low $0.15$-$0.20$ mag scatter in the relation seen today \\citep[\\eg,][]{Jha07}. We expect future observational and theoretical studies of ULP Cepheids will further decrease the already fairly small scatter found in this first analysis.  We strove to find all the known ULP Cepheids in the literature, but our sample is likely not a complete census of ULP Cepheids. We note that the Araucaria Project found ULP Cepheids in three of the five galaxies they observed at the time of this analysis, so the ULP Cepheid sample should grow at a reasonable rate as Cepheid studies are extended to more galaxies and previous data sets are reanalyzed with longer period searches. There is also evidence ULP Cepheids exist in a broad range of metallicities, as preliminary data analysis of a M81 variability survey with the Large Binocular Telescope has already discovered at least one ULP Cepheid (Kochanek, private communication).  In addition to following a fairly tight PL relation, the ULP Cepheids are also very luminous, with a median absolute magnitude in $I$($V$) for our sample of $-7.86(-6.97)$. Using WFPC3, {\\em HST}\\/ can obtain $10\\%$ photometry at $V=28$ or $I=27$ (DM $= 35$ --- distance of $100\\;$Mpc for the median ULP Cepheid) in about 10 orbits, while only two orbits are needed to reach a DM of $34$. As one would only need a few orbits per epoch, one could detect the median ULP Cepheid (with a period of $\\sim121$ days) at 100 Mpc and the brightest ULP Cepheids at $\\sim150$ Mpc in a reasonable amount of time. Since the luminosity of an ULP Cepheid is a weak function of its period, relatively accurate distances would not require as precise period measurements as is needed for classical Cepheids. We encourage future variability proposals to search for Cepheids with periods greater than 100 days.  We note two concerns in using ULP Cepheids for distance measurements. It is obvious from our sample size that ULP Cepheids are far less common than classical Cepheids. The relatively small ULP Cepheid population will make precision distance measurements less practical. As the sample size grows and the ULP Cepheid PL relations become established, a single ULP Cepheid observation may yield distance measurements accurate to $10-20\\%$. Blending is always a concern in Cepheid studies. Blending can compromise Classical Cepheid measurements as stars of comparable brightness are likely to lie within a single PSF \\citep{Mochejska00}, especially at larger distances. We expect blending to be less of a problem for our ULP Cepheids even at large distances simply because they are so bright and there are very few stars of comparable brightness in a given galaxy. The effect of blending on ULP Cepheid observations will need to be investigated further as the sample is increased.  \\begin{figure*}[p] \\figurenum{7} \\plotone{f8.eps} \\caption{\\label{fig:dpdt} $V$-band photometry of SMC HV829 taken from 1970-1976 (left panel) and 2000-2004 (right panel; see text for data references). Each set of data was phased with the period shown at the top of each pane. The pulsation period has clearly decreased, from 85.9 days to 84.4 days over approximately 30 years.} \\end{figure*}  \\begin{figure*}[p]% \\epsscale{0.80} \\figurenum{8} \\plotone{f7.eps} \\caption{\\label{fig:mass} CMD of OGLE SMC Cepheids (small blacks dots) and ULP Cepheids (large open symbols) overlaid with 12 M$_{\\odot}$, 15 M$_{\\odot}$, and 20 M$_{\\odot}$ evolutionary models of \\citet{Lejeune01}. The metallicity chosen for the models is similar to the SMC. For each mass, models incorporating ``normal'' (solid line) mass loss and ``double'' (dashed line) mass loss are plotted.  The two mass loss models produce essentially the same evolutionary track for each mass plotted. The large arrow denotes that the pulsation period of HV829 is becoming smaller; therefore the star is undergoing its second crossing of the instability strip.} \\end{figure*}  The data set and analysis herein provides meaningful constraints on theoretical work on Cepheids in this mass and period range. We have shown that ULP Cepheids show strong evidence for a different, flatter PL relation than their lower period cousins. Several papers have modeled Cepheid pulsations and mapped these to theoretical PL relations \\citep[\\eg,][]{Bono02}. However, this work has not been reliably extended to the period range of our sample. Pulsation models in this period range would also help determine the intrinsic scatter to the PL$_{ULP}$ relation.  Our current sample is not large enough to constrain the sensitivity of the ULP Cepheid PL to metallicity. The PL residuals as a function of metallicity are consistent with the results for shorter period Cepheids \\citep[$\\gamma= -0.29 \\pm0.09\\pm 0.05$ mag dex$^{-1}$;][]{Macri06}. Note that the median $12 + \\log($O/H$)$ value for a galaxy in our sample is about $0.5\\;$dex lower than the corresponding value for the {\\em HST}\\/ Key Project Cepheid hosts.  Period changes of ULP Cepheids have powerful implications for stellar evolutionary models in this mass regime. We examined photometry from two epochs separated by thirty years to check for evidence of period changes in seven Magellanic Cloud ULP Cepheids. Only one, SMC HV829, showed a significant period change--- from $85.9\\pm0.3$ to $84.4\\pm0.4$ days. The negative period derivative indicates that the Cepheid is undergoing its second crossing. Assuming the metallicity of the SMC, many current evolutionary models do not predict second crossings for Cepheids in this mass range (See Figure~\\ref{fig:mass}). Future models should incorporate this observational result as it should place limits on several key input parameters. In addition to their potential as distance indicators, ULP Cepheids provide observational constraints on high mass stellar evolution models.  To summarize, ULP Cepheids, while often dismissed in the past, are potentially a powerful distance indicator, probe of PL metallicity sensitivity, and also a probe of massive star evolutionary models. ULP Cepheids could provide the first direct stellar distance measurements to galaxies in the 50-150 Mpc range, extending the cosmological distance ladder well into the Hubble flow."
    },
    "0807/0807.4102_arXiv.txt": {
        "abstract": "{Within the context of investigating possible differences between the mechanisms at play in Radio Loud AGN and those in Radio Quiet ones, we study the spectral characteristics of a selected sample of Intermediate-Luminosity Broad-Line Radio Galaxies in X-rays, optical, IR and radio. Here, we present the radio spectra acquired with the 100-m radio telescope in Effelsberg between 2.6 and 32 GHz. These measurements reveal a large variety of spectral shapes urging for radio imaging that would disclose the source morphology. Such studies could potentially discriminate between different mechanisms. ",
        "introduction": "A great deal of our current understanding of the processes at play in AGNs is coming from studies of chiefly Radio-Quiet (RQ) objects. In fact, it appears that different mechanisms become important at different luminosities so that the emission from luminous RQ systems is dominated by accretion flow \\citep{Haardt1991ApJ}, whereas in lower luminosity systems the accretion flow may be radiatively less efficient \\citep{Ho1999ApJ}.  On the other hand, little is known about the mechanisms determining the behavior of the Radio-Loud (RL) AGNs, mostly because they comprise only a small fraction of the complete AGN population. Studies of the Broad-Line Radio Galaxy (BLRGs) sub-class have been focusing on the brightest and most luminous objects. This rough picture leaves several questions unanswered such as whether there are differences between the physics of RL and RQ systems. A uniform coverage of the parameter space of the studied AGNs could shed light on many of these issues. In particular, the detailed comparison of RL and RQ AGNs of similar luminosity can provide significant insight.  To make up for the bias towards the most luminous sources we have selected a sample of 6 Intermediate-Luminosity BLRGs (IL-BLRGs) and studied their radio spectrum. The ultimate goal of this attempt would be to determine their Spectral Energy Distribution from radio to X-rays utilizing the {\\sl XMM-Newton} observations that have been performed for these sources (see Sambruna et al. in prep.). That would potentially reveal the relative importance of the various emission mechanisms. Further, radio observations in conjunction with X-ray measurements can provide information about the relative contribution of the disk and the jet and how that depends on luminosity. Here, we present the results of the radio study only. ",
        "conclusions": "Clearly, the 6 BLRGs discussed here show a variety of radio spectra that could be interpreted in different ways. Interferometric observations would be essential for the correct interpretation of the observed phenomenology. All these results in conjunction with the X-ray and optical data are in preparation for publication (Sambruna et al., in prep.)."
    },
    "0807/0807.1727_arXiv.txt": {
        "abstract": "We present results from a low resolution (R~$\\simeq$~300) near-infrared spectroscopic variability survey of actively accreting T Tauri stars (TTS) in the Taurus-Auriga star forming region.  Paschen and Brackett series H~{\\scshape i} recombination lines were detected in 73 spectra of 15 classical T Tauri systems.  The values of the Pa$n_{up}$/Pa$\\beta$, Br$n_{up}$/Br$\\gamma$, and Br$\\gamma$/Pa$n_{up}$ H~{\\scshape i} line ratios for all observations exhibit a scatter of~$\\la$~20\\% about the weighted mean, not only from source to source, but also for epoch-to-epoch variations in the same source.  A representative or `global' value was determined for each ratio in both the Paschen and Brackett series as well as the Br$\\gamma$/Pa$n_{up}$ line ratios.  A comparison of observed line ratio values was made to those predicted by the temperature and electron density dependent models of Case~B hydrogen recombination line theory.  The measured line ratios are statistically well-fit by a tightly constrained range of temperatures ($T$~$\\la$~2000~K) and electron densities  (10$^{9}$~$<$~n$_{\\rm e}$~$\\la$~10$^{10}$~cm$^{-3}$).  A comparison of the observed line ratio values to the values predicted by the optically thick and thin local thermodynamic equilibrium cases rules out these conditions for the emitting H~{\\scshape i} gas.  Therefore, the emission is consistent with having an origin in a non-LTE recombining gas.  While the range of electron densities is consistent with the gas densities predicted by existing magnetospheric accretion models, the temperature range constrained by the Case~B comparison is considerably lower than that expected for accreting gas.  The cooler gas temperatures will require a non-thermal excitation process (e.g., coronal/accretion-related X-rays and UV photons) to power the observed line emission. ",
        "introduction": "\\label{intro}  In the decades since their discovery, photometric variability continues to be a defining characteristic of T Tauri stars (TTS) and has been observed at most wavelengths with timescales as short as hours \\citep[e.g.,][]{joy1945,bouv1993}.  In addition to the widely studied photometric variations of TTS, astronomers have begun, more recently, to probe the variability of the unique spectral emission features associated with these sources using multi-epoch spectroscopic observations to note variations in the shapes and line strengths of the emission features \\citep[e.g.,][]{john1995,basr1997,alen2001,bouv2007}.  An investigation of the physical conditions that give rise to these emission features will illuminate many of the processes involved in the formation of these stars.  Single-epoch observations of H~{\\scshape i} Balmer series line emission from TTS reveal a complicated kinematic line structure, with features often including blue-shifted and/or red-shifted absorption components.  As evident from the catalog of H$\\alpha$ features presented in \\cite{reip1996}, a variety of complex profiles exist for this feature and several other Balmer emission lines.  The same spectral line studies reveal quite large full-width at half of maxima (FWHM) for the emission components of these features and suggest that the gas is moving at hundreds of km~s$^{-1}$ relative to the sources.  The complexity of these emission features alone has led to different conclusions about the origins of this emitting gas.  However, the blue-shifted absorption components or P Cygni profiles for these emission features are most likely produced by gas entrained in an outflowing wind \\citep{hart1990,grin1991,calv1992}.  As a result, the observation of a large H$\\alpha$ equivalent width (EW) for an individual source was initially interpreted as evidence of and arising entirely from an active outflowing wind.  In another model of the star-disk interaction, the H~{\\scshape i} lines and other TTS phenomena were formed in a viscous boundary layer, where the accretion disk touched the surface of the star \\citep{lynd1974,bert1989}.  However, the inability of both the boundary layer and stellar wind models to explain the shapes and widths of H~{\\scshape i} emission features led to the application of the magnetospheric accretion model to TTS \\citep{ghos1979,koen1991}.  Previously developed to explain disk accretion onto compact objects, the magnetospheric accretion model established a new paradigm for the star-disk interaction in these pre-main sequence stars.  This model involves a stellar magnetosphere that truncates the disk, opening a gap with a size of a few stellar radii between the disk and the surface of the star.  Looping stellar magnetic field lines thread the inner most regions of these disks, acting as accretion funnels through which disk matter may free fall onto the stellar surface.  The free falling material, confined to a magnetic channel, naturally explains the large velocities observed for the H~{\\scshape i} emission features, as well as the appearance of red-shifted absorption components.  In addition, the accretion shocks formed at the point of impact at the stellar surface provide the heated gas necessary to produce the `blue' continuum excesses characteristic of actively accreting TTS.  Several authors have employed the magnetospheric accretion model to predict the fluxes and shapes of a few of the optical and infrared H~{\\scshape i} emission features \\citep{hart1994,muze1998a,muze1998b}.  These modeling efforts support the idea that the bulk of the emitting H~{\\scshape i} gas is confined to the columns of accreting gas, rather than an outflowing wind.  Advances in infrared spectroscopy have allowed astronomers to observe at high spectral resolution a number of H~{\\scshape i} Paschen and Balmer series emission lines, predominantly Pa$\\beta$ and Br$\\gamma$, from TTS \\citep{naji1996,muze1998a,muze1998b,folh1998,folh2001,natt2006}. These infrared H~{\\scshape i} lines, which possess lower opacities relative to the Balmer series, `conspicuously' lack the blue-shifted absorption features present in H$\\alpha$ and H$\\beta$ profiles \\citep{folh1998}.  The infrared lines also appear to correlate with the mass accretion rates determined from measurements of the `blue' continuum excess \\citep{gull1998,gull2000,muze2001} suggesting that the H~{\\scshape i} emission is related to accretion activity.  Therefore, as a result of the successes of the magnetospheric accretion model in explaining the shapes of the H~{\\scshape i} emission features and the correlation of these features with other accretion diagnostics, the interpretation of the H~{\\scshape i} emission lines and the location of the emitting gas has shifted in the last 15 years.  Where once the presence of a strong H$\\alpha$ emission line in the spectrum of a TTS was widely interpreted as an indicator of an outflowing wind and the `blue' continuum excesses were evidence of a hot boundary layer, today the H~{\\scshape i} features and the `blue' excesses are both interpreted as evidence of magnetospheric accretion.  More recently, spectro-astrometric observations of several classical TTS (cTTS; TTS possessing evidence for actively accreting circumstellar disks) and Herbig AeBe stars possessing symmetric (Type~I) Pa$\\beta$ emission features have shown that some of the emission in the line-wings have measurable offsets with respect to the central sources.  These results suggest that outflowing gas is a likely source for this component of the emission \\citep{whel2004} and potentially explain the difficulty \\cite{folh2001} encountered modeling the line-wings of these emission features with a magnetospheric accretion scenario that did not include a wind.  In addition, AMBER/VLTI observations of Herbig Ae star HD~104237 revealed that the Br$\\gamma$ and {\\it K}-band continuum emission have similar spatial distributions \\citep{tatu2007}.  \\cite{tatu2007} successfully explain the visibility of the Br$\\gamma$ emission by modeling it with an outflowing wind and places the emission at the base of the jet with a radial extent of 0.2 to 0.5~AU.  Observations such as these indicate that a portion of the near-infrared H~{\\scshape i} emission features can be produced in outflowing gas.  Despite these apparent advancements in understanding the origin of the H~{\\scshape i} emission features, the physical conditions of the emitting gas (e.g., ionization state, level populations, temperature, density) remain poorly constrained \\citep{alen2000,whel2004,tatu2007,eisn2007a}.  \\cite{mart1996} calculated a temperature and density profile for gas traveling along an accretion funnel from the disk to the star by modeling several sources of gas heating (e.g., ambipolar diffusion, adiabatic compression, photoionization, and external radiative excitation) and cooling (e.g., Ca {\\scshape ii} and Mg {\\scshape ii} ions).  Assuming an initial gas temperature of 3000~K for gas located at the inner edge of the disk where it is loaded into the accretion funnel, for a fiducial mass accretion rate of $\\dot{M}$~=~10$^{-7}$~M$_\\odot$~yr$^{-1}$, \\cite{mart1996} found that the gas temperature profile peaks at $T$~$\\sim$~6500~K.  In another investigation seeking to constrain the temperature of magnetospherically accreting gas, \\cite{muze2001} calculated H$\\alpha$ and H$\\beta$ line fluxes and profiles dependent upon $\\dot{M}$ and a range of gas temperatures.  The authors found `optimal' temperatures of 6000~$<$~$T$~$<$~12000~K for a range of accretion rates 10$^{-6}$~$\\ge$~$\\dot{M}$~$\\ge$~10$^{-10}$~M$_\\odot$~yr$^{-1}$, with the lower gas temperatures corresponding to sources with higher mass accretion rates.  This study suggested that the gas temperatures predicted by \\cite{mart1996} were too low to produce the observed line luminosities for the $\\dot{M}$ values considered.  These temperatures agree with the standard thermal gas temperature of $T$~$\\sim$~10$^4$~K observed for most H~{\\scshape ii} emission regions.  A direct measurement of the physical conditions of the gas will provide a valuable test of the models and may help to further constrain the origin of the H~{\\scshape i} line emission.  In optically thin gases, line ratios provide valuable information about the physical conditions.  However, at high densities, optical depth effects require radiative transfer modeling and thus complicate the interpretation of these line measurements.  In cTTS, H$\\alpha$ emission lines often appear to be optically thick and may be formed by the contributions of emitting gas in multiple regions of the star-disk system, not limited to the accretion flow alone \\citep{kuro2006}.  The near-infrared H~{\\scshape i} lines, with higher $n$-values, have lower optical depths and appear to be less complicated than the Balmer transitions \\citep{folh2001}. Observational attempts to compare the measured values of the Pa$\\beta$/Br$\\gamma$ H~{\\scshape i} infrared line ratio to the values predicted for optically thin and thick hydrogen gas in LTE, in the past, have failed to determine a temperature for the emitting gas \\citep{muze2001}.  Similarly, comparisons of the Pa$\\beta$/Br$\\gamma$ ratio to the values predicted using the Case~B approximation of hydrogen recombination line theory (\\cite{stor1995} and references therein) have also failed to find a meaningful result.  In fact, the lack of agreement between the observed and predicted line ratios lead to the conclusion that the emitting H~{\\scshape i} gas is neither optically thin or thick nor can it be described as a recombining gas \\cite{muze2001}.  However, in most cases where comparisons were made, the values for this line ratio were determined from spectra collected at different times.  Given that the variability of infrared H~{\\scshape i} line fluxes has timescales as short as hours \\citep{john1995b,alen2001,bouv2007}, a line ratio calculated from two distinct, non-simultaneous spectral observations will likely be misleading.  Therefore, it is not surprising that previous comparisons of a single line ratio have failed to meaningfully constrain the physical conditions of the emitting gas.  Simultaneous observations of not only Pa$\\beta$ and Br$\\gamma$ \\citep[e.g.,][]{gatt2006}, but including several other Paschen and Brackett series emission lines will improve upon these previous comparisons.  We present an analysis of H~{\\scshape i} spectral line data {\\it simultaneously} acquired as part of a multi-epoch near-infrared spectroscopic variability study of cTTS in the Taurus-Auriga star forming region.  Multiple Paschen and Brackett emission features were detected in the spectra of our sample of actively accreting cTTS collected during the three year survey.  We determine the values for 16 distinct line ratios that are representative of the whole sample of TTS.  In \\S\\ref{analysis}, we compare the set of observed ratios, broken into three groups, Paschen series, Brackett series, and Br$\\gamma$/Pa$n_{up}$, where $n_{up}$ is the upper level of the transition, with the temperature and electron density sensitive Case~B models as well as the temperature-only sensitive optically thick and thin LTE scenarios.  The results of these comparisons and their predictions of the physical conditions of the emitting H~{\\scshape i} gas are discussed in \\S\\ref{disc}. ",
        "conclusions": "In this multi-epoch spectroscopic survey, the line ratios from 17 Paschen and Brackett H~{\\scshape i} emission features were measured from a total of 73 spectra of 15 actively accreting cTTS.  The values of the line fluxes measured for the strongest features (e.g., Pa$\\beta$) over the entire sample span almost two orders of magnitude.  In spite of this large variation in line fluxes, the measured values of {\\em all line ratios} were approximately constant, within the observational uncertainties, from source to source and epoch to epoch.  These results are consistent with earlier work that also found little variation in the value of the Pa$\\beta$/Br$\\gamma$ ratio \\citep{muze2001,natt2006}.  Due to the little variation observed in the values of the line ratios, a single value, given by the uncertainty-weighted average of all the data for each ratio, was adopted.  \\subsection{Level Populations and Optical Depth}   The measured values for the H~{\\scshape i} line ratios and those expected for both the optically thick and thin LTE cases for any temperature were not in agreement.  This result is evidence that the $n, l$-level populations of the H~{\\scshape i} gas are {\\bf not} in LTE.  On the other hand, the observed ratios were statistically well-fit by the predictions of the recombination line theory Case~B approximation, tightly constraining the values of $T$ and n$_{\\rm e}$.  This is evidence that the characteristic assumptions of Case~B theory hold for the emitting gas.  In addition to the success of the Case~B models, the measured value of the temperature- and density-independent ratio of Br$\\gamma$/Pa$\\delta$ is consistent with the optically thin value (\\S \\ref{sec_ratio}).  \\subsection{Density}  The observed line ratios are statistically well-fit by Case~B models with an electron density of n$_{\\rm e}$~$=$~$10^{10}$~cm$^{-3}$ (\\S \\ref{sec_contours}).  Furthermore, electron densities of $\\ga~10^{11}$~cm$^{-3}$ and $\\la~10^{9}$~cm$^{-3}$ are statistically ruled out at the 99\\% level or higher (depending on the temperature).  Thus it appears that the range of electron densities allowed for all 15 stars in this survey and for all epochs is constrained to lie within only a factor of a few of $10^{10}$~cm$^{-3}$.  The electron density favored by the Case~B analysis is consistent with that predicted by magnetospheric accretion models \\citep{mart1996, muze1998a}, for typical cTTS parameters and assuming an ionization fraction of $\\sim10^{-2}$$-$10$^{-3}$ (\\S \\ref{sec_flowdensity}).  However, it is a surprise that the allowed range in electron densities is so narrow, given that the range of accretion rates exhibited by the stars in our sample spans more than a factor of 100.  If the H~{\\scshape i} emission arises primarily from material in a magnetospheric accretion flow, this seems to require systematic differences from source to source that regulate the electron density (see \\S\\ref{disc}).   \\subsection{Temperature}   The Case~B statistical analysis rules out electron temperatures of $T \\ge 5000$~K for the emitting H~{\\scshape i} gas, at the 99.9\\% level (\\S \\ref{sec_contours}).  Temperatures of $\\la 2000$~K are favored, and temperatures as low as a few hundred Kelvin cannot be ruled out.  These temperatures are well below the predictions of the magnetospheric accretion models of $6000 \\la T \\la 12000$~K \\citep{muze2001, mart1996}.  However, these models currently do not include all important sources of heating, cooling, and ionization (\\S \\ref{sec_flowtemperature}).  Therefore, the lower gas temperatures we find may still be consistent with H~{\\scshape i} arising from accreting gas.  Independent of the location of the H~{\\scshape i} gas, the Case~B temperature range appears to be too low to produce the observed emission line luminosities, unless a non-thermal process is ionizing the gas (\\S \\ref{sec_ionizing}).  We propose that a significant flux of ionizing photons, arising from the stellar coronae and/or accretion shocks near the stellar surfaces, may simultaneously explain the high observed line luminosities and low temperatures of the emitting H~{\\scshape i} gas."
    },
    "0807/0807.2662_arXiv.txt": {
        "abstract": "{ The mounting evidence for violent nuclear star formation in Seyfert  galaxies has led us to consider the hydrodynamics of the matter reinserted by massive stars through strong stellar winds and supernovae, under the presence of a central massive BH. We show that  in all  cases there is a bimodal solution strongly weighted by the location of the stagnation radius ($R_{st}$), which splits the star cluster into two different zones. Matter reinserted within the stagnation volume is to be accreted by the BH while its outer counterpart would  composed a star cluster wind. The mechanical power of the latter, ensures that there is no accretion of the ISM into the BH and thus the BH accretion and its luminosity is regulated by the star formation feedback. The location of the stagnation  radius is a function of three parameters: the BH mass, the mechanical power (or mass) of the star formation event and the size of the star forming region. Here we present our self-consistent, stationary solution, discuss the accretion rates and BH luminosities and show that our model predicts the intrinsic link between the BH activity and the starburst parameters. ",
        "introduction": "The imprints of massive starbursts around the central supermassive black hole (BH) have been revealed in a number of Seyfert galaxies. Levenson et al. (2001), Jim\\'enez-Bail\\'on et al. (2005) found that in many cases the spectral fit to the X-ray emission from Seyfert galaxies requires a complementary thermal component associated with an extended starburst. Alexander et al. (2005) provided ultra deep Chandra observations of 20 radio-selected submillimeter-emitting galaxies selected from the deep SCUBA surveys and combined this with Keck deep spectroscopy. They found strong evidences that a substantial fraction ($\\sim 75$\\%) of the selected star forming galaxies harbor also an active galactic nucleus (AGN). They also found the average X-ray to the far-IR luminosity ratio an order of magnitude smaller than that in the case of a typical quasar (QSO). This led them to conclude that vigorous star formation activity ($\\sim 1000$ \\Msol \\, yr$^{-1}$) dominates the bolometric luminosity of their AGN-classified sources and that submillimeter-emitting galaxies are associated with a massive black hole (BH) and stellar component rapid growth stage which ends up with the emergence of an un-obscured quasar.  Heckman et al. (1997) and Gonz\\'alez Delgado et al. (1998) found absorption line features associated with photospheres of O and B stars and their stellar winds in the ultraviolet and optical spectra of four Seyfert 2 galaxies that presented direct evidence for the existence of nuclear starbursts in the AGN galaxies.  Terlevich et al. (1990) proposed to use the Ca II triplet as the near-IR indicator of star formation in AGN galaxies. Schweitzer et al. (2006) detected the polycyclic aromatic hydrocarbon (PAH) features, associated with soft UV radiation emerging from the star forming region, in 11 from 26 selected quasars. They found that at least 30\\% but likely much more of the far-IR luminosity in their sample QSOs arises from star formation and concluded that there is a strong connection between the AGN and starburst activity in the QSOs. These results have been confirmed in Netzer et al. (2007) who presented the supplementing spectral energy distributions.  Seth et al. (2007) presented evidences on the presence of super massive black holes in galaxies which host compact nuclear star clusters. Many of these galaxies present a mixed AGN-starburst optical spectra and are classified as composite.  Strong evidences for intense high velocity outflows in the composite Seyfert 2/starburst ultra luminous infrared galaxies have been presented and thoroughly discussed in  Gonz\\'alez Delgado et al. (1998) and Rupke et al. (2005).  Thus star formation occurs at different space and energy scales around the SMBH in many AGN galaxies and QSOs. The mechanical power of young nuclear starbursts might prevent through the cluster winds the accretion of interstellar matter from the bulges and disks of their host galaxies onto the central BHs (see, for example, McLaughlin et al. 2006). Here we show that in such cases the BHs are fed with the matter injected by numerous stellar winds and SNe explosions that result from the multiple evolving sources. This implies that nuclear starbursts must strongly affect and even control the power of the central BH at the stage of vigorous star formation. In fact, it  may be the dominant factor to be consider in order to understand the physics of the BH growth and the relation between the BH and star burst luminosity in AGN galaxies and QSOs. ",
        "conclusions": "Here we discuss the link between the supermassive BH activity and star formation in the central zones of AGN galaxies and QSOs. We briefly formulate the input physics and major results from our (Silich et al. 2008) self-consistent, stationary, spherically symmetric solution for gaseous flows which are formed inside starburst regions with a central supermassive BH.  We show that the thermalization of the kinetic energy released by massive stars inside young stellar clusters results into a bimodal solution with an accretion of the injected matter onto a central BH in the inner zones of the cluster, and a fast outflow of the returned matter from the outer zones of the star forming region.  Our model predicts the intrinsic link between the BH activity and parameters of the starburst region. Specifically, we show that the BH luminosity, $L_{BH}$,  depends on the mass and size of the stellar cluster hosting a BH. It grows with mass of the cluster and is larger in the case of more compact clusters. In the case of extended starbursts, the BH luminosities fall well below the Eddington limit. However, $L_{BH}$ approaches the Eddington value in the case of very compact star forming regions.  Recent results dealing with far-IR emission of QSO and spectral properties of SCUBA galaxies suggest that a continuous mode of star formation is relevant in order to understand the intrinsic links between AGNs and star formation activity. This will be a subject for our forthcoming communication."
    },
    "0807/0807.0983_arXiv.txt": {
        "abstract": "PSR J0205+6449 is a X-ray and radio pulsar in supernova remnant 3C 58.  We report on observations of the central region of 3C 58 using the 4.2-m William Herschel Telescope with the intention of identifying the optical counterpart of PSR J0205+6449 and characterising its pulsar wind nebula. Around the pulsar position we identified extended emission with a magnitude of $B = 23 \\fm 97 \\pm 0.10$, $V = 22 \\fm 95 \\pm 0.05$ and $R = 22 \\fm 15 \\pm 0.03$ consistent with a pulsar wind nebula. From the R-band image we identified three knots with $m_R$ = $24 \\fm 08 \\pm 0.07$ (o1), $24 \\fm 15 \\pm 0.07$ (o2) and $24 \\fm 24 \\pm 0.08$ (o3). We confirm the presence of an optical pulsar wind nebula around PSR J0205+6449 and give an upper limit of $m_R \\approx$ 24 for the optical magnitude of the pulsar. Furthermore we make the tentative suggestion that our object o1, with an $m_R \\approx$ 24.08 is the optical counterpart. If confirmed the pulsar would have an $L_R/L_x\\approx 0.004$ and an optical efficiency of about 5\\% of the Crab pulsar. Such a low efficiency is more consistent with the characteristic age of the pulsar rather than that of SN 1181. ",
        "introduction": "PSR J0205+6449 in supernova remnant (SNR) 3C58  is a recently discovered X-ray and radio pulsar \\citep{ca02,mu02} with a period of 65 ms.  3C 58  was thought to be young due to a possible  association with SN 1181 \\citep{vb78}. Consequently it should share many of the characteristics of the Crab nebula, including the presence of a young pulsar. However an embedded pulsar  defied detection for over twenty years. PSR J0205+6449 has a considerably weaker flux than the Crab pulsar. Its X-ray emission is 1000 times lower than the Crab and its radio emission 120 times lower. Some of this discrepancy can be attributed to the mounting evidence that the characteristic age of the pulsar ($P / 2\\dot{P} $ $\\approx$ 5400 years) is near to its true age and hence it is not associated with SN 1181 \\citep{go07}.  VLA observations at 1.4GHz show a nebula expansion rate of $0.014\\% \\pm 0.003$ yr$^{-1}$ \\citep{bi06},  inconsistent with SN 1181 unless substantial deceleration of the remnant has occurred.  A comprehensive optical study of 3C 58 showed no evidence of an optical pulsar at m$_R \\approx $22.5 \\citep{Fesen08}. Furthermore based upon the relatively low proper motion of knots within 3C 58 these authors also cast doubt on the association between 3C 58 and SN 1181. Deeper optical observations of 3C 58 show evidence of an optical nebulosity at the same location as the X-ray counterpart to PSR J0205+6449 \\citep{Shibanov08}. These authors interpret this nebulosity as a pulsar wind nebula (PWN).  PSR J0205+6449  has the third highest spin down energy flux [$\\dot{E}/d^2 \\approx 2.6~10^{36}$ ergs/s/kpc$^2$] after the Crab and Vela pulsars. The X-ray determined hydrogen column density [$N_H\\approx 4.1~10^{21} cm^{-2}$], \\citep{go07} and optical extinction, E(B-V)  $\\sim$ 0.68 \\citep{Fesen08}, are similar to the Crab pulsar.  From H{\\small I} observations a kinematic distance of 3C 58  has been established at  3.2 kpc  \\citep{ro93}, consistent with its Dispersion Measure \\citep{ca02}. These combine to make  PSR J0205+6449 a likely candidate for optical emission studies. If we assume that the optical luminosity scales with light cylinder magnetic field,  $L_{opt} \\propto B_{lc}^{1.6}$ \\citep{sg02}, then we estimate that the pulsar should have a visual magnitude in the range 23--25, depending on interstellar absorption and effects of beaming geometry \\citep{sh08b}. This paper reports on a William Herschel telescope (WHT) service time observation of  PSR J0205+6449 to look for evidence of an optical pulsar and its possible PWN. ",
        "conclusions": "From these observations we confirm the presence of an optical PWN around PSR J0205+6449 and identify possible structure within the nebula. Although from these observations we can not definitively identify the optical counterpart of PSR J0205+6449 we can give a deep upper limit and make a tentative suggestion that our object o1 is  the counterpart. If it is shown that o1 is the counterpart then we can estimate, using \\citet{fg94}, the flux from the pulsar to be 0.71 $\\mu$Jy in the range $\\lambda\\lambda$ 6000-7000 \\AA, corresponding to a R-band luminosity of $7.3~10^{30}$ ergs sec$^{-1}$ based upon an E(B-V)=0.68 giving A$_R \\approx$ 1.8. This luminosity gives an efficiency  $L_R/\\dot{E}$ of $\\approx 3~10^{-7}$ and a $L_R/L_x\\approx 0.004$. The optical efficiency is about 5\\% of the Crab pulsar's optical efficiency and comparable to the efficiency of older optical pulsars such as Vela and PSR B0656+14 \\citep{sg02}. The optical -- X-Ray luminosity ratio is relatively high at 0.004, but this is more of a reflection of the low X-ray luminosity. From these tentative observations we can conclude that the optical emission, if confirmed, is consistent with a pulsar with an age similar to its spin down age rather than that of a very young pulsar. Final optical identification will have to wait for observations sensitive to flux variations on millisecond time scales using  instruments such as Optima \\citep{st01} or GASP \\citep{co08}."
    },
    "0807/0807.2639_arXiv.txt": {
        "abstract": "Cosmological constant problem (in its various versions) is arguably the deepest gap in our understanding of theoretical physics, the solution to which may very likely require revisiting the Einstein theory of gravity. In this letter, I argue that the simplest consistent way to decouple gravity from the vacuum energy (and hence solve the problem) is through the introduction of an incompressible {\\it gravitational aether} fluid. The theory then predicts that gravitational constant for radiation is $33\\%$ larger than that of non-relativistic matter, which is preferred by most cosmological observations (with the exception of light element abundances), but is not probed by current precision tests of gravity. I also show that slow-roll inflation can happen in this theory, with only minor modifications.  Finally, interpreting gravitational aether as a thermodynamic description of gravity, I propose a finite-temperature correction to the equation of state of gravity, which would explain the present-day acceleration of the cosmic expansion as a consequence of the formation of stellar mass black holes. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0807/0807.2313_arXiv.txt": {
        "abstract": "{Galactic open clusters are since long recognized as one of the best tools for investigating  the radial distribution of iron and other metals.} {We employed FLAMES at VLT to collect UVES spectra of bright giant  stars in a large sample of open clusters, spanning a wide range of Galactocentric distances, ages, and metallicities. We present here the results for four clusters: Berkeley~20 and Berkeley~29, the two most distant clusters in the sample;  Collinder~261, the oldest and the one with the minimum  Galactocentric distance; Melotte~66.} {Equivalent width analysis was carried out using the spectral code MOOG and Kurucz model atmospheres to derive abundances of Fe, Al, Mg, Si, Ca, Ti, Cr, Ni, Ba; non-LTE Na abundances were derived by direct line-profile fitting.} {We obtain subsolar metallicities for the two anticenter clusters Be~20 ([Fe/H]=$-$0.30, rms=0.02) and Be 29 ([Fe/H]=$-$0.31, rms=0.03), and for Mel~66 ([Fe/H]=$-$0.33, rms=0.03), located in the third Galactic quadrant, while Cr~261, located toward the Galactic center, has higher metallicity ([Fe/H]=+0.13, rms=0.05 dex). The $\\alpha$-elements Si, Ca and Ti, and the Fe-peak elements Cr and Ni are in general close to solar; the s-process element Ba is enhanced. Non-LTE computations of Na abundances indicate solar scaled values, suggesting that the enhancement in Na previously determined in giants in open clusters could be due to neglected non-LTE effects.} {Our results support the presence of a steep negative slope of the Fe radial gradient up to about 10--11 kpc from the Galactic center, while in the outer disk the [Fe/H] distribution seems flat. All the elemental ratios measured are in very good agreement with those found for disk stars of similar metallicity and no trend with Galactocentric distance seems to be present.} ",
        "introduction": "The radial metallicity gradient -- i.e. the behaviour of [Fe/H] with Galactocentric radius ($R_{\\rm gc}$) --  provides a powerful observational constraint on models of Galactic disk formation and chemical evolution  (e.g., Tosi \\cite{tosi96}; Twarog, Ashman, \\& Anthony-Twarog \\cite{twarog}; Freeman \\& Bland-Hawthorn \\cite{freeman}). At the same time, the ratios of the abundances of different elements to Fe, are fundamental for understanding the role of stars with different masses in the Galactic chemical enrichment, and thus to get insights on the initial  mass function (IMF) at the epoch of disk formation.  Open clusters are widely recognized as ideal targets for the investigation of chemical abundances and their evolution, due to their internal homogeneity in age and composition, to the fact that their distance is in principle  easy to measure, and to their broad coverage of ages and Galactocentric distances. For this reason, many spectroscopic determinations of abundances in open clusters have been carried out during the last years. Among them we cite  those based on low resolution spectroscopy by Friel (\\cite{friel95}) and Friel et al.~(\\cite{friel02}) which focused on metallicity ([Fe/H]),  and those based on high resolution, by Friel et al. (\\cite{friel03}; \\cite{friel05})  Carretta et al. (\\cite{carretta04}, \\cite{carretta05}), Carraro et al. (\\cite{carraro04}, \\cite{carraro6791}, \\cite{carraro07}),  Yong, Carney, \\& de Almeida (\\cite{yong05}), most of which present the analysis of several other elements besides iron.  In spite of these observational efforts, various problems/discrepancies persist, especially regarding the abundance gradients. For example, a consensus on the shape of the gradient has not been reached yet. Early works suggested an unique negative slope (e.g., Friel \\cite{friel95}; Carraro, Ng, \\& Portinari \\cite{carraro98}, Friel et al.~\\cite{friel02}); but Twarog et al.~(\\cite{twarog}) and Corder \\& Twarog (\\cite{corder})  favor a step-like distribution of the Fe content with Galactocentric distance. Very recent results (e.g., Carraro et al. \\cite{carraro04}, \\cite{carraro07}; Yong et al. \\cite{yong05}) suggest that the gradient might become flat for $R_{\\rm gc}$ larger than about 14 kpc.  In this context we have carried out two VLT/FLAMES programs, aimed at deriving the gradient, based on the homogeneous analysis of a large sample of open clusters. With the fiber link to UVES we observed giant stars in order to obtain accurate element abundances based on high resolution data. The first program was based on Guaranteed Time Observations (GTO; Pallavicini et al. \\cite{pallavic06}) and included three open clusters, while within the other one (Randich et al.~\\cite{messenger}) nine clusters were observed. The targets span a wide range of ages (from $\\sim$1 Gyr up to $\\sim$7 Gyr), metallicities ([Fe/H] from $\\sim-0.5$ up to $\\sim$+0.40)  and Galactocentric radii ($\\sim$6--22 kpc);  In Sestito et al. (2006 -- hereafter Paper~{\\sc i}) and Bragaglia et al. (2008 -- hereafter Paper~{\\sc ii}), the analysis of Fe and other elements in five of the observed clusters (NGC~2660, NGC~2324, NGC~2477, NGC~3960, and Be~32) was presented; in   Sestito et al. (\\cite{ngc6253} -- hereafter S07) we reported on abundances in the very metal rich cluster NGC~6253. In this paper we present the analysis of the two far, old, anticenter clusters in the sample by Randich et al.~(\\cite{messenger}; Berkeley~20, and Berkeley~29), and of two old clusters observed within the GTO program (Cr~261 and Mel~66; Pallavicini et al. \\cite{pallavic06}).  Be 20 and Be 29 are the most distant clusters in the sample, while Cr 261 is the one with the lowest  $R_{\\rm gc}$. For all these clusters there are already papers based on high resolution spectra (see next Section), but  we improve on past studies both on statistics (for most of them only two stars were analyzed) and S/N ratio.  The paper is organized as follows: in Sect.~2 we describe the observational samples, data reduction, and abundance analysis; the results are reported in Sect.~3 and discussed in Sect.~4; a summary and conclusion are given in Sect.~5. ",
        "conclusions": "\\label{discussion}  In Table ~\\ref{cluster_summary} we give a summary of the properties and abundances of some key elements in the open clusters analyzed by our group (this paper, Paper~{\\sc i},~{\\sc ii} and S07).   The Galactocentric radii are from Friel et al. (2002), when present; otherwise they were taken from Friel (1995), Friel (2006).  Ages are from Bragaglia \\& Tosi (\\cite{bt06}), Bragaglia et al. (\\cite{bragaglia06}), Tosi, Bragaglia, \\& Cignoni (\\cite{tosi07}). The left panel of Fig.~\\ref{gradiente} shows the distribution of [Fe/H]  as a function of the Galactocentric distance: the sample includes our data and other clusters analyzed with high resolution spectroscopy\\footnote{ Literature high resolution studies: Smith \\& Suntzeff~(\\cite{smith87}: NGC~2420), Gratton \\& Contarini~(\\cite{gratton94}: NGC~2243), Brown et al.~(\\cite{brown96}: Mel~71, To~2), Hamdani et al.~(\\cite{hamdani00}: NGC~2360, NGC~2447), Gonzalez \\& Wallerstein (\\cite{gonzalez00}: M11), Bragaglia et al. (\\cite{bragaglia01}: NGC~6819), Randich et al.~(\\cite{randich01}: IC~2602, IC~2391; \\cite{randich03}: NGC 188; \\cite{randich06}: M67), Sestito et al.~(\\cite{sestito03}: NGC~6475), Paulson et al.~(\\cite{paulson03}: Hyades), Schuler et al.~(\\cite{schuler03}: M34; \\cite{schuler04}: Pleiades), Carretta et al.~(\\cite{carretta04}: IC~4651, NGC~2506, NGC~6134), Carraro et al.~(\\cite{carraro04}: Saurer~1; \\cite{carraro07}: Be~25, Be~73, Be~75, Rup~4, Rup~7), Ford et al.~(\\cite{ford05}: Blanco~1), Tautvaisiene et al. (\\cite{taut05}: NGC~7789), Villanova et al.~(\\cite{villanova05}: Be~22, Be~66),Yong et al.~(\\cite{yong05}: Be 31, NGC 2141), Pace et al. (in prep.: Praesepe). The point representative of NGC~6791 (the most metal rich cluster) is an average between the [Fe/H] values obtained by Carraro et al.~(\\cite{carraro6791}), Gratton et al.~(\\cite{gratton06}), and Origlia et al.~(\\cite{origlia}). We included in the plot only clusters older than $\\sim$ 100 Myr; in the cases of clusters in common between us and other authors, we chose our own metallicity determination. Most [Fe/H] values are based on the analysis of giant stars, except those for the young clusters IC~2602, IC~2391, Pleiades, Blanco~1, NGC~6475, M34, Hyades, and for the old clusters M67 and NGC~188.}. The right panel of Fig.~\\ref{gradiente} shows the radial [Fe/H] distribution for open clusters observed at low resolution by Friel at al.~(\\cite{friel02}). The [Fe/H] values in the two panels of  Fig.~\\ref{gradiente} are taken directly from the literature sources: thus, whereas the low resolution data by Friel et al. represents  a homogeneous sample,  our data and the other high resolution ones from the literature could be on different scales. Whereas this might introduce some spurious scatter, some general features are evident in the figure. If one considers high resolution data, there appears to be a quite steep metallicity gradient up to a Galactocentric distance $R_{\\rm{gc}}\\lesssim10-11$ kpc, and then a flattening at larger distances. More in detail, using a weighted linear fit, we find a slope of $-0.17\\pm0.02$ dex kpc$^{-1}$ (with a correlation coefficient $-0.83$) in the  disk region with $R_{\\rm{gc}}\\leq11$ kpc, while in the outer disk ($R_{\\rm{gc}}>11$ kpc) the slope is, within the errors, consistent with zero. The sample by Friel et al. (2002) is instead limited to $R_{\\rm{gc}}\\lesssim16$ kpc, with only one cluster in the distance range $R_{\\rm{gc}}\\sim14-16$ kpc: these authors found a negative slope of $-0.063\\pm0.010$ dex kpc$^{-1}$ between $\\sim$7 and 16 kpc, and they did not discuss the possibility of a flattening of the gradient outside of a given radius. Considering only the clusters at $R_{\\rm{gc}}\\leq11$ kpc, the [Fe/H] gradient from low resolution data has a slope of $-0.09\\pm0.02$ dex kpc$^{-1}$. The difference between this value of the slope in the inner disk and that found from high resolution data  within the same radius ($R_{\\rm{gc}}\\leq11$ kpc) is important, since even small differences in the slope might lead to changes  in the infall rate and star formation efficiency adopted in the models. This point will be discussed in more detail in a future paper of the series, where we plan to compare our empirical results with updated Galactic chemical evolution models. On the other hand, considering only high resolution data, our sample and the one from the literature give similar results;  in particular, the possibility of a  plateau in the Fe gradient for clusters in the outer disk was already evidenced by Yong et al.~(\\cite{yong05}) who studied Be~20 and Be~29 (also in our sample), and two clusters at $R_{\\rm{gc}}\\sim12$ kpc, Be~31 and NGC~2141, as well as by Carraro et al.~(\\cite{carraro04}, \\cite{carraro07}).  Our clusters significantly improve the statistics, confirming on a more solid basis the change in the slope of the gradient, and in particular of a flattening for radii larger than $\\sim$11 kpc; the average metallicity outside of this Galactocentric distance is $-0.27\\pm0.13$. However, we have to remember that the \"inner\" slope is based on a  short baseline (about 4 kpc),  since the OC closest to the Galactic centre has $R_{\\rm GC}$ slightly less than 7 kpc, and that the number or \"outer\" OCs is still scarce. To improve the situation, we should concentrate on deriving the metallicity of clusters   both nearer to the Galactic centre and in the outer part of the disk, and in the transition region between the two slopes. The growing number of  photometric studies dedicated to overlooked  OCs (e.g., Carraro et al. \\cite{carraro06}, \\cite{carraro07}, Ortolani et al. \\cite{ortolani}) and the new catalogues of candidate clusters (e.g., Froebrich et al.  \\cite{froebrich}, Kronberger et al. \\cite{kronberger})  will provide a good database for future abundance studies.    Figure \\ref{confronto_disco} shows [X/Fe] ratios vs. [Fe/H] for the open clusters in our sample and for field disk stars. We show the $\\alpha$-elements Si, Ca, Ti and the Fe-peak elements Cr and Ni. We considered the disk star catalogue by  Soubiran \\& Girard~(\\cite{soubiran}), which is a compilation and rehomogenization of several recent literature studies, but we also included the data by Bensby et al.~(\\cite{bensby}), since the latter authors derived also Cr abundances,  not included in the other catalogue. It is evident that the open clusters follow the same distribution of chemical abundances with [Fe/H] of disk stars, with average values close to (but slightly larger than) solar. Note that the $\\alpha$-elements  have a certain amount of scatter, in particular Ti.  Figure \\ref{grad_altri} shows the radial gradients for $\\alpha$-elements, Fe-peak elements and the s-process element Ba, based on our sample and the other high resolution data from the literature. The plots indicate some amount of dispersion at any given  Galactocentric distance (this is especially true for Ba),  with no obvious trends, with the only exception of  Ca: in this case there seems to be a gradient with a positive slope (0.08$\\pm$0.02, with correlation coefficient 0.64)  in the inner disk.  The interpretation of the distributions of these elements is rather complex and it is strictly related to the role of SNe{\\sc I}a and SNe{\\sc II} in the Galactic chemical enrichment (see, e.g. Yong et al. 2005). The comparison between our data and Galactic chemical evolution models,  planned for a future paper of the series,  will be aimed to distinguish the main contribution  in element production by different nucleosynthesis process.  Whereas we find nearly solar [X/Fe] ($\\alpha$ and Fe-peak elements)  in all the clusters, Yong et al. (2005) and Carraro et al. (2004) claim that the outer disk clusters Sau~1, Be~20 and Be~29 (the two latter in common with the present study) have  an enhancement in the $\\alpha$-elements. Note, however, that Yong et al. (2005) derive [Si+Ca/Fe]$\\sim$+0.1, i.e.  values similar to those we find  for some stars, and which are not much higher than solar, and have been also found in  field stars of similar metallicity ; on the other hand, they derive [Mg+Ti/Fe]$\\sim$+0.2--0.3. Also the values  for the $\\alpha$-elements found by Carraro  et al. (2004) are actually only slightly oversolar  (average [$\\alpha$/Fe] ratios of less than 0.1 for Be~29 and about 0.15  for Sau~1);  anyway, in the most recent analysis of five outer disk clusters Carraro et al. (\\cite{carraro07}) do seem to converge on solar-scaled [$\\alpha$/Fe] ratios, similar to open clusters in the solar vicinity and thin disk stars. With the present data, we do not think there is a strong argument for a significantly increased [$\\alpha$/Fe] at large Galactocentric distances.  The only s-process element analyzed by us is Ba, which is clearly enhanced for our open clusters. At the Fe abundances  of our clusters, the Ba production is dominated by  the s-processes in low mass stars  (see e.g. Travaglio et al.~\\cite{travaglio01} and Cescutti  et al.~\\cite{cescutti06} for  a complete discussion about this element).  The Ba enhancement, mainly in the younger clusters, might   be explained by its high production in low mass stars, mainly from 1 to  2 M$_{\\odot}$, that were formed from $\\sim$ 7 to 2 Gyr ago (e.g. Charbonnel et al.~\\cite{charbonnel96}).  These stars contribute to the Ba enrichment exactly in the  metallicity range covered by the examined clusters,  while at lower metallicities the main contribution  is from r-processes which are less effective  (see also Fig.~4 by Cescutti et al.~\\cite{cescutti06}). However, Ba abundances should be treated with caution, since the hyperfine structure could be important, as a preliminary analysis on available data seems to show (D'Orazi and Randich, in preparation), contrary to what found e.g., by Mashonkina and Gehren (2001); this would decrease the abundances. Furthermore,  Ba abundances (ours and literature ones) show a scatter much higher than for any other element here considered, suggesting that  we need a careful appraisal of their reliability before attempting meaningful interpretations."
    },
    "0807/0807.0857_arXiv.txt": {
        "abstract": "Organic compounds are ubiquitous in space: they are found in diffuse clouds, in the envelopes of evolved stars, in dense star-forming regions, in protoplanetary disks, in comets, on the surfaces of minor planets, and in meteorites and interplanetary dust particles.  This brief overview summarizes the observational evidence for the types of organics found in these regions, with emphasis on recent developments. The {\\it Stardust} sample-return mission provides the first opportunity to study primitive cometary material with sophisticated equipment on Earth. Similarities and differences between the types of compounds in different regions are discussed in the context of the processes that can modify them.  The importance of laboratory astrophysics is emphasized. ",
        "introduction": "Organic matter is defined as chemical compounds containing carbon-hydrogen bonds of covalent character, i.e., with the carbon and hydrogen forming a true chemical bond. Observations over the last century have established that these molecules are ubiquitous throughout the universe, not only in our Galaxy (\\cite{kwok07b}) but even out to high redshifts (\\cite{yan05}). With the detection of more than 200 exoplanetary systems, a major question is whether these organic compounds can be delivered in tact to new planetary systems where they could form the basis for the origin of life. The answer requires a good understanding of the entire lifecycle of organic molecules from their formation in the outflows of evolved stars to the diffuse interstellar medium (ISM) and subsequently through the star-forming clouds to protoplanetary disks.  Ehrenfreund \\& Charnley (2000) and Ehrenfreund et al.\\ (2002) summarized our understanding of this cycle several years ago, but since then new observations (e.g., with the {\\it Spitzer Space Telescope}), laboratory experiments, and in-situ space missions (in particular the {\\it Stardust} mission) have occurred. Thus, a symposium such as this reviewing these recent developments is timely.  In this opening paper, a broad overview will be given where organic matter is found in space and which species have been identified. Subsequently, observational evidence of how and where these molecules are modified will be summarized. Finally, a number of questions to address at the symposium and in the future are raised.  The importance of laboratory astrophysics in providing the basic data to interpret astronomical and solar-system observations and analyse meteoretic and cometary samples is emphasized. ",
        "conclusions": ""
    },
    "0807/0807.2549_arXiv.txt": {
        "abstract": " ",
        "introduction": "Scaling relations between the intrinsic properties of galaxies provide clues to the physical mechanisms which operate within these systems.  In general, the tighter a relation is, i.e.\\ the less scatter it has, the more fundamental the relation is expected to be.  Therefore, it is perhaps not surprising that there has been a huge interest in the $M_{\\rm bh}$-$\\sigma$ relation (Ferrarese \\& Merritt 2000; Gebhardt et al.\\ 2000) which was reported to have very little or no intrinsic scatter.  In addition to providing an indirect means to measure supermassive black hole (SMBH) masses in many galaxies, the $M_{\\rm bh}$-$\\sigma$ relation (Merritt \\& Ferrarese 2001a; Tremaine et al.\\ 2002; Ferrarese \\& Ford 2005; Novak et al.\\ 2006), along with the equally strong $M_{\\rm bh}$-$n$ relation (Graham \\& Driver 2007a) and $M_{\\rm bh}$-$L$ relation (Kormendy \\& Richstone 1995; Magorrian et al.\\ 1998; McLure \\& Dunlop 2002; Marconi \\& Hunt 2003; Graham 2007), provides insight into the joint formation process of SMBHs and their host bulges.  Such relations can also be applied to volume-limited galaxy samples, providing an estimate of the SMBH mass density in the Universe (e.g., Salucci et al.\\ 1999; Graham \\& Driver 2007b, and references therein).  As the number of galaxies with direct SMBH mass measurements has increased, it has become possible to explore the demographics of the SMBH population within the $M_{\\rm bh}$-$\\sigma$ diagram.  Rather than delineating a single line, Graham (2008) and Hu (2008) have revealed a tendency for SMBHs in barred galaxies, or perhaps equivalently pseudobulges, to reside below the $M_{\\rm bh}$-$\\sigma$ relation defined by non-barred galaxies.  In addition, Hu (2008) has noted that there may be a third subdivision in the $M_{\\rm bh}$-$\\sigma$ diagram such that ``core galaxies'' (Ferrarese et al.\\ 1994; Faber et al.\\ 1997; Trujillo et al.\\ 2004) define a steeper relation than non-core galaxies.  Such departures from a single unifying expression offer the promise of further valuable clues into the coevolution of galaxies and the million to billion solar mass black holes which reside at their centres.  This paper presents the largest sample of galaxies for which direct SMBH mass estimates are available.  While the structure within the updated $M_{\\rm bh}$-$\\sigma$ diagram is explored here, it is additionally hoped that this database will be a helpful resource, or rather stepping stone, for future investigations. ",
        "conclusions": ""
    },
    "0807/0807.2852_arXiv.txt": {
        "abstract": "The NASA Discovery mission EPOXI, utilizing the Deep Impact flyby spacecraft, comprises two phases: EPOCh (Extrasolar Planet Observation and Characterization) and DIXI (Deep Impact eXtended Investigation). With EPOCh, we use the 30-cm high resolution visible imager to obtain ultraprecise photometric light curves of known transiting planet systems. We will analyze these data for evidence of additional planets, via transit timing variations or transits; for planetary moons or rings; for detection of secondary eclipses and the constraint of geometric planetary albedos; and for refinement of the system parameters. Over a period of four months, EPOCh observed four known transiting planet systems, with each system observed continuously for several weeks. Here we present an overview of EPOCh, including the spacecraft and science goals, and preliminary photometry results. ",
        "introduction": "The EPOXI mission, led by PI Michael A'Hearn and Deputy PI Drake Deming, is a NASA Discovery Program mission of opportunity, designed to be a low cost, low risk, scientifically focused investigation. The mission comprises of two distinct science projects: Extrasolar Planet Observation and Characterization (EPOCh) and the Deep Impact eXtended Investigation (DIXI). Both investigations make complementary use of the Deep Impact flyby spacecraft, an existing and flight-proven space asset that was previously used to observe comet Tempel 1. The EPOCh project is currently observing known transiting planet systems to obtain ultraprecise photometric light curves and is the focus of this contribution.  The aim of the EPOCh project is to observe a carefully selected set of known, bright, transiting planet systems for several weeks at a time, obtaining light curves with very high precision, phase coverage and cadence. Using these data we can investigate a variety of science goals, including: refining system parameters; searching for additional bodies in the system via transit timing variations in the transits of the known planet (\\cite[Agol et al. 2005]{Agol05}; \\cite[Holman \\& Murray 2005]{Holman05}) or detections of additional transits (\\cite[Croll et al. 2007a,b]{Croll2007a}); searching for reflected light at secondary eclipse and constraining the planetary albedo (Rowe et al. and Matthews et al., these proceedings); and searching for the signatures of moons (\\cite[Doyle \\& Deeg, 2004]{Doyle04}, \\cite[Brown et al. 2001]{Brown01}) or rings (\\cite[Barnes \\& Fortney, 2004]{Barnes04}) associated with the known planet. Another aspect of the EPOCh project was the characterization of Earth as an extrasolar planet (Deming et al. 2007); only the transiting planet observations and results are discussed here. ",
        "conclusions": "Our preliminary analysis has thus far focused on the HAT-P-4 dataset; the initial investigations and further work are briefly described below.  \\subsection{Refinement of system parameters}  We have used the analytic equations describing transiting light curves of \\cite[Mandel \\& Agol (2002)]{Mandel02} and the Monte Carlo Markov chain (MCMC) method outlined for radial velocity measurements by \\cite[Ford (2005)]{Ford05} and adapted to transit light curve measurements by, e.g., \\cite[Holman et al. (2006)]{Holman06} to analyze the light curve of HAT-P-4. We used the four-parameter non-linear limb-darkening law of \\cite[Claret (2004)]{Claret04}, with the limb-darkening coefficients derived from a weighted fit across the instrument bandpass, interpolated across the closest theoretical stellar models\\footnote{\\texttt http://kurucz.cfa.harvard.edu/} (\\cite[Kurucz 1994]{Kurucz94}, \\cite[Kurucz 2005]{Kurucz05}). The phase-folded binned light curve is shown in Figure \\ref{fig:phase} with the best-fit solution over-plotted. Further investigation of the correlation of the errors in the parameters is required, however our results are thus far consistent with those presented in the discovery paper (\\cite[Kov{\\'a}cs et al. 2007]{Kovacs07}).  \\begin{figure} \\begin{center} \\includegraphics[angle=90,width=5.7in]{figs/phase.eps} \\caption{Phase-folded HAT-P-4 data, overplotted with the best-fit analytic solution from \\cite[Mandel \\& Agol (2002)]{Mandel02}, with fixed limb-darkening coefficients derived for the instrument bandpass. The data are binned by a factor of 20 and the error bars are calculated using the out-of-transit standard deviation. The residuals to the solution are shown underneath, with a standard deviation of 0.4~mmag.} \\label{fig:phase} \\end{center} \\end{figure}  \\subsection{Reflected light analysis}  We have analyzed the data around secondary eclipse in order to search for the signature of reflected light and to put constraints on the geometric planetary albedo. The {\\em MOST} satellite has been used very successfully to put tight constraints on this parameter for HD~209458b and HD~189733b (\\cite[Rowe et al. 2006, see also the Rowe et al. contribution to these proceedings]{Rowe06}). The upper panel of Figure \\ref{fig:eclipse} shows the phased HAT-P-4 data around the predicted times of secondary eclipse (set to a phase of 0.0), binned by a factor of 80. The best-fit eclipse model gives an upper limit on the geometric albedo of 0.346, below that of Jupiter, 0.6. Further work allowing the predicted time of eclipse to vary owing to the range in eccentricity allowed by the radial velocity measurements is still required.  \\begin{figure} \\begin{center} \\includegraphics[angle=90,width=5.4in]{figs/eclipse.eps} \\caption{Searching for the secondary eclipse of HAT-P-4. {\\em Upper}: The data around secondary eclipse (set to a phase of 0.0) are shown here, binned by a factor of 80. The solid line shows a model with a secondary eclipse depth inconsistent with the data at a 3$\\sigma$ significance, giving an upper limit on the planet-to-star flux ratio of $7.1\\times10^{-5}$. {\\em Lower}: The change in chi-squared from the best-fit value, with the vertical dashed line delineating the allowed upper limit on the depth.} \\label{fig:eclipse} \\end{center} \\end{figure}  \\subsection{Transit timing variations}  The timestamps for the HAT-P-4 data have been corrected to the barycentric Julian date (BJD), however the relative offset between the clock on-board the spacecraft and UT on Earth is yet to be applied. As a result, we can conduct an independent period search and transit timing analysis within the EPOCh dataset, however at this stage we cannot combine these results with those of ground-based observations. We find a period of 3.056370$\\pm$0.000094 days, and Figure \\ref{fig:resids} shows the residuals of the individual transit times to this period. When the timing offset has been applied, we will analyze the long-term period and transit timing behaviors.  \\begin{figure} \\begin{center} \\includegraphics[angle=90,width=5.5in]{figs/resids.eps} \\caption{Residuals of the measured center of each individual transit of HAT-P-4 to the best-fit period.} \\label{fig:resids} \\end{center} \\end{figure}"
    },
    "0807/0807.2255_arXiv.txt": {
        "abstract": "We describe the Submillimeter Array (SMA) Polarimeter, a polarization converter and feed multiplexer installed on the SMA. The polarimeter uses narrow-band quarter-wave plates to generate circular polarization sensitivity from the linearly-polarized SMA feeds. The wave plates are mounted in rotation stages under computer control so that the polarization handedness of each antenna is rapidly selectable. Positioning of the wave plates is found to be highly repeatable, better than 0.2 degrees. Although only a single polarization is detected at any time, all four cross correlations of left- and right-circular polarization are efficiently sampled on each baseline through coordinated switching of the antenna polarizations in Walsh function patterns. The initial set of anti-reflection coated quartz and sapphire wave plates allows polarimetry near 345 GHz; these plates have been have been used in observations between 325 and 350 GHz. The frequency-dependent cross-polarization of each antenna, largely due to the variation with frequency of the retardation phase of the single-element wave plates, can be measured precisely through observations of bright point sources. Such measurements indicate that the cross-polarization of each antenna is a few percent or smaller and stable, consistent with the expected frequency dependence and very small alignment errors. The polarimeter is now available for general use as a facility instrument of the SMA. ",
        "introduction": "\\label{sec:intro} As the first dedicated submillimeter interferometer, the Submillimeter Array (SMA) provides unique capabilities for detailed study of the dusty sites of star formation. In particular, measurements of polarized dust emission, which traces the structure of the magnetic field in the plane of the sky ($B_\\perp$), benefit greatly from higher frequencies and finer angular resolution. The flux density from optically thin warm dust emission rises like $\\nu^{3-4}$, multiplying the emission by factors of several between the 230~GHz band of millimeter interferometers such as PdBI and CARMA and the 345~GHz band of the SMA. Furthermore, structure is expected in the magnetic field on scales as small as $\\sim$100~AU in star forming regions, corresponding to $\\sim$1\\arcsec\\ in even the nearest objects. Previous observations of polarized dust emission have generally been made at lower frequency and few arcsecond resolution or higher frequencies with no better than 10\\arcsec\\ resolution. The SMA provides the first opportunity to examine polarization at both high frequencies and high angular resolution. Illuminating the important and poorly understood role of the magnetic field in the star formation process provides a strong motivation for adding a polarimetric capability to the SMA. In addition, there are other targets, such as the polarization signatures of low-luminosity accretion\\cite{MarroneE06}, that are accessible only near submillimeter frequencies.  The SMA was designed to allow polarimetry, but requires additional hardware to enable these observations. Most importantly, the SMA can measure only a single linear polarization at a given frequency, although this will change for some receivers in future upgrades. Therefore, modulation of the polarization sensitivity of each antenna is required in order to recover complete polarization information. Moreover, the interferometric measurement of linear polarization is best made with circularly-polarized receivers (see, e.g., ref.~\\citenum{URP2}), so a linear-to-circular polarization conversion is desirable. The instrument described here provides the polarization conversion and rapid modulation that is required for these observations. We outline the instrument design, control system, calibration procedures and results, and a provide examples of the science enabled by the polarimeter. ",
        "conclusions": "\\label{sec:concl} We have presented the 345~GHz polarimeter instrument for the Submillimeter Array. This instrument now provides a fully-supported polarization observing mode for the SMA, used in roughly 10\\% of observations with the array. Several upgrades to the polarimeter are planned or have been completed since its deployment. As part of a new calibration system, the installation and removal of the wave plates was automated in 2007. Additional wave plates for 230/690~GHz and 240/400~GHz have been added; commissioning of these plates is ongoing. Over the coming year, the SMA will gain a new 340-420~GHz receiver that will provide dual-polarization observations near 345~GHz. This obviates the need for rapid switching of the antenna polarization state, although some polarization modulation is still desirable to prevent errors in the orientation of the derived polarization due to inter-receiver phase drift. The decrease in the switching overhead and the instantaneous measurement of all four polarized cross correlations will improve the sensitivity by a factor of more than $\\sqrt{5}$ for this band, greatly increasing the number of sources accessible to the array."
    },
    "0807/0807.0299_arXiv.txt": {
        "abstract": "It is widely known that cosmic magnetic fields, i.e. the                % fields of planets, stars, and galaxies, are produced by the hydromagnetic dynamo effect in moving electrically conducting fluids. It is less well known that cosmic magnetic fields play also an active role in cosmic structure formation  by enabling outward transport of angular momentum in accretion disks via the magnetorotational instability (MRI). Considerable theoretical and computational progress has been made in understanding both processes. In addition to this, the last ten years have seen tremendous efforts in studying both effects in liquid metal experiments. In 1999, magnetic field self-excitation was observed in the large scale liquid sodium facilities in Riga and Karlsruhe. Recently, self-excitation was also obtained in the French ''von K\\'{a}rm\\'{a}n sodium'' (VKS) experiment. An MRI-like mode was found on the background of a turbulent spherical Couette flow at the University of Maryland. Evidence for MRI as the first instability of an hydrodynamically stable flow was obtained in the ''Potsdam Rossendorf Magnetic Instability Experiment'' (PROMISE). In this review,  the history of dynamo and MRI related experiments is delineated, and some directions of future work are discussed. ",
        "introduction": "Magnetism has been known for approximately 3000 years. There is some evidence that a hematite bar, found close to Veracruz (now Mexico), had served the Olmecs as a simple compass \\cite{CARLSON}. In any case, the Chinese have built, probably in the first century B.C., a compass  in the form of a lodestone spoon that was freely turnable on a polished bronze plate \\cite{NEEDHAM}. The old Greek philosophers, starting with Thales of Miletus \\cite{ARISTOTLE}, were well aware of the attracting forces of lodestone, and the Roman philosopher Lucretius (95?-55 B.C.) described its action in an atomistic way \\cite{LUKREZ}: \"First, stream there must from off the lode-stone seeds. Innumerable, a very tide, which smites by blows that air asunder lying betwixt the stone and iron. And when is emptied out this space, and a large place between the two is made a void, forthwith the primal germs of iron, headlong slipping, fall conjoined into the vacuum, and the ring itself by reason thereof doth follow after and go thuswise with all its body.\"   \\begin{vchfigure}[] \\begin{center} \\includegraphics[width=0.7\\linewidth]{fig1_new.eps} \\vchcaption{Lodestone experiments of Petrus Peregrinus. (a) Experiments showing the attracting and repelling forces of a broken piece of loadstone. Upper configuration: \"naturale\", lower configuration: \"non naturale\". (b) A simple compass. Figures from \\cite{PEREGRINUS}. } \\label{fig1} \\end{center} \\end{vchfigure}  As early as 1269, a first systematic experimental study of the attracting and repelling forces of lodestone was published  by Petrus Peregrinus in his \"Epistola de magnete\" \\cite{PEREGRINUS}. For the first time, he defined the concept of polarity and distinguished the north and south poles of the magnet. He was the first to formulate the law that poles of opposite polarity attract while poles of the same polarity repel each other (cf. Figure 1a). Besides the construction of several compasses (cf. Figure 1b), he also proposed a magnetic perpetuum mobile.  Three centuries later, Peregrinus' work inspired William Gilbert to make his own experiments with small spheres of lodestone (``terrellae''), which led him, in 1600, to the conclusion that ``...that the terrestrial globe is magnetic and is a loadstone \\cite{GILBERT1600}''  However, this lodestone theory soon ran into trouble when the westward drift of the Earth's magnetic field declination was described by Gellibrand in  1635 \\cite{GELLIBRAND},  and the detection of abrupt polarity reversals by David and Brunhes in 1904/05 \\cite{BRUNHES} has dealt it the ultimate deathblow.   Interestingly, it was not the well-studied magnetic field of the Earth, but the observation of magnetic fields in sunspots \\cite{HALE}, that put Larmor on the right track speculating  \\cite{LARMOR} that it could be \"...possible for the internal cyclic motion to act after the manner of the cycle of a {\\it self-exciting dynamo}, and maintain a permanent magnetic field from insignificant beginnings, at the expense of some of the energy of the internal circulation.''  This one-page communication, in which a natural process was explained in terms of a technical device \\cite{SIEMENS,SIMONYI,WHEATSTONE}, was the birth certificate of the modern hydromagnetic theory of cosmic magnetic fields. ",
        "conclusions": "The last ten years have seen tremendous progress in liquid metal experiments on the origin and the action of cosmic magnetic fields. With the success of the complementary sodium experiments in Riga and Karlsruhe in 1999 it was shown that self-excitation works not only in computer programs on the basis of rather smooth flows, but also in real-world turbulent flows. Kinematic dynamo theory has been shown to be robust with respect to low levels of turbulence and complicated boundary and interface conditions. The saturation mechanism in the Riga experiment is non-trivial as it results not only from a global pressure increase but also from a significant redistribution of the flow.   While the Riga and Karlsruhe dynamos are characterized by an amazing predictability, in some respect one can learn even more from the efforts to make the VKS dynamo (and also the MDX dynamo) running. After some modifications of the original concept, the VKS experiment was eventually successful in self-exciting a magnetic field. The fact that the observed eigenfield in VKS is essentially an axisymmetric dipole, in contrast to the original prediction of an equatorial dipole, is an inspiring challenge to understand better the intricate field amplification loop the dynamo mechanism relies on. The ''little brother'' of the VKS experiment, the VKG experiment in Lyon, has yielded fascinating field reversals when complemented by an external amplification loop that mimics a sort of $\\alpha$ effect.   On the way to the final rotating torus experiment, the Perm group has obtained important results concerning the mean-field coefficients in turbulent flows.  After dynamo action has thus been proven, one observes presently some tendency to take a breath and turn back to somewhat smaller machines in order to study MHD instabilities and wave phenomena. The identification of the MRI is only one aspect in this direction, though an important one due to the enormous astrophysical implications of this instability. Apart from standard MRI (Maryland, Princeton, New Mexico) and helical MRI (Dresden-Rossendorf) there are other magnetic instabilities that are capable of destabilizing hydrodynamically stable flows or even fluids at rest. Among them we have to note the Taylor-Vandakurov instability (with a current flowing through the liquid) \\cite{VANDAKUROV,TAYLER,SPRUIT} and the ''azimuthal MRI'' (AMRI) \\cite{RUEDIGERHOLLERBACHSCHULTZ} based on an purely azimuthal field. Future experiments on those instabilities are very desireable.  There are many wave phenomena in rotating fluids under the influence of (externally applied or self-excited) magnetic fields which still  deserve a deeper understanding. At this point we observe a revival of the  activities in the sixties on  Alfv\\'{e}n wave studies with liquid metals. In this respect, it might also be interesting that presently magnetic fields are available \\cite{HOCHFELD} which are so strong that for potassium and sodium the Alfv\\'{e}n velocity exceeds the sound velocity. The small scale experiments in Maryland and Grenoble have provided a wealth of data, and the 3 m sphere in Maryland will be a fascinating tool for extending those investigations into extreme parameter regions.  An interesting direction of future research could be the experimental investigation of precession driven dynamos, for which a water test experiment had been carried out by L\\'{e}orat et al. \\cite{LEORAT1,LEORAT2}. Numerical work by Tilgner \\cite{TILGNERPRECESSION} points to a critical $Rm$ of around 200 which makes a laboratory experiments not unrealistic. Of course, much more optimization would be necessary before such an experiment could be designed.   If one had a wish for free, one could  think of constructing a huge self-sustaining nonlinear dynamo \\cite{RINCON} (or ''self-creating dynamo'' \\cite{RAEDLERFUCHS}). One could start, for example, in the hydrodynamic stable regime of a Taylor-Couette flow (mimicking a Keplerian flow) which could be destabilized, via MRI, by an externally applied magnetic field. Suppose now that the resulting flow would act as a dynamo, the resulting eigenfield could possibly replace the initially applied magnetic field as a trigger for the instability.   Having another wish for free, one could also think about the experimental realization of a fluctuation dynamo, based on more or less homogeneous isotropic turbulence \\cite{SCHEKOCHIHIN}. For those flows, critical $Rm$ around 200 have been found in recent simulations.   % However, simple Kolmogorov scaling arguments tell us that such a ''James-Bond-dynamo'' (''...shaken, not stirred...'') is only possible with a giant input power of many Megawatts.   In spite of the fact that  liquid metal experiments should not be expected to be perfect Bonsai models of any real astrophysical systems, experimental work has already started to change some views on those natural systems. Maybe that the Riga dynamo will once become a model of the double helix nebula close to the galactic center \\cite{MORRIS,SHUKUROVNEBULA}, maybe that the partitioning  of $\\alpha$ and $\\Omega$ effects in the VKS dynamo will re-animate the old Babcock-Leighton theory of the solar dynamo \\cite{BABCOCK}, maybe that  the PROMISE experiment will illuminate the possible  action of helical MRI in cold parts of accretion disks were standard MRI cannot work.             \\begin{acknowledgement} Support from Deutsche Forschungsgemeinschaft in the framework of SFB 609, from European Commission under contract 028679, and from Leibniz-Gemeinschaft in the framework of the SAW project PROMISE is gratefully acknowledged. We are indebted to Raul Avalos-Zu\\~{niga}, Michael Christen, Andr\\'{e} Giesecke, Thomas Gundrum, Uwe G\\\"unther, Rainer Hollerbach, Sa\\v{s}a Kenjere\\v{s}, Jacques L\\'{eorat}, Olgerts Lielausis, Ernests Platacis, J\\={a}nis Priede, Karl-Heinz R\\\"adler, G\\\"unther R\\\"udiger, Luca Sorriso-Valvo, Jacek Szklarski, and Mingtian Xu for the fruitful collaboration over the course of many years.  \\end{acknowledgement}"
    },
    "0807/0807.0250_arXiv.txt": {
        "abstract": "Observing extrasolar planetary transits is one of the only ways that we may infer the masses and radii of planets outside the Solar System. As such, the detections made by photometric transit surveys are one of the only foreseeable ways that the areas of planetary interiors, system dynamics, migration, and formation will acquire more data. Predicting the yields of these surveys therefore serves as a useful statistical tool. Predictions allows us to check the efficiency of transit surveys (``are we detecting all that we should?'') and to test our understanding of the relevant astrophysics (``what parameters affect predictions?''). Furthermore, just the raw numbers of how many planets will be detected by a survey can be interesting in its own right. Here, we look at two different approaches to modeling predictions (forward and backward), and examine three different transit surveys (TrES, XO, and Kepler). In all cases, making predictions provides valuable insight into both extrasolar planets and the surveys themselves, but this must be tempered by an appreciation of the uncertainties in the statistical cut-offs used by the transit surveys. ",
        "introduction": "Planetary transits are one of the many methods by which extrasolar planets have been discovered, and are also the one that provides the most complete set of information about the planetary system. Only planets with very specific orbital characteristics have a transit visible from Earth, because the orbital plane has to be aligned to within a few degrees of the line of sight. Therefore transiting planets are rare. Nevertheless, a transiting extrasolar planet offers the opportunity to determine the mass of the planet - when combined with radial velocity (RV) measurements - since the inclination is now measurable, as well as the planetary radius, the density, the composition of the planetary atmosphere, the thermal emission from the planet, and many other properties (see \\cite[Charbonneau et al. (2007)]{charbonneau2007} for a review). Additionally, and unlike RV surveys, transiting planets should be readily detectable down to $1 R_\\oplus$ and beyond, even for relatively long periods.  Having accurate predictions of the number of detectable transiting planets is immediately important for the evaluation and design of current and future transit surveys. For the current surveys, predictions allow the operators to judge how efficient are their data-reduction and transit detection algorithms. Future surveys can use the general prediction method that we describe here to optimize their observing set-ups and strategies. More generally, such predictions allow us to test different statistical models of extrasolar planet distributions. As more transiting planets are discovered, these statistical properties are increasingly becoming the frontier of research, shifting the focus of the field away from individual detections. ",
        "conclusions": "The field of transiting planets is transitioning from pure discovery, wherein every new transiting planet is newsworthy, to the details. Individual planets, aside from the smallest ones being discovered, are becoming less interesting in and of themselves, and more interesting for what they can tell us about the statistical properties of planets. Given that transiting planets are so far one of the only ways that we may infer the radius of extrasolar planets and their exact orbital properties, the statistical characteristics of this group of objects is one of the only foreseeable ways that the areas of planetary interiors, system dynamics, migration, and formation will acquire more data.  That being said, it is often hard to draw specific statistical conclusions from the results of the transit surveys. In B\\&G we were able to adopt hard cut-offs on our simulated detections, but the transit surveys typically do not adopt the strict detection criteria that we have used. Promising planet candidates are often followed up even if they are beyond the stated S/N or magnitude limits of the survey. Understanding and quantifying how the survey teams select candidates is vital to appropriately deriving the statistical properties of extrasolar planets. Indeed, the results of calculations shown here and in B\\&G demonstrate that the actual predictions depend crucially on the specific magnitude and S/N threshold used.  In the future, besides using the results of B\\&G to design more efficient transit surveys, I would hope that a more systematic approach towards transit surveys will allow the model to be used to make more specific statistical comparisons. As the number of known transiting planets grows, we would also expect that the model will be used to test different distributions of planet frequencies, periods, and radii against those observational results. This will allow us to better understand the statistics of extrasolar planetary systems, improve our ability to find new planets, and help to understand the implications of the ones that have already been detected."
    },
    "0807/0807.3703_arXiv.txt": {
        "abstract": "Quasars are thought to be powered by supermassive black holes accreting surrounding gas\\cite{Salpeter64,Lynden-Bell69,Shields78}. Central to this picture is a putative accretion disk which is believed to be the source of the majority of the radiative output\\cite{Lynden-Bell69,Shields78,Malkan83may}.  It is well known, however, that the most extensively studied disk model\\cite{Shakura73} --- an optically thick disk which is heated locally by the dissipation of gravitational binding energy --- is apparently contradicted by observations in a few major respects\\cite{Antonucci99,Koratkar99}. In particular, the model predicts a specific blue spectral shape asymptotically from the visible to the near-infrared\\cite{Shakura73,Hubeny00}, but this is not generally seen in the visible wavelength region where the disk spectrum is observable\\cite{Neugebauer87,Cristiani90,Francis91,Zheng97,VandenBerk01}. A crucial difficulty was that, toward the infrared, the disk spectrum starts to be hidden under strong hot dust emission from much larger but hitherto unresolved scales, and thus has essentially been impossible to observe.  Here we report observations of polarized light interior to the dust-emiting region that enable us to uncover this near-infrared disk spectrum in several quasars.  The revealed spectra show that the near-infrared disk spectrum is indeed as blue as predicted.  This indicates that, at least for the outer near-infrared-emitting radii, the standard picture of the locally heated disk is approximately correct. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0807/0807.3229_arXiv.txt": {
        "abstract": "{ We present results of simultaneous dual-frequency (2 GHz and 8 GHz) very long baseline interferometry (VLBI) observations of 12 active galactic nuclei with prominent jets.  Spectral properties of the jets and evolution of their brightness temperature are discussed. Measured sizes and brightness temperatures of VLBI features are found to be consistent with emission from relativistic shocks dominated by adiabatic losses. Physical scenarios with different magnetic field orientation in the jets are discussed. ",
        "introduction": " ",
        "conclusions": "Measured sizes and brightness temperatures of VLBI components in BL Lac object $1823+568$ and quasar CTA~102 are consistent with emission from relativistic shocks dominated by adiabatic losses."
    },
    "0807/0807.1706_arXiv.txt": {
        "abstract": "Prompted by high resolution observations, I propose an explanation for the 40+ year old problem of structure and energy balance in the solar transition region. The ingredients are simply cross-field diffusion of neutral atoms from cool threads extending into the corona, and the subsequent excitation, radiation and ionization of these atoms via electron impact. The processes  occur whenever chromospheric plasma is adjacent to coronal plasma, and are efficient even when ion gyro-frequencies exceed collision frequencies. Cool threads - fibrils and spicules perhaps -  grow slowly in thickness as a neutral, ionizing front expands across the magnetic field into coronal plasma. Radiative intensities estimated for H L$\\alpha$ are within an order of magnitude of those observed, with no ad-hoc parameters - only thermal parameters and geometric considerations are needed.  I speculate that the subsequent dynamics of the diffused material might also explain observed properties of trace elements. ",
        "introduction": "\\label{sec:introduction}  The upper transition region (henceforth, ``TR'') - plasma with electron temperatures in the range $2\\times 10^5 \\la T_e \\la 10^6$ K, is adequately described by field-aligned thermal conduction down from the corona.  The lower TR ($10^4$ K $< T_e < 2\\times 10^5$ K) however, is not so easily understood \\citep{Gabriel1976,Jordan1980b}.  Models dominated by field-aligned heat conduction produce too little emission from the lower TR by orders of magnitude, a problem already evident in work by \\citet{Athay1966}. Neither could such models radiate away the downward directed conductive flux of $F_{cond}\\sim 10^6$ \\flxu{} \\citep[e.g.][]{Jordan1980b,Athay1981}. \\nocite{Fontenla+Avrett+Loeser2002} Fontenla {\\em et al.} (2002 and earlier papers in the series), henceforth ``FAL'' showed that energy balance can be achieved through field-aligned (1D) diffusion of neutral hydrogen and helium atoms.  The neutral atoms diffuse into hot regions, radiate away much of the coronal energy, and can reproduce the H and He line intensities.  The problem might be considered by some as solved, in principle. But there exists the serious and nagging problem of the peculiar spatial relationship between the observed corona, TR and chromosphere \\citep{Feldman1983}. Feldman and colleagues have since analyzed many observations, concluding that the lower TR is thermally disconnected from the corona \\citep[e.g.][and references therein]{Feldman+Dammasch+Wilhelm2001}. Yet \\citet{Fontenla+Avrett+Loeser1990} declared that ``The above [i.e. their] scenario explains why (as noted by Feldman 1983) the \\nocite{Feldman1983} structure of the transition region is not clearly related to the structures in the corona''. That the debate still rages is evidenced by advocates for  ``cool loop'' models in which lower TR radiation originates from loops never reaching coronal temperatures and having negligible conduction \\citep[][and references therein, henceforth ``PGV'']{Patsourakos+Gouttebroze+Vourlidas2007}. Here I propose a different scenario, prompted by new data and analyses which show that neither cool loops nor field-aligned processes adequately describe the L$\\alpha$ chromospheric network. I speculate that other TR lines might also be accounted for. ",
        "conclusions": "Based upon observations of spicules and other fine, thread-like structures on the solar disk, it is clear that non-planar thermal interfaces exist at the base of the corona, and that the morphology of the TR emission from such interfaces cannot be explained by field-aligned particle transport at the base of coronal loops, in contrast to the  claims by \\citet{Fontenla+Avrett+Loeser1990}. The picture proposed here uses unspecified mechanisms in the chromosphere to maintain a reservoir of cool mostly neutral plasma directly adjacent to hot coronal plasma.  The cylindrical geometry, inspired by observations, presents a large surface area (per unit volume) of contact between cool and hot plasma.  The chromosphere supplies mass via neutral diffusion across the surface to a thermal boundary layer, and the corona supplies energy to the neutral particles. The originally neutral particles drain energy from the corona by latent heat of ionization and by inelastic collisions leading to strong L$\\alpha$ emission. The diffusing layer propagates outwards, emitting radiation like the boundary of a wild fire\\footnote{Secchi in 1877 described the chromosphere as a ``burning prairie'', but in a different sense.}, into the corona until either the supply of neutral mass or coronal energy dries up.  The present proposal is related to models invoking cross-field heat conduction \\citep{Rabin+Moore1984,Athay1990}.  This effect is included here (via $d_s^E$), but it is far less efficient at moving heat to cool plasma than  diffusion is at moving neutral atoms to the coronal heat.  The calculations presented here fall short of accounting for the large radiative flux of L$\\alpha$, by factors of $\\approx 10$.  However, the calculations miss important additional sources of energy in the corona: thermal and gravitational potential energy.  The cool threads extend only a few Mm into the corona, and form just the lower parts of a much larger coronal structure.  The diffused cool material is thus subject to parallel transport (heat conduction, diffusion) which will transfer heat from the overlying coronal plasma to the diffused material.  Spicules formed by ejection from the chromosphere will have their entire length exposed to this energy flux, because the lowest parts of the spicules diffuse first into the corona- the diffusion fronts are not exactly parallel to field lines.  Coronal plasma along connected field lines contains $L \\frac{3}{2} nkT$ erg~cm$^{-2}$, where $L$ is the pressure scale height ($\\sim 50$Mm) or loop length. Since $L \\gg l_c$ the energy available for L$\\alpha$ radiation would be $L/l_c \\ge 10\\times$ larger than computed above, more if the tubes expand with height.  I speculate that cross-field diffusion and subsequent parallel conduction {\\em might bring theoretical and observed intensities values into agreement.}  The time needed to conduct this energy must lie between the electron sound speed $c_e$ as $L/c_e \\sim 13$ s, and $\\sim 10^3$ s, an upper limit obtained from the thermal energy divided by the conductive flux for a uniform temperature gradient.  Gravitational potential energy might contribute to the heating and dynamics of the sheath as the diffused material cools the corona and adds mass, such that vertical pressure balance no longer is expected.  It may be that larger red-shifts would be expected where magnetic fields are more vertical, i.e. directly over the magnetic network.  This expectation is not in disagreement with results found by \\citet{McIntosh+others2007}.  However, little more can be said without solving the 2D multi-fluid conservation equations including parallel heat conduction and cross-field diffusion, beyond the scope of this letter.  Such calculations will also show if the emission lines of trace species (ions of carbon, oxygen etc. in the TR) can be explained.  Cool threads are observed in different coronal environments (PGV)- their intensities appear to vary relatively little compared with the embedding coronal intensities. This fact is part of Feldman's (1983) \\nocite{Feldman1983} claim that TR emission is energetically disconnected from the corona.   The calculations presented here indeed produce a non-linear relationship between L$\\alpha$ and coronal brightness. The L$\\alpha$ intensities scale with the local coronal energy density and with the diffusion speed. But the EUV and X-ray radiation emitted by the corona itself vary with density$^2$ and peaked functions of temperature along lines of sight different from the direction of field lines into the sheath. The scenario might therefore explain most of the observed puzzling facets noted by Feldman and colleagues, yet still maintain a strong energetic link between the corona and TR, and thereby resolve a long-standing debate (see the different perspectives of \\citealp{Feldman+Dammasch+Wilhelm2001} and \\citealp{Wikstol+Judge+Hansteen1998}, for example).  To see if the scenario survives scrutiny, more observations of chromospheric fine structure and its relation with the corona and TR would be as important as numerical modeling work.  \\noindent I am grateful to Tom Holzer, Scott McIntosh and the referee for comments.  \\def\\aspcs{{ASP Conf.\\ Ser.}}"
    },
    "0807/0807.1530_arXiv.txt": {
        "abstract": "The recently discovered transiting very hot Jupiter, HAT-P-7b, a planet detected by the telescopes of HATNet, turned out to be among the ones subjected to the highest irradiation from the parent star. As known, the combination of photometric and spectroscopic data for such an object yields the stellar, orbital and planetary parameters. In order to best characterize this particular planet, we carried out a complex analysis based on a complete and simultaneous Monte-Carlo solution using all available data. We included the discovery light curves, partial follow-up light curves, the radial velocity data, and we used the stellar evolution models to infer the stellar properties.  This self-consistent way of modeling provides the most precise estimate of the a posteriori distributions of all of the system parameters of interest, and avoids making assumptions on the values and uncertainties of any of the internally derived variables describing the system. This analysis demonstrates that even partial light curve information can be valuable. This may become very important for future discoveries of planets with longer periods -- and therefore longer transit durations -- where the chance of observing a full event is small. ",
        "introduction": "The transiting extrasolar planet candidate HTR154-011 has been identified by the telescopes of the HATNet (HAT-7, HAT-8 and HAT-6, HAT-9) -- see \\cite{bakos2002} for a detailed description about the instrumentation. The reduction of the frames followed the standard frame calibration procedures, followed by star detection, astrometry and aperture photometry which was supported by various detrending algorithms. We search the light curves for box-shaped transit signals using the BLS algorithm (\\cite{kovacs2002}): a periodic dip was detected in the light curve of 2MASS~19285935+4758102. The initial ``rejection mode'' spectroscopy with CfA Digital Speedometer (\\cite{latham1992}) at FLWO eliminated the possibility of an eclipsing binary star system. The analysis of Keck HIRES spectra resulted the radial velocity data, stellar atmospheric parameters and the bisector information (which ruled out a ``false positive'' hierarchical triple). We also carried out a partial photometry follow-up in the Sloan $z$-band with KeplerCam on the FLWO~1.2m telescope. See \\cite{pal2008} for more details about the observational details and data reduction.  The basics of the characterization of transiting extrasolar planets are the following. The period of the orbital motion can be very accurately determined because of the long timespan of the discovery (HATNet) observations. The mass of the planet is proportional to the amplitude of the RV curve, $K$. The radius of the planet relative to the star determines the depth of the transit (which is resulted both by the follow-up light curves and the HATNet data). The inclination of the orbit can be derived from the shape of the transit light curve, i.e. the impact parameter is proportional to the cosine of the inclination. To calculate the absolute mass and radius, we need the same data or the host star too, therefore we need at least one luminosity indicator for a good stellar evolution model. Because of the lack of parallax, this indicator can be $a/R_\\star$, see \\cite{sozzetti2007}. The bottleneck in this case is the partial follow-up, i.e. the shape parameters of the light curve -- namely the ingress/egress length relative to the total duration -- cannot be derived directly to obtain a reasonable value of $a/R_\\star$. ",
        "conclusions": "The outlined method provides a straightforward way for characterizing planetary systems since (i) it omits any kind of intermediate values which otherwise would distort the error estimations, (ii) automatically results the uncertainties and correlations of the system parameters, including the stellar, planetary and orbital parameters too, (iii) it can utilize even the partial light curves which otherwise would not provide enough information for further steps  -- these partial follow-up light curves become even more important when systems with longer periods will be discovered (i.e. the chance to observe a complete transit on a single night will be definitely smaller). The planet HAT-P-7b is also very interesting since (i) it has a very high surface temperature, somewhere between $2200-2700$~K, depending on the heat redistribution, (ii) it happens to fall in the field of view of the upcoming Kepler Mission, yielding an opportunity for further continuous observations."
    },
    "0807/0807.2828_arXiv.txt": {
        "abstract": "The formation of resonant planets pairs in exoplanetary systems involves planetary migration inside the protoplanetary disc\\,: an inwards migrating outer planet captures in Mean Motion Resonance an inner planet. During the migration of the resonant pair of planets, the eccentricities are expected to rise excessively, if no damping mechanism is applied on the inner planet. We express the required damping action to match the observations, and we show that the inner disk can play this role. This result applies for instance to the system GJ~876\\,: we reproduce the observed orbital elements through a fully hydrodynamical simulation of the evolution of the resonant planets. ",
        "introduction": "Most of the known exoplanets are giant planets on short period orbits around their stars. In the core-accretion model, such planets should have formed much further away, beyond the so-called ``snow-line'', where water is solid and contributes to the formation of a massive solid core. This suggests that inward planetary migration occured in the protoplanetary disc. Through mutual gravitationnal interactions, a planet and the gaseous disc exchange energy and angular momentum. For a giant planet, this lead to the opening of a gap and to the migration of the planet (\\cite{LinPapa86}).  A stronger evidence for migration is that a large fraction of multiplanet systems are engaged a Mean Motion Resonance (MMR)\\,: the ratio of the orbital periods of the outer to the inner planet, equals that of two small integers. This cannot happen by chance. It betrays a convergent migration of the planets and a dissipative process to enter this dynamical configuration.  For instance, in GJ~876, the two most massive planets are in a 2:1 MMR, with periods of $\\approx 30$ and $\\approx 60$ days. To model the formation of this system, \\cite{LeePeale} performed 3-body simulations of a central host star and two planets, with additional (dissipative) forces that reproduced the effects of disc-planet interaction. Only the outer planet was forced to migrate inward (assuming that the inner planet has no ambient gas). Capture into resonance occured when the outer planet crosses the location at which the mean orbital periods have a ratio of 2. After the resonant capture, however, the eccentricities of the two planets rise dramatically (except if an irrealistically high eccentricity damping is applied to the outer planet). ",
        "conclusions": "This work shows that the eccentricity of resonant giant exoplanets is governed by the presence of a gas disc inside the orbits of the innermost one. For given eccentricities, resonances and planet masses, the required action of the inner disc is provided analytically. In the GJ~876 system, the radius of the inner edge of the disc can be estimated thanks to the presence of a third planet\\,; a numerical simulation of this system reproduces very well the observed parameters, which confirms the validity of the model.  In conclusion, the observed moderate eccentricities of the exoplanets in resonance suggest that in these systems, an inner disc was present during the migration.  This work has also led to the publication of an article in A\\&A \\cite{CSK08}."
    },
    "0807/0807.0646_arXiv.txt": {
        "abstract": "The clustering properties of sterile neutrinos are studied within a simple extension of the minimal standard model, where these neutrinos are produced via the decay of a gauge singlet scalar. The distribution function after decoupling is strongly out of equilibrium and features an enhancement at small comoving momentum $\\propto 1/\\sqrt{p}$. Dark matter abundance and phase space density constraints from dwarf spheroidal galaxies constrain the mass in the $\\mathrm{keV}$ range consistent with a Yukawa coupling to a gauge singlet with mass and vacuum expectation value in the range $\\sim 100\\,\\mathrm{GeV}$ and a decoupling temperature of this order. The dark matter transfer function and power spectrum are obtained from the solution of the non-relativistic Boltzmann-Vlasov equation in the matter dominated era. The small momentum enhancement of the non-equilibrium distribution function leads to long range memory of gravitational clustering and a \\emph{substantial enhancement of the power spectrum at small scales as compared to a thermal relic  or sterile neutrino produced via non-resonant mixing with active neutrinos}. The scale of suppression of the power spectrum for a sterile neutrino with $m\\sim \\mathrm{keV}$ produced by scalar decay that decouples at $\\sim 100 \\mathrm{GeV}$ is $\\lambda \\sim 488 \\,\\mathrm{kpc}$. At large scales   $T(k)\\sim 1-C \\,  k^2/k^2_{fs}(t_{eq})  +\\cdots$ with $C \\sim \\mathrm{O}(1)$. At small scales $65\\,\\mathrm{kpc} \\lesssim \\lambda \\lesssim 500\\,\\mathrm{kpc}$  corrections to the fluid description and memory of gravitational clustering become important, and we find $T(k) \\simeq 1.902\\,e^{-k/k_{fs}(t_{eq})}$, where $k_{fs}(t_{eq}) \\sim 0.013/\\mathrm{kpc}$ is the free streaming wavevector at matter-radiation equality. The enhancement of power at small scales may provide a possible relief to the tension between the constraints from X-ray and Lyman-$\\alpha$ forest data. ",
        "introduction": "In the \\emph{concordance} $\\Lambda\\mathrm{CDM}$ standard cosmological model  dark matter (DM) is composed of primordial particles which are cold and collisionless\\cite{primack}.  In this  cold dark matter (CDM) scenario structure formation proceeds in a hierarchical ``bottom up'' approach: small scales become non-linear and collapse first and their merger and accretion leads to structure on larger scales. CDM particles feature negligible small velocity dispersion  leading to a power spectrum that favors small scales. In this hierarchical scenario, dense clumps that survive the merger process form satellite galaxies.  Large scale simulations seemingly yield  an overprediction of satellite galaxies\\cite{moore2} by almost an order of magnitude larger than the number of satellites that have been observed in Milky-Way sized galaxies\\cite{kauff,moore,moore2,klyp,will}. Simulations within the $\\Lambda$CDM paradigm also yield a density profile in virialized (DM) halos that increases monotonically towards the center\\cite{frenk,dubi,moore2,bullock,cusps} and features a cusp, such as the Navarro-Frenk-White (NFW) profile\\cite{frenk} or more general central density profiles $\\rho(r) \\sim r^{-\\beta}$ with $1\\leq \\beta \\lesssim 1.5$\\cite{moore,frenk,cusps}. These density profiles accurately describe clusters of galaxies but there is an accumulating body of observational evidence\\cite{dalcanton1,van,swat,gilmore,salucci,battaglia,cen,wojtak} that seem to indicate that the     central regions of (DM)-dominated dwarf spheroidal satellite (dSphs) galaxies feature smooth cores instead of cusps as predicted by (CDM). More recently\\cite{cen} a ``galaxy size'' problem has been reported, where large scale simulations at $z=3$ yield galaxies that are too small, this problem has been argued to be related to that of the missing dwarf galaxies.  Warm dark matter (WDM) particles were invoked\\cite{mooreWDM,turokWDM,avila} as possible solutions to these discrepancies both in the over abundance of satellite galaxies and as a mechanism to smooth out  the cusped   density profiles predicted by (CDM) simulations into the  cored profiles that fit the observations in   (dShps). (WDM) particles feature a range of velocity dispersion in between the (CDM) and hot dark matter (HDM) leading to free streaming scales that smooth  out small scale features and could be consistent with core radii of the (dSphs). If the free streaming scale of these particles is  smaller than the scale of galaxy clusters, their large scale structure properties are indistinguishable from (CDM) but may affect the  \\emph{small} scale power spectrum\\cite{bond} so as to provide an explanation of the smoother inner profiles of (dSphs), fewer satellites and the size of galaxies at $z=3$\\cite{cen}.  Sterile neutrinos with masses $\\sim \\mathrm{keV}$ may be suitable (WDM) candidates\\cite{dw,este,shapo,kusenko,kusepetra,petra}. The main property that is relevant for structure formation  of any dark matter candidate is its distribution function after decoupling\\cite{hogan,coldmatter}, which depends on the production mechanism and the (quantum) kinetics of its evolution from production to decoupling. There is a variety of mechanisms of sterile neutrino production\\cite{dw,colombi,este,shapo,kuse2}, and mixing between sterile and active neutrinos can be one of them\\cite{dw,colombi,este}. There is considerable tension between the X-ray\\cite{xray} and Lyman-$\\alpha$ forest\\cite{lyman,lyman2} data if sterile neutrinos are produced via the Dodelson-Widrow (DW)\\cite{dw} non-resonant mixing  mechanism, leading to the suggestion\\cite{palazzo} that these cannot be the dominant (DM) component. Constraints from the Lyman-$\\alpha$ forest spectra are particularly important because of its sensitivity to the suppression of the power spectrum by free-streaming in the linear regime\\cite{lyman,lyman2}. The most recent constraints from the Lyman-$\\alpha$ forest\\cite{lyman2} improve upon previous ones, but rely on the Dodelson-Widrow\\cite{dw} model for the distribution function of sterile neutrinos, leaving open the possibility of evading these tight constraints with non-equilibrium distribution functions from other production mechanisms.  The gravitational clustering properties of collisionless (DM) in the linear regime are described by the   power spectrum of gravitational perturbations. Free streaming of collisionless (DM)  leads to a suppression of the transfer function on length scales smaller than the free streaming scale via Landau damping\\cite{bond,bondszalay,kt}. This scale is   determined by the decoupling temperature, the particle's mass and the distribution function at decoupling\\cite{freestream}.  In this article we study the gravitational clustering properties of sterile neutrinos as (WDM) candidates within a model in which these are produced via the decay of a gauge singlet scalar. Such model has been advocated recently in ref.\\cite{kusepetra,petra,kuse2}, and is based on a   phenomenologically appealing extension of the minimal standard model\\cite{shapo} consistent with the observed neutrino masses and mixing. This model also shares many features in common with  models of gravitino production\\cite{rubakov}, hence it provides a   viable extension of the standard model to study in detail the production and clustering  properties of potential (WDM) candidates. In this model sterile neutrinos   decouple a temperatures much larger ($\\sim 100 \\,\\mathrm{GeV}$ )  than in the Dodelson-Widrow scenario ($\\sim 150 \\,\\mathrm{MeV}$)\\cite{dw}, therefore they are   \\emph{colder} at matter-radiation equality (and today), being dubbed, for this reason,  ``chilled'' neutrinos in refs.\\cite{kusepetra,petra}.  Clustering properties of active neutrinos in non-standard cosmology, for example quintessence have been reported in ref.\\cite{pettorino}.  A program that yields a quantitative assessment of a particle physics candidate for (DM) in the linear regime implements the following steps:  \\begin{itemize} \\item{Establish the quantum kinetic equations that describe the production of these particles and follows the evolution of their distribution function through their decoupling from the cosmological plasma.}  \\item{The distribution function after decoupling becomes the \\emph{unperturbed} distribution, which determines the abundance,   the primordial phase space densities\\cite{hogan,coldmatter} and free streaming lengths\\cite{freestream}. The generalized  Tremaine-Gunn\\cite{TG} constraints  obtained in\\cite{coldmatter} in combination with the recent photometric observations of the phase space densities of (dSphs) combined with the  DM  abundance lead to bounds on the mass, couplings and decoupling temperature\\cite{coldmatter}.   }  \\item{The unperturbed distribution function is input in the Boltzmann-Vlasov equation for density and gravitational perturbations\\cite{ma}.  The solution of which yields the transfer function, and the power spectrum. }    \\end{itemize}     We follow this program in the   model of sterile neutrino production proposed in refs.\\cite{shapo,kusepetra,petra,kuse2}. In principle, in order to obtain the transfer function and the power spectrum of density perturbations the coupled set of Boltzmann equations for baryons, photons, dark matter and gravitational perturbations must be solved\\cite{ma,dodelson}. Photons and baryons are coupled by Thompson scattering and dark matter only couples to the gravitational perturbations that are sourced by all the components. In practice this is a computationally daunting task because popular codes\\cite{lewis} based on the set of coupled Boltzmann equations for photons, baryons and dark matter\\cite{ma} need to be modified to input arbitrary non-equilibrium distribution functions, masses and couplings.  Recently a simple analytic framework to obtain the dark matter transfer function, and consequently the power spectrum during matter domination has been presented\\cite{gildan}. The main premise of this formulation is that the contribution from baryons and photons modifies the  DM  transfer function at most by a few percent\\cite{EH,HS} during matter domination and that a preliminary robust assessment of the clustering properties of a DM candidate can be systematically established by neglecting in first approximation the contribution from baryons and photons. The influence of baryons on the DM power spectrum is more prominent on the scale of baryon acoustic oscillations (BAO), corresponding to the scale of the sound horizon at recombination, or   $\\sim 150\\,\\mathrm{Mpc}$ today\\cite{eisenstein}. On this scale the DM power spectrum does not distinguish between (CDM) or (WDM), and at smaller scales, of interest for the satellite and cusp problems, the (BAO) features are not prominent and can be safely neglected.  The main ingredient to study the (DM) transfer function in absence of baryons is the non-relativistic Boltzmann-Vlasov equation for  DM  density and gravitational perturbations. The non-relativistic limit is warranted for particles that decoupled early and became non-relativistic prior to matter-radiation equality and for perturbations that entered the horizon prior to matter-radiation equality, these describe all the relevant scales for structure formation.  The method developed in ref.\\cite{gildan} yields a simple analytic approximation to the transfer function that is remarkably accurate in a wide range of scales relevant to structure formation. An important ingredient is a non-local kernel that depends on the unperturbed distribution function of the decoupled particles. This kernel describes memory of gravitational clustering and is a \\emph{correction to the fluid description} which become important at small scales. Distribution functions that feature larger support at small momentum yield longer range memory kernels thereby enhancing the transfer function at small scales\\cite{gildan}placing  greater importance on non-equilibrium aspects of the distribution function.  In this article we study the clustering properties of sterile neutrinos in the model proposed in references\\cite{kusepetra,petra,shapo,kuse2} by implementing all the steps described above, from obtaining and solving the quantum kinetic equation for production that establishes the distribution function at decoupling, narrowing the range of parameters, masses and couplings with the observational constraints from (DM) abundance and coarse grained phase space densities of (DM) dominated (dSphs)\\cite{gilmore,coldmatter}, and   solving the Boltzmann-Vlasov equation obtaining the transfer function and power spectrum which we compare to the case of thermal relics or Dodelson-Widrow-type\\cite{dw} distribution functions.  \\vspace{2mm}  {\\bf Summary of results:}  The production of  sterile neutrinos of $m\\sim \\mathrm{keV}$ via the decay of scalar gauge singlet with $M\\sim 100 \\,\\mathrm{GeV}$    leads to decoupling at a temperature $\\sim 100 \\,\\mathrm{GeV}$ and a distribution function that is strongly out of equilibrium and  behaves as $1/\\sqrt{p}$ for small comoving momenta $p$.  The constraints from (DM) abundance and coarse-grained phase space density from the latest compilation of photometric data from (dSphs) lead to a narrow window in the $\\mathrm{keV}$ range for the value of the mass of the sterile neutrino, consistently with the phenomenologically motivated extension beyond the standard model studied.  The (DM) transfer function and power spectrum are obtained from the solution of the non-relativistic Boltzmann-Vlasov equation for (DM) density and gravitational perturbations during matter domination. We implement a simple analytic approximate method\\cite{gildan} to obtain the density and gravitational perturbations and the transfer function that is remarkably accurate in a wide range of cosmologically relevant scales, as confirmed by the exact solution. This approach yields a wealth of information    that relates the small scale behavior of the transfer function to the range of memory of gravitational clustering, which is determined by the small (comoving) momentum region of the distribution function.  The enhancement of the non-equilibrium distribution at small momentum leads to a long range memory of gravitational clustering and slower fall off of the free-streaming solution. Both features lead to an enhancement of the transfer function and power spectrum at small scales.  We compare the transfer function and power spectrum from sterile neutrinos produced via gauge singlet decay to that of relativistic fermions decoupled in local thermodynamic equilibrium (LTE) (thermal relic) and sterile neutrinos produced by non-resonant mixing with active neutrinos \\emph{a la} Dodelson-Widrow\\cite{dw}. Thermal relics and (DW)-produced sterile neutrinos feature the \\emph{same} transfer function for similar ratios  of the decoupling temperature to mass.  The transfer function and power spectrum for sterile neutrinos produced by scalar decay is \\emph{substantially enhanced} with respect to that of (DW)-sterile neutrinos (and thermal relics) at small scales $ \\lambda \\lesssim 500 \\,\\mathrm{kpc}$. Whereas for (DW)-sterile neutrinos with  $m\\sim \\mathrm{keV}$ the transfer function is suppressed on scales $\\lambda \\lesssim 900 \\,\\mathrm{kpc}$, the scale of suppression for  $m\\sim \\mathrm{keV}$ sterile neutrinos produced by scalar decay at a scale $\\sim 100\\,\\mathrm{GeV}$ is $\\lambda \\lesssim 488 \\,\\mathrm{kpc}$ with a large enhancement of power at smaller scales.  For sterile neutrinos produced by scalar decay we find the following   behavior for the transfer function: at long wavelengths,  \\be T(k) \\simeq 1- C \\bigg(\\frac{k}{k_{fs}(t_{eq})}\\bigg)^2 + \\cdots  ~~;~~ k\\ll k_{fs}(t_{eq}) \\ee with $C \\sim \\mathrm{O}(1)$,   and at small scales  where the corrections to the fluid description and the memory of gravitational clustering becomes important   \\be T(k) \\simeq 1.902\\,e^{-k/k_{fs}(t_{eq})}~~;~~k \\geq k_{fs}(t_{eq})\\,   \\ee valid for scales $65\\,\\mathrm{kpc} \\lesssim \\lambda \\lesssim 500\\,\\mathrm{kpc}$ where $k_{fs}(t_{eq}) $ is the free streaming wavevector at matter radiation equality.  For $m\\sim \\mathrm{keV}$ and decoupling temperature $\\sim 100\\,\\mathrm{GeV}$ we obtain $k_{fs}(t_{eq}) \\sim 0.013/\\mathrm{kpc}$.     The smaller suppression scale may relieve the tension between the X-ray\\cite{xray} and Lyman-$\\alpha$ forest\\cite{lyman,lyman2} data and may provide the necessary enhancement of power at small scales to smooth out the inner profile of  (dSphs). ",
        "conclusions": "In this article we have implemented a program that begins with the microphysics of production and decoupling of a dark matter particle candidate, constrains the mass and couplings from the observed DM abundance and phase space density of DM dominated satellite galaxies (dSphs) and obtains the DM transfer function and power spectrum by solving the Boltzmann-Vlasov equation for density and gravitational perturbations during matter domination.  The model studied is a phenomenologically appealing extension of the minimal standard model proposed in references\\cite{shapo,kusepetra,petra,kuse2} in which sterile neutrinos are produced by the decay of a gauge singlet scalar. With the scale of the expectation value and mass of this scalar $\\sim 100\\,\\mathrm{GeV}$ a consistent description of $\\sim \\mathrm{keV}$ sterile neutrinos decoupled strongly out of equilibrium at a decoupling temperature $\\sim 100\\,\\mathrm{GeV}$ emerges and satisfies the DM abundance and phase space constraints from (dSphs).  The distribution function after decoupling for sterile neutrinos produced by gauge singlet decay features a strong enhancement at small comoving momentum $\\propto 1/\\sqrt{p}$ in contrast to sterile neutrinos produced via non-resonant mixing with active neutrinos (\\emph{a la} Dodelson-Widrow) for which   the distribution function is that  of ultrarelativistic thermal fermions multiplied by a suppression factor. Such distribution function was used in the hydrodynamical simulations performed in ref.\\cite{lyman2} to analyze the Lyman-$\\alpha$ forest data.  We have implemented an accurate analytic approximation to the solution of the Boltzmann-Vlasov equation and the transfer function introduced in ref.\\cite{gildan} and  obtained the power spectrum. This approximation allows to identify which features of the distribution function determine the small scale behavior of the transfer function. Distribution functions that favor small (comoving) momentum lead to longer range memory of gravitational clustering and slower fall-off of the free streaming solution. Both features lead to \\emph{small scale enhancement} of the transfer function and power spectrum.  We compare the transfer function and power spectrum in the case of sterile neutrinos produced by gauge singlet decay and by non-resonant mixing with active neutrinos, and find that the former is substantially enhanced over  latter at small scales. The transfer function and power spectrum of sterile neutrinos produced via non-resonant mixing is the \\emph{same} as that for fermionic thermal relics. The suppression factor in the distribution (see eqn. \\ref{fdw}) modifies the abundance an primordial phase space densities but \\emph{not} the transfer function or power spectrum.  While at large scales $\\lambda \\gg 1\\,\\mathrm{Mpc}$ the transfer function in both cases are nearly indistinguishable and the same as the (CDM) case, the power spectrum for $m\\sim \\mathrm{keV}$ (DW) sterile neutrinos produced by non-resonant mixing is suppressed below a scale $\\lambda \\lesssim 900 \\,\\mathrm{kpc}$ whereas the transfer function for sterile neutrinos produced via scalar decay is suppressed below a scale $\\lambda \\lesssim 488 \\,\\mathrm{kpc}$ and substantially enhanced at smaller scales when compared to the (DW) case.  We find the simple fits to $T(k)$ in the limits of large and small scales. For large scales we find  \\be T(k) \\sim 1-C\\bigg(\\frac{k}{k_{fs}(t_{eq})}\\bigg)^2 + \\cdots ~~;~~ {k} \\ll k_{fs}(t_{eq})\\,, \\ee with $C \\sim \\mathrm{O}(1)$. In this long wavelength limit the fluid description is valid and the contribution from the memory kernel is subleading.  At small scales the corrections to the fluid description in terms of the non-local kernel that includes memory of gravitational clustering  becomes important, in the small scale regime ${k} \\geq k_{fs}(t_{eq})$  a   simple and   accurate numerical fit yields:  \\be T(k) \\simeq 1.902\\,e^{-k/k_{fs}(t_{eq})}\\,, \\ee where $k_{fs}(t_{eq})$ is the free streaming wave vector at matter-radiation equality. For  a sterile neutrino with $m\\sim 1\\,\\mathrm{keV}$ decoupling at $T_d\\sim 100 \\,\\mathrm{GeV}$ we find \\be k_{fs}(t_{eq}) \\simeq 0.013/\\mathrm{kpc}\\,.\\ee  This fit is remarkably accurate in the wide range of scales $60\\,\\mathrm{kpc} \\lesssim \\lambda \\lesssim 500\\,\\mathrm{kpc} $ and is   different from the often quoted result of ref.\\cite{BBKS}.       We  have given arguments that show that the results presented above are an   \\emph{upper bound} to the small scale properties of $T(k)$, since the evolution of WDM perturbations during   (RD) leads to further suppression of $T(k)$ with a \\emph{larger} suppression for the case of sterile neutrinos with distribution functions of the form (\\ref{fdw}) as compared to those with (\\ref{f0}). This is because the distribution function (\\ref{f0}) favors the small momentum region leading to shorter free streaming lengths and larger free streaming wavevectors, allowing more power at small scales.  A  more  detailed analysis of the initial conditions obtained by including the evolution during the (RD) will be reported elsewhere.    The substantial difference between the suppression scales in the transfer function and power spectrum at small scales between sterile neutrinos produced by gauge singlet decay and those produced by the (DW) mechanism suggest that sterile neutrinos produced by scalar decay \\emph{may } relieve the tension between the constraints from X-ray \\cite{xray} and Lyman-$\\alpha$ forest data\\cite{lyman,lyman2}.  In order to assess whether sterile neutrinos produced by scalar decay may explain the cored profiles of (dSphs), a full N-body simulation with the power spectrum obtained above must be carried out.   Whereas in this article we have focused on the production mechanism from gauge singlet decay, the model includes sterile-active mixing via a see-saw (Majorana) mass matrix. Therefore there is also a complementary mechanism of sterile neutrino production via active-sterile mixing akin to the (DW) mechanism, which is effective at a much lower temperature $\\sim 150\\,\\mathrm{MeV}$. The wide separation of decoupling scales, $\\sim 100\\,\\mathrm{GeV}$ for scalar decay, vs. $\\sim 150\\,\\mathrm{MeV}$ for (DW) suggests that once the distribution function has been established after decoupling at the higher temperature, the non-equilibrium effects at the lower scale may not modify the small momentum region significantly. We conjecture this to be the case because the wide separation of scales \\emph{suggests} that the distribution function obtained from scalar decay, may be taken as an \\emph{initial condition} for the kinetics of production via active-sterile mixing, in which case for small mixing angle and neglecting again the build up of the population, the final distribution function would be the \\emph{sum} of (\\ref{f0}) and (\\ref{fdw}).  We have not included  this possibility in this article, postponing the more detailed kinetic description, a study of the consequences and different initial conditions to a forthcoming article.  In obtaining the transfer function and power spectrum by solving the Boltzmann-Vlasov equation for (DM) and gravitational perturbations during matter domination, we have neglected the contribution from baryons and photons. Although these will only affect the results for $T(k)$ and $P(k)$ at the few percent level, as discussed in the introduction, a precise assessment of $P(k)$ to few percent accuracy requires solving the full set of Boltzmann-equations by modifying publicly available codes.  The study presented in this article provides a reliable preliminary assessment of   (DM) candidates, allows a systematic comparison  and highlights the important small-scale aspects, perhaps eventually justifying the more numerically demanding task of solving the full set of coupled Boltzmann equations for photons, baryons, gravitational and (DM) perturbations."
    },
    "0807/0807.0652_arXiv.txt": {
        "abstract": " ",
        "introduction": "Oxygen is important for the chemistry of many of the current Solar System bodies since it readily bonds with many other atoms. There are three prominent atomic oxygen lines in the optical region of the spectrum. The transition from the first excited state of the singlet branch, $^1$D, to the ground state, $^3$P, produces a doublet in the red at 6300.304~\\AA\\ and 6363.776~\\AA.  The second excited singlet state, $^1$S, decays to the $^1$D state 90--95\\% of the time with a transition at 5577.339~\\AA; the remaining 5-10\\% of the time, the $^1$S state decays directly to the ground state ($^1$S -- $^3$P) through two UV lines at 2958.365 and 2972.288~\\AA\\ (Slanger {\\it et al.} 2006a). \\nocite{slangbranch06} These are all forbidden lines. The green line and the red doublet are observed in the Earth's atmosphere and are affected by auroral activity. The green line has been observed in spectra of Venus and appears to be quite variable there (Slanger {\\it et al.} 2006b). \\nocite{slangvenus06} It is assumed that these lines will also be present in Mars' atmosphere, but only the UV line at 2972.288~\\AA\\ has been detected definitively (Barth {\\it et al.} 1971; Leblanc {\\it et al.} 2006). \\nocite{barthetal71,leblanc06}  The three forbidden lines are also seen in the spectra of comets. Photodissociation of oxygen-bearing parent species produces oxygen atoms in the ground $^3$P state or directly into the excited $^1$D or $^1$S states, depending on the parent molecule and the nature of the solar photodissociation. The red doublet is substantially stronger than the green line and the red lines are generally much stronger than the telluric red lines. The cometary green line is generally a small fraction of the intensity of the telluric green line.  To observe the oxygen in the spectra of comets, it is desirable to observe with moderately high spectral resolution in order to resolve the Doppler-shifted cometary lines from the telluric lines.   The red doublet occurs in a part of the cometary spectrum which is relatively uncrowded, so contamination from other cometary species is not an issue at high spectral resolution. The spectral resolution allows for the separation of the cometary and telluric lines. The situation for the green line is quite different.  In addition to the telluric green line which must be separated, the cometary green line sits in the middle of the C$_{2}$ (1,2) P-branch; the P$_1$(27) and P$_2$(26) lines flank the oxygen line.  The C$_{2}$ spectrum is quite dense.  Thus, for many cometary spectra, it is impossible to derive the intensity of the green line.  However, for some comets, the oxygen line is stronger than the C$_{2}$ bands and can be successfully deblended from the cometary C$_{2}$ and the telluric oxygen if the spectral resolution is sufficient.  Some comets have been observed to have almost no C$_{2}$, making detection of the green line simple (Cochran and Cochran 2001). \\nocite{coco01}  The green line and the red doublet are not produced by fluorescence or dissociative recombination (Festou and Feldman 1981) \\nocite{fefe81} but instead arise from atoms produced directly in the $^1$S or $^1$D states (i.e.  prompt emissions). The lifetimes of the upper states of the oxygen lines are short, with the $^1$D state lasting about $\\sim110$\\,s and the $^1$S state lasting $\\sim1$\\,s. Indeed, the lifetime of the $^1$S state is so short that its spatial distribution traces the parent's distribution directly, since the oxygen atom will not have time to travel far from its direct parent prior to its decay to the ground state.  The atomic oxygen observed in the spectra of cometary comae did not sublime from pure oxygen ice. Instead, oxygen is stored in the nucleus as an oxygen-bearing ice.  Jackson and Cochran (2007) \\nocite{jaco08} observed the green line in the spectrum of comet Tempel 1 following the impact of the Deep Impact spacecraft with the comet.  They saw the line increase in brightness quickly after the impact and then decay.  Figure~\\ref{tempel1} shows the intensity of the green line after the impact as a function of the cometocentric distance. The initial increase shows the impulsive nature of the impact. The decay started at $\\sim2000$s after the impact.  After that time, the green line flux decayed as 1/$\\rho$, where $\\rho$ is the projected cometocentric distance. This behavior is characteristic of prompt emission produced by a single-step dissociation.  Since, with its short lifetime, the emission is tracing the direct parent, the behavior shown in Fig.~\\ref{tempel1} is indicative that the $^1$S state is a daughter, and not a granddaughter, product of some species. But what is this parent? Complex oxygen-bearing species (e.g. H$_2$CO or HCOOH) can not be the parent, as they can only produce oxygen via a two-step process.  In addition, Jackson and Cochran showed that the green line flux decayed much faster than the OH flux, indicating that the oxygen line did not come primarily from OH.  Thus, the green line is likely to be directly produced from H$_{2}$O, CO$_2$, or CO. It is generally assumed that the red lines come from the dissociation of H$_{2}$O as this is consistent with the distribution of  O ($^1$D) in the inner coma of a comet (Combi and McCrosky 1991; Fink and Johnson 1984; Magee-Sauer {\\it et al.} 1988). \\nocite{comc91,maetal88,fijo84} We have used high spectral resolution observations of 8 comets to detect the green line and the red doublet of oxygen in order to try to understand the production of the atomic oxygen in cometary spectra. \\begin{figure} \\vspace{4in} \\special{psfile=\"tempel1oxy-rot.ps\" vscale=60 hscale=60 voffset=-40 hoffset=-20} \\caption[fig1]{The intensity of the green oxygen line at 5577.339~\\AA\\ was measured immediately after the impact of comet 9P/Tempel~1 and the Deep Impact spacecraft using the HIRES instrument at Keck I.  The green line initially increased in intensity and then decayed as it flowed outwards.  This figure shows the intensity of the line as a function of cometocentric distance, $\\rho$.  The line decayed as $1/\\rho$, the signature of a parent.  Since the lifetime of the $^1$S state is $\\sim1$\\,s at 1\\,{\\sc {\\sc au}}, this points to a single-step process from parent to O ($^1$S) state.}\\label{tempel1} \\end{figure} ",
        "conclusions": "Spectra of comets tell us about the physical and chemical processes taking place within the cometary comae.  In addition, comets are excellent laboratories for testing our understanding of chemistry since the long path lengths through the coma allow relatively small gas densities to produce significant column densities.  We are interested in the oxygen lines for several reasons.  First, oxygen is one of the most abundant elements in the Universe after hydrogen and helium and it readily bonds with other atoms and molecules.  Indeed, 80\\% of cometary ices are H$_{2}$O, so oxygen is a critical component of comets.  Water is hard to detect spectrally from the ground so daughter products of water have traditionally been used as a proxy to study the water. Since the OH fluorescent bands are in the UV region of the spectrum, where there is low detector quantum efficiency and high atmospheric extinction, they are difficult to study.  The three forbidden oxygen lines discussed in this paper are in the optical and have been used as a water proxy.  Unfortunately for studies of atomic oxygen in comets, the telluric spectrum has strong atomic oxygen emissions which are always present and are quite variable.  Thus, until the advent of the era of high resolution spectroscopy of comets, contamination of the cometary lines by the telluric lines was always an issue. Even with high resolving powers, the green line in cometary spectra has proved to be a difficult target since it sits within the spectral region of the C$_{2}$ (1,2) P-branch. Despite that, there have been some observations of the detection of the green line in cometary spectra previously reported. These are listed in Table~\\ref{otherdata}. \\begin{table} \\renewcommand{\\baselinestretch}{1} {\\small \\caption{Previously Reported Green Line Detections}\\label{otherdata} \\vspace*{5pt} \\centering \\begin{tabular}{lccccl} \\hline \\multicolumn{1}{c}{Comet} & Date(s) & Heliocentric & Geocentric & $\\frac{I_{5577}}{(I_{6300}+I_{6364})} $ & Ref.\\\\ & & Distance & Distance & flux & \\\\ & & (AU) & (AU) & ratio & \\\\ \\hline C/1983 H1 & 6, 8 May 1983 & 1.02 & 0.13, 0.09 & 0.022--0.034 & 1 \\\\ \\hspace*{2em}(IRAS-Araki-Alcock) & \\\\ 1P/Halley & 5 Apr 1986 & 1.25 & 0.45 & 0.05 -- 0.1 & 2 \\\\ C/1996 B2 (Hyakutake) & 23, 27 Mar 1996 &  1.08, 1.00 & 0.12 &  0.13 -- 0.16 & 3 \\\\ C/1995 O1 (Hale-Bopp) & 26, 28 Mar, 22 Apr 1997 & 0.92, 0.99 & 1.33, 1.63 & 0.18 -- 0.22 & 4 \\\\ 153P/Ikeya-Zhang & 20 Apr 2002 & 0.89 & 0.43 & 0.12$\\pm$0.1 & 5 \\\\ 9P/Tempel 1 & 2 -- 6,9 Jul 2005 & 0.89--0.92 & 1.51 & 0.056-0.13$\\pm$0.1 & 6 \\\\ \\hline \\multicolumn{6}{p{5.5in}}{References -- 1: Cochran (1984) - red and green lines observed different nights; 2: Smith and Schempp (1989) - 6364 line not observed; 3: Morrison {\\it et al.} (1997) - 6300 line not observed; 4: Zhang {\\it et al.} (2001); 5: Capria {\\it et al.} (2005); Capria {\\it et al.} (2008)} \\\\ \\hline \\nocite{co84,smsc89,moetal97,zhzhhu01,caetal2005,caetal2008} \\end{tabular}  } \\end{table}  The three oldest observations all suffered from non-simultaneity of the observations of the three lines.  Indeed, the Halley and Hyakutake observations reported in Table~\\ref{otherdata} each contain observations of only 1 line of the red doublet.  Thus, these observations should be used cautiously.  The ratios of the green line to the sum of the red line fluxes for Hale-Bopp, Ikeya-Zhang and Tempel 1 are consistent with the values we found for our 8 comets reported in this paper.  Hale-Bopp seems to have the highest ratio.  This might be because Hale-Bopp had extremely high water production rates and a very large collisional zone.  However, there are not a great deal of details in the Zhang {\\it et al.} (2001) paper so it is difficult to judge why this ratio might be higher. Hale-Bopp had a very strong continuum and C$_{2}$ band  so there may have been issues correcting for these factors.  Comparing these remaining three comets with the 8 observations presented here, it is obvious that the green line is typically about 10\\% the summed intensity of the red doublet. If there is collisional quenching, as implied in Fig.~\\ref{hyatrend}, the red line fluxes might be lower than without quenching so the $^1$S/$^1$D ratio in the absence of quenching might be even lower. All of the authors referenced in Table~\\ref{otherdata}, and indeed our previous paper on 1999 S4 (LINEAR) (Cochran and Cochran 2001), have used this ratio along with Table~3 of Festou and Feldman (1981) to conclude that water is the source of the green atomic oxygen line.  In this paper, we reported on the widths of the oxygen lines in addition to their intensities and found that the green line was consistently wider than the red lines.  Since 90--95\\% of the photons in the $^1$S state decay to the ground state via the $^1$D state, the higher velocity photons should widen the red lines. We examined how this would affect our measurements by generating two Gaussian line profiles with FWHM values appropriate for the green line and assuming the red line FWHM was the ``correct\" value.   These two lines were then combined with the same average flux ratio observed in the comets reported in this paper.  We then measured the properties of this artificially generated line against a red line generated by assuming none of the red line comes from photons which were first excited to the $^1$S state. We find the combined line has 2\\% lower flux than a pure line but that the FWHM of the lines agree to 0.001~\\AA, or at least as good as our measurement accuracy.  Therefore, based on the relative intensity of the green and red lines, we would not have been able to measure any broadening of the red lines and the measured red line widths can be assumed to be indicative of excitation in the $^1$D state only.  The wider green line leads to two possibilities for the production of photons in the $^1$S state.  The first is that there is a parent other than H$_{2}$O that dissociates into photons in the $^1$S state.  As discussed above, CO can be eliminated based on its bonding and on observations of comet 1999~S4.  OH can be eliminated because it decays more slowly than the green line.  Thus, CO$_2$ would be the likely alternative parent.  The other possibility is that H$_{2}$O is the parent for both the $^1$S and $^1$D states.  This would imply that the energy of the photons responsible for the production of the $^1$S state is higher than that for the production of the $^1$D state. As Huestis (2006) pointed out, these excitation rates are not calibrated.  The laboratory measurement of the production of the green line from H$_{2}$O has never been made. Nor has the production of the red lines from CO$_2$.  Until the point that these rates are known with any certainty, all we can definitely conclude is that the green line flux is about 10\\% that of the combined red line flux in all comets.  However, this in itself may offer a clue. The observations of the 8 comets reported here span 11 years, or a full solar activity cycle.  They also span heliocentric distances from 0.66 to 1.54\\,{\\sc au}.  We see no difference in either the relative line intensities nor in the line widths from one comet to another.  Thus, whatever photodissociative process(es) produce these lines, their relative importance does not change with solar UV flux or from comet to comet. Inspection of Table~I of Bockel\\'{e}e-Morvan {\\it et al.} (2004) \\nocite{bocrmuwe2004} shows that the abundances of minor species such as CO and CO$_2$ relative to water are quite variable.  The constancy of the green line to the red line flux thus suggests a single parent for both states with different energies responsible for the production of the different states. Clearly, more laboratory work is needed but comets are lending clues to the processes involved.  \\vspace*{0.75in} \\begin{center}Acknowledgments\\end{center} This paper includes data taken at The McDonald Observatory of The University of Texas at Austin.  This work was supported by NASA Grant NNG04G162G. I thank D. Lambert, M. Endl, and W. Cochran for obtaining some of the spectra.  \\newpage"
    },
    "0807/0807.2461_arXiv.txt": {
        "abstract": "This is the second paper of a series where we study the clustering of LRG galaxies in the latest spectroscopic SDSS data release, DR6, which has 75000 LRG galaxies covering over 1 $Gpc^3/h^3$ for $0.15<z<0.47$. Here we focus on modeling redshift space distortions in $\\xisp$, the 2-point correlation in separate line-of-sight and perpendicular directions, at small scales and in the line-of-sight. We show that a simple Kaiser model for the anisotropic 2-point correlation function in redshift space, convolved with a distribution of random peculiar velocities with an exponential form, can describe well the correlation of LRG at all scales. We show that to describe with accuracy the so called \"fingers-of-God\" (FOG) elongations in the radial direction, it is necessary to model the scale dependence of both bias $b$ and the pairwise rms peculiar velocity $\\sigma_{12}$ with the distance. We show how both quantities can be inferred from the $\\xisp$ data. From $r \\simeq 10$ Mpc/h to $r \\simeq 1$ Mpc/h, both the bias and $\\sigma_{12}$ are shown to increase by a factor of two: from $b=2$ to $b=4$ and from $\\sigma_{12}=400$ to $800$ Km/s. The later is in good agreement, within a 5 percent accuracy in the recovered velocities, with direct velocity measurements in dark matter simulations with $\\Omega_m=0.25$ and $\\sigma_8$=0.85.\\ ",
        "introduction": "\\label{sec:intro}   The luminous red galaxies (LRGs) are selected by color and magnitude to obtain intrinsically red galaxies in Sloan Digital Sky Survey (SDSS) \\cite{eisenstein2001}. These galaxies trace a big volume, around $1Gpc^3h^{-3}$, which make them perfect to study large scale clustering or smaller scales with larger statistics. In this paper we focus on the 2-point correlation function $\\xisp$ on the smaller scales, where the non-linear bias and the random peculiar velocities can complicate the analysis. Our motivation is to provide a model that can explain the observed $\\xisp$ of the LRG galaxies. Such a model does not exist right now and we are not aware of any attempt to reproduce the $\\xisp$ in the detail we will explore here. LRG galaxies seem to display larger \"Fingers-of-God\" \\cite{jackson} than regular galaxies, is this evidence for larger velocities? If so, is this evidence that LRG trace stronger gravitational potentials? We will show that a simple Kaiser model for the anisotropic 2-point correlation function in redshift space, convolved with a distribution of random peculiar velocities with an exponential form, can describe well the correlation of LRG at all the scales. To describe with accuracy the so called \"fingers-of-God\" (FOG) elongations in the radial direction, it is necessary to model the scale dependence of both bias $b$ and the pairwise rms peculiar velocity $\\sigma_{12}$ with the distance.  In Paper I \\cite{paper1} of this series, we have analyzed larger scales and present the basis for this work, including more detailed theory, error and systematic effects.  On large scales, the density fluctuations are small enough to be linearized, and can be used to constrain cosmological parameters, since we can assume that the clustering is well described by (linearly biased) dark matter (see Paper I). On smaller scales, we can learn about the relation of galaxies to dark matter through the biased clustering of halos. This range of distances can be fitted by a power law, but there are small deviations that can be understood in the theory of the halo occupation distribution. The transition between galaxy pairs of the same halo and galaxy pairs that belong to different halo, occurs around 1Mpc/h. When moving to scales smaller than 1Mpc/h, in the 1-halo term, we can see processes more complex that modify the galaxy clustering, such as dynamical friction, tidal interactions, stellar feedback, and other dissipative processes.  We use the most recent spectroscopic SDSS data release, DR6 \\cite{dr6}, to perform the study of small scales in the anisotropic 2-point correlation function, and its derivatives projected correlation function and real-space correlation function. LRGs are supposed to be red old elliptical galaxies, which are usually passive galaxies, with relatively low star formation rate. They have steeper slopes in the correlation function than the rest of galaxies, since they are supposed to reside in the centers of big halos, inducing non-linear bias dependent on scale, for small scales.  The same LRGs (but with reduced area) have been studied from different points of view. Zehavi et al (2005) study LRGs at intermediate scales (0.3 to 40Mpc/h), where they calculate the projected correlation function, the monopole and real-space correlation function to study mainly linear bias, non-linear bias and the differences between luminosities. They find that there are differences from a power law for scales smaller than 1Mpc/h and they find no strong evolution for LRG. At smaller scales, Eisenstein et al (2005) did a cross-correlation between spectroscopic LRG with photometric main sample in order to reduce shot-noise in small scales clustering. They conclude that the clustering is higher for most luminous galaxies. Moreover, LRG have a scale dependent bias on luminosity, while for normal galaxies the luminosity bias is scale independent. LRG galaxies are surrounded by other red galaxies near them, which seem to be approaching the center LRG galaxy. Finally, Masjedi et al (2006) deal with very small scale clustering to scales smaller than 55'' by cross-correlating the spectroscopic LRG sample and the targeted imaging sample and find that the correlation function from 0.01-8Mpc/h is really close to a power law with slope -2, but there are still some features that diverge from the power law.    The small scale slope depends on the interplay between two factors which control how the correlation function of galaxies is related to that of the underlying matter : the number of galaxies within a dark matter halo (HOD) and the range of halo masses which contain more than one galaxy \\cite{benson}. LRGs are to be found in halos with the median of the distribution occurring at $3\\;10^{13}M_\\odot/h$, estimated using weak lensing measurements \\cite{mandelbaum}. Almeida et al (2008) find with simulations that 25\\% of LRGs at z = 0.24 are satellite galaxies, which play an important role in the form of the small scale correlation function, as well as the pairwise velocity dispersion, and also provide information about galaxy formation and evolution. Zheng et al (2008), who study HOD in LRG clustering, find that the satellite fraction is small (5-5\\% for $M_g<-21.2$) and decreases with LRG luminosity.  Slosar et al (2006) show that pairwise velocity distribution in real space is a complicated mixture of host-satellite, satellite-satellite and two-halo pairs. The peak value is reached at around 1Mpc/h and does not reflect the velocity dispersion $\\sigma_{12}$ of  a typical halo hosting these galaxies, but is instead dominated by the sat-sat pairs in high-mass clusters. Different groups have found this dependence of $\\sigma_{12}$ on the scale (Zehavi et al 2002, Hawkins et al 2003, Jing and Borner 2004, Li et al 2006, Van den Bosch 2007, Li et al 2007). They also find that $\\sigma_{12}$ depends on luminosity, with higher values for both large and small luminosities, but this tendency can not be predicted by the halo model. Li et al (2006, 2007) add that redder galaxies have stronger clustering and larger velocities at all scales, so they move in strongest gravitational fields. Tinker et al (2007) use the halo occupation distribution framework to make robust predictions of the pairwise velocity dispersion (PVD). They assume that central galaxies move with the center of mass of the host halo and satellite galaxies move as dark matter. The pairs that involve central galaxies have a lower dispersion, so the fraction of satellites strongly influences both the luminosity and scale dependence of the PVD in their predictions. At r between 1 and 2 Mpc/h, the PVD rapidly increases with smaller separation as satellite-satellite one-halo pairs from massive halos dominate. At r $<$ 1 Mpc/h, the PVD decreases towards smaller scales because central-satellite one-halo pairs become more common. At r $>$ 3Mpc/h, the PVD is dominated by two-halo central galaxy pairs, and reach a constant value.  LRGs have also been analyzed at higher redshifts (z=0.55) with the 2dF-SDSS LRG and QSO Survey (2SLAQ, Cannon et al 2007). Ross et al (2007), Da \\^Angela et al (2008) and Wake et al (2008) analyze the redshift distortions in the LRGs and quasars for this catalog. For part of our analysis we have followed the method explained in Hawkins et al (2003), an extensive analysis of the redshift distortions in the 2dF catalog.  Here we define the parameters that we assume during all this work, which are motivated by recent results of WMAP, SNIa and previous LSS analysis: $n_s=0.98$, $\\Omega_b=0.045$, $h=0.72$ and flat geometry. Unless otherwise said, we use $\\Omega_m=0.25$ and $\\sigma_8=0.85$, both obtained from paper I. We will use the power spectrum analytical form for dark matter by Eisenstein and Hu (1998), and the non-linear fit to halo theory by Smith et al (2003). ",
        "conclusions": "\\label{sec:discuss}   We have studied small scales in $\\xisp$  measured from luminous red galaxies, which have both a strong non-linear bias and are affected by large random peculiar velocities in redshift space. We have compared our results for the monopole in redshift space $\\xi(s)$, the real-space correlation function $\\xi(r)$ and the perpendicular projected correlation function $\\Xi(r)$ with Zehavi et al (2005) and Masjedi et al (2006) results. Our analysis agree at all the scales except for very small scales, below 1Mpc/h, where we find some minor differences with respect to previous analysis.  Here we also recover, for the first time, the pairwise velocity dispersion $\\sigma_{12}$ of LRG galaxies as a function of separation by fitting $\\xisp$. Our method is shown to work in realistic simulations. We find that the scale variation of $\\sigma_{12}$  in LRG galaxies is similar to that in  dark matter simulations. We show in section \\S\\ref{sec:consistency} that a simple Kaiser model convolved with an exponential distribution of pairwise velocities, can explain well the complicate shape of $\\xisp$ at small scales, once we add the scale dependent bias and the scale dependent $\\sigma_{12}$. The $\\chi^2$ per degree of freedom reduces by over a factor of two when we allow $b$ and $\\sigma_{12}$ to change with scale. We notice that if we attribute all the distortion at small scales to the peculiar velocities, without taking into account the non-linear bias, velocities appear   to be artificially larger. On small scales the errors in $\\xisp$ are relatively small compare to the signal, see section \\S\\ref{sec:errorpisigma}. Thus the agreement of the $\\xisp$ data to our simple modeling shown in Fig.\\ref{fig:pisigmar2} is quite remarkable and significant. We believe that this is the first time this simple $\\xisp$ model is shown to agree in detail with small scale LRG clustering."
    },
    "0807/0807.5077_arXiv.txt": {
        "abstract": "We study a set of universe models where the dark sector  is described by  a perfect fluid with  an affine equation of  state $P=P_0+\\alpha \\rho$, focusing specifically on cosmological perturbations in a flat universe. We perform a Monte Carlo Markov Chain analysis  spanning the full parameter space of the model using the WMAP 5 years data and the SDSS LRG4 survey. The affine fluid can either  play the role of a unified dark matter (UDM), accounting for both dark matter and a cosmological constant, or work alongside cold dark matter (CDM), as a form of dark energy. A key ingredient is the sound speed, that depends on the nature of the fluid and that, for any given background model, adds a degree of freedom to the perturbations: in the barotropic case the square of the sound speed is simply equal to the affine parameter $\\alpha$; if entropic perturbations are present the effective sound speed has to be specified as an additional parameter. In addition to the barotropic case, we consider the two limiting cases of effective sound speed equal to 0 or 1. For $\\alpha=c_s^2=0$ our UDM model is equivalent to the standard $\\Lambda$CDM  with adiabatic perturbations. Apart of a trivial subcase, all models considered  satisfy the data constraints, with quite standard values for the usual cosmological parameters. In general our analysis confirms that cosmological datasets require both a collisionless massive and cold component to form the potential wells that lead to structure formation, and an effective cosmological constant that drives the late accelerated expansion. ",
        "introduction": "In the last few years cosmological observations have become increasingly accurate, allowing various models to be tested or even ruled out. The one that currently seems to satisfy most observational requirements is the so-called concordance or $\\Lambda$CDM model \\cite{Tegmark:2004,Spergel:2003}: a flat universe, with $\\sim 4\\%$ of baryons,  $\\sim20\\%$ of an unknown weakly-interacting heavy component (or dark matter), leaving the remaining $\\sim 76\\%$ in the form of a cosmological constant (or vacuum energy density) responsible for the late time acceleration of the universe  \\cite{Riess:astro-ph/0611572,Perlmutter:astro-ph/9812133}. However, despite its simplicity, this model lacks solid theoretical motivations and actually seems to require \\emph{ad hoc} assumptions, both on dark matter and the cosmological constant. Many hypotheses have been proposed to alleviate the problem of the cosmological constant and the related one of why the density of the two unknown components are of the same order of magnitude in the present universe (the so-called coincidence problem). For an incomplete list see \\cite{bento03,Kamenshchik00,2008arXiv0807.1020B, sahniwang00,mainini05,giannakis05,beca05,chimento04,scherrer04} and the recent review on dark energy \\cite{DErevCopeland07}.  Observationally, the dark sector is degenerate: by definition, dark components can be probed only through gravitational effects, leaving open a wide range of possibilities regarding their nature and possible interactions \\cite{2007astro.ph..2615K}. In this paper we investigate an effective model for the dark sector, based on the affine parameterisation of the equation of state \\cite{Ananda:astro-ph/0512224}: \\begin{equation} \\label{affine_pressure} p_{X}=P_0+\\alpha\\rho_{X}, \\end{equation} where $p_X$ is the pressure, $\\rho_X$ is the energy density, and $P_0$ and $\\alpha$ are constant parameters; this leads to a time dependent equation of state parameter \\begin{equation} \\label{w_eq} w_{X}=\\frac{P_0}{\\rho_{X}}+\\alpha. \\end{equation} An interesting property of this parameterisation is that it results in a constant energy density term mimicking an effective cosmological constant, with $\\Omega_\\Lambda = -P_0/\\left[\\rho_c(1+\\alpha)\\right]$, plus an evolving term that can reproduce a dark matter behaviour for certain choices of the parameter $\\alpha$. This allows one to either treat the affine fluid as a single unified dark component, or to use it to model dark energy alone.  As shown in \\cite{quercellini_fields07}, when $\\alpha$ is negative, this description can be seen as the attractor solution for a quintessence scalar field dynamics. Alternatively, when treating perturbations, a barotropic affine fluid can be interpreted as a k-essence scalar field (naturally describing an effective cosmological constant plus dark matter), while a scalar field with sound speed $c_s^2=1$ act as a dark energy component. In addition, an affine fluid description can also be interpreted as the result of two interacting dark components (one of them being a cold dark matter component), as we discussed in detail in \\cite{QuercelliniEtAl2008_coupled}. In the present work we explore several different cases resulting from the affine fluid description. We consider two classes of models: one where the affine fluid describes a unified dark component, the other containing a cold dark matter component as well. For each class, we also study three separate subcases, identified by the value of the speed of sound: the barotropic case, with $c^2_{\\eff}=\\alpha$, the case $c^2_{\\eff}=1$, and the ``silent'' case \\cite{Bruni:1995b,Bruni:1995a} with $c^2_{\\eff}=0$.  To study the properties of the model, we calculate the evolution of scalar perturbations in the affine fluid by modifying the publicly available CAMB code, and set constraints to the parameters of the model by performing  a Monte Carlo Markov Chain analysis using the cosmic microwave background (CMB) anisotropy WMAP 5 year data \\cite{WMAP5Komatsu2008} and the large-scale matter distribution derived from the Sloan Digital Sky Survey (SDSS) Luminous Red Galaxy (LRG) 4 year data \\cite{sdsslrg4}.  In the next section we describe our model and the theoretical framework adopted, in Sec.~\\ref{results} we discuss the results of our numerical calculations and comparison with observations, and in Sec.~\\ref{conclusions} we draw our final conclusions. ",
        "conclusions": "\\label{conclusions} We studied the effect of an affine EoS fluid model applied to the dark sector, both as a unified description of dark matter and an effective cosmological constant, and as a pure dark energy component. Our model makes use of a dynamical parameterisation relating $p$ and $\\rho$, as opposed to the usual cynematical description of the EoS parameter in terms of its current value and its first derivative. In a previous paper \\cite{LADM07} we carried on a comparison of the background evolution of this model with existing cosmological observations. In the present work, we focused on the behaviour of cosmological perturbations, and compared the theoretical predictions with the CMB WMAP 5 year data, and with the SDSS large scale structure data.  \\begin{table*}[h!] \\caption[]{Best fit parameter values for $\\alpha$DM.} \\label{final_table} \\begin{center} \\begin{tabular}{|c|c|c|c|c|c|} \\hline models/params  & $\\Lambda$CDM +    & \\multicolumn{2}{|c|}{$\\alpha DM$ - bar}& \\multicolumn{2}{|c|}{$\\alpha DM$ - $c^2_{\\eff}=0$} \\\\ & SNe \\& BAO (WMAP5)  & CMB               &  MPS               & CMB                & MPS               \\\\ \\hline $\\Omega_b h^2$ & $0.02265\\pm0.00059$ & $0.0223\\pm0.0007$ &$0.0223\\pm0.0006$ & $0.0224\\pm0.0006$  & $0.0224\\pm0.0003$ \\\\ \\hline $\\Omega_c h^2$ & $0.1143\\pm0.0034 $ & -                 & -                  & -                   & -                 \\\\ \\hline $H_0$          & $70.1\\pm1.3$       & $75\\pm5$          &  $71\\pm2$          & $69\\pm6$            & $67\\pm4$          \\\\ \\hline $\\tau$         & $0.084\\pm0.016$    & $0.090\\pm0.018$   &  $0.083\\pm0.016$   & $0.086\\pm0.017$     & $0.085\\pm0.017$   \\\\ \\hline $n_s$          & $0.960\\pm0.014$    & $0.99\\pm0.03$     &  $0.960\\pm0.014$   & $0.959\\pm0.015$     & $0.957\\pm0.014$   \\\\ \\hline $\\log(10^{10}A_s)|_{k=0.002}$ & $3.20\\pm0.08$ & $3.10\\pm0.08$ &$3.19\\pm0.04$ & $3.19\\pm0.05$       & $3.19\\pm0.04$     \\\\ \\hline $\\Omega_\\Lambda$     &$0.721\\pm0.015$     & $0.76\\pm0.04$  &$0.73\\pm0.02$    & $0.70\\pm0.09$       & $0.69\\pm0.05$     \\\\ \\hline $\\alpha$       & -                 & $(8\\pm11)\\times10^{-4}$& $(0.2\\pm4)\\times10^{-7}$ & $(-1.5\\pm3)\\times10^{-3}$ & $(-2\\pm2)\\times10^{-3}$ \\\\ \\hline \\end{tabular} \\end{center} \\end{table*}    \\begin{table*}[h!] \\caption[]{Best fit parameter values for $\\alpha$CDM.} \\label{final_table2} \\begin{center} \\begin{tabular}{|c|c|c|c|c|c|c|} \\hline models/params  &  \\multicolumn{2}{c|}{$\\alpha CDM$ - bar}& \\multicolumn{2}{c|}{$\\alpha CDM$ - $c^2_{\\eff}=0$} & \\multicolumn{2}{c|}{$\\alpha CDM$ - $c^2_{\\eff}=1$}\\\\ & CMB                &  MPS               & CMB                & MPS               & CMB                & MPS               \\\\ \\hline $\\Omega_b h^2$ &  $0.0224\\pm0.0007$ &  $0.0220\\pm0.0007$ & $0.0223\\pm0.0006$  & $0.0224\\pm0.0006$ & $0.0224\\pm0.0006$ & $0.0224\\pm0.0006$\\\\ \\hline $\\Omega_c h^2$ &  $0.07\\pm0.03$      &  $0.106\\pm0.004$   & $0.109\\pm0.006$    & $0.107\\pm0.006$   & $0.109\\pm0.006$ & $0.107\\pm0.006$\\\\ \\hline $H_0$          &  $74\\pm4$           &  $72\\pm2$          & $71\\pm6$           & $73\\pm3$          & $73\\pm5$ & $74\\pm3$\\\\ \\hline $\\tau$         &  $0.087\\pm0.017$    &  $0.083\\pm0.017$   & $0.085\\pm0.018$    & $0.088\\pm0.017$    & $0.086\\pm0.016$ & $0.085\\pm0.018$\\\\ \\hline $n_s$          &  $0.975\\pm0.019$    &  $0.964\\pm0.015$   & $0.962\\pm0.014$    & $0.964\\pm0.014$   & $0.963\\pm0.014$ & $0.963\\pm0.015$\\\\ \\hline $\\log(10^{10}A_s)|_{k=0.002}$ &  $3.15\\pm0.06$ &$3.19\\pm0.05$   & $3.18\\pm0.05$      & $3.17\\pm0.04$     & $3.18\\pm0.05$ & $3.17\\pm0.05$\\\\ \\hline $\\Omega_\\Lambda$ &  $0.76\\pm0.03$      &  $0.74\\pm0.02$     & $0.5\\pm0.2$        & $0.3\\pm0.2$     & $0.5\\pm0.2$ & $0.2\\pm0.2$\\\\ \\hline $\\alpha$       &  $(6\\pm9)\\times10^{-3}$ & $(1.9\\pm1.4)\\times10^{-2}$ & $-1.1\\pm0.4$ & $-1.0\\pm0.3$      & $-1.2\\pm0.4$ & $-1.1\\pm0.2$\\\\ \\hline \\end{tabular} \\end{center} \\end{table*}  As a first result, we obtained much tighter constraints on the parameters of the model with respect to the analysis carried on the background observables in \\cite{LADM07}, confirming that perturbations should be properly included in the calculations when  developing effective models for the dark sector \\cite{ValiviitaMajerottoMaartens2008}.  In the case when the fluid is treated as a unified dark component, we get values of the effective cosmological constant $\\Omega_\\Lambda\\simeq 0.7$, essentially independent of the speed of sound. For the equation of state parameter $\\alpha$, the constraints vary when the fluid is treated as barotropic (resulting in a slightly positive $\\alpha$) or a vanishing speed of sound is assumed (resulting in a slightly negative $\\alpha$). Both cases are however compatible with $\\alpha=0$ at one sigma confidence level. The inclusion of the matter power spectrum in the analysis has generally the effect of shrinking the confidence interval on the parameters, in particular in the barotropic case, due to the effect of $\\alpha$ on the Jeans length of the perturbations \\cite{quercellini_fields07}.  When standard dark matter is included, the effects of $\\alpha$ on the clustering process is less relevant, because the matter-like component of the unified fluid is forced to mimic the cosmological constant behaviour. This is apparent from the fact that the $\\alpha$ best fit value moves to $\\alpha\\sim -1$ which is typical of a cosmological constant. The constraints in the barotropic case remain quite tight, but get larger when the sound speed is set to zero. We also considered the case with a sound of speed equal to unity, which describes a scalar field behaviour. Also in this case the constraints on $\\alpha$ are rather loose."
    },
    "0807/0807.0464_arXiv.txt": {
        "abstract": "We analyse the local variance effect in the standard method for detecting the integrated Sachs-Wolfe effect (ISW) via cross-correlating the cosmic microwave background (CMB) with the large-scale structure (LSS). Local variance is defined as the systematic noise in the ISW detection that originates in the realization of the matter distribution in the observed Universe. We show that the local variance contributes about 11 per cent to the total variance in the standard method, if a perfect and complete LSS survey up to $z \\approx 2$ is assumed. Due to local variance, the estimated detection significance and cosmological parameter constraints in the standard method are biased. In this work, we present an optimal method of how to reduce the local variance effect in the ISW detection by working conditional on the LSS data. The variance of the optimal method, and hence the signal-to-noise ratio, depends on the actual realization of the matter distribution in the observed Universe. We show that for an ideal galaxy survey, the average signal-to-noise ratio is enhanced by about 7 per cent in the optimal method, as compared to the standard method. In the framework of our method, it is straightforward to correct for the magnification bias coming from gravitational lensing effects. Furthermore there is no need to estimate the covariance matrix by Monte Carlo simulations as in the standard method, which saves time and increases the accuracy. Finally, we derive the correct joint likelihood function for cosmological parameters given CMB  and LSS data within the linear LSS formation regime, which includes a small coupling of the two datasets due to the ISW effect. ",
        "introduction": "\\label{intro}  The integrated Sachs-Wolfe (ISW) effect \\citep{isw} is an important probe of the existence and nature of dark energy \\citep{crittenden_turok} and the nature of gravity \\citep{lue_gravity,zhang_gravity}. Spatial curvature also gives rise to an ISW effect, but is well constrained to be very close to zero by cosmic microwave background (CMB) experiments such as the {\\it Wilkinson Microwave Anisotropy Probe} \\citep{wmap_5}. However, the detection of the ISW signal remains challenging, for it is obscured by primordial fluctuations in the CMB. In recent years, substantial effort has been made to detect the ISW effect via cross-correlation of the CMB with large-scale structure (LSS) surveys, such as optical galaxy and quasar surveys (Sloan Digital Sky Survey, \\cite{sdss_6}, and Two-Micron All-Sky Survey, \\cite{2mass}), radio surveys (NRAO VLA Sky Survey, \\cite{nvss}), and X-ray surveys (High Energy Astrophysics Observatory, \\cite{heao}). Such cross-correlation studies have, for example, been done by \\cite{x_ray_boughn}, \\cite{boughn_crittenden_nature}, \\cite{ boughn_crittenden_detection}, \\cite{afshordi}, \\cite{rassat}, \\cite{ reassessment}, \\cite{mcewen}, \\cite{pietrobon_NVSS}, \\cite{ fosalba_1}, \\cite{fosalba_2}, \\cite{vielva}, \\cite{liu}, \\cite{ ho} and \\cite{giannantonio}, just to name a few of them.  The standard method for detecting the cross-correlation between the LSS and the CMB involves comparing the observed cross-correlation function with its theoretical prediction for a given fiducial cosmological model. The theoretical prediction is by construction an ensemble average over all possible realizations of the universe given the fiducial parameters, including all possible realizations of the local matter distribution. Assuming ergodicity, this ensemble average can also be thought of as an average over all possible positions of the observer in the Universe ('cosmic mean').  However, the ISW effect is created by the decay of the gravitational potential coming from the structures on the largest scales, i.e. from structures that have not yet undergone significant gravitational collapse and are still not decoupled from the expansion of the Universe. These largest scales are most affected by cosmic variance. Therefore, when comparing the observed (local) cross-correlation function to its cosmic mean value, the realization of the matter distribution in our vicinity acts as a source of systematic noise in the estimation of the cross-correlation, hence leading to a biased detection significance, due to cosmic variance.  In this work, we estimate the contribution of this local variance effect to the total variance in the detected signal under the simplifying assumption that there is no shot noise in the galaxy distribution. We find that the local variance in the detected signal amounts to 11 per cent in the case of an ideal LSS survey going out to about redshift 2 and covering enough volume to include the large scales relevant for the ISW. This agrees with \\cite{cabre}, who compare different methods to estimate the error in standard cross-correlation studies. They find that their MC1 error, which ignores the variance coming from the realization of the matter field and only considers the variance in the CMB fluctuations, systematically underestimates the error by about 10 per cent.  From the above-mentioned surveys, the local matter distribution is known to a certain degree, and hence the local variance effect can be reduced by working conditional on that information. We present a generic technique of how to include the knowledge of the matter distribution into the detection of the ISW via cross-correlation, hence reducing the sources of noise to the unknown part of the matter distribution and the primordial CMB fluctuations. We define the systematic noise that comes from the known part of the matter distribution as bias, for it can be removed by working conditional on the LSS data. Our method is referred to as optimal method, in contrast to the standard method for ISW detection mentioned above. The main idea of the optimal method is to create an ISW template from a Wiener filter reconstruction of the LSS. We then use this template to detect the amplitude of the ISW template rather than of the theoretical cross-correlation function. This makes the variance in the estimated amplitude, and hence the signal-to-noise ratio both depend on the actual realization of the matter in the Universe. For an ideal LSS survey, we show that the average variance in the detected amplitude is reduced by 13 per cent in the optimal method. Rephrased in terms of the signal-to-noise ratio, this reduction of the noise leads on average to a higher detection significance by about 7 per cent. As we show in this work, in the framework of the optimal method it is straightforward to correct for the magnification bias due to gravitational lensing, as described by \\cite{LoVerde}. Furthermore, there is no need to estimate the covariance matrix by Monte Carlo simulations as in the standard method. This saves time and increases the accuracy of the method.  To our knowledge, the only suggested methods for ISW detection besides this work that also do not suffer from the local variance effect are by \\cite{zhang} and \\cite{carlos}. The former involves combining lensing-LSS cross-correlation measurements with the ISW-LSS cross-correlation, and thereby relies on the nowadays still-difficult lensing measurements \\citep{hu_lensing,hu_okamoto}. The latter follows an approach very similar to ours. However, \\cite{carlos} does not use a Wiener filter reconstruction of the LSS distribution but works directly on the sphere. In contrast to our work he neglects the shot noise, which could be easily included in his analysis, though. In this work we will go one step further than \\cite{carlos} and derive the correct way of including the information encoded in the ISW for cosmological parameter estimation.  Many of the above-mentioned cross-correlation studies have attempted to constrain cosmological parameters using a likelihood function for the cosmological parameters $p$ given the observed cross-correlation function. Just like the detection significance, these parameter estimates are biased by local variance. Furthermore, to our knowledge, there is no straightforward way of combining the likelihood function for the cross-correlation with the likelihoods for CMB  and LSS data. In this work, we derive the correct joint likelihood function $P(T,\\delta_g \\,|\\, p)$ for cosmological parameters, given the CMB map $T$ and the LSS data $\\delta_g$, from first principles for the linear LSS formation regime. This joint likelihood consistently includes the coupling between the two datasets introduced by the ISW effect, which so far has been neglected in analyses deriving cosmological parameter constraints by combining CMB and LSS data \\citep{sdss_wmap,wmap_3}.  The article is organized as follows. We start by briefly describing the integrated Sachs-Wolfe effect in section \\ref{isw}, and explain in detail the different stochastic processes that are relevant for our analysis and the correction for the magnification bias in section \\ref{stochastic}. In section \\ref{standard} we review the standard method for detecting the ISW via cross-correlation and estimate the contribution of the local variance to the total variance of the detected signal. Section \\ref{optimal_method} is devoted to presenting the optimal method of ISW detection we developed, and to comparing it in detail to the standard method. We discuss the role of the biasing effect due to local variance in parameter constraints and derive the joint likelihood function $P(T,\\delta_g \\,|\\, p)$ in section \\ref{params}. Concluding remarks on our work are given in section \\ref{conclusions}. ",
        "conclusions": "\\label{conclusions}  Due to the obscuration by primordial CMB fluctuations, the detection of the integrated Sachs-Wolfe effect remains a challenge, and has to be performed by cross-correlating the CMB signal with the large-scale structure. The standard method for doing so involves comparing the observed cross-correlation function to its theoretical prediction, which is by construction an ensemble average over all realizations of the primordial CMB fluctuations and matter distributions. Hence, the realization of the matter distribution in our Universe acts as a source of systematic noise in the estimate of the cross-correlation function, an effect that we have named the local variance.  Since the ISW is only present on the largest scales, the effect of the local variance is quite notable, amounting to about 11 per cent of the total variance in the standard method for an ideal LSS survey. This leads to a biased estimated detection significance of the cross-correlation, and when moving from the pure ISW detection to parameter estimation, it also biases the parameter constraints. We note that even if the local variance contributes only about 11 per cent to the total variance of the detected signal, we could be unlucky and live in an unlikely realization of the matter distribution, given the power spectrum. This would enhance the effect of the bias on the detection significance and parameter constraints.   Given that information about the matter distribution can be inferred from the LSS survey, the local variance can be reduced by working conditional on this information. In  this work, we have presented a generic technique of how to include the knowledge of the matter distribution into ISW detection in an optimal way, hence reducing the effect of the local variance. This optimal method requires a three-dimensional Wiener filter reconstruction of the LSS, including an estimator of the full reconstruction uncertainty covariance matrix. Note that also other reconstruction techniques that provide an estimator of the uncertainty covariance can easily be included into our method. The reduction of the local variance stresses the importance to measure and reconstruct the LSS to the highest possible accuracy, as aimed by \\cite{paco} and Kitaura et al. (in preparation).  The conditionality on the LSS data results in a dependence of the variance in the detected signal on the actual realization of the matter distribution in the observed Universe. The average variance in the optimal method is reduced by about 13 per cent as compared to the standard method, again in the case of an ideal LSS survey. The reduction of the noise translates into an average enhancement of the signal-to-noise or detection significance by about 7 per cent for the optimal method. However, note that also the signal-to-noise ratio depends on the actual realization of the matter distribution.  We would also like to point out that in the optimal method, there is no need to estimate the covariance matrix by Monte Carlo simulations, which safes time and increases the accuracy of the method (using 1000 Monte Carlo simulations to estimate the standard covariance matrix of the cross-correlation function only reaches an accuracy of about 5 per cent, as stated by \\cite{cabre}).  In order to consistently include the information encoded in the ISW effect for deriving cosmological parameter constraints, we have derived the joint likelihood $P(T,\\delta_g \\,|\\, p)$ for the cosmological parameters $p$, given CMB  and LSS data within the linear regime of structure formation. If one wishes to use the ISW effect for constraining cosmological parameters, one should include the additional CMB-galaxy data coupling term $Q(T,\\delta_g \\,|\\, p)$, which we have factored out in eq. (\\ref{likelihood}), into the usual likelihood analysis."
    },
    "0807/0807.0008_arXiv.txt": {
        "abstract": "Duration and period of transits in extrasolar planetary systems can exhibit long-term variations for a variety of reasons. Here we investigate how systemic proper motion, which steadily re-orients planetary orbit with respect to our line of sight, affects the timing of transits. We find that in a typical system with a period of several days proper motion at the level of 100 mas yr$^{-1}$ makes transit duration vary at a rate $\\sim 10-100$ ms yr$^{-1}$. In some isolated systems this variation is at the measurable level (can be as high as $0.6$ s yr$^{-1}$ for GJ436) and may exceed all other transit timing contributions (due to the general relativity, stellar quadrupole, etc.). In addition, proper motion causes evolution of the observed orbital period via the Shklovskii effect at a rate $\\gtrsim 10$ $\\mu$s yr$^{-1}$ for the nearby transiting systems ($0.26$ ms yr$^{-1}$ in GJ436), which in some cases exceeds all other contributions to $\\dot P$. Earth's motion around the Sun gives rise to additional periodic timing signal (even for systems with zero intrinsic proper motion) allowing a full determination of the spatial orientation of the planetary orbit. Unlike most other timing effects the proper motion signatures persist even in systems with zero eccentricity and get stronger as the planetary period increases. They should be the dominant cause of transit timing variations in isolated wide separation (periods of months) systems that will be sought by {\\it Kepler}. ",
        "introduction": "\\label{sect:intro0}   Planetary transits have provided us with a unique opportunity to get a handle on the physical properties of the extrasolar planets such as their radii and densities. Recently it has been suggested (Miralda-Escud\\'e 2002; Heyl \\& Gladman 2007) that precision timing of the moments at which transits occur can give us additional information about the transiting systems. Various physical effects cause orbit of the planet precess in space leading to the changes in transit geometry, which can be measured through the timing of transits. Among these effects are the general relativistic precession of the orbit, gravitational influence of other planets in the system or companion stars, torques due to the spin-induced quadrupole moment of the star and due to the tidal deformations of both the star and the planet (Miralda-Escud\\'e 2002; Heyl \\& Gladman 2007; Ribas \\etal 2008; P\\'al \\& Kocsis 2008; Jord\\'an \\& Bakos 2008).  Another obvious reason for the re-orientation of the planetary orbit with respect to observer at Earth is the proper motion of the exoplanetary system with respect to the Solar System. Stars in the Solar neighborhood move at velocities of tens of km s$^{-1}$ and some of them exhibit proper motion at the level of $1$ mas yr$^{-1}$. Also, the distance to stars constantly changes as a result of their relative motion with respect to the Sun and this affects transit timing because of the finite speed of light. At some level proper motion is a characteristic of any star, including those with transiting exoplanets, and it is thus important to understand its implications for transit timing.  Proper motion is well known to be important in the timing of isolated and binary radio pulsars (Shklovskii 1970; Kopeikin 1996). In these systems proper motion affects the pulsar spin and orbital periods through the so-called Shklovskii effect (Shklovskii 1970) while the re-orientation of the binary orbit can be (and has been) measured via the variation of the projected size of the orbit (Kopeikin 1996; Arzoumanian \\etal 1996). Pulsar acceleration in external gravitational field can also be important especially for pulsars in globular clusters (Edwards \\etal 2006).  Of course, there are significant differences between the timing of pulsars and of planetary transits: accuracy with which some millisecond pulsars can be timed is at the $\\mu$s level (Manchester 2008) while a single planetary transit can only be timed to several seconds at best (Knutson \\etal 2007). Also, the whole idea of timing is different in the two cases: for binary pulsars one is usually able to trace the whole orbit of the neutron star in time domain while in the case of planetary transits only two narrow time windows --- primary and secondary transits --- are available to play with. Nevertheless, some of the ideas developed in pulsar timing may be applied to the timing of planetary transits.  Previously, Kopeikin \\& Ozernoy (1999) have utilized a post-newtonian relativistic approach for the precision Doppler measurements of the binary star orbits and discussed some of the relevant effects of the proper motion. Here we aim at investigating the role of the proper motion in timing of planetary transit duration and period in extrasolar planetary systems. We lay out the basics of the orbital element evolution due to the proper motion in \\S \\ref{sect:intro} . We discuss the evolution of the transit duration in \\S \\ref{sect:transit} and the evolution of the period between transits in \\S \\ref{sect:period}. Comparison with other transit timing effects and application to real systems can be found in \\S \\ref{sect:disc}. ",
        "conclusions": "\\label{sect:concl}  We investigated the effect of the proper motion on timing of transiting exoplanetary systems in which the gravitational effect of other possible companions can be neglected. Proper motion re-orients planetary orbit with respect to our line of sight and changes the distance to the system. Short period transiting systems having proper motion at the level of $100$ mas yr$^{-1}$ should exhibit variation of the transit duration at the level of $\\sim 100$ ms yr$^{-1}$, which may be comparable to or exceed the timing signatures produced by the general relativity or stellar quadrupole and which should not be hard to detect. Proper motion also causes variation of the observed orbital period through the Shklovskii effect which dominates $\\dot P$ for high proper motion systems. Orbital motion of the Earth around the Sun gives rise to periodically varying transit timing signal even in systems having zero intrinsic proper motion. Timing effects induced by the proper motion become especially important in systems with zero eccentricity and in wide separation systems with periods longer than a month which should be discovered by {\\it Kepler}."
    },
    "0807/0807.0232.txt": {
        "abstract": "Stochastic microlensing is a central tool in probing dark matter on galactic scales. From first principles, we initiate the development of  a mathematical theory of stochastic microlensing. Beginning with the  random time delay function and associated lensing map, we determine exact expressions for the  mean and variance of these transformations. In addition, we derive the probability density function (p.d.f.) of a random point-mass potential, which form the constituent  of a stochastic microlens potential.  We characterize the exact p.d.f. of a normalized random time delay function at the origin, showing that it is a shifted gamma distribution, which also holds  at leading order in the limit of a large number of point masses if  the normalized time delay function was at a general point of the lens plane. For the large number of point masses limit, we also prove that the asymptotic p.d.f. of the random lensing map under a specified scaling converges to a bivariate normal distribution. We show analytically that the p.d.f. of the random scaled lensing map at leading order depends on the magnitude of the scaled bending angle due purely to point masses as well as demonstrate explicitly how this radial symmetry is broken at the next order. Interestingly, we found at leading order a formula linking the expectation and variance of the  normalized random time delay function to the first Betti number of its domain. We also determine an asymptotic p.d.f. for the  random bending angle vector and find an integral expression for the probability of a lens plane point being near a fixed point. Lastly, we show explicitly how the results are affected by location in the lens plane. The results of this paper are relevant to the theory of random fields and provide a platform for further generalizations as well as analytical limits for checking astrophysical studies of stochastic microlensing. ",
        "introduction": "Stochastic microlensing is the study of the deflection of light by a random collection of stars. In recent years, this subject has become  a central tool in understanding the nature and distribution of dark matter on galactic scales (e.g.,  Schechter and Wambsganss \\cite{Schechter1}).  This has raised the need for a rigorous framework for stochastic microlensing.  Research on stochastic microlensing has been mainly numerical, or at least semi-analytical, in flavor. Work in the early 1980's included an analytical study of certain stochastic effects  of microlensing by a single random lens  (e.g., Vietri and Ostriker \\cite{Vietri}). This was later extended to the case of multiple lenses  using numerical simulations (e.g., Kayser, Refsdal, and Stabell \\cite{Kayser}; Paczynski \\cite{Pac}).  While successful in furthering our picture of stochastic microlensing, these simulations often cannot be checked analytically and are time-consuming. Some of the first analytical counterparts of the numerics were focused  on the study of  random surface brightnesses, bending angles, magnifications, and redshifts (e.g., Deguchi and Watson \\cite{Deguchi,Deguchi2}; Katz, Balbus, and Paczynski  \\cite{Katz};  Schneider \\cite{Schneider}; Turner, Ostriker, and Gott \\cite{Turner};  Vietri \\cite{Vietri1}). This was followed by numerical and semi-rigorous work characterizing the  magnification probabilities and the expected number of microimages in a variety of  lensing situations (e.g., Rauch {\\it et al.} \\cite{Rauch}; Wambsganss \\cite{Wambsganss}; Wambsganss, Witt, and Schneider \\cite{wamb}). Recently, similar techniques were employed to study the distributions of microimages, magnifications, and  time delay differences between images, with the last two   applied directly to probe dark matter using flux-ratio anomalies (e.g., Granot, Schechter, and Wambsganss \\cite{GSW};  Keeton, and Moustakas \\cite{Keeton}; Schechter, and Wambsganss \\cite{Schechter1};  Schechter, Wambsganss, and Lewis \\cite{Schechter2}; Tuntsov {\\it et al.} \\cite{Tuntsov}). In a more general context,  Berry and collaborators also explored the statistical aspects of caustics in optical systems for the case of Gaussian random fields (see \\cite{Berry} and references therein).    This series  aims to  develop a mathematical theory of stochastic microlensing, starting from first principles. This fits in the general theory of random fields, namely, the study of random functions and  their critical point structure (for more on random fields, see Adler and Taylor \\cite{Adlerjon} and references therein). We derive new stochastic results, for a general point in the lens plane,  about random time delay functions and lensing maps, which are transformations forming the core of microlensing theory. We determine the exact (as opposed to asymptotic) expectation and variance of both the lensing map and the point mass potential. Using the expected value of the components of the  lensing map, we deduce, for a fixed image in the lens plane,  the corresponding average position of the point it is mapped to in the light source plane. Furthermore, since  point mass potentials form the density components of the  time delay function, we use the  exact formulas of their first and second moments to derive the asymptotic expectation and variance of the normalized time delay function in the large number of stars limit. We also reveal  an interesting topological link, namely, at leading order we show that the expectation and variance of the normalized random microlensing time delay function is related to  the first Betti number of the time delay functions's domain.  Continuing our study of the normalized random  time delay function, we establish that at the origin its exact p.d.f. is a shifted gamma distribution, and  that this is asymptotically the case for an arbitrary  point within the lens plane  in the large number of stars limit. In the same limit, we prove that the p.d.f of a scaled lensing map converges to a bivariate Gaussian, with once again consistent asymptotic mean and  variance. This analytical work shows  that the p.d.f. of the random scaled lensing map at leading order  term depends on the square of  magnitude of the scaled bending angle due purely to point masses. It also illustrate how  this radial symmetry is broken at the next order. Similarly, we note that the p.d.f. of the random time delay function at the origin depends only on the potential time delay due purely to point masses, a result also holding at leading order for an arbitrary point in the lens plane.   Taking our analysis further, we compute the probability that the image of  a given point under the lensing map will be  in a given disk  in light the source plane, which provides information about the light source position.  We deduce from these results the expectation, variance, and asymptotic p.d.f. of the bending angle vector at an arbitrary position in the lens plane. We also use these results to compute the probability that a  point in the lens plane will be mapped in a certain neighborhood of ``itself'' in the light source plane. This study relates to the theory of fixed points, which was studied earlier by  Petters and Wicklin \\cite{Fixe} in the case of a deterministic lens system.  Overall, these first steps lay the foundation in terms of concepts, results, and tools for future analytical work on random time delay functions and lensing mappings, including lensing observables like magnification, arising from more complex stochastic mass distributions. The work can also help as a guide on tractable aspects of the general theory of random mappings.  We  state the basic random microlensing framework in Section~\\ref{basic}. In Section~\\ref{time}, we present the results of our study of the   random time delay function, obtaining its expectation, variance, exact p.d.f. at the origin, and asymptotic p.d.f. at an arbitrary lens plane point.  A similar presentation of results is done  for  the random lensing map   in Section~\\ref{lens}. In Section~\\ref{angle}, we apply our results about random lensing maps to random bending angle vectors and fixed points. The detailed proofs of these results are given in the Appendix.  Finally, the second paper in this program will explore the random shear and expected number of microimages produced by stochastic microlensing. ",
        "conclusions": "We presented first steps in the development of a mathematical theory  of stochastic microlensing that focused on the building blocks of the theory, namely, the random time delay function and random lensing map.  We derived exact analytical formulas for the expectation and  variance of the random time delay function  and random lensing map about any point. In the large limit of stars, we found  a simple asymptotic expression for the expectation and variance of  a normalized time delay function at an arbitrary point. For the same limit, we highlighted an interesting link between the leading order term of the expectation and variance of our  normalized random time delay function and the first Betti number of its domain. In addition, the asymptotic p.d.f.s of both the  normalized time delay function and the scaled lensing map were characterized in the large number of stars limit: The asymptotic p.d.f. of the former is a shifted gamma density at leading order, while for the latter the p.d.f. is a bivariate Gaussian distribution. The p.d.f. of the random scaled lensing map is also shown to depend on the magnitude of the scaled bending angle due purely to point masses at leading order and we illustrated explicitly how this radial symmetry fails at the next order. We also derived the asymptotic p.d.f. of the random bending angle vector and gave an estimate of the probability  of a lens plane point being close to a fixed point. Overall, the paper determined and illustrated new analytical results about the microlensing behavior of random time delay functions and random lensing maps about an arbitrary point of the lens plane.  In Paper II of this series, we shall explore the microlensing random shear and expected number of  microimages."
    },
    "0807/0807.4889_arXiv.txt": {
        "abstract": "{The solar chromosphere and transition region are highly structured and complex regimes. A recent breakthrough has been the identification of dynamic fibrils observed in \\Ha\\ as caused by field-aligned magnetoacoustic shocks.} {We seek to find whether such dynamic fibrils are also observed in \\Lya.} {We used a brief sequence of four high-resolution \\Lya\\ images of the solar limb taken by the Very high Angular resolution ULtraviolet Telescope (VAULT), which displays many extending and retracting \\Lya\\ jets. We measured their top trajectories and fitted parabolas to the 30 best-defined ones.} {Most jet tops move supersonically. Half of them decelerate, sometimes superballistically, the others accelerate. This bifurcation may arise from incomplete sampling of recurrent jets.} {The similarities between dynamic \\Lya\\ jets and \\Ha\\ fibrils suggest that the magnetoacoustic shocks causing dynamic \\Ha\\ fibrils also affect dynamic \\Lya\\ jets.} ",
        "introduction": "The solar chromosphere observed in \\Ha\\ is a bewildering mass of elongated features, but a breakthrough concerning so-called dynamic fibrils (henceforth \\Ha\\ DFs) has been made in the studies of \\citet{DePontieuetal2004,DePontieuetal2007a}, \\citet{Hansteenetal2006}, \\citet{RouppevanderVoortetal2007}, \\citet{Hegglandetal2007}, and \\citet{Langangenetal2008a, Langangenetal2008b} following earlier work by \\citet{Suematsuetal1995}. These studies have established that \\Ha\\ DFs, which are rows of dark fibrilar features jutting out from plage and network with periodic extension and retraction, display repetitive mass loading by upward propagating magnetoacoustic shock waves driven by the global solar oscillations. Reduction of the effective gravity along tilted magnetic channels lowers their cutoff frequency and lets them propagate into the chromosphere, steepen into shocks, and repetitively lift the chromospheric-transition region interface.  These studies are all based on optical observations, but \\citet{deWijnandDePontieu2006} have studied transition-region jets in \\ion{C}{iv} in ultraviolet TRACE images and found remarkable morphological similarities between \\Ha\\ DFs and their \\ion{C}{iv} counterparts. The DFs may represent the injection of cool material postulated by \\citet{Judge2008} as the source of hot sheaths making up the transition region.  In this paper we report the presence of similar features in solar \\Lya\\ images taken with the Very high Angular resolution ULtraviolet Telescope \\citep[VAULT,][]{Korendykeetal2001} during its second rocket flight. VAULT acquired \\Lya\\ images with much higher spatial and temporal resolution than the EUV imagery from the TRACE, SoHO, and STEREO satellites but only for a few minutes. The \\Lya\\ fine structure in these images is discussed by \\cite{Patsourakosetal2007} and \\cite{JudgeandCenteno2008}. Here we locate and study extending or retracting \\Lya\\ brightness structures in images taken near the limb. We call these dynamic \\Lya\\ jets, abbreviated to \\Lya\\ DJs, and compare them to \\Ha\\ DFs.   \\begin{figure} \\centering \\resizebox{\\hsize}{!}{\\includegraphics{small-0710fig1}} \\caption[]{ \\Lya\\ image taken at 18:16:56\\,UT covering the east limb. We invite the reader to inspect this image at large magnification with a pdf viewer. \\Lya\\ jets are grouped in hedgerows at network boundaries that are best seen in projection against the dark internetwork. The rectangle selects the hedgerow enlarged in Fig.~\\ref{fig:closeup}. The contrast is enhanced by logarithmic brightness scaling. Tickmark spacing: 10\\,arcsec. }  \\label{fig:image} \\end{figure}  \\begin{figure} \\centering \\resizebox{\\hsize}{!}{\\includegraphics{small-0710fig2}} \\caption[]{ Enlargement of the subfield outlined in Fig.~\\ref{fig:image}. The rectangle outlines the smaller subfield shown in Fig.~\\ref{fig:evolution} with \\Lya\\ DJs 3 and 11. Tickmark spacing: 1\\,arcsec. }  \\label{fig:closeup} \\end{figure}  \\begin{figure} \\centering \\resizebox{\\hsize}{!}{\\includegraphics{small-0710fig3}} \\caption[]{ The upper part of Fig.~\\ref{fig:image} with 55 \\Lya\\ DJs marked by plus signs. The 30 DJs selected for further analysis are labeled with their reference number. Tickmark spacing: 10\\,arcsec. }  \\label{fig:55DJs} \\end{figure}  \\begin{figure} \\centering \\resizebox{\\hsize}{!}{\\includegraphics{small-0710fig4}} \\caption[]{ Time evolution of extending (nrs.~22, 19, 20, 3) and retracting (nrs.~11, 10, 9, 24) \\Lya\\ DJs. The time increases with 17-s intervals between the columns. The arrows define the jet axis. The tickmarks have 1-arcsec spacing. }   \\label{fig:evolution} \\end{figure}  \\begin{figure} \\centering \\resizebox{0.967\\hsize}{!}{\\includegraphics{small-0710fig5}} \\caption[]{ Intensity profiles along the longitudinal axes of the \\Lya\\ DJs shown in Fig.~\\ref{fig:evolution}. The solid, dotted, dashed, and dashed-dotted curves correspond to the 1st, 2nd, 3rd, and 4th panel in the corresponding row of Fig.~\\ref{fig:evolution}, respectively, and show the temporal evolution. The thin-line markers specify the ranges of intensity thresholds defining each jet top. The reference footpoints with $\\Lambda=0$ are fixed in time. }  \\label{fig:profiles} \\end{figure} ",
        "conclusions": "\\label{sec:discussion}  Obviously, longer \\Lya\\ image sequences at the VAULT-II resolution are needed to gain better information than these four images provide. Nevertheless, they do show the existence of dynamic jets that fan out from network and extend or retract during the one-minute image sequence.  How do these \\Lya\\ DJs compare to the now well-studied and fairly well-explained \\Ha\\ DFs? They have similar shapes, occupy similar locations on the solar surface, show similar hedge-row morphology, and extend or retract along their length the same way. The measurements for the 16 decelerating ones in Table~\\ref{tab:measurements} show fairly good correspondence with the \\Ha\\ DF measurements listed at the bottom of the table. Figure~\\ref{fig:velocities} suggests a positive correlation between $v_{\\rm max}$ and $|a|$ comparable to the linear relations for \\Ha\\ DFs. Our estimate of the mean \\Lya\\ DJ excursion duration is about one and half minutes shorter than the average \\Ha\\ DF periodicity reported by \\citet{DePontieuetal2007a}, but represents only a lower limit. Three of the sixteen DJs show substantially greater superballistic deceleration than the maximum of 320\\,m\\,s$^{-2}$ for DFs reported by \\citet{DePontieuetal2007a} and the maximum of 400\\,m\\,s$^{-2}$ reported by \\citet{RouppevanderVoortetal2007} and \\citet{DePontieuetal2007c} for quiet-Sun mottles and for type-1 spicules, respectively, but with large uncertainty for one of them. We conclude that, overall, our decelerating \\Lya\\ DJs are rather similar to \\Ha\\ DFs.  However, in contrast to the \\Ha\\ DF studies, we found a roughly equal number of \\Lya\\ DJs that significantly accelerate rather than decelerate. These are collected in the lower half of Fig.~\\ref{fig:trajectories}. The two sets of trajectories together suggest that they might actually be incomplete samples of sinusoidal motion or, as the referee has pointed out, sample the end of one shock-driven feature and the onset of the next one if these \\Lya\\ DJs occur in succession as \\Ha\\ DFs do. Note that our length estimates do not differ significantly between the two sets, but these are rather arbitrary.  An obvious difference between \\Lya\\ DJs and \\Ha\\ DFs is that the former appear bright against dark backgrounds, the latter dark against bright backgrounds. \\Lya\\ emissivity indicates the presence of gas that is sufficiently hot for collisional excitation of its upper level at 10.2~eV. Sufficient population of that level is also required to make fibrils opaque in \\Ha. The recent study by \\citet{Leenaartsetal2007a} suggests that cool fibrils may maintain much greater \\Ha\\ opacity than their temperature would suggest if they have recently undergone shock heating. Hence, \\Ha\\ DFs may remain opaque also when they are cool during part of their lifetime. They will be much more opaque in \\Lya. Can the same fibril appear bright in \\Lya\\ and dark in \\Ha?  Perhaps through hot-sheath topology of the transition region, with the latter surrounding cool DFs as hot shells \\citep[e.g.,][]{Rutten2007,DePontieuetal2007a}. The neutral-hydrogen diffusion mechanism recently proposed by \\citet{Judge2008} seems a viable candidate. Detailed time-dependent numerical MHD simulation may demonstrate this. However, realistic evaluation of \\Ha\\ formation from such simulations remains a formidable challenge that may be helped by constraints derived from \\Lya.  Finally, the \\Lya\\ DJs studied here are likely related to chromospheric phenomena such as regular spicules \\citep[e.g.][]{Sterling2000}, straws \\citep{Rutten2006,Rutten2007}, and type-2 spicules \\citep{DePontieuetal2007c,Langangenetal2008c}."
    },
    "0807/0807.3651_arXiv.txt": {
        "abstract": "{We have performed a 70 billion dark-matter particles N--body simulation in a 2 $h^{-1}$  Gpc periodic box, using the concordance, cosmological model as  favored by the latest WMAP3  results. We have computed  a full-sky convergence  map with  a resolution  of $\\Delta \\theta \\simeq 0.74$ arcmin$^{2}$,  spanning 4 orders of magnitude in angular dynamical  range.  Using various high-order  statistics on a realistic  cut sky, we  have characterized  the transition  from the linear to the nonlinear regime  at $\\ell \\simeq 1000$ and shown that realistic  galactic masking affects  high--order moments  only below $\\ell  < 200$.   Each domain  (Gaussian and  non--Gaussian)  spans 2 decades in angular  scale.  This map is therefore  an ideal tool for testing map--making  algorithms on  the sphere. As  a first  step in addressing the full map  reconstruction problem, we have benchmarked in this paper two denoising  methods: 1) Wiener filtering applied to the  Spherical Harmonics  decomposition  of  the map  and  2) a  new method,  called MRLens,  based on  the modification  of  the Maximum Entropy  Method on  a Wavelet  decomposition.  While  the  latter is optimal  on large  spatial  scales, where  the  signal is  Gaussian, MRLens outperforms the Wiener  method on small spatial scales, where the  signal is  highly  non--Gaussian.  The  simulated full-sky convergence map  is freely  available to the  community to  help the development  of new  map--making  algorithms dedicated  to the  next generation of weak--lensing surveys.} ",
        "introduction": "Weak  gravitational lensing,  or ``cosmic  shear'', provides  a unique tool for mapping the matter  density distribution in the Universe (for reviews,     see     \\cite{refregier03},     \\cite{hoekstra03}     and \\cite{munshi06}).  Current weak-lensing surveys cover altogether about 100 square degrees and have been  used to measure the amplitude of the matter   power  spectrum  and   other  cosmological   parameters  (see \\cite{benjamin07}, \\cite{fu08}  and references therein).   A number of new  instruments are  being planned  to carry  out these  surveys over wider  sky areas (PanSTARRS,  DES, SNAP  and LSST)\\footnote{PanSTARRS: http://pan-starrs.ifa.hawaii.edu,                                DES: https://www.darkenergysurvey.org,   SNAP:   http://snap.lbl.gov  and LSST: http://www.lsst.org}  or even over the  full extragalactic sky (DUNE\\footnote{DUNE:  http://www.dune-mission.net}).  These wide-field surveys  will yield  cosmic-shear measurements  on both  large scales, where  gravitational  dynamics is  in  the  linear  regime, and  small scales,  where the dynamics  is highly  nonlinear.  The  comparison of these measurements  with theoretical predictions of  the density field evolution  will place strong  constraints on  cosmological parameters, including dark energy  parameters (e.g. \\cite{hu99}, \\cite{huterer01}, \\cite{amara06}  and \\cite{albrecht07}).  On  small scales,  the highly nonlinear nature  of the density field ensures  that predictions based on analytic calculations are  prohibitively difficult and requires the use of numerical simulations.   N-body simulations have thus been used to simulate weak-lensing  maps across small patches of  the sky, using the  flat sky approximation  (e.g. \\cite{jain00},  \\cite{hamana01} and \\cite{vale04}).  The  simulation of  full-sky maps in  preparation for future surveys involve a wide range of both mass and length scales and is challenging  for current N-body  simulations.  The range  of scales involved  also requires  the development  of efficient  algorithms for deriving a mass map from  true noisy data sets.  These algorithms need to  be  well--suited  to   both  the  large--scale  signal,  which  is essentially a Gaussian random field, and those on small--scales, where it is highly non--Gaussian and exhibits localized features.  \\begin{figure*} \\centering \\resizebox*{!}{9cm}{\\includegraphics{0657fig1.eps}} \\vskip -0.1cm \\caption{Full-sky simulated convergence map derived from the {Horizon} Simulation.   Its   resolution  of  200  million   pixels  has  been downgraded  to fit  the page.   The various  inserts display  a zoom sequence into smaller and smaller  areas of the sky.  The pixel size is 0.74 arcmin$^{2}$.  } \\label{fullsky} \\end{figure*}  In  this  paper,  we  used  a high  resolution  N-body  simulation  to construct   a    full-sky   weak-lensing   map   and    test   a   new map-reconstruction method based  on a multi-resolution technique.  For this purpose,  we use  the Horizon simulation,  a 70  billion particle N--body simulation, featuring more than  $140$ billion cell in the AMR grid of  the RAMSES code \\citep{teyssier02}.  The  simulation covers a sufficiently  large  volume   ($L_{box}=2  h^{-1}$Gpc)  to  compute  a full-sky convergence  map, while  resolving Milky-Way size  halos with more than 100 particles, and exploring small scales into the nonlinear regime  (see   Sect.~\\ref{simu}).   This  unprecedented  computational effort allows us, for the first  time, to close the gap between scales close to  the cosmological horizon  and scales deep  inside virialized dark-matter  haloes.   A  similar   effort  at  lower  resolution  was presented by \\cite{fosalba08}.  The  dark-matter distribution in  the simulation  was integrated  in a light  cone to  a redshift  of 1,  around an  observer located  at the centre of the simulation  box (see Sect.~\\ref{cone}).  This light cone was  then  used  to   calculate  the  corresponding  full-sky  lensing convergence field, which  we mapped using the Healpix\\footnote{HeaPix: http://healpix.jpl.nasa.gov}  pixelisation  scheme  \\citep{gorski05} with a  pixel resolution of  $\\Delta \\theta \\simeq  0.74$ arcmin$^{2}$ ($n_{\\rm side}=4096$), and added  ``instrumental'' noise for a typical all--sky  survey with 40  galaxies per  arcmin$^{2}$, as  expected for example  for the  DUNE mission  \\citep{dune06}.  Using  an Undecimated Isotropic   Wavelet   Decomposition   of  this   realistic   simulated weak-lensing map  on the  sphere, we analyzed  the statistics  of each wavelet plane  using second, third and fourth  order moments estimator (Sect.~\\ref{stat}).   We   then  applied,  in   Sect.~\\ref{filter},  a multi-resolution algorithm  to filter a  fictitious simulated $\\kappa$ data set based  on an extension of the  wavelet filtering technique of \\citet{pires06}.  We  characterised the quality  of the reconstruction using the  power spectrum  of the  error map and  compare this  to the result  of standard  Wiener  filtering on  the  sphere.  Our  results, summarised  in Sect.~\\ref{conclusion}, illustrate  the virtue  of high resolution simulations  such as the  one reported here to  prepare for future weak-lensing surveys and to design new map--making techniques. ",
        "conclusions": "\\label{conclusion} Using the 70  billion particles of the Horizon  N--body Simulation, we have computed for the first  time a realistic full-sky convergence map with   a   pixel   resolution   of   $\\Delta   \\theta   \\simeq   0.74$ arcmin$^{2}$.   We    have   analyzed   the    resulting   map   using multi-resolution  statistics (variance,  skewness,  and kurtosis)  and angular power-spectrum  analysis.  We  have shown that  this simulated map spans 4 decades of useful  signal in angular scale, with 2 decades within  the  linear, Gaussian  regime  and  2  decades well  into  the nonlinear, non--Gaussian regime.  We have shown that, when considering a realistic sky cut, we  can reliabily estimate high--order moments of the map above  $\\ell \\simeq 200$.  Using even  higher resolution maps, angular scales smaller than $\\theta \\simeq 1^{'}$ could be explored in future works, although  the mass ditribution on these  scales might be affected by  baryons physics \\citep{jing06},  so that the  present map might already cover all cosmologically relevant scales.  As a  first step  towards a realistic  map--making procedure,  we have tested  two  de-noising schemes  on  a  simplified fictitious  dataset derived from the full-sky map,  namely Wiener filtering and the MRLens method  \\citep{pires06}.   We have  shown  quantitatively that  Wiener filtering is the most effective  method on large scales, although some signal is lost  on small scales.  MRLens performs  more effectively on small scales and recovers the dense clumps associated with dark matter halos, but deals less accurately  with low frequency waves in the map. Hence,  this work  demonstrates the  need for  hybrid multi-resolution approach,   e.g.,  by  combining   spherical  harmonics   and  wavelet coefficients.  The present analysis will be extended in future work to map--making  algorithms  dealing  directly  with galaxy  shears.   The simulated convergence  map may prove to  be an effective  tool for the design of new map--making methods  and for the preparation of the next generation  weak--lensing   surveys\\footnote{The  convergence  map  is freely          available          for          download          at \\texttt{http://www.projet-horizon.fr}.  }.  \\begin{acknowledgement} We  would  like  to  thank  Julien Devriendt,  Pierre  Ocvirk,  Arthur Petitpierre and Philippe Lachamp  for their unvaluable help during the course  of this  project. The  Horizon Simulation  presented  here was supported by  the ``Centre de Calcul Recherche  et Technologie'' (CEA, France) as a  ``Grand Challenge project''. This work  was supported by the  Horizon Project.  Some  of the  results in  this paper  have been derived using the HEALPix \\citep{gorski05} package. \\end{acknowledgement}  \\appendix"
    },
    "0807/0807.2288_arXiv.txt": {
        "abstract": "We calculate the jet power of the classical Blandford-Znajek (BZ) model and hybrid model developed by Meier based on the global solutions of advection dominated accretion flows (ADAFs) surrounding Kerr black holes. We find that the jet power of the hybrid model is larger than that of the pure BZ model. The jet power will dominate over the accretion power, and the objects will enter into ``\\emph{jet-power-}dominated advective systems,\" when the accretion rate is less than a critical value $\\dot{m}_{\\rm c}=\\dot{M}_{\\rm c}/\\dot{M}_{\\rm Edd}$, where $3\\times10^{-4}\\lesssim\\dot{m}_{\\rm c}\\lesssim 5\\times10^{-3}$ is a function of black hole spin parameter. The accretion power will be dominant when $\\dot{m}\\gtrsim \\dot{m}_{\\rm c}$ and the objects will enter into ``\\emph{accretion-power-}dominated advective systems.\" This is roughly consistent with that constrained from the low/hard-state black hole X-ray binaries (e.g., Fender et al.).  We calculate the maximal jet power as a function of black hole mass with the hybrid jet formation model, and find it can roughly reproduce the dividing line of the Ledlow-Owen relation for FR I/FR II dichotomy in the jet power-black hole(BH) mass plane ($Q_{\\rm jet}-M_{\\rm BH}$) if the dimensionless accretion rate $\\dot{m}\\sim0.01$ and BH spin parameter $j\\sim0.9-0.99$ are adopted. This accretion rate $\\dot{m}\\sim0.01$ is consistent with that of the critical accretion rate for the accretion mode transition of a standard disk to an ADAF constrained from the state transition of X-ray binaries. Our results imply that most FR I galaxies may be in the ADAF accretion mode similar to the low/hard-state XRBs. ",
        "introduction": "Black hole accretion is thought to power active galactic nuclei (AGNs) and  X-ray binaries (XRBs). Both the UV/optical bumps observed in luminous quasars and the soft X-rays observed in the high/soft state XRBs can be naturally interpreted as blackbody emission from a cold, optically thick, and geometrically thin standard disk \\citep*[SSD; e.g.,][]{ss73,sm89}. A hot, optically thin, geometrically thick advection dominated accretion flow model has been developed in the last several decades (ADAF, or ``radiative inefficient accretion flows''; e.g., Ichimaru 1977; Rees et al. 1982; Narayan \\& Yi 1994, 1995; Abramowicz et al. 1995; see Kato et al. 2008 and Narayan \\& McClintock 2008 for reviews), which can successfully explain most features of  the nearby low-luminosity AGNs and low/hard-state XRBs (see Remillard \\& McClintock 2006; Done et al. 2007; Yuan 2007; Ho 2008 for recent reviews). Their spectral energy distributions (SEDs) can be well reproduced by the ADAF model. It was found that the hard X-ray photon indices of both XRBs and AGNs are anti-correlated with the Eddington ratios when the Eddington ratio is less than a critical value, while they become positively correlated when the Eddington ratio is higher than the critical value \\citep*[][and references therein]{wg08,gc08}. These results provide evidence for the accretion mode transition near the critical Eddington ratio.   It is widely believed that the radio emission of both the low/hard state XRBs \\citep*[e.g.][and references therein]{fe06} and the low luminosity AGNs \\citep*[LLAGNs, e.g.,][]{wc05} comes from the jets. However, the detailed physical mechanism for the jet formation is still unclear for either AGNs or XRBs. The currently most favored jet formation mechanisms include the Blandford-Znajek (BZ) process \\citep{bz77} and the Blandford-Payne (BP) process \\citep{bp82}. In the BZ process, energy and angular momentum are extracted from a rotating black hole and transferred to a remote astrophysical load by open magnetic field lines. In the BP process, the magnetic fields threading the disk extract energy from the rotation of the accretion disk itself to power the jet/outflow. \\citet{me99} showed that even if the BZ process is neglected entirely, the jet power contributed by the field threading the disk alone will still be a function of the black hole spin since the rotating metric contributes to the rotation of the magnetic field. Recent magnetohydrodynamic(MHD) simulations also showed that both the BH spin and the accretion rate in the underlying accretion disk play important roles for the jet formation \\citep*[][]{ko00,kg04,hi04,de05,hk06}. The relative importance of these two mechanisms was explored by many authors \\citep*[e.g.,][]{ga97,li99,me01,ca02,ne07,wang08}. Assuming the poloidal magnetic field component at the disk surface to be of the same order as the toroidal field component, the maximal jet power extracted from the accretion disk (BP process) may dominate over the maximal power extracted by the BZ process \\citep*[e.g.,][]{li99,me01,ca02,ne07}. However, \\citet{re06} found that the dynamics of the accretion disk in plunging region within the innermost stable circular orbit can greatly enhance the trapping of large scale magnetic field on the BH, and therefore increase the importance of the BZ mechanism effect compared to previous estimates that ignore the plunge region \\citep*[e.g.][]{ga97}.   Radio galaxies are usually classified as FR I or FR II sources depending on their radio morphology. FR I radio galaxies (defined by edge-darkened radio structure) have lower radio power than FR II galaxies (defined by edge-brightened radio structure due to compact jet terminating hot spots; Fanaroff \\& Riley 1974). What causes the morphological difference between FR I and FR II radio galaxies is still unclear. The theoretical models fall into two different groups: (1) the morphological differences arise of the different physical conditions in their ambient medium \\citep*[see][for a summary]{gk00}; (2) their intrinsic difference of their central engines, i.e., different accretion modes and/or jet formation processes \\citep*[e.g.,][]{bi95,re96a,me99,gc01,ma04,wo07,ha07}.  Most previous works on the jet power extracted from ADAFs were based on self-similar solutions of ADAFs \\citep*[e.g.,][]{an99,me01,cr04,ne07}. The self-similar solution can reproduce the global solution quite well at large radii, while it deviates significantly near the black hole \\citep*[e.g.,][]{na97}, where the relativistic jets are supposed to be formed. In this paper we calculate the jet power, incorporating some recent MHD simulation results, based on the global ADAF solutions surrounding Kerr black holes. We then compare our model calculations with the accretion power and jet power of XRBs and the radio galaxy FR I/FR II dichotomy. ",
        "conclusions": "The main conclusions of this work can be summarized as follows:  (1) We investigate the jet power of the BZ mechanism and hybrid mechanism for the ADAFs surrounding rotating black holes based on our global ADAF solutions in Kerr metric. We find that the jet power of the hybrid model is about 1 order of magnitude higher than that of the BZ model, and the jet efficiency of the hybrid model is roughly consistent with that of the numerical MHD simulations for Kerr black holes\\citep{hk06}.   (2) The jet power dominates over the accretion power when the accretion rate is less than a critical value $\\dot{m}_{\\rm c}$, while the accretion power will be dominant when the accretion rate is larger than this critical value (Fig. 2), which is roughly consistent with that constrained from the observations of XRBs \\citep*[e.g.,][]{fe03,mf06}.  (3)  The dividing line of the Ledlow-Owen relation in $Q_{\\rm jet}-M_{\\rm BH}$ plane for the FR I/FR II dichotomy can be well reproduced by the hybrid jet model, provided that the ADAFs surrounding Kerr black holes are accreting at $\\dot{m}\\la 0.01$ in FR I sources(Fig. 3). The ionization luminosity corresponding to the dividing line plane also suggests the critical accretion rate to be $\\sim 0.01$ if the empirical relation between the photo-ionizing luminosity and the radio luminosity is adopted. These strongly imply that most FR I radio galaxies may have a different accretion mode from FR II sources."
    },
    "0807/0807.2594_arXiv.txt": {
        "abstract": "{} {To understand low- to intermediate-mass star-formation in the nearby R\\,Cr\\,A molecular cloud, we try to identify the stellar content that is accessible with near-infrared observations. } {We obtained a $JHK_{s}$ band mosaic of $\\sim 10\\arcmin \\times 60\\arcmin$ covering the entire R\\,CrA molecular cloud with unprecedented sensitivity. We present a catalogue of about 3500 near-infrared sources fainter than the saturation limit $K_{s} \\sim 10\\,$mag, reaching  $K_{s} \\sim$ 18\\,mag. We analysed the extended sources by inspecting their morphology and point sources by means of colour-colour and colour-magnitude diagrams. Additionally, we compared the extinction inferred from the NIR data with the line-of-sight dust emission at 1.2\\,mm. Sources towards high dust emission but relatively low $H-K_{s}$ show a projected mm-exces; these sources are either immediately surrounded by cold circumstellar material or, if too red to be a true foreground object, they are embedded in the front layer of the 1.2\\,mm emitting dust cloud. In both cases they are most likely associated with the cloud.} {By means of the projected mm-excess technique we find 33 new faint near-infrared sources deeply embedded in the Coronet cluster around R\\,CrA, for which so far about 20 bright infrared stars have been known. In contrast to the Coronet region, both the northwestern dust ridge and the southeastern cloud condensation \"C\" appear to be devoid of associated stars detectable with our near-infrared data. Furthermore, about a dozen sources, which are spread over the entire molecular cloud region, exhibit a possible $K$-band excess, but only with marginal statistical significance ($<3\\sigma$), so that we do not consider the indicated $K$-band excess as real. Finally, while the Herbig-Haro-like objects seen on our maps are concentrated around the Coronet, we find four new nebulae also located farther down to the southeast. At the position of IRAS\\,18595-3712, an X-shaped bipolar nebula is resolved; its exciting star is hidden behind an edge-on disc. } {The deep near-infrared survey of the entire R\\,CrA molecular cloud strengthens the evidence for the Coronet being the region where most of the young stars are found. Our results are consistent with earlier predictions that the R\\,CrA cloud has fragmented into sub-condensations at different star-forming stages. } ",
        "introduction": "At a distance of about 170\\,pc (Knude \\& Hog 1998), the R\\,Cr\\,A molecular cloud (Rossano's cloud A, 1978) is among the closest star-forming regions. Surveys of this cloud have revealed a population of young intermediate- to low-mass stars, through the presence of H$\\alpha$ emission (e.g. Marraco \\& Rydgren 1981, Fern\\'andez \\& Comer\\'on 2001) or as X-ray sources (e.g. Forbrich \\& Preibisch 2007). Near-infrared (NIR) observations detected the {\\it Coronet}, a small cluster of embedded sources around R\\,Cr\\,A itself (Taylor \\& Storey 1984). Further NIR observations by Wilking et al. (1997) focussed on the {\\it Coronet} region and analysed sources down to $K = 15\\,$mag, which -- due to the large extinction in most parts of the observed region -- yields a mass detection limit of about 1\\,M$_{\\odot}$. More deeply embedded sources have been identified by radio continuum (Brown 1987), mid-infrared (Olofsson et al. 1999), and millimetre (Chini et al. 2003, Nutter et al. 2005) observations. The young (10$^{\\rm 3}$ -- 10$^{\\rm 5}$~yrs) stars drive various outflows that are observed either as Herbig-Haro objects (e.g. Hartigan \\& Graham 1987, Wang et al. 2004) or as molecular outflows (e.g. Anderson et al. 1997). Using the Swedish ESO Submillimetre Telescope (SEST), the gas and dust content of the entire R\\,Cr\\,A cloud having an extent of 10$\\arcmin$\\,$\\times$\\,50$\\arcmin$ has been mapped in C$^{\\rm 18}$O J=1-0 by Harju et al. (1993) and at 1.2\\,mm with SIMBA (Chini et al. 2003). At galactic coordinates $l = 0^\\circ, b=-18^\\circ$ any fore- and background contamination should be negligible.  The R\\,Cr\\,A molecular cloud lies at the apex of a several degrees large cometary-shaped flocculent reflection nebula. As discussed by Harju et al. (1993), the cloud may have been influenced about 10-13\\,Myr ago by the expanding H\\,I shell of the Upper Centaurus-Lupus OB association located west-northwest of R\\,CrA. Then a rotating disc-like cloud seen roughly edge-on could have fragmented into several sub-condensations that further contract to form stars. The central condensation A, the {\\it Coronet}, houses the currently observed young sources, while the neighbouring regions northwest (NW), and southeast (C) show virtually no ongoing star formation. The dense region C was predicted to be a site for future star formation (Harju et al. 1993), and in fact a cold cloud core has been detected there (Chini et al. 2003). On the other hand, it is puzzling why the dust ridge northwest of the {\\it Coronet} lacks signs of recent star formation. This stimulated our aim of identifying the stellar content along the entire cloud down to at least 0.1 M$_{\\odot}$ (spectral type M7V - M8V, see Landolt-B\\\"ornstein 1982).  Because of the expected high extinction at optical wavelengths, we performed a deep and spatially complete NIR survey of the R\\,Cr\\,A cloud. We gave special emphasis on covering the outermost regions NW and C, which in the NIR had only been surveyed by 2MASS, and to map the {\\it Coronet} with at least 2 mag higher sensitivity than was achieved by Wilking et al. (1997). Here we present the data and perform a first analysis with respect to the 1.2\\,mm dust continuum map. Further comparison of the NIR data with MIR data will be provided in a follow-up publication. ",
        "conclusions": "\\label{section_conclusions}  We performed a deep near-infrared survey of the entire R\\,CrA molecular cloud and analysed the data with respect to the 1.2\\,mm dust continuum map.  \\begin{itemize} \\item[1)] Applying the projected mm-excess technique we find 33 new sources deeply embedded in the front layer towards the {\\it Coronet}. Nine sources could be detected at $H$ and $K_{s}$, 24 only have $K_{s}$-band data. In contrast to the {\\it Coronet}, the NW dust ridge and the C core region appear to be devoid of associated stars detectable in our near-infrared images. This strengthens the evidence for the {\\it Coronet} being the region where most of the Class I-III stars are found. Obviously, core C contains at most very young protostars, which are too faint and embedded to be visible at near-infrared wavelengths. Stringent estimates about the contamination from background objects may be derived from Spitzer MIR data.  \\item[2)] Only marginal evidence of $K$-excess sources can be found from our data ($K_{s} > 10\\,$mag). If these $K$-excess candidates are associated with the R\\,CrA cloud, then they are very low-mass stars. In general, it seems that our sources do not exhibit circumstellar discs with {\\it prominent} hot dust emission showing up already in the $K$-band.  \\item[3)] The four new nebulae found towards the southeastern part corroborate the picture that star formation has also started in these regions.   \\item[4)] The dust and CO peaks A {\\it Coronet} and B (MMS\\,23) with ongoing star formation exhibit relatively sharp eastern borders and shallow western tails; the same holds for cloud C. If the star formation in the diverse R\\,CrA subcondensations were directly triggered by the wind from the Upper Centaurus-Lupus OB association, such a morphology is not expected at first glance. Instead in a scenario where an external wind hits an originally symmetric gas clump, one would expect that the front side, where the wind hits the clump, appears compressed and houses the first and stronger star formation. Because this is not observed, the picture is supported that the expanding Upper Centaurus-Lupus HI shell just may have led to fragmentation of the R\\,CrA cloud into subcondensations, which further collapse to form stars mainly under the control of their own gravitation.  \\end{itemize}  Our results are in line with the predictions by Harju et al. (1993) that the R\\,CrA cloud has fragmented into subcondensations that are at different star-forming stages."
    },
    "0807/0807.0302_arXiv.txt": {
        "abstract": "{Accurate measurements of lithium abundances in young low-mass stars provide an independent and reliable age diagnostics. Previous studies of nearby star forming regions have identified significant numbers of Li-depleted stars, often at levels inconsistent with the ages indicated by their luminosity. } {We aim at a new and accurate analysis of Li abundances in a sample of $\\sim$100 pre-main sequence stars in Taurus-Auriga using a homogeneous and updated set of stellar parameters and model atmospheres appropriate for the spectral types of the sample stars. } {We compute Li abundances using published values of the equivalent widths of the Li $\\lambda$6708~\\AA~doublet obtained from medium/high resolution spectra. } {We find that the number of significantly Li-depleted stars in Taurus-Auriga is greatly reduced with respect to earlier results. Only 13 stars have abundances lower than the interstellar value by a factor of 5 or greater. All of them are weak-lined T Tauri stars drawn from X-ray surveys; with the exception of four stars located near the L1551 and L1489 dark clouds, all the Li-depleted stars belong to the class of dispersed low-mass stars, distributed around the main sites of current star formation. If located at the distance of Taurus-Auriga, the stellar ages implied by the derived Li abundances are in the range 3-30 Myr, greater than the bulk of the Li-rich population with implication on the star formation history of the region. } {In order to derive firm conclusions about the fraction of Li-depleted stars of Taurus-Auriga, Li measurements of the remaining members of the association should be obtained, in particular of the group of  stars that fall in the Li-burning region of the HR diagram.} ",
        "introduction": "The use of lithium as an alternative and secure method to derive the ages of young, very low-mass (M$\\sim$0.1--0.6~\\msun) stars is now well established on theoretical and observational grounds. These stars are fully convective; thus they have a simple internal structure and the physics of their Li depletion during the pre-main sequence (PMS) phases is well understood and not affected by major uncertainties. Theory predicts that they start depleting their surface Li after $\\sim$2~Myr from  birth, and eventually exhaust all the internal supply in about 10--50~Myr (e.g. D'Antona \\& Mazzitelli~\\cite{dantona94}; Baraffe et al. \\cite{baraffe98}; Siess et al. \\cite{siess}). The timescale of Li depletion depends on mass, with lower mass stars being the slowest ones. Therefore, in a cluster or association of a given age there is a clear/sharp transition, the so-called lithium depletion boundary (LDB), between fully depleted objects and those with the initial lithium content. The mass at which the boundary occurs reveals the age of the cluster. Observationally, the LDB test has been successfully applied to a number of young open clusters in the age range 30-120~Myr (e.g., Stauffer et al.~\\cite{stauf98}, \\cite{stauf99}; Barrado y Navascu\\'es et al.~\\cite{byn}; Jeffries \\& Oliveira~\\cite{jo05}; Manzi et al.~\\cite{manzi08}).  The physics of lithium depletion is considerably more complex for higher mass stars in the range M$\\sim$0.6--1.2~\\msun that develop a radiative core during PMS evolution. The exact amount of depletion in these stars depends on a number of parameters and assumptions entering the models (e.g., Jeffries \\cite{jeff06}); hence, use of Li abundances to derive exact ages of stars in this mass range is considerably uncertain, although it still allows distinguishing between very young and more evolved PMS stars.  In addition to providing absolute age estimates, Li abundances can also be used to obtain information on the magnitude of the age spread within less evolved system that are too young for the development of the LDB. Recently, Palla et al. (\\cite{pallaONC05}, \\cite{pallaONC07}) have measured Li abundances in a large sample of low-mass members of the Orion Nebula Cluster (ONC), whose bulk population is only $\\sim$1--2 Myr old. They found that, in addition to the overwhelming population of stars with Li abundances consistent with the initial interstellar value (log N(Li)=3.1$\\pm$0.2, e.g. Jeffries~\\cite{jeff06} and references therein), six high probability members are strongly Li depleted, with Li depletion ages greater than 10 Myr. For four of these stars the agreement with the isochronal ages derived from the position of stars in the HR diagram is excellent. Similarly, Li depleted stars were found in Upper Scorpius (Mart\\'{\\i}n \\cite{martin98}) and $\\sigma$~Orionis (Sacco et al. \\cite{sacco}). The presence of older stars mixed with younger members reveals an extended phase of star formation that lasted for much more than a few dynamical times.  Another region that has been the object of extensive Li studies in the past is the Taurus-Auriga (Tau-Aur) association. The largest collection of Li data has been presented in a series of papers by Basri et al. (1991, hereafter B91), Magazz\\`u et al. (1992, hereafter M92), and Mart\\'{\\i}n et al. (1994, hereafter M94). These studies demonstrated for the first time that the majority of young stars have  Li abundances consistent with the initial interstellar value, but that a significant number (about 30) stars are characterized by low or very low abundances. If real, the observed depletion level would imply large  Li depletion ages ($\\gesssim$10~Myr), in excess of the bulk of the undepleted population ($\\gesssim$1--3~Myr). However, the interpretation of the data remained controversial, since the measured Li abundances were not entirely consistent with the predictions of theoretical models. In particular, the lowest mass stars of the samples indicated too much Li destruction at the observed luminosity. In other words, these T Tauri stars were too young, as indicated by the isochronal ages, to have undergone significant nuclear burning and Li destruction. The interpretation of this result was not straightforward and the possible explanations invoked large uncertainties in the stellar parameters (effective temperature, luminosity, surface gravity), problems with model atmosphere used to derive Li abundances, and limitations of the PMS evolutionary models. More recently, Wichmann et al. (2000; hereafter W00) have presented a detailed analysis of the young population discovered by ROSAT in and near the Tau-Aur complex. Using high resolution echelle spectra, they measured lithium equivalent widths and used these measurements as a selection criterion for youth, expanding in this way the sample of candidate low-mass TTS of Tau-Aur. Since the knowledge of the properties of the stellar population of Tau-Aur has greatly increased in the last decade (e.g., Kenyon \\&  Hartmann 1995; Brice\\~no et al. 2002; Luhman et al. 2003; G\\\"udel et al. 2007), and new atmospheric (e.g., Pavlenko et al. 2001) and evolutionary (Siess et al. 2000; Baraffe et al. 2002; D'Antona \\& Montalb\\`an 2003; Montalb\\`an \\& D'Antona 2006) models have been developed, the time is ripe for a careful re-analysis of the Li abundances in T Tauri stars in this important star forming region.  The structure of the paper is the following: In Sect.~2 we describe the sample, stellar parameters,  method of analysis, and results obtained by B91, M92, M94 and W00; in Sect.~3 we present our new analysis of the sample and the resulting  Li abundances based on updated stellar parameters and curves of growth (COGs). In Sect.~4 we address issues related to the origin of the Li-depleted stars, to the completeness of the sample and the possibility of a new survey aimed at discovering new Li depleted TTS. Sect.~5 closes the paper. ",
        "conclusions": "We have computed Li abundance in a sample of $\\sim$100 T Tauri stars of Tau-Aur with published values of EW(Li) obtained from high resolution spectra. The Li abundance was derived using updated atmospheric models and a set of homogeneous stellar data as inputs. The most significant results are the following: \\begin{itemize} \\item The serious inconsistency found in previous studies on the relatively high number of Li-depleted of TTS at the wrong luminosity is removed. Only 13 stars have Li abundance below a value \\nli$=$2.6, corresponding to depletion factors greater than five. We attribute the early findings to incorrect Li-abundance analysis and/or to the use of inconsistent stellar parameters. \\item The majority of the Li-depleted stars lie in the HR diagram either within or close to the boundary of the theoretical region of Li-burning in low-mass stars. Conversely, most of the remaining stars with Li abundance consistent with the initial interstellar value are distributed in the appropriate portion of the HR diagram. \\item Unlike previous results, none of the 21 CTTS included in the sample shows evidence for Li-depletion. On the other hand, all of the 13 Li-depleted stars are WTTS identified in X-ray surveys. Most of them belong to the widely dispersed population distributed away from (or at the periphery of) the main star forming sites. In several cases the radial velocity and proper motions are consistent with membership to the Tau-Aur moving group. \\item If located at the distance of Tau-Aur, the Li depletion and isochronal ages of the Li-depleted stars are in the range $\\gesssim$3--15~Myr, greater than the bulk of the undepleted population. \\item In order to draw firm conclusions on the Li-content of TTS in Tau-Aur, it is necessary to obtain high quality measurements of the remaining known members of the complex, which outnumber the group studied here by a factor $\\sim$3. In particular, a group of $\\sim$25 stars that lie inside the Li-burning region constitute an interesting subsample that will be the subject of a future study. \\end{itemize}"
    },
    "0807/0807.2900_arXiv.txt": {
        "abstract": "Dimension-reduction techniques can greatly improve statistical inference in astronomy. A standard approach is to use Principal Components Analysis (PCA). In this work we apply a recently-developed technique, diffusion maps, to astronomical spectra for data parameterization and dimensionality reduction, and develop a robust, eigenmode-based framework for regression. We show how our framework provides a computationally efficient means by which to predict redshifts of galaxies, and thus could inform more expensive redshift estimators such as template cross-correlation.  It also provides a natural means by which to identify outliers (e.g., misclassified spectra, spectra with anomalous features). We analyze 3835 SDSS spectra and show how our framework yields a more than 95\\% reduction in dimensionality. Finally, we show that the prediction error of the diffusion map-based regression approach is markedly smaller than that of a similar approach based on PCA, clearly demonstrating the superiority of diffusion maps over PCA for this regression task. ",
        "introduction": "\\label{sect:intro}  Galaxy spectra are classic examples of high-dimensional data, with thousands of measured fluxes providing information about the physical conditions of the observed object. To make computationally efficient inferences about these conditions, we need to first reduce the dimensionality of the data space while preserving relevant physical information. We then need to find simple relationships between the reduced data and physical parameters of interest. Principal Components Analysis (PCA, or the Karhunen-Lo\\`eve transform) is a standard method for the first step; its application to astronomical spectra is described in, e.g., \\citet{BorosonGreen1992}, \\citet{Connolly1995}, \\citet{Ronen1999}, \\citet{Folkes1999}, \\citet{Madgwick2003}, \\citet{Yip2004a}, \\citet{Yip2004b}, \\citet{Li2005}, \\citet{Zhang2006}, \\citet{VDB2006}, \\citet{Rogers2007}, and \\citet{ReFiorentin2007}. In most cases, the authors do not proceed to the second step but only ascribe physical significance to the first few eigenfunctions from PCA (such as the ``Eigenvector 1\" of \\citeauthor{BorosonGreen1992}). Notable exceptions are \\citeauthor{Li2005}, \\citeauthor{Zhang2006}, and \\citeauthor{ReFiorentin2007} However, as we discuss in {\\S}\\ref{sect:app}, these authors combine eigenfunctions in an ad hoc manner with no formal methods or statistical criteria for regression and risk (i.e., error) estimation.  In this work we present a unified framework for regression and data parameterization of astronomical spectra. The main idea is to describe the important structure of a data set in terms of its {\\em fundamental eigenmodes}. The corresponding eigenfunctions are used both as coordinates for the data and as orthogonal basis functions for regression. We also introduce the {\\em diffusion map} framework (see, e.g., \\citealt{Coifman:Lafon:06}, \\citealt{LafonLee2006}) to astronomy, comparing and contrasting it with PCA for regression analysis of SDSS galaxy spectra.  PCA is a global method that finds linear low-dimensional projections of the data; it attempts to preserve Euclidean distances between all data points and is often not robust to outliers. The diffusion map approach, on the other hand, is non-linear and instead retains distances that reflect the (local) connectivity of the data. This method is robust to outliers and is often able to unravel the intrinsic geometry and the natural (non-linear) coordinates of the data.  In {\\S}\\ref{sect:diff} we describe the diffusion map method for data parameterization. In {\\S}\\ref{sect:regress} we introduce the technique of {\\em adaptive regression} using eigenmodes. In {\\S}\\ref{sect:app} we demonstrate the effectiveness of our proposed PCA- and diffusion-map-based regression techniques for predicting the redshifts of SDSS spectra. Our PCA- and diffusion-map-based approaches provide a fast and statistically rigorous means of identifying outliers in redshift data. The returned embeddings also provide an informative visualization of the results.  In {\\S}\\ref{sect:summary} we summarize our results. ",
        "conclusions": "\\label{sect:summary}  The purpose of this paper is two-fold. First, we introduce the diffusion map method for data parametrization and dimensionality reduction. We show that for the types of high-dimensional and complex data sets often analyzed in the astronomy, diffusion map can yield far superior results than commonly-used methods such as PCA.  Moreover, the simple, intuitive formulation of diffusion map as a method that preserves the local interactions of a high-dimensional data set makes the technique easily accessible to scientists that are not well-versed in statistics or machine learning.  Second, we present a fast and powerful eigenmode-based framework for estimating physical parameters in databases of high-dimensional astronomical data.  In most astrophysical applications, PCA is used as a data-explorative tool for dimensionality reduction, with no formal methods and statistical criteria for regression, risk estimation and selection of relevant eigenvectors. Here we propose a statistically rigorous, unified framework for regression and data parameterization.  Our proposed regression model combines basis functions in a simple and statistically-motivated manner while our clear objective of risk minimization drives the estimation of the model parameters.  Again, the simplicity of the proposed method will make it appealing to the non-specialist.  We apply the proposed methodology to predict redshift for a sample of SDSS galaxy spectra, comparing the use of the proposed regression model with PCA basis functions versus diffusion map basis functions. We find that the prediction error for the diffusion-map-based approach is markedly smaller than that of a similar framework based on PCA. Our techniques are also more robust than commonly used template matching methods because they consider the structure of the entire high-dimensional data set when reparametrizing the data. Statistical inferences are based on this learned structure, instead of considering each data point separately in an object-by-object matching algorithm as is currently used by SDSS and commonly employed throughout the astronomy literature. Work in progress extends our approach to photometric redshift estimation and to the estimation of the intrinsic parameters (e.g., mean metallicities and ages) of galaxies."
    },
    "0807/0807.4467_arXiv.txt": {
        "abstract": "{  Deep polarimetric Westerbork observations of the galaxy cluster \\object{Abell~2256} are presented, covering a frequency range of \\modified{325--377~MHz}. The central halo source has a diameter of the order of 1.2~Mpc (18\\arcmin), which is somewhat larger than at 1.4~GHz. With $\\alpha=-1.61\\pm0.04$, the halo spectrum between 1.4~GHz and 22.25~MHz \\modified{is less steep} than previously thought. The centre of the ultra steep spectrum source in the eastern part of the cluster exhibits a spectral break near 400~MHz.  It is estimated to be \\modified{at least 51 million years old, but possibly older than 125 million years. A final measurement requires observations in the 10--150~MHz range.} It remains uncertain whether the source is a radio tail of Fabricant galaxy 122, situated in the northeastern tip of the source. Faraday rotation measure synthesis revealed no polarized flux at all in the cluster.  The polarization fraction of the brightest parts of the relic area is less than 1\\%.  The RM-synthesis nevertheless revealed 9 polarized sources in the field enabling an accurate measurement of the Galactic Faraday rotation ($-33\\pm2$~rad~m$^{-2}$ in front of the relic). \\modified{Based on its depolarization on longer wavelengths, the line-of-sight magnetic field in relic filament G is estimated to be between 0.02 and 2~$\\mu$G. A value of $0.2$~$\\mu$G appears most reasonable given the currently available data.} } ",
        "introduction": "\\label{brentjens_abell2256_sec:introduction}   Abell~2256 is one of the most massive clusters in the nearby universe. It has been studied extensively at X-ray, optical, and radio wavelengths. The galaxies associated with the cluster can be divided into three distinct populations, based on their kinematics \\citep{BerringtonLuggerCohn2002}. X-ray images of the cluster reveal several substructures superimposed on a fairly smooth main source \\citep{SunEtAl2002,MolendiEtAl2000, RoettigerBurnsPinkney1995, BrielEtAl1991}, indicating that Abell~2256 is currently involved in a major merger event.  The cluster has some interesting radio properties. It contains the largest number of head-tail galaxies of all known clusters \\citep{MillerOwenHill2003}. One of them has a straight, almost 1~Mpc long tail.  There is a complex steep spectrum source in the eastern part of the cluster \\citep{RottgeringEtAl1994,MassonMayer1978,BridleEtAl1979}.  The cluster is embedded in a large, diffuse, unpolarized radio halo \\citep{BridleFomalont1976, Kim1999, ClarkeEnsslin2006} that dominates the total flux at decametric wavelengths \\citep{CostainEtAl1972}.  The most striking feature is the large, bright relic source in the northwestern part of the cluster. Although it is \\modified{20--40\\%} linearly polarized at 1.4~GHz \\citep{BridleEtAl1979,ClarkeEnsslin2006}, \\citet{Jaegers1987} did not detect any linear polarization at 608.5~MHz, establishing a $3\\sigma$ upper limit of 20\\% fractional polarization.  There are strong indications that relic sources are caused by large scale structure formation (LSS) shocks compressing buoyant bubbles of magnetized plasma that have been emitted by active galactic nuclei in the past \\modified{\\citep{EnsslinGopalKrishna2001, EnsslinBrueggen2002, Brueggen2003, HoeftBrueggen2007}}.  Shortly after their injection into the intergalactic medium (IGM), these bubbles have faded beyond detection limits due to synchrotron losses and adiabatic expansion. An encounter with a LSS shock wave would then compress the plasma, increasing and aligning the magnetic field, and accelerate the electrons to energies enabling synchrotron emission at radio wavelengths. This model explains both the high fractional polarization that is usually observed in these relic sources and their peripheral position.   De Bruyn \\& Brentjens (2005) \\nocite{DeBruynBrentjens2005} have discovered several highly polarized structures in the direction of the \\object{Perseus cluster} that were tentatively attributed to the cluster itself. Some of the structures resembled buoyant bubbles in the IGM of the Perseus cluster, while others looked like long, straight shock fronts. One of these \"fronts\" was located at the western edge of the cluster at the interface to the Perseus-Pisces supercluster, exactly where such shocks are expected to occur. However, these structures have not been detected in Stokes $I$, hence the exact fractional polarization is unknown. The two most likely reasons for the non-detection of Stokes $I$ are: \\begin{itemize} \\item The sources are highly polarized and the Stokes $I$ brightness is below the confusion limit of the Westerbork Synthesis Radio Telescope (WSRT) in these observations; \\item The structures have Stokes $I$ structure at much larger scales than the Stokes $Q$ and $U$ structure, rendering it invisible at even the shortest interferometer baseline. \\end{itemize} The latter situation is common in the Galactic synchrotron foreground \\citep{Wieringa1993}. A Galactic interpretation can therefore not be ruled out. However, \\citet{GovoniEtAl2005} have observed similar structures in Stokes $I$ near \\object{Abell~2255} with the VLA at 1.4~GHz. Abell~2256 was chosen to search for similar objects because of its proximity, its dynamic nature, its relic source, and its moderate Galactic latitude of $32\\degr$.  Wide field polarimetric images with noise levels well below 1~mJy~beam$^{-1}$ are required in order to detect such sources. Additionally, the images can be used to find more polarized background sources in order to establish the Galactic contribution to the Faraday rotation measure (RM) towards the cluster, and \\modified{to} determine the decline of the fractional polarization with increasing wavelength, enabling a measurement of the thickness of the relic sources in RM-space.  \\modified{The spectrum} of the diffuse halo source has been rather uncertain. \\citet{BridleEtAl1979} favoured a spectral index ($S_\\nu\\propto\\nu^\\alpha$) \\modified{of $\\alpha < -1.2$ between 610 and 1415~MHz and $\\alpha \\approx -1.8$ between 151 and 610~MHz}, based on maps where the halo was only marginally detected. More recently \\citet{ClarkeEnsslin2006} determined the 1.4~GHz halo flux in more sensitive VLA D \\modified{observations, but could not derive the spectrum of the halo because of the large uncertainties in the halo flux measurements available to them}. They \\modified{assumed} that $-1.7 \\lesssim \\alpha \\lesssim -1.25$. \\citet{Kim1999} determined that $\\alpha=-2.04\\pm0.04$, but this estimate was partly based on an erroneous value of the total cluster flux at 81.5~MHz, and partly on maps in which the halo was only marginally detected. New measurements of the 351~MHz flux of the halo and the entire cluster are presented in Sect.~\\ref{brentjens_abell2256_sec:maps}.  These values are combined with corrected literature values in order to obtain accurate spectra for the cluster as a whole and the diffuse halo source itself.  The redshift of \\object{Abell~2256} is 0.058 \\citep{BerringtonLuggerCohn2002}. Assuming $H_0 = 71$~km~s$^{-1}$~Mpc$^{-1}$ \\citep{SpergelEtAl2003}, 1\\arcmin\\ on the sky corresponds to 67.8~kpc at the cluster. ",
        "conclusions": "The spectral index of the diffuse central halo of \\object{Abell~2256} is $-1.61\\pm0.04$. \\modified{Although this value is steeper than the 610/1415~MHz value of $-1.2$ in \\citet{BridleEtAl1979}, it is somewhat flatter than their 150/610~MHz spectral index ($-1.8$) and considerably flatter than the value of $-2.04$ derived by \\citet{Kim1999}.} The reason for this is twofold: \\begin{itemize} \\item in \\citet{ClarkeEnsslin2006} and in this work, the halo was detected over a larger extent than in previous observations by \\citet{BridleEtAl1979} and \\citet{Kim1999} due to the increased sensitivity of the observations; \\item \\citet{Kim1999} did not correct the 81.5~MHz flux \\citep{Branson1967} as proposed by \\citet{MassonMayer1978}. \\end{itemize}  The \\modified{observed spectral index} puts the classical minimum energy magnetic field in the range of 2--4~$\\mu$G, and the hadronic minimum energy field between 4 and 9~$\\mu$G \\citep{ClarkeEnsslin2006}.  Despite the deep polarization images obtained after RM-synthesis, no polarized emission was detected that could be attributed to Abell~2256. The complete depolarization of filaments G and H was expected because of the relatively high density environment in which they reside. The upper limit to the fractional polarization of less than 1\\% is consistent with the 608.5~MHz upper limit by \\citet{Jaegers1987} and the depolarization between 1705~MHz and 1395~MHz that \\citet{ClarkeEnsslin2006} found.   \\begin{figure} \\resizebox{\\hsize}{!}{\\includegraphics{a2256fg16}} \\caption{\\modified{Estimates of the magnetic field parallel to the line of sight in filament G.}} \\label{brentjens_abell2256_fig:b-field-filament-g} \\end{figure}   \\modified{ Reliable measurements of the magnetic field in relic sources are relatively rare. \\citet{BagchiPislarLimaNeto1998} used inverse Compton scattering to derive the cluster scale magnetic field in the relic of \\object{Abell~85}. They found a value of $0.95\\pm0.10$~$\\mu$G. \\citet{ChenHarrisHarrisonMao2008} derived classical equipartition field strengths of 0.63 and 1.3~$\\mu$G for the relics in \\object{0917+75} and \\object{1401-33} respectively.  Their $3\\sigma$ inverse Compton lower limits are 0.81 and 2.2~$\\mu$G. \\citet{EnsslinEtAl1998} listed an equipartition field of $3 h_{71}^{2/7}$~$\\mu$G for the entire relic in Abell~2256. Although the estimate of the line-of-sight field in filament G presented in this work is not very accurate ($0.02$--$2$~$\\mu$G with $0.2$~$\\mu$G in case of reasonable assumptions about the shock compression ratio and path length along the line of sight), it is consistent with the quoted literature values for magnetic fields in relic sources in general. An actual measurement of the depolarization of filament G can improve the estimate significantly, although the major uncertainty is in the physical thickness of the relic along the line of sight, which is very difficult to obtain.}  Polarization observations covering frequencies between 1.4~GHz and 400~MHz, where the fractional polarization decreases rapidly, would make it possible to recover the (linearly polarized) emissivity of filaments G and H as a function of Faraday depth along the line of sight.  The high fractional polarization at 1.4~GHz ensures that the magnetic field has no reversals along the line of sight. Therefore, the Faraday depth of a certain location in the relic sources should be a monotonic function of the geometric distance between the front of the sources and that location, enabling the first 3D reconstruction of an extragalactic synchrotron source.  The fact that no evidence was found for polarized emission from LSS shocks or buoyant bubbles further away from the cluster centre than filaments G and H could have several reasons: \\begin{itemize} \\item they are depolarized due to internal Faraday dispersion, just like sources G and H; \\item there are no shocks oriented edge-on, only face-on, resulting in an undetectably low surface brightness; \\item apart from the relic area, there are simply no other bubbles or large scale shocks in Abell~2256 that emit radio waves. \\end{itemize}  The first point can be tested by applying RM-synthesis to deep polarimetry at 1.4~GHz. Discriminating between the other two possibilities would be difficult at best in a single cluster. One way to move forward is by observing a larger sample of nearby clusters at high Galactic latitude using RM-synthesis around 350~MHz. A fractional bandwidth of 20\\% is then sufficient to separate cluster emission from Galactic foreground emission. If RM-synthesis does not reveal any bubble-like polarized structures, similar to the ones towards the Perseus-cluster \\citep{DeBruynBrentjens2005}, it must be concluded that the structures detected towards the Perseus cluster really belong to the Galactic ISM. At this moment, however, there is not enough evidence to draw a final conclusion."
    },
    "0807/0807.1845_arXiv.txt": {
        "abstract": "{The improvement in observational facilities requires refining the modelling of the geometrical structures of astrophysical objects. Nevertheless, for complex problems such as line overlap in molecules showing hyperfine structure, a detailed analysis still requires a large amount of computing time and thus, misinterpretation cannot be dismissed due to an undersampling of the whole space of parameters.} {We extend the discussion of the implementation of the Gauss--Seidel algorithm in spherical geometry and include the case of hyperfine line overlap.} {We first review the basics of the short characteristics method that is used to solve the radiative transfer equations. Details are given on the determination of the Lambda operator in spherical geometry. The Gauss--Seidel algorithm is then described and, by analogy to the plan--parallel case, we see how to introduce it in spherical geometry. Doing so requires some approximations in order to keep the algorithm competitive. Finally, line overlap effects are included.} {The convergence speed of the algorithm is compared to the usual Jacobi iterative schemes. The gain in the number of iterations is typically factors of 2 and 4 for the two implementations made of the Gauss--Seidel algorithm. This is obtained despite the introduction of approximations in the algorithm. A comparison of results obtained with and without line overlaps for N$_2$H$^+$, HCN, and HNC shows that the $J=3-2$ line intensities are significantly underestimated in models where line overlap is neglected. } {} ",
        "introduction": "Our understanding of the evolution of the molecular complexity in astrophysical media relies both on the quality of the observations and on their interpretation. The increasing spatial resolution that will be provided by the new generation of ground-- and space--based telescopes (Herschel, ALMA, JWST) will enable finer geometrical structures to be infered in the objects under study. Hence, the best way to interpret the origin of the molecular line emission and to derive physical and chemical conditions is to use non local radiative transfer calculations with realistic geometries adapted to the angular resolution. While the Sobolev approximation can be safely used when a large velocity gradient is present, interstellar clouds often show linewidths that are thermal in nature or that involve low--velocity fields, so the use of such an approximation is rather doubtful, and a widely used method in the field of molecular astrophysics is the Monte--Carlo technique \\citep{bernes1979}. This method is particularly attractive due to its conceptual simplicity and because it is easily implemented to treat multidimensional geometries. Nevertheless, an inherent drawback of the method is that it suffers from numerical noise and presents a slow convergence rate when optically thick lines are modelled. Although line overlaps can be easily implemented \\citep{gonzalez1993}, computing time becomes very expensive when treating a case such as N$_2$H$^+$ where the two nitrogen atoms contribute to the hyperfine structure.  Since the 90s, various techniques have been introduced to solve the coupled set of radiative transfer and statistical equilibrium equations. These methods lead to substantial gains in term of calculation speed and accuracy of the solutions. Among them, the preconditioning of statistical equilibrium equations entails drastic improvement in the convergence properties of the iterative scheme for problems with large optical depths \\citep[see][]{rybicki1991,rybicki1992}. A crucial issue is that the Lambda iteration scheme does not always converge to the true solution for large optical depths \\citep[see][]{mihalas1978,auer1991}. Another progress consisted in the introduction by \\citet{trujillo1995} of the Gauss--Seidel algorithm for the solution of radiative transfer problems.  The description of this method was first made for the case of one-dimensional plane-parallel stellar atmospheres \\citep{trujillo1995,paletou2007} and was further generalised to the case of two and three-dimensional radiative transfer problems with multilevel atoms in Cartesian coordinates \\citep{bendicho1997,leger2007} and to the case of scattering line polarization \\citep{trujillo1999}. It was then described \\citep{asensio2006} and used \\citep{cernicharo2006} in spherical geometry. Although the geometry of interstellar clouds cannot always be assumed to be as simple as spherical, it enables a first guess more appropriate than plan--parallel geometry in order to model centrally condensed clouds. In this article, we give more details on the implementation in spherical geometry, and we extend the discussion to the case of line overlap. We find that in spherical geometry, the requirements of the original Gauss--Seidel algorithm cannot be fulfilled without increasing the computational time (even in the case of isolated lines). To address this problem, we introduced several approximations into the algorithm. Interestingly, despite these approximations, the usual increase in reaching convergence in terms of number of iterations persists. Moreover, these approximations do not entail false convergence, in comparison to results obtained within the $\\Lambda$ iteration scheme.  Several molecules used to trace the physical conditions in molecular clouds harbour a hyperfine structure that complicates the analysis of the observed emission \\citep[ see e.g.][]{daniel2007}. Indeed, the influence of the overlap is usually not considered when interpreting observations despite the fact that for a molecule with hyperfine structure, the line blending increases with the rotational quantum number. For high rotational transitions, all the hyperfine lines are usually overlapped in astrophysical objects. Previous works have shown that, for HCN, line overlaps play an important role in the pumping of molecular levels and have to be considered in order to interpret the relative intensities of hyperfine transitions \\citep[cf.][]{guilloteau1981,gonzalez1993}. One goal of this paper is to extend the discussion to other molecules, i.e. HNC and N$_2$H$^+$, and to derive the general radiative transfer effects that are common to the different species.  This article is organised as follows. In section \\ref{section:th1}, we first overview the short characteristics (SC) method. In section \\ref{theorie:ALI}, the main equations used in the preconditioning of statistical equations are given. In section \\ref{theorie:GS}, the implementation of the Gauss--Seidel algorithm in spherical geometry is discussed. In section \\ref{app_num}, a test case for H$_2$CO gives the main properties of the various algorithms. Finally, in section \\ref{app_LO}, the influence of introducing line overlap in the calculations is discussed for the N$_2$H$^+$, HCN, and HNC molecules. ",
        "conclusions": "We have presented technical details concerning the determination of the Lambda operator. Moreover, we gave details on the implementation of the Gauss--Seidel algorithm in spherical geometry. Finally, the numerical code was applied to the treatment of line overlap. The main conclusions are: \\begin{enumerate} \\item The diagonal of the Lambda operator, as obtained within the short characteristics method, has to be approximated in practical applications. The use of approximated diagonal terms still enable speeding the convergence. Moreover, the tests carried out did not show false convergence by comparison with LI calculations. \\item The implementation of the Gauss--Seidel algorithm in spherical geometry also requires some approximations. These are introduced by analogy with the plan--parallel case. \\item The convergence rate of the GS down and GS up--down are similar to what is obtained in plan--parallel geometry, i.e. they reduce the number of iterations to reach convergence by factors of 2 and 4 respectively. The gain in terms of computing time strongly depends on the numerical implementation, and in the present code, the GS up--down reduces the calculation time by a factor $\\sim$ 2 by comparison to the ALI scheme. \\item Line overlaps can lead to a large increase in the predicted intensities, the effect being greater for higher rotational lines. Especially, differential effects for HCN and HNC can alter the estimate of the abundance ratio X(HNC)/X(HCN), and models where line overlap is neglected will overestimate it.  \\end{enumerate}"
    },
    "0807/0807.1508_arXiv.txt": {
        "abstract": "{We revisit indirect detection possibilities for neutralino dark matter, emphasizing the complementary roles of different approaches. While thermally produced dark matter often requires large astrophysical \"boost factors\" to observe antimatter signals, the physically motivated alternative of non-thermal dark matter can naturally provide interesting signals, for example from light wino or higgsino dark matter.  After a brief review of cosmic ray propagation, we discuss signals for positrons, antiprotons, synchrotron radiation and gamma rays from wino annihilation in the galactic halo, and examine their phenomenology.  For pure wino dark matter relevant to the LHC, PAMELA and GLAST should report signals.} ",
        "introduction": "While a host of astrophysical measurements have precisely determined the amount of dark matter in our universe, we do not yet know its identity.  At present one could imagine that the dark matter is a weakly interacting massive particle (WIMP), an axion, or something more exotic.  This situation should change, perhaps soon.  If the dark matter is indeed a WIMP, evidence for it could be found both at the Large Hadron Collider (LHC) and a host of dark matter detection experiments, both direct and indirect.  In this paper, we will assume that the dark matter is a WIMP, in particular the lightest supersymmetric particle (LSP). The identity of the LSP depends on the details of supersymmetry breaking.  Determining its identity will be a necessary step towards understanding the cosmological history of our universe, and an important clue towards the determination of the underlying theory. A phenomenologically attractive candidate is the lightest neutralino.  We concentrate on a case that is both physically well-motivated and potentially gives large signals for dark matter indirect detection: a non-thermally produced LSP with large annihilation cross section. This scenario does not require additional anomalously large astrophysical ``boost factors'' to produce interesting signals.  By now, a large literature on the indirect detection of dark matter exists.  For reviews, see \\cite{JKG, HooperBaltz}.  We will place particular emphasis on a dark matter interpretation of positrons, for earlier work on this subject see, e.g., \\cite{BaltzPositron,Kane1and2,Freese1and2,ProfumoUllio}.  For the LHC to provide complementary data on the dark matter\\cite{Brhlik,PeskinBaltz}, it must be kinematically accessible.  Often, the dark matter is most efficiently searched for in the cascade decays of colored particles.  However, there can be a large gap between the dark matter mass and the lightest colored particle.  In models with gaugino mass unification, there is roughly a factor of seven between the WIMP candidate mass and the gluino mass.  In anomaly mediated models of supersymmetry breaking, the ratio is a factor of nine; in other models with non-universal gaugino masses, it can be a factor of a few.  Thus, if the gluino is to be produced copiously (say with a mass less than 2 TeV), the dark matter should not be too heavy.  In this paper, we will focus on a light mass region where the LSP is a wino with a mass of a few hundred GeV.  \\subsection{Thermal vs. Non-Thermal Production} Often, SUSY dark matter candidates are assumed to be produced from the primordial thermal plasma in the early stages of the universe (see e.g. \\cite{JKG} for a review). Under this assumption, the relic density of the LSP depends inversely on the annihilation cross section.  For a neutralino, $\\chi$, one finds \\cite{RPP,KolbTurner}: \\be \\label{omega} \\Omega_{\\chi} h^{2} \\approx 0.1 \\left(\\frac {3 \\times 10^{-26} \\rm{cm}^{3} \\rm{s}^{-1}}{\\langle \\sigma_{A} v \\rangle}\\right) \\ee For the case of a light neutralino LSP (a few hundred GeV or less), this typically restricts the neutralino to have a substantial bino component as pure wino and Higgsino states (co)-annihilate very effectively to weak gauge bosons.  But precisely because of the smaller annihilation cross section, the annihilation signals from bino-like dark matter can be disappointingly small unless one appeals to large ``boost factors.''  This issue is further exacerbated by the fact that bino annihilations are p-wave suppressed in the early universe, and are thus suppressed by powers of the final state masses today. If, as is often the case,  the final state is $b$-quarks, the annihilation rate in our galaxy can be very small.  Models with gaugino mass unification often do typically give rise to a bino LSP, with its associated small annihilation cross section.  One is then challenged to reduce the relic density to the observed value.  However, if one does not assume a simple unification of gaugino masses at the high scale, other possibilities arise, well-motivated by top-down models of supersymmetry breaking.  One attractive possibility is a wino LSP.  This naturally occurs in theories where anomaly mediation gives the dominant contribution to the gaugino masses\\cite{AnomalyMediation}.  It also occurs in string compactifications, see, e.g. \\cite{Acharya:2008zi}.  This type of dark matter can also occur in the simplest models of split-supersymmetry \\cite{Split,WellsPeV}, where the gauginos get anomaly mediated masses (with attendant loop suppression), but scalars receive large masses, suppressed only by the Planck scale.  A light wino has a large annihilation cross section, which is good for indirect detection, but also implies a small thermal relic abundance.  The solution to recovering the correct relic abundance comes from non-thermal production.  Often, the very same models that predict a wino LSP also provide mechanisms by which the LSP is produced non-thermally.  If particles decay into the wino below its freeze-out temperature, this can provide the correct relic abundance\\cite{WatsonKane}.  Excellent candidates for the decaying particle include gravitinos and weakly coupled moduli \\cite{Moroi:1999zb,Giudice:2000ex,Acharya2008:bk}.  Non-thermal production of dark matter leads to WIMPs with larger cross sections, since the standard thermal relic abundance calculation no longer applies.  Since the flux of anti-particles coming from dark matter annihilations depends linearly on the cross section, this implies that non-thermal production of dark matter may lead to larger fluxes that may be detectable in future indirect experiments.\\footnote{While we will concentrate on wino dark matter, the results are a bit more general.  In the region of interest, the winos annihilate nearly exclusively to $W$ bosons.  So, basically what we are probing is a dark matter candidate that annihilates to $W$'s with a given cross section.}  In the remainder of the paper, we review elements that enter any discussion of the indirect detection of dark matter.  First, we briefly review basics of cosmic ray propagation, as well as the form of the source term arising from dark matter annihilation. We then discuss constraints from both anti-protons and synchrotron radiation.  We then discuss prospects for observations of non-thermally produced wino dark matter in positrons and gamma rays.  With both PAMELA (a Payload for Antimatter Matter Exploration and Light-nuclei Astrophysics) and GLAST (Gamma Ray Large Area Space Telescope) in orbit, these two signals are particularly timely.  Throughout, we attempt to point out where astrophysical assumptions enter.  Finally, we comment on implications for the LHC, and briefly discuss implications for direct detection and indirect searches for dark matter via neutrinos. ",
        "conclusions": "A non-thermally produced wino is a well-motivated candidate for the dark matter observed in our universe.  Its large annihilation cross section could potentially allow it to explain the suggestion of an excess from HEAT/AMS-01, which could be confirmed soon by PAMELA.  However, to avoid conflict with bounds from gamma rays (and perhaps synchrotron radiation), the dark matter distribution cannot be too highly peaked towards the center.  There is already some tension in the case of an NFW profile.  This fact suggests that if PAMELA were to observe an excess in positrons that comes from dark matter, GLAST should follow with a confirmation.  Any candidate detection by PAMELA and GLAST will need to be examined in the context of direct detection experiments.  We do not do that here, since the pure wino LSP suggested by the present positron excess gives signals well below the current sensitivity of the current direct detection experiments. However, adding an admixture of Higgsino to the neutralino allows an increase in the direct detection cross section (via the $\\tilde{w}-\\tilde{h}-h$ coupling).  An increased Higgsino content also increases the capture cross section on the sun, allowing for a possible indirect detection via neutrinos. Thus, signals in these types of experiments could help to probe the Higgsino content of the LSP.  At the LHC, a pure wino of a few hundred GeV by itself may be difficult to observe via direct production.  However, it may be possible to find it in decays or associated production.  The sensitivity of this modes depends on the mass of the lightest colored mode.  In minimal models of anomaly mediation \\cite{AnomalyMediation}, the ratio of the wino mass to the gluino mass is a factor of nine.  So a 200 GeV wino implies a 1.8 TeV gluino, which might preclude an early discovery.  However, if the mass difference is smaller, as occurs in some models of non-thermal production then it might be possible to determine the wino nature of the LSP by looking for charged tracks, as recently studied in \\cite{Yanagida}. More generally, several LHC signatures will depend on the mass and type of the LSP, so we expect that careful studies will be able to test whether a candidate seen in indirect data is also present in LHC data."
    },
    "0807/0807.3027_arXiv.txt": {
        "abstract": "We present an approximation for the average density profile of dark matter haloes in the $\\Lambda$CDM cosmological model, which is accurate to within 10--15\\% even for large radial distances from $0.05R_{\\rmn {vir}}$ up to $10R_{\\rmn {vir}}$ for halo masses ranging from $10^{11.5}$ to $10^{15.0}$ $h^{-1}M_{\\odot}$. We propose a modified form of the Navarro, Frenk \\& White (NFW) approximation: $\\rho(r)=\\rho_{\\rmn{NFW}}(r)+ A (r/R_{\\rmn {vir}})^{-1}+B (1+r/R_{\\rmn {vir}})^{-1}$. This generalized expression, which is applicable to the external regions of dark matter haloes, only very slightly affects the density in the inner regions of haloes. The strong correlation among the different parameters in the model allows us to describe the profile in terms of just one parameter: the virial mass. We integrate our density profile to derive the enclosed mass in a sphere of a given radius and compare it with the NFW results. We find that the NFW underestimates the enclosed mass by more than 50\\% at $10R_{\\rmn {vir}}$, whereas our model reproduces the results from numerical simulations to within 2\\% accuracy even at this distance. We also use this new approximation to study the weak gravitational lensing and to obtain an analytic expression for the tangential shear. This allows us to quantify the contribution to the shear from the outer regions of the density profile. For the first time we calculate the difference between the tangential shear calculated via the NFW profile and the corresponding result when the external regions of haloes in cosmological simulations are taken into account. We find a 4\\% difference for all the mass ranges under study. ",
        "introduction": "\\label{sec:intro}  Measurements of the mass distribution of dark matter (DM) haloes are essential for the theory of structure formation. This issue has been addressed in two ways: observational and theoretical. The observational measurements have profited from recent gravitational lensing studies and also from measurements of galactic dynamics using large scale and massive galaxy surveys (e.g. \\citealt{Ma06}, \\citealt{Pr03}). Theoretical estimates made substantial progress thanks to recent advances in both top-hat gravitational collapse models and fully non-linear evolution via $N$-body simulations (e.g. \\citealt{Sa07}, \\citealt{As07}, \\citealt{Di07}).  Cosmological $N$-body simulations have been extensively used in order to provide predictions for the structure of DM haloes in the hierarchical clustering scenario. The Navarro, Frenk \\& White (NFW) density profile has become a convenient formula for the description of DM haloes in a broad mass range. This profile provides a reasonable fit for the halo density distribution, despite that many other functions have been suggested in order to improve accuracy (\\citealt{Mo98}, \\citealt{Se68}, \\citealt{Ei65} profiles, see also \\citealt{Me05} for reference). Nevertheless, the NFW density profile \\citep{Na97} was proposed and contrasted to fit the inner regions of DM haloes roughly up to the virial radius, where haloes are expected to be in equilibrium \\citep{Co96}. Therefore, it was not expected that the NFW analytical profile will describe the external density profile properly. Our aim here is to extend the validity of this approximation to describe the outer regions of DM haloes. Moreover, recent works (see \\citealt{Pr06}, \\citealt{Cu07}, \\citealt{Di07}) have shown that the DM virialized regions extend well beyond the formal virial radius especially for low-mass haloes. This reinforces our interest in studying haloes beyond the virial radius.  On the other hand, current observational techniques allow us to measure the mass distribution around galaxies and clusters at large distances. The dynamics of satellite galaxies (e.g. \\citealt{Za94}; \\citealt{Za97}; \\citealt{Pr03}; \\citealt{Br04}; \\citealt{Co07}) provide strong constraints on the shape of the density profiles of DM haloes in the outer regions. Another technique which has proven its usefulness despite its recent development, is the analysis of the weak lensing effect in gravitational lensing theory (e.g. \\citealt{Me99}; \\citealt{Ba01}; \\citealt{Ho04} \\citealt{He06}; \\citealt{Mas07}). In both observational techniques, the external regions under study are often beyond the virial radius. Yet, the use of the NFW profile to model observations is still very common in the determination of DM halo mass. Although the NFW profile is a great improvement as compared to the isothermal profile to describe for example the weak lensing observations (e.g. \\citealt{Ma06b}), it is not suitable for extrapolation at large distances beyond the virial radius. The study of the physical processes occurring in the external regions of DM haloes will provide unique insight into our understanding of the formation of DM haloes. A good example is the recent observational work by \\citealt{Ge08} who measured temperature and density profiles beyond the virial radius of a cluster and this allowed them to produce improved constraints on the mass and gas fraction profiles. A first step in this direction is the study of average density profiles in the outer regions of DM haloes in cosmological simulations.  In this paper, we present an approximation for the density profile of DM haloes well beyond the virial radius, obtained from high resolution cosmological simulations. This approximation focuses on the external regions of the haloes as an extension of the NFW formula that fits well the inner regions. We took the NFW model as our starting point because of its simplicity and with the aim to derive an expression which is easy to implement in many of the applications mentioned above. In order to describe representative DM halo density profiles and compare them with our approximation, we average over many individual density profiles for a given mass range. Hence, our conclusions apply for the ensemble of haloes of a given mass, with a halo-to-halo mass dependent scatter. As it is already known from the early work of \\citet{Na96}, individual haloes at a given range of mass usually show a similar profile. Thus, the analysis of average halo density profiles happens to be a reasonable option for our purposes.  The approximation presented in this paper includes free parameters which have to be fixed for a given mass range. However, we will also study some correlations among them and prove that only one parameter (for example the virial mass) is enough in order to provide a high quality fitting function for the inner and outer regions. Moreover, our approximation can be useful to obtain a suitable formula to describe the projected mass density of average DM haloes. We determine with this formula the tangential shear associated to gravitational lensing. We also have compared our results with those obtained using the NFW formula, in order to assess its validity at large distances from the halo centre.  The paper is organized as follows. In Section 2, we introduce our approximation that describes the density profiles of DM haloes well beyond $R_{\\rmn{vir}}$. In Section 3, we fit the numerical density profiles drawn from cosmological simulations. In Section 4, we show the application of our model to the measurement of the tangential shear in gravitational lensing. Finally, we discuss and summarize the main conclusions in Section 5. ",
        "conclusions": "In this paper we have presented a simple approximation for the DM density profiles of haloes with masses ranging from $10^{11.5}$ to $10^{15.0}h^{-1}M_{\\odot}$, which is valid even beyond the virial radius up to $\\sim 10R_{\\rmn{vir}}$. The expression we are suggesting here is $\\rho(s)=\\frac{\\rho_{s}}{cs\\left(1+cs\\right)^2}+\\left(\\frac{b_{1}}{s}+\\frac{b_{2}}{s+1}+1\\right)\\bar{\\rho}$, where $s=r/R_{\\rmn{vir}}$. This approximation is an extension of the NFW formula but includes two additional parameters. We have shown that these parameters are very well correlated with the virial mass, so that the density profile is effectively just a function of $M_{\\rmn{vir}}$. Other approximations found in the literature will prove to be useful to describe, with high accuracy, the halo density profile up to the virial radius. Yet, most of them fail when they are applied to fit the outer regions, where the mass predicted by extrapolation of these profiles is far below the actual mass in these regions.  The description of DM distribution far away from the halo centre is especially interesting. In particular, the halo--dark matter correlation function is related in a straightforward way to the \\textit{average} density profile. Although the transition from the one-halo term to the two-halo term in this correlation function is present before $10R_{\\rmn{vir}}$, our approximation has turned out to be a reasonable description of the DM distribution even at these distances. In order to build the numerical mean density profiles, we averaged over many hundreds of haloes from high-resolution cosmological simulations so that the profiles corresponding to most of our mass bins are entirely unaffected by statistics. This procedure for averaging density profiles is analogue to the stacking method, used in observational studies like the one by \\citet{Ma08} to infer the density profile of a cluster of galaxies. This similarity is useful for the comparison of the results from cosmological simulations to the real data.  This parametrization for the average density profile is accurate to within 10--15\\% in the range from $0.05R_{\\rmn{vir}}$ to $10R_{\\rmn{vir}}$. There are two main discrepancies from the numerical density profile which have a different origin: whereas the overestimation around $r=r_s$ is inherited from NFW profile, the overestimation just beyond the virial radius suggests that our model is not able to reproduce the steepest region. This steep region just outside $R_{\\rmn{vir}}$ is more pronounced for most massive haloes, suggesting a depletion of the halo outskirts due to dark matter infall (\\citealt{Pr06}, \\citealt{Cu07}). The presence of our additional terms with respect to the NFW formula has only a very small influence on the inner regions of the density profile, so that our approximation can also be considered as an extension of the NFW profile. At larger distances our model shows deviations around 20\\% in the range 10--30$R_{\\rmn{vir}}$ just before entering the asymptotic regime. These deviations are caused by our additional $(r/R_{\\rmn{vir}})^{-1}$ terms which improve the fit in the interesting region below $10R_{\\rmn{vir}}$ where the density is much higher. In any case, we must remark that our approximation implements the correct asymptotic behaviour: the density profile tends to the asymptotic value of the mean matter density of the Universe $\\bar{\\rho}$.  The cumulative mass inside a sphere of a given radius is underestimated by more than 50\\% at $10R_{\\rmn{vir}}$ by the NFW formula. On the contrary, it is much better approximated (to within 12\\%) when the NFW profile is modified by addition of the mean matter density, although with our model the difference with numerical density profiles is reduced even up to 1\\% in the range 1--9$R_{\\rmn{vir}}$. This is especially interesting for new measurements of the enclosed mass beyond virial radius in X-ray clusters \\citep{Ge08}, where plain NFW is still used even at $r>1R_{\\rmn{vir}}$. While current observations cannot distinguish between modified NFW and our approximation, in the near future they should be able to find the need for adding the mean matter density term to the density profile.  We have also presented an application for our approximation in the context of mass estimation using gravitational lensing effect. We derived expressions for tangential shear corresponding to different regions around the halo, which are in turn related to the different terms in our approximation. The contribution from the outer regions is small as compared to the contribution of the inner region, as expected. We calculated the difference between this tangential shear and the one derived from the NFW profile as a function of distance, showing that the inclusion of the outer regions produces a difference around 4\\% . This small difference could provide an observational test for the validity of our approximation, which has been derived from the results of cosmological $N$--body simulations. Although present resolution of weak lensing experiments prevents us from drawing a robust conclusion, the stacking of different observations should prove that this approximation, which includes the contribution of external regions, is more realistic than most of the so far proposed density profiles, which do not account for them.\\\\  HT wants to thank the I.E.S. Jos\\'e Cadalso for allowing him to combine this research with his teaching duties. FP, AJC and MASC thank the Spanish MEC under grant PNAYA 2005-07789 for their support. AK acknowledges support from NASA and NSF grants to NMSU. AJC acknowledges the financial support of the MEC through Spanish grant FPU AP2005-1826. MASC acknowledges the financial support of the CSIC through Spanish grant I3P."
    },
    "0807/0807.2398_arXiv.txt": {
        "abstract": "We have modeled optical to far infrared images, photometry, and spectroscopy of the object known as Gomez's Hamburger. We reproduce the images and spectrum with an edge-on disk of mass $0.3M_\\odot$ and radius 1600~AU, surrounding an A0 III star at a distance of 280~pc. Our mass estimate is in excellent agreement with recent CO observations. However, our distance determination is more than an order of magnitude smaller than previous analyses which inaccurately interpreted the optical spectrum. To accurately model the infrared spectrum we have extended our Monte Carlo radiation transfer codes to include emission from polycyclic aromatic hydrocarbon (PAH) molecules and very small grains (VSG). We do this using pre-computed PAH/VSG emissivity files for a wide range of values of the mean intensity of the exciting radiation field. When Monte Carlo energy packets are absorbed by PAHs/VSGs we reprocess them to other wavelengths by sampling from the emissivity files, thus simulating the absorption and re-emission process without reproducing lengthy computations of statistical equilibrium, excitation and de-excitation in the complex many level molecules. Using emissivity lookup tables in our Monte Carlo codes gives the flexibility to use the latest grain physics calculations of PAH/VSG emissivity and opacity that are being continually updated in the light of higher resolution infrared spectra. We find our approach gives a good representation of the observed PAH spectrum from the disk of Gomez's Hamburger. Our models also indicate the PAHs/VSGs in the disk have a larger scaleheight than larger radiative equilibrium grains, providing evidence for dust coagulation and settling to the midplane. ",
        "introduction": "The object known as Gomez's Hamburger \\citep{Ruiz87} exhibits the spectral and imaging characteristics of an edge-on circumstellar disk. High resolution HST imagery shows a bipolar reflection nebula of size 5.5\\arcsec by 2.5\\arcsec bisected by an optically thick dust lane, reminiscent of models and images of disks surrounding low mass Classical T~Tauri stars \\citep[e.g.,][]{WH92, Burrows96, Grosso2003}. The HST data support the circumstellar disk (or torus) interpretation first proposed by \\citet{Ruiz87} based on lower resolution optical imaging polarimetry and the object's broad band spectral energy distribution (SED) which is faint in the optical and bright at mid- to far-infrared wavelengths. Such SEDs are characteristic of edge-on circumstellar disks \\citep[e.g.,][]{Wood2002}. The disk hypothesis for Gomez's Hamburger has been strengthened by recent CO observations  \\citep{Bujarr2008} that are reproduced with a model of a Keplerian disk of mass $0.9 M_\\odot (d/500{\\rm pc})^2$, where $d$ is the distance to the source.  Optical spectroscopy of Gomez's Hamburger indicates the central source is an A0 III star. Using parameters typical to this spectral type and the integrated broad band spectrum, \\citet{Ruiz87} estimated the object to be at a distance of 2.9~kpc with a circumstellar disk of radius 8000~AU. However, as we will show below the disk is smaller and closer than estimated by \\citet{Ruiz87}. This is because they did not take into consideration that the integrated spectrum from an edge-on disk is strongly dependent on inclination \\citep{Whitney2003a}. At optical wavelengths edge-on disks are only seen in scattered light and can be several orders of magnitude fainter than systems where the central star is directly seen. In addition to the error in distance determination, the spectral type of the central source may not be accurate because the optical spectrum is entirely scattered light. Therefore, the combination of high resolution optical spectroscopy and radiation transfer modeling of the scattered starlight from the disk is required for more accurate spectral typing. However, in the absence of higher resolution optical spectroscopy, in our modeling below we adopt the A0 III spectral type from \\citet{Ruiz87}.  More recent Spitzer IRS observations of Gomez's Hamburger show strong mid infrared emission features typical of emission from polycyclic aromatic hydrocarbon (PAH) molecules and very small grains (VSGs). Such emission features are common in massive star forming regions \\citep[e.g.,][]{Churchwell2006} as the emission mechanism is sensitive to the strength of the exciting UV radiation from O and B stars. However, charged PAH molecules exhibit significant opacity at optical and infrared wavelengths, so that the PAH excitation may arise from softer photons \\citep{LD2002}, thus enabling PAH emission to be observed from the circumstellar environment of lower mass stars. Indeed, PAH emission has been detected from disks and envelopes around A stars and classical T~Tauri stars \\citep[e.g.,][]{Habart2004, Pontop2007, Geers2006}.  In this paper we present models for the HST images, broad band spectral energy distribution, and high resolution Spitzer IRS spectrum of Gomez's Hamburger. In doing so, in \\S2 we introduce extensions to our Monte Carlo radiation transfer codes that allow us to model emission from PAHs and VSGs. \\S3 presents the radiation transfer modeling of Gomez's Hamburger, and in \\S4 we summarize our findings and briefly discuss future applications of our new code developments. ",
        "conclusions": "We have presented a straightforward method to include the absorption and re-emission from PAH molecules and VSGs into Monte Carlo radiative equilibrium codes. Our technique utilizes pre-computed intensity-dependent emissivities from which we sample the emission spectrum of Monte Carlo energy packets that have been absorbed by PAHs/VSGs. Our codes are readily adaptable to include a range of sizes for the PAHs/VSGs, but in this first paper we adopt a grain size cutoff of 200~{\\AA}, above which we assume the dust grains attain radiative equilibrium temperatures. For grain sizes smaller than 200~{\\AA}, absorbed energy packets are reprocessed using the pre-computed emissivity files as described in \\S2. Incorporating the PAH/VSG emission in this manner allows us to compute model spectra and images very quickly without repeating lengthy calculations of statistical equilibrium and transient heating and cooling of the PAH/VSG component --- these calculations are in effect already incorporated in the pre-computed PAH/VSG opacities and emissivities presented in figures~1 and 2.  In a future paper we will extend our technique for sampling emissivities to be dependent not only on the intensity of the exciting radiation field, but also on the spectrum of the absorbed energy in each grid cell of our Monte Carlo simulation. Similar to the mean intensity calculation, determining the spectrum of the absorbed energy throughout our grid is straightforward to implement in Monte Carlo radiation transfer simulations. This more accurate treatment of the PAH/VSG reprocessing will be possible using planned calculations of PAH/VSG emissivity files for a range of exciting spectra comprising model stellar atmosphere spectra and black-bodies for a range of temperatures and reddening.  As a first application of our new code we modeled the object known as Gomez's Hamburger. We find that the broad band SED, high resolution images, and {Spitzer} IRS spectrum can be reproduced with an edge-on circumstellar disk with a component of PAHs/VSGs comprising 1.5\\% of the dust mass. Our Monte Carlo radiation transfer using the adopted dust opacities ensures that very few high energy photons penetrate into regions of large optical depth. Therefore there is little or no PAH/VSG emission from such regions and the PAH/VSG emission is largely confined to a surface layer exposed to the direct stellar radiation. This is what is observed in the limb brightened bubbles in the GLIMPSE survey \\citep{Churchwell2006}. In future papers we will use our code to model the bubbles observed in the GLIMPSE survey to address the question of PAH/VSG survival around hot stars and in H~{\\sc ii} regions in addition to modeling the dust content within the H~{\\sc ii} regions and the effects this may have on determining the ionizing luminosity of the central sources. We will also investigate emission from PAHs/VSGs in protostars, thus extending the radiative equilibrium dust modeling presented by \\citet{Whitney2003b}.   \\begin{figure} \\plotone{f1.eps}  \\caption{Solid line shows the wavelength dependent opacity for the PAH/VSG (dust plus gas) model of \\citet{DL2007} for grains with sizes $a < 200$~{\\AA}. Dot-dashed line is for the radiative equilibrium dust model of \\citet{Wood2002}. For comparison the dashed line shows the opacity for an ISM dust plus gas mixture.} \\end{figure}  \\begin{figure} \\plotone{f2.eps} \\caption{Emissivity for PAHs/VSGs with size $a < 200${\\AA} for different values of $U$, the energy density of the exciting radiation field relative to the average interstellar value. The emissivities shown are in the range $1 \\le U\\le 10^7$. The solid line shows the emissivity for $U=1$ and the long dashed line is for $U=10^7$.} \\end{figure}  \\begin{figure} \\plotone{f3.eps} \\caption{SED data (dots and red curve) and our best disk model (upper solid line) for Gomez's Hamburger. The lower solid line at mid-IR wavelengths shows the same disk model without the PAH/VSG components.} \\end{figure}   \\begin{figure} \\plottwo{f4a.eps}{f4b.eps} \\caption{Three-color images at $B$ (blue), $V$ (green), and $R$ (red) bands showing HST data (left) and model (right) images for Gomez's Hamburger} \\end{figure}"
    },
    "0807/0807.0162_arXiv.txt": {
        "abstract": "In a continuing survey of active galactic nuclei observed by the $\\it Far~Ultraviolet~Spectroscopic~Explorer$, we provide a deeper analysis of intrinsic absorption features found in 35 objects. Our survey is for low-redshift and moderate-luminosity objects, mostly Seyfert galaxies. We find a strong correlation between maximum radial velocity and luminosity. We also examine the relationships between equivalent width (EW), full width at half maximum, velocity and continuum flux. The correlation between velocity and luminosity has been explored previously by Laor \\& Brandt, but at a significantly higher redshift and heavily weighted by broad absorption line quasars. We also have examined each object with multiple observations for variability in each of the aforementioned quantities, and have characterized the variation of EW with the continuum flux. In our survey, we find that variability of O VI $\\lambda$$\\lambda$1032, 1038 is less common than of the UV doublets of C IV and N V seen at longer wavelengths, because the O VI absorption is usually saturated. Lyman $\\beta$ absorption variability is more frequent. In a target-by-target examination we find that broad absorption line and narrow absorption line absorbers are related in terms of maximum outflow velocity and luminosity, and both can be exhibited in similar luminosity objects. We also find one object that shows radial velocity change, seven objects that show equivalent width variability and two objects that show either transverse velocity changes or a change in ionization. ",
        "introduction": "It is commonly known that intrinsic UV absorption occurs in a large fraction ($\\sim$ 50\\%) of Seyfert galaxies. First noticed by Anderson \\& Kraft (1969), several surveys have been done in recent years (Crenshaw et al. 1999; Kriss 2002; and Dunn et al. 2007, hereafter Paper I) to characterize these features. In the first step of our survey we examined low-z active galactic nuclei (AGN) (Paper I). Seyfert galaxies are moderately luminous, relatively nearby, and mostly spiral galaxies, which make up a large portion of our survey, while the remainder are low-redshift quasars. We concentrate only on AGN that show broad permitted emission lines and narrow emission lines (i.e., Seyfert 1's, Kachikian \\& Weedman 1974). This is because intrinsic absorption is only readily detectable if seen against the broad emission lines and central continuum source characteristic of type 1 objects. Seyfert 2 galaxies have only narrow permitted and forbidden lines, and thus intrinsic absorption cannot be easily detected.  Crenshaw et al. (1999) performed a survey of Seyfert galaxies that had been observed with the Faint Object Spectrograph (FOS) and investigated the properties of the intrinsic absorption lines. They found that nearly 60\\% of the targets in the survey showed narrow C IV absorption features (FWHM $<$ 500 km s$^{-1}$) and outflow velocities up to 2000 km s$^{-1}$. Also in this survey, they found a link between warm X-Ray absorbers and UV intrinsic absorption, in the sense that Seyfert galaxies that show one also show the other.  Broad absorption line (BAL) quasars have been studied even more extensively. These objects show lines with much larger velocity widths, up to $\\sim$ 10,000 km s$^{-1}$ (Weymann et al. 1981, Turnshek 1988), which appear in $\\sim$ 10\\% of all quasars (Foltz et al. 1990). This is a much smaller percentage than the percentage that exhibit narrow absorption lines only, so the global covering factor is an important consideration in determining their nature.  Laor \\& Brandt (2002) looked for various trends concerning intrinsic absorption in data from the {\\it Hubble Space Telescope (HST)} and the {\\it International Ultraviolet Explorer (IUE)}. Their survey was derived from the PG quasar sample (Boroson \\& Green 1992). One of the more notable trends they found was a dependence of maximum outflow velocity on luminosity (v$_{max}$  $\\propto$ L$^{0.62 \\pm 0.08}$). Their plot contains both BAL and narrow absorption lines (NALs). They found that the equivalent width (EW) also followed an increase with increasing luminosity. This trend is based on Soft X-Ray Weak Quasars (SXWQs), which show larger values of EW and v$_{max}$ than any other quasar across the luminosity span. They furthermore showed that $[$O III$]$ EW and v$_{max}$ are correlated. SXWQs showed a significant correlation of EW versus v$_{max}$. Of the 56 objects they investigated, 28 showed intrinsic UV absorption. This agrees well with the $\\sim$60\\% frequency Crenshaw et al. (1999) found in their lower-luminosity sample and the $\\sim$50\\% frequency we found in Paper I.  For a purely BAL sample, Ganguly et al. (2007) recently used data from the Sloan Digital Sky Survey (SDSS) to investigate their properties. They conclude that the overall fit between v$_{max}$ and luminosity found by Laor \\& Brandt applies to their much larger sample of over 5000 objects. In Ganguly \\& Brotherton (2008), they compile all intrinsic absorption features for all known surveys and show that the fit still holds.  Our knowledge of the variability for intrinsic absorption is limited. Variability is important because it is our only link to several key physical parameters concerning the AGN. Cromwell \\& Weymann (1970) were the first to detect variability, in the Seyfert 1 NGC 4151. Later, in data from the {\\it IUE}, more absorption variability was seen (Crenshaw et al. 1999). Nearly all of the variability has been due to changes in the ionic column densities, as a result of changes in the ionizing flux from the continuum source. However, in four Seyfert galaxies (NGC 3516, NGC 3783, NGC 4151, and NGC 5548), bulk transverse motion of gas across the line of sight was identified as a likely culprit (Crenshaw et al. 2003). In the bright Seyfert galaxy NGC 4151, both ionization change and bulk transverse motion were detected. Kraemer et al. (2006) performed a detailed analysis of the absorption from the Space Telescope Imaging Spectrograph (STIS) data. For the component they dubbed D+Ed, they found a transverse velocity lower limit of v$_t$ $\\ge$ 2100 km s$^{-1}$.  NGC 3783 has shown previously to be variable (Gabel et al. 2005, Maran et al. 1996). In 1993 February data from STIS, NGC 3783 showed a spectrum with no narrow intrinsic absorption. By 1994 January, two obvious absorption features arose for each member of the C IV doublet. By 2000 February another pair of absorption lines had emerged. Gabel et al. (2003) showed that the radial velocity changed as well. In a series of several spectra the radial velocity of the line changed in two time spans. They first found a change of $\\sim$35~km~ s$^{-1}$ and in the second span they saw a change of $\\sim$ 55 km s$^{-1}$. This is the first well-documented case of a change in radial velocity for an intrinsic absorber in a Seyfert galaxy. Transverse velocity was also measurable, where they found a lower limit of v$_t$ $\\ge$ 540 km s$^{-1}$, which is comparable to the radial velocity ($\\sim$$-$450 km s$^{-1}$). Motivated by their findings, we undertook a survey of O VI and Ly$\\beta$ absorption in $\\it Far Ultraviolet Sepctroscopic Explorer (FUSE)$ spectra of low-redshift AGN.  The penultimate goal of our research is to understand the general dynamics of mass outflow and to ultimately help complete the picture of the AGN central engine. Several dynamical scenarios have been offered as explanations for the outflow (see Crenshaw et al. 2003 for a review of the models). One method, Compton heated winds assumes that EUV and X-ray photons from the accretion disk are irradiating the gas at larger distances and creating a thermal wind (Begelman et al. 1983). Radiative driving is also a promising method to explain fast-moving outflows such as BALs (deKool \\& Begelman 1995, Arav et al. 1994). Another possibility is that magnetic fields from the accretion disk are lifting plasma from the disk in magnetohydromagnetic (MHD) flows, where the plasma follows a helical magnetic field line (Blandford \\& Payne 1982). Depending on how far from the central core these clouds can follow the field lines, this could also be an explanation for the high transverse velocities.  In Paper I we cataloged 35 AGN with intrinsic absorption and presented approximate velocities. We found 11 new UV intrinsic absorption objects in a sample of 90 nearby AGN observed by $\\it FUSE$. We calculated the frequency of O VI absorption to be on the order of 50\\%, slightly lower than the value Crenshaw et al. (1999) found for C IV. The global covering factor, based on an estimate of the average covering factor in the line of sight ($<$C$_f$$>$) was similar to previous results of $\\sim$0.4. We continue our study of that absorption measurements of the absorption lines we found in the 35 objects from Paper I. ",
        "conclusions": "Our goal was to explore the Dunn et al. (2007) $\\it FUSE$ survey in depth for correlations and variability. We sought to characterize the intrinsic absorption seen in nearly 50\\% of all nearby active galaxies in order to help characterize the outflow from the nuclear engine.  As shown by Laor \\& Brandt (2002), there is a relation between the maximum outflow velocity and the luminosity. Their survey contained moderate redshift objects and BAL quasars, while our survey was for smaller redshifts and luminosities. However, as shown previously, our relation between v$_{max}$ and L is consistent with theirs. In both our surveys there is a distinct upper limit on the maximum velocity dependent upon the luminosity of the object. This leads us to believe that both classical narrow intrinsic absorption lines and broad absorption lines are governed by the same relationship, thus they may be related, in the sense that they may share a common driving mechanism. Due to their simultaneous presence in WPVS 007 and PG 1411+442 however, we feel that the nature of their origins may be different.  Laor \\& Brandt (2002) also found for higher redshift galaxies that EW \\& luminosity are correlated for C IV. We examined correlations between the observed lines and the AGN properties. We found no link between the EW and the luminosity. Because EW increases with column density of the particular ion in question, it would seem that an increase in flux would increase the O VI EW in cases where the ionization parameter is low and vise versa for a highly ionized cloud, at least until the line reaches heavy saturation. We do not see this trend for either O VI or Ly$\\beta$. This is most likely due to the heavy saturation in O VI and even Ly$\\beta$ in many cases while C IV and N V do not saturate as easily.  It has been seen in the past that some objects show bulk transverse motion along with the more easily detectable radial motion (Crenshaw et al. 2003). Of the 35 objects available we have 22 with more than one observation. Two of these objects have only one usable observation, leaving 20 total useful data sets. In the survey we find 8 that show change. The frequency of variability events seems to be low, in the $\\it FUSE$ data. With 22 available objects, only 8 show variability. This could mean the events are rare, that only weak absorption is subject to this change and/or the sample is biased by small number statistics, or the O VI saturation makes detecting variability difficult.  None of our objects with EW changes can be explained with transverse bulk motion. We have 7 that can be explained with ionization changes. In the case of Mrk 509 with modeling of the subcomponents, we can find a distance and density (n$_H$) of the third subcomponent of 0.3 pc and 3.8$\\times$10$^{8}$ cm$^{-1}$, respectively. Also, we find one object that shows evidence for radial velocity changes (PG 0804+761); we currently have no explanation for this.  We find in Mrk 817 that the absorption appears in relatively weak lines, and furthermore even the changes in the absorption features for Mrk 509 are in relatively weak lines. This is the case for most transversely moving clouds seen in the past. With the low signal-to-noise $\\it FUSE$ data it may be that these weak lines are swamped by the surrounding noise and only strong absorption features are noticeable, or it may be that bulk transverse motion is not a common event. Given that these changes are in weak absorbers, it suggests that strong absorbers appear more static in the far UV than weak absorbers due to saturation and blending in many of the lines."
    },
    "0807/0807.1091_arXiv.txt": {
        "abstract": "{The evolution of quasar host galaxies is still hardly studied at high redshifts ($z>2$), although this is a very interesting redshift range as both the quasar activity and the star formation rate density have their peak at $z\\approx2$--$3$. This makes is especially interesting to study properties of quasar host galaxies, such as the star formation rate or the black hole mass at this redshift. A proper classification of quasar host galaxies at high redshift would help to answer the question which role quasars play in galaxy evolution.} {In this paper we study different properties of quasars and their host galaxies at high redshifts up to $z\\approx3.4$. We compare our results to those of other authors and discuss the correlation between galaxy evolution and quasar activity.} {We analysed broad-band images in eight filters (from U to K) of eight quasars in the FORS Deep Field with redshifts between $z=0.87$ and $z=3.37$. A fully 2-dimensional decomposition was carried out to detect and resolve the host galaxies. Using the magnitudes in different filters, we investigated the presumed galaxy type, galaxy age, star formation rate and the mass of the central black hole of the host galaxies. In addition, the masses of the central black hole for the whole sample were derived using the corresponding quasar spectra.} {We were able to resolve the host galaxies of two out of eight quasars between $z=0.87$ and $z=2.75$. Additionally, two host galaxies were possibly resolved.} {The resolved low-redshift quasar ($z=0.865$) was identified as a late type galaxy with a moderate star formation rate of $1.8 M_{\\odot}/yr$ hosting a supermassive black hole with a mass of $\\lesssim10^{8}M_{\\odot}$. The resolved high redshift host galaxy ($z=2.7515$) shows moderate star formation of $4.4$--$6.9 M_{\\odot}/yr$, for the black hole mass we found a lower limit of $>10^{7}M_{\\odot}$. All quasars host supermassive black hole with masses in the range $\\sim10^{7}$--$10^{9}M_{\\odot}$. Our findings are well consistent with those of other authors.} ",
        "introduction": "The quasar activity and the star formation rate density apparently peak at $z\\approx2$--$3$ (see e.g. Nandra et al. \\cite{Na05}, Appenzeller et al. \\cite{Ap04}). This implies a strong link between quasar activity and galaxy formation. Additionally, this is the redshift range at which the luminosity of quasar host galaxies seems to turn over (Falomo et al. \\cite{Fa08}), even though this turnover is poorly constrained due to a small number of resolved host galaxies at high redshifts. Although the formation and evolution of nuclear activity is not yet understood it seems to be clear that there is a strong link between the evolution of galaxies and nuclear activity. After the observational confirmation of supermassive black holes in nearby early-type galaxies, correlations between the properties of the host galaxy and its central black hole were found (e.g. Novak et al. \\cite{No06}). This affirmed the assumption that the evolution of early-type galaxies and the central black hole, and thus nuclear activity, are linked.  The evolution of host galaxies of quasars has been studied over a wide redshift range from very low (e.g. Dunlop et al. \\cite{Du03}; Percival et al. \\cite{Pe01}) up to very high redshifts of $z\\sim3$ (e.g. Ridgway et al. \\cite{Ri01}; Falomo et al. \\cite{Fa04}, \\cite{Fa08}; Jahnke et al. \\cite{Ja04}; Kotilainen et al. \\cite{Ko07}). It is nowadays established that luminous quasars typically reside in luminous massive early-type galaxies. The luminosity of host galaxies seems to increase with redshift, with the host galaxies of radio-quiet quasars apparently being on average less luminous than the ones of radio-loud quasars (e.g. Kukula et al. \\cite{Ku01}). Quasar host galaxies at high redshifts of $z=2\\sim3$ have not yet been studied elaborately, since due to high QSO to host light ratios, extremely deep high resolution images are required to assure a sufficiently high signal to noise ratio. Consequently, most studies of quasar hosts were carried out using the Hubble Space Telescope (HST) (e.g. Kukula et al. \\cite{Ku01}; Ridgway et al. \\cite{Ri01}; Jahnke et al. \\cite{Ja04}). Jahnke et al. (\\cite{Ja04}) found rest-frame near-UV absolute magnitudes of --$20.0<M_{200nm}<$--$22.2$, but did not study morphological properties. Kukula et al. (\\cite{Ku01}) found absolute V-band magnitudes around --$22<M_{V}<$--$24$, early-type morphologies and differing evolution of the host galaxies of radio-quiet quasars (RQQ) and radio-loud quasars. In their samples, host galaxies of RQQ are fainter than those of RLQ on average with a growing gap towards higher redshifts. Adaptive Optics (AO) observations are beginning to play a major role as well. Recent studies of high redshift quasars (Kotilainen et al. \\cite{Ko07}, Falomo et al. \\cite{Fa08}) using AO showed no deviation from former findings.  Since quasar host galaxy studies are typically carried out in 1--2 filters, galaxy types are derived via their morphological appearance. Thus, investigations of stellar populations, star formation rates or ages of the host galaxy are difficult or even impossible. Recently, a spectroscopic study of quasar host galaxies at low redshift ($z\\sim0.3$) was performed by Letawe et al. (\\cite{Le06}). Most host galaxies were found to have Sc-like populations, showing large amounts of ionised gas and signs of interaction. Canalizo \\& Stockton (\\cite{CS01}) analysed a sample of low-redshift radio-quiet quasars ($z\\leqslant0.4$) spectroscopically. They found tidal interaction-induced starbursts in almost all of their objects. For 5/9 sources they found star-forming regions in the centre which is consistent with merger-induced starbursts. They found clues that central starbursts and quasar activity were triggered simultaneously. Hyv\\\"{o}nen et al. (\\cite{Hy07}) analysed the colours of 18 low redshift ($z\\leqslant0.3$) BL Lac objects and found old evolved stellar populations with an additional about 1 Gyr old young stellar population. Absence of signs of interaction implies that star formation and nuclear activity where not triggered simultaneously. Schramm et al. (\\cite{Sc08}) studied high redshift ($z\\approx3$) high luminosity quasars in restframe optical and found host with stellar populations indicating recent massive star formation. The different results seem to indicate that stellar populations of quasar host galaxies depend on the luminosity of the quasar itself, the radio-loudness and the redshift.  In this paper we present the analysis of eight quasars in the FORS Deep Field (see Heidt et al. \\cite{heifdf}) with a redshift range from $z\\approx0.9$-$3.4$. The field was observed with the VLT and HST from near UV to near IR (UBgRIJKs,F814W), allowing multiband analysis of the host galaxies from rest-frame UV to restframe optical. Due to very long exposure times (up to 44400s in U), the images are extremely deep in the optical with 50\\% completeness limits up to 27.7 in the B-band. Additionally, the seeing is excellent, with the best seeing being 0.''53 in the I-band. This gives us a unique possibility to study the probable galaxy type, star formation rate and mass of central black hole of the host galaxies.  The paper is organised as follows. In Section 2 we shortly present the quasar sample and summarise the observations. In Section 3 we describe the data analysis. In Section 4 we discuss the results. We investigate the galaxy types, host galaxy formation timescales, star formation rates and masses of the central black holes and compare our results to those of other authors. We discuss our results regarding the link between galaxy formation and quasar activity in Section 5 and compare our results to those of other authors. Our results are summarised in Section 6.  All magnitudes are given in the Vega-system. The cosmology used is $H_{0}=70\\textrm{km/s/Mpc}, \\Omega_{\\Lambda}=0.7, \\Omega_{m}=0.3$, unless otherwise stated. ",
        "conclusions": "In this paper we have presented a multiband (eight filters from near UV to NIR) analysis of eight high-redshift quasars with redshifts between $z=0.87$ and $z=3.37$. For two of the quasars (FDF0809 at $z=0.87$ and FDF6007 at $z=2.75$) the host galaxies could be resolved. For two quasars the host galaxy is regarded as marginally resolved, therefore magnitudes found are merely upper limits. For the resolved and marginally resolved host galaxies we analysed their age, stellar populations and SFR and estimated their central black hole masses. The results can be summarised as follows:  By analysing the colours of the resolved host galaxies we tried to constrain the galaxy type and formation redshift of the galaxy. FDF0809 is consistent with an either E/S0 or a Sb-type galaxy, while extreme bursts can clearly be rejected. Additionally, using simulation of spectra we were able to confirm stellar populations consistent with a spiral type galaxy. For the other three host galaxies, the colours are consistent with either an E/S0 or Sb-type galaxy with a formation at $z=5$ or $z=3$ or a long burst with a formation at $z=3$.  The determination of the star formation rate is critical, as the used model breaks down if the star formation rate changes on small time scales ($\\lesssim10^{8}$yr). This limitation of the model was verified by calculating SFRs for model galaxies. Thus the results should be handled with care, as the SFR is potentially overestimated. For all objects, the SFR is estimated to be between 0.03 and 33 $M_{\\odot}/yr$, which is actually consistent with values derived by Jahnke et al. (\\cite{Ja04}) using the same method at similar redshifts.  The mass of the central black hole was calculated using the C IV line width for 6/8 sources. Within the uncertainties we derived BH masses between $\\sim 10^{7}$ and $10^{9} M_{\\odot}$. The results are consistent with those of Vestergaard et al. (\\cite{Ve08}) and Kuhlbrodt et al. (\\cite{Ku05}), considering the fact that both studies observed quasars more luminous than those from our sample.  Both clearly resolved host galaxies are consistent with those studied by other authors. Being a late-type galaxy, FDF0809 fits well into the picture of quasar host galaxies at low redshift being hosted by low luminosity quasars. FDF6007, showing moderate star formation rate and a UV luminosity consistent with those of similar host galaxies studied by Jahnke at al. (\\cite{Ja04}), seems to be no expection either. Both fit into the general picture of quasar host galaxies being either late-type galaxies, like FDF0809, interacting galaxies or galaxy that have recently undergone strong star formation (FDF6007 might fit into this class).  Analysis of high-redshift quasar hosts on the one hand requires extremely deep images to guarantee high signal-to-noise. Otherwise the fainter host galaxies can not be detected. On the other hand high resolution imaging is necessary to be able to resolve high redshift hosts, which are just merely more extended than a point-like source. Further exploration of the evolution of the mass of the central black hole and properties of the galaxies like SFR or stellar populations will help to understand triggering and evolution of quasars, their host galaxies and galaxies in general."
    },
    "0807/0807.4950_arXiv.txt": {
        "abstract": "We have employed laser guide star (LGS) adaptive optics (AO) on the Keck II telescope to obtain near-infrared (NIR) images in the Extended Groth Strip (EGS) deep galaxy survey field. This is a continuation of our Center for Adaptive Optics Treasury Survey (CATS) program of targeting $0.5<z<1$ galaxies where existing images with the {\\it Hubble Space Telescope} (HST) are already in hand.  Our AO field has already been imaged by the Advanced Camera for Surveys (ACS) and the Near Infared Camera and Multiobject Spectrograph (NICMOS).  Our AO images at 2.2$\\mu$m ($K'$) are comparable in depth to those from HST, have Strehl ratios up to 0.4, and FWHM resolutions superior to that from NICMOS.  By sampling the field with the LGS at different positions, we obtain better quality AO images than with an immovable natural guide star.  As examples of the power of adding LGS AO to HST data we study the optical to NIR colors and color gradients of the bulge and disk of two galaxies in the field with $z=0.7$. ",
        "introduction": "Although the bulges of disk-dominated galaxies can appear to be the remnants of long-ago mergers, their kinematics may point to close ties with ongoing formation processes in the disk \\citep[for reviews, see][and references therein]{Kormendy2004, Wyse1997}. The strong correlation of bulge and disk color seen in local galaxies \\citep{deJong1996, Peletier1996} suggests a shared history, with similar ages and metallicities in their stellar populations. So it is important to probe bulge and disk colors at significant look-back times.  Of particular interest might be the link between bulge and central black hole mass, for example, as that correlation plausibly develops at an epoch when bulges are first undergoing collapse or mergers, providing fuel for the black hole \\citep{Maggorian1998, Ferrarese2000, Gebhardt2000}.  Separating the bulge from disk in images of high-redshift field galaxies is a difficult task though, because bulges are small; typically less than a few kpc across in the local universe \\citep{Andredakis1995}. At redshift $z=0.5$ each kpc corresponds to only 0\\farcs16 in a concordance cosmology ($H_0=70$ km ${\\rm s}^{-1}$ Mpc, $\\Omega_{\\rm m}=0.3$, $\\Omega_{\\Lambda}=0.7$) which we adopt throughout. The {\\it Hubble Space Telescope} (HST) has been used to decompose many $z>0.5$ field galaxies into bulges and disks in the optical \\citep{Schade1995, Simard2002, Koo2005}. But equally high resolution near-infrared (NIR) color maps would also be valuable, as they should lessen contamination by dust and probe older stellar populations, especially for galaxies with $z>0.7$ where the $V$ filter desirably samples rest-frame $U$, which is below the 4000~\\AA~break. Our initial efforts to do high-resolution NIR photometry of $0.5<z<1$ galaxies with natural-guide-star (NGS) AO were hampered by poor control over the point-spread function (PSF) \\citep{Steinbring2004}. Better results are now possible with the new generation of laser guide-star (LGS) AO systems.  The motivation for LGS AO is to increase sky coverage by providing a pointable artificial beacon on which to guide. LGS AO will also limit anisoplanatism, and thus improve image quality, by allowing the wavefront sensor and deformable mirror to be aligned with the target, instead of a bright nearby star. One complication is that the laser spot can only be used to sense high-order aberrations, not phase tip and tilt. Thus, a nearby star, which can be several magnitudes fainter than that used in full NGS mode, is still needed to account for image motion. A separate sensor guides on this ``tip-tilt star\" and drives a fast steering mirror.  Despite this added complexity, tests with the Lick LGS AO system show that the resulting point-spread function (PSF) is no harder to constrain than with a conventional NGS system \\citep{Steinbring2005}.  Here we discuss our ongoing effort to employ the latest AO technology to obtain images of high-redshift field galaxies. This work is part of the Center for Adaptive Optics (CfAO) Treasury Survey (CATS) \\citep{Melbourne2005, Koo2006}. We have employed Keck LGS AO in the Extended Groth Strip (EGS), an HST field extensively imaged in both the optical and near infrared.  The data are presented in Section~\\ref{data}, including a description of the improved PSF provided by the laser and the comparison of AO and HST images. Details of photometry and image decomposition follow for two $z=0.7$ galaxies in the field in Section~\\ref{photometry}, and color gradients between the bulge and disk are discussed further in Section~\\ref{discussion}. ",
        "conclusions": "\\label{discussion}  The integrated ACS $I$ magnitude of GSS 044\\_6719 is 21.1, which gives an absolute rest-frame $B$ magnitude of -21.0 using the $K$-corrections of \\cite{Simard2002}. For GSS 044\\_6712 ($I=20.3$) this is $M_B=-21.4$. These $M_B$, along with a disk scale length $R_{\\rm scale}\\approx 2.4$ kpc found by GIM2D fitting are similar to those of the three DEEP galaxies we observed in Keck NGS mode reported in \\cite{Steinbring2004}. These are seemingly normal spirals, near the peak of the luminosity function at $0.6<z<0.8$ as determined in the DEEP2 survey \\citep{Willmer2006}. The Keck DEIMOS spectrum from that survey also provides velocity dispersion for GSS 044\\_6719 of $\\sigma = 24\\pm8$ km $s^{-1}$, consistent with a disk-dominated galaxy. From our GIM2D fitting, we can see that while the bulge is weak in the optical (restframe UV), it may comprise 20\\% of the $K'$ (restframe~$\\sim J$) light. Interestingly this galaxy is detected at $24\\mu{\\rm m}$ (restframe 14 $\\mu$m) with Spitzer \\citep{Davis2007}, indicating significant dust may be present.  In an attempt to understand the very red optical to NIR colors of these galaxies, and especially the bulge of GSS 044\\_6719, we used the Galaxev code of \\citet{Bruzual2003} to compare population synthesis models with the data. Galaxev constructs spectra for arbitrary star formation histories by adding together the single-stellar-population spectra from various ages, weighted appropriately. We produced $V-I$ versus $I-K'$ colors from these spectra by scaling them according to the \\citet{Calzetti1997} dust attenuation curve, shifting them to $z=0.7$, and then convolving this with the appropriate filter bandpasses. We constructed two models, both with solar metalicity ($Z=0.02$). The first has an exponential ($\\tau=3$ Gyr) star formation rate, and the other passively evolves - that is, there is no active star formation - after a brief (50 Myr) burst. These were also run assuming a large amount of dust ($A_V=10$), for a total of four models. The results are shown in Figure~\\ref{figure_bulge_colors}, for ages of 10 Myr, 100 Myr, 1 Gyr, 2 Gyr, 4 Gyr, and 7 Gyr (the age of the universe at $z=0.7$ for our chosen cosmology) after the onset of star formation. The total colors of GSS 044\\_6719 from aperture photometry are well fit by an old $\\tau$ model with no dust; a value of $\\tau>3$ Gyr will just shift the track blueward. The core (0\\farcs4 aperture) is redder than the total light, and for a uniform age throughout the galaxy more closely matched by the single-burst+passive model. By adding dust (and keeping $\\tau=3$) the same models can also hold for the other galaxy, GSS 044\\_6712. However, the bulge of GSS 044\\_6719 can only be fit by either a $\\tau$ model or a single-burst+passive model after invoking significant dust. One way to reconcile this bulge color with the total light, which seems essentially unreddened, would be a clumpy distribution of dust in the galaxy. We could be seeing an old bulge through a dust clump in a thick disk. The bulge is very small, and so any obscuration would affect its color much more dramatically than the disk. The same might be true for an AGN which, highly obscured, could contribute to these red colors, as well as produce significant mid-IR emission.  The combination of the high resolution and long wavelength of our NIRC2 AO imaging, along with existing HST imaging, give us an opportunity to measure and interpret the optical to NIR color gradients of the bulge and disk of two galaxies at $z=0.7$.  Aperture photometry shows that the centers of the galaxies are significantly redder than the outer regions.  While we have only explored a few stellar population models in our interpretation of this, we find that the color gradient can be explained by two different star-formation histories for bulge and disk, with similar ages. However, aperture photometry is much more difficult to interpret in terms of galaxy subcomponents, as the light in any given aperture is a combination of light from the galaxy bulge and disk. Furthermore, our GIM2D analysis shows that the bulge of GSS 044\\_6719 is extremely small, so only the central aperture we used is likely to contain a significant amount of bulge light.  A much larger sample of galaxies is needed, but it seems we have established a path towards obtaining it.  Employing the Keck laser in our NIRC2 AO observations has resolved the PSF issues that troubled our NGS work and made it particularly difficult to measure and analyze galaxy subcomponent colors.  We have demonstrated a method of dithering the laser guide star that stabilizes the PSF across the field of view during science observations.  Furthermore, our method of interleaving science and PSF calibration frames, and of imaging a crowded stellar field for further calibration, allows us to much better characterize the PSF in our LGS science images than in the NGS images we analyzed in our previous work.  This is important because LGS AO is quickly becoming a standard feature of large telescopes, opening up the sky to deep NIR imaging surveys in archival HST faint galaxy fields. This is the path that the CfAO Treasury Survey is pursuing."
    },
    "0807/0807.3577_arXiv.txt": {
        "abstract": "In this paper we consider the relation between the volume deceleration parameter obtained within the Buchert averaging scheme and the deceleration parameter derived from the supernova observation. This work was motivated by recent findings that showed that there are models which despite $\\Lambda=0$ have volume deceleration parameter $q^{vol} < 0$. This opens the possibility that backreaction and averaging effects may be used as an interesting alternative explanation to the dark energy phenomenon.  We have calculated $q^{vol}$ in some Lema\\^itre--Tolman models. For those models which are chosen to be realistic and which fit the supernova data, we find that $q^{vol} > 0$, while those models which we have been able to find which exhibit $q^{vol} < 0$ turn out to be unrealistic. This indicates that care must be exercised in relating the deceleration parameter to observations. ",
        "introduction": "Accelerated expansion, modeled by a positive cosmological constant, is an essential element of the current standard cosmological model of the Universe. The accelerated expansion was originally motivated by supernova observations \\cite{sn} and is supported by many other types of cosmological observations. Observational data is, in modern cosmology, analyzed almost exclusively within the framework of homogeneous and isotropic Friedmann models \\cite{grs}. This analysis leads to the Concordance model, which provides a remarkably precise fit to cosmological observations. In this situation, if the Ehlers-Geren-Sachs theorem \\cite{EGS} and `almost EGS theorem' \\cite{SME} are invoked\\footnote{These theorems imply that if anisotropies in the cosmic microwave background radiation are small for all fundamental observers then the Universe is locally almost spatially homogeneous and isotropic. However, as shown in \\cite{NUWL99} the almost Robertson--Walker geometry also requires the smallness of the Weyl curvature.}, then it seems that an assumption of large scale homogeneity of the Universe can be justified. This on the other hand implies that the Universe must be filled with dark energy which currently drives the acceleration of the Universe.  However the Concordance model is not the only one which can fit cosmological observations. Anti-Copernican inhomogeneous models which assume the existence of a local Gpc scale void also fit cosmological observations \\cite{ide} (see \\cite{C07} for a review). Moreover, on small and medium scales our Universe is not homogeneous. Therefore, one may ask whether Friedmann models can describe our Universe correctly. In particular, it is important to ask what is the best way to fit a homogeneous model to a realistic and inhomogeneous Universe. This problem, known as the fitting problem, was considered by Ellis and Stoeger \\cite{ES87}. In considering the fitting problem, it becomes apparent that a homogeneous model fitted to inhomogeneous data can evolve quite differently from  the real Universe. The difference between evolution of homogeneous models and an inhomogeneous Universe is caused by backreaction effects, due to the nonlinearity of the Einstein equation. Unfortunately, in the standard approach, the backreaction is rarely taken into account -- in most cases when modelling our Universe on a local scale Newtonian mechanics is employed and on large scales the Friedmann equations (or linear perturbations of Friedmann background) are used \\cite{P93}. Such an approach to cosmology is often encouraged by the ``no--go'' theorem which states that the Universe can be very accurately described by the conformal Newtonian metric perturbed about a spatially flat background, even if $\\delta \\rho / \\rho \\gg 0$. In such a case the backreaction is negligible \\cite{IW06,KAF06}. However, the results obtained by van Elst and Ellis \\cite{EE98} and recently by Kolb, Marra and Matarrese \\cite{KMM08} show that the application of ``no-go'' theorem is limited. Therefore, one should be aware that in the absence of an analysis of the backreaction and other effects caused by inhomogeneities in the universe, there remains the possibility that the observed accelerated expansion of the Universe is only apparent \\cite{El08}. The direct study of the dynamical effects of inhomogeneities is difficult. Due to the nonlinearity of the Einstein equations, the solution of the Einstein equations for the homogeneous matter distribution leads in principle to a different description of the Universe than an average of a inhomogeneous solution to the exact Einstein equations (even though inhomogeneities when averaged over a sufficiently large scale might tend to be zero).  Neither the analysis of the evolution of a general matter distribution nor the numerical evolution of cosmological models employing the full Einstein equations are available at the level of detail which would make them useful for this problem. There are currently several different approaches which attempt to take backreaction effects into account. One approach is based on exact solutions -- see for example \\cite{ha}. Another, and more popular approach is based on averaging.  In the averaging approach to backreaction, one considers a solution to the Einstein equations for a general matter distribution and then an average of various observable quantities is taken. If a simple volume average is considered then such an attempt leads to the Buchert equations \\cite{B00}. The Buchert equations are very similar to the Friedmann equations except for the backreaction term which is in general nonvanishing, if inhomogeneities are present. For a review on backreaction and the Buchert averaging scheme the reader is referred to \\cite{SR2006,B08}. Within this framework and using spherically symmetric inhomogeneous models Nambu and Tanimoto \\cite{NT}, Paranjape and Singh \\cite{PS}, Kai, Kozaki, Nakao, Nambu, and Yoo \\cite{KKNNY}, Chuang, Gu, and Hwang \\cite{CGH}, provided explicit examples that one can obtain negative values of the volume deceleration parameter even if $\\Lambda = 0$. Another interesting example was presented by R\\\"as\\\"anen \\cite{SR2006,R06b} where it was shown that the total volume deceleration parameter of two isolated and locally decelerating regions can also be negative.   There are however important ambiguities in the application of an averaging procedure. The average itself not only depends on a choice of volume but also on a choice of time slicing. This is very crucial in cosmology. Once inhomogeneities are present the age of the Universe is not everywhere the same. Namely, the big bang in inhomogeneous models is not a single event, so the average taken over a hypersurface of constant cosmic time $t$ is different from the average taken over a hypersurface of constant age of the Universe $t - t_B$ \\cite{SR2004b}. Moreover, the results of the averaging procedure vary if the discrepancy between the average cosmic time and the local time is introduced (the local time is the time which is measured by local clocks; the cosmic time is the time which appears in the averaged homogeneous model). This phenomenon was studied by Wiltshire \\cite{Wa}, and has been used in an ambitious alternative concordance model. The model proposed by Wiltshire introduces some additional assumptions which allow to some extent a comparison of averaged quantities with observations. Such a comparison shows quite good agreement with observations, \\cite{LNW08}. Thus, while serious fundamental questions remain concerning Wiltshire's approach, it is another example of an approach where one does not need  dark energy to fit cosmological observations.  The averaging procedure is also gauge-dependent. For example using different gauge one can obtain that the backreaction mimics not dark energy but dark matter \\cite{KKM}. The averaging schemes, therefore, in the literature have been criticized, and their inherent ambiguities (and in some cases obscurity) have been discussed, cf. e.g. \\cite{IW06}.  A key point is that it is far from obvious if the average quantities, such as the acceleration of the averaged universe are really the quantities which are measured in astronomical observations. In particular, an operational analysis is to a large extent lacking in the discussions of averaging. Thus, it is important to test the averaging procedures with the exact and inhomogeneous solutions of the Einstein equations. Within exact models each quantity can easily be calculated and then compared with its averaged counterpart.  This paper aims to perform such an analysis within the Lema\\^itre--Tolman model.   The structure of this paper is as follows. Buchert's averaging procedure is presented in section \\ref{BSsec}, and some background on the Lema\\^itre--Tolman model is given in section \\ref{LTsec}. The volume and distance deceleration parameters are introduced in section \\ref{DPsec}. Finally, in section  \\ref{QAsec}, we discuss the relation between the deceleration parameters, supernova observations and models of cosmic structures. ",
        "conclusions": "In this paper we have studied the relation between the volume deceleration parameter obtained within the Buchert averaging scheme and the deceleration parameter derived from the observations of supernovae. This work was motivated by recent results showing there there are models which despite $\\Lambda=0$ and average expansion rate is accelerating, i.e. $\\ddot{a} > 0$ [where $a$ is defined by relation (\\ref{aave})]. This opens the possibility that backreaction and averaging effects may be used as an interesting alternative explanation to the dark energy phenomenon.  We have compared the quantities obtained within the exact and inhomogeneous models with their average counterparts. We focused on the supernova observations and models of cosmic structures. For this purpose the Lema\\^itre--Tolman model was employed. It was showed numerically that the averaging of models which fit the supernova observations does not lead to volume acceleration ($\\ddot{a} < 0$ for these averaged models and hence $q^{vol} >0$). It was also shown that realistically evolving models of cosmic structures have also $q^{vol} >0$. It was possible to modify these model in such a way that after the averaging $q^{vol} <0$. This was obtained by choosing $E$ function of positive amplitude - as was recently proved by Sussman \\cite{S08} this is a necessary condition to obtain $q^{vol} <0$. However, in models with realistic density distribution, in such a cases, $E \\gg 1 \\gg M/R \\approx 10^{-7} - 10^{-6}$, hence as seen from (\\ref{evo}) $t_B \\approx t$ (to remind $c \\times 10^{10}$y $\\approx 3$ Gpc). Thus, within such models the age of the Universe is unrealistically small.  Our analysis has been performed in the limited class of Lema\\^itre--Tolman models, which due to their spherical symmetry are arguably too simple to give a full understanding of averaging and backreaction problems. However, within this class, we conclude that the volume  deceleration parameter $q^{vol}$ is not a quantity which can be directly related to observations.  It is possible that the volume deceleration parameter $q^{vol}$ becomes negative only after averaging over the scales which are larger than 100 Mpc. On such large scales the structure of the Universe becomes too complicated to be fully described by spherically symmetric models. However, it is intriguing that models which fit the supernova observations and for which the distance deceleration parameter, $q^{dis}$, is negative have still $q^{vol}>0$. This suggest that the volume  deceleration $q^{vol}$ does not have a clear interpretation in terms of observable quantities. It does not, of course, mean that averaging and backreaction effects cannot potentially be employed to explain the phenomenon of dark energy. However, our work here indicates that such a potential solution of the dark energy problem should be based upon different methods than those related to volume deceleration parameter. Rather than showing that $q^{vol}<0$ the averaging approach should explain observations -- reproduce correct values of distance to supernovae, correct shape of the CMB power spectrum, etc. An interesting,  quasi-Friedmannian approach, was recently suggested in \\cite{LABKC08}. In this approach backreaction is modeled in terms of the morphon field \\cite{BLA06}. In such a case a Universe is describe by a homogeneous model with the spatial curvature being just a function of time. As shown in \\cite{LABKC08} such approach lead to an agreement with supernova and CMB data without the need for dark energy, but requires $q^{vol}<0$.   \\ack  We would like to thank Henk van Elst and Thomas Buchert for useful discussions and comments. KB would like to thank Peter and Patricia Gruber and the International Astronomical Union for the PPGF Fellowship, the support of the Polish Astroparticle Network (621/E-78/SN-0068/2007) is also acknowledged. KB is also grateful to LA and the Albert Einstein Institute, where part of this research was carried out, for their hospitality. LA acknowledges the support of the NSF with grants DMS-0407732 and DMS-0707306 to the University of Miami.  \\appendix"
    },
    "0807/0807.4161_arXiv.txt": {
        "abstract": "We consider a troublesome form of non-isoplanatism in synthesis radio telescopes: non-coplanar baselines. We present a novel interpretation of the non-coplanar baselines effect as being due to differential Fresnel diffraction in the neighborhood of the array antennas.  We have developed a new algorithm to deal with this effect. Our new algorithm, which we call ``W-projection'', has markedly superior performance compared to existing algorithms. At roughly equivalent levels of accuracy, W-projection can be up to an order of magnitude faster than the corresponding facet-based algorithms.  Furthermore, the precision of result is not tightly coupled to computing time.  W-projection has important consequences for the design and operation of the new generation of radio telescopes operating at centimeter and longer wavelengths. ",
        "introduction": "Wide-field imaging with synthesis radio telescopes can be limited by a number of effects:  \\begin{itemize} \\item The intrinsic performance of the deconvolution algorithms. \\item Time and frequency averaging. \\item The finite size of the primary beam. \\item The individual, angular, frequency, polarization, and temporal variations in the antenna primary beams. \\item Non-isoplanatic atmospheres. \\item The non-coplanar baselines effect. \\end{itemize}  Effective and reasonably efficient algorithms exist to correct for many of these effects \\citep[see][Chapter 19]{Tayloretal1999}.  Early interferometric arrays were often designed to be aligned East-West, in which case the baselines are coplanar and no error arises (for an excellent review of the development of radio interferometry see chapter 1 in the standard text by Thompson, Moran, and Swenson \\citep[][]{Thompsonetal2001}). Early {\\em non-coplanar} arrays had limited sensitivity and so only the brightest sources at the centre of a small field of view were of interest. As the sensitivity improved over the years with lower system temparatures, it became necessary to find and remove the effects of the other bright sources in the antenna primary beam.  Only as non-coplanar arrays become sufficiently sensitive did the need for a full field of view correction algorithm become apparent.  For the Very Large Array (VLA), this transition occured with the commissioning of the 327MHz observing system in the early nineties. At that point, Cornwell and Perley developed a faceted algorithm \\citep{CornwellPerley1992} which has been used in one form or another since then. Reaching the thermal noise sensitivity limit on the VLA required using the faceted algorithm to correct for the sidelobes from the myriad background sources in any field. However, the faceted algorithm is typically 100 to 1000 times slower than the simple two dimensional inversion, and so the need for a faster algorithm remained.  In this paper, we briefly re-examine various algorithms for dealing with the non-coplanar baselines effect before presenting our new algorithm, its implementation, performance and implications for the new generation of radio telescopes. ",
        "conclusions": "We have demonstrated a new algorithm for correcting the non-coplanar baselines effect. This has superior performance in both speed and error control, at the cost of greater memory usage. These advantages grow more substantial as the non-coplanarity grows more severe. It does not make it necessary to use only pixel based conventional CLEAN deconvolution step and thus is well-suited to wide-field imaging of very extended emission (for example, the Galactic plane) where multi-scale or Maximum Entropy methods are required.  The facets algorithms remain roughly competitive with W-projection only at low dynamic range, for example that occuring in VLA observations at 74MHz.  At high sensitivities, W-projection will be decisively superior.  Hence use of this algorithm (and similar algorithms for other problems) will have a substantial impact on the predicted computing costs for new radio synthesis arrays such as the EVLA, LOFAR, and SKA."
    },
    "0807/0807.4816_arXiv.txt": {
        "abstract": "{In the standard formulation of the disc instability model for soft X-ray transient and dwarf nova outbursts, the mass transfer rate from the secondary is assumed to be constant. This may seem natural since the $L_1$ point remains shielded from the accretion luminosity by the accretion disc. However an indirect heating could take place and could lead to an enhancement of the mass transfer rate. It is still debated whether such an enhancement of the mass transfer rate during outbursts is a missing ingredient of the current model.} {We discuss two mechanisms that could result in an enhancement of the mass transfer rate during outbursts: the hot outer disc rim itself could significantly heat the $L_1$ point; scattered radiation by optically thin outflowing matter could also heat $L_1$ significantly. We determine quantitatively the increase of the mass transfer rate resulting from an extra heating.} {We estimate the heating of the $L_1$ point by the disc rim by analytical and numerical arguments. The fraction of the luminosity scattered by matter above the disc, here modeled as a spherically symmetric outflow, is determined analytically. We finally solve in a numerical way the vertical structure equations for the secondary star and calculate the mass transfer enhancement.} {During outbursts, the temperature at the outer edge of the disc reaches $10^4$ K. The disc edge heats up the upper layer of the secondary with a flux of the order of the intrinsic stellar flux. This probably has no large effect on the mass transfer rate. In soft X-ray transients, the environing medium of the disc (corona+wind) could back-scatter a certain fraction of the accretion luminosity toward $L_1$. In dwarf novae, the same effect could be due to the wind present during outburst. Since soft X-ray transients reach high luminosities, even a low efficiency of this effect could yield a significant heating of $L_1$, whereas we show that in dwarf novae this effect is negligible. Initially the incoming radiation does not penetrate below the photosphere of the secondary. Depending on the heating efficiency, which has to be determined, the mass transfer rate could be significantly increased.} {} ",
        "introduction": "Dwarf novae (DNe) and Soft X-ray transients (SXTs) are semi-detached binary systems that undergo regular outbursts. SXTs in outbursts reach luminosities of order of $10^{38}$ erg.s$^{-1}$ in the soft X-ray band, whereas they are faint X-ray sources during quiescence (with luminosities $\\lesssim 10^{33}$ erg.s$^{-1}$). Their outbursts last for a few months and have a recurrence time in the range 1 -- 50 years \\citep[see][]{1997xrb..book.....L}. DNe are $1-5$ magnitude brighter during outbursts than in quiescence. Their outbursts last for a few days and recur every few weeks \\citep[see][]{2003cvs..book.....W}. In both types of systems, a compact object (i.e. a white dwarf in DNe and a neutrons star/black hole in SXTs) accretes from a low mass secondary star via an accretion disc. It is now believed that DNe and SXTs outbursts are both due to a thermal/viscous instability of the accretion disc triggered when hydrogen becomes partially ionized \\citep[see][]{1996PASP..108...39O,2001NewAR..45..449L}. In this picture, the accretion disc performs a limit cycle between a cold, quiescent state of low accretion rate and a hot, viscous state of high accretion rate corresponding to the outburst. The difference between DNe and SXTs timescales is thought to be due both to the irradiation of the disc by the central source in SXTs \\citep[see][]{1998MNRAS.293L..42K,2001A&A...373..251D}, and to a higher primary mass.  \\begin{table*}[t] \\caption{Physical parameters for the four models condidered in this paper. $P_{\\rm hr}$ is the orbital period in hours, $a$ the orbital separation, $R_2$ the secondary radius, $M_1$ and $M_2$ the primary and secondary masses, $T_\\star$ the effective temperature of the secondary, $\\dot{M}_{\\rm tr}$ the average mass transfer rate from th esecondary, and $\\dot{M}_{\\rm acc,max}$ is the maximum accretion rate onto the compact object. $T_\\star$ for model $1-3$ is taken from \\cite{2004AcA....54..181S} and from \\cite{1986A&A...162...71H} for model $4$.}            % \\label{params_phys}      % \\centering % \\begin{tabular}{c c c c c c c c c c c}        % \\hline % Model & System & $\\mathrm{P_{orb}}$ (h) & $a\\ (R_\\odot)$  & $R_2/a$& $q=M_2/M_1$ & $M_2(M_\\odot)$ & $T_\\star$ (K) & $\\dot M_\\mathrm{tr}$ (10$^{16}$ g.s$^{-1}$) & $\\dot M_\\mathrm{acc,max}$ (10$^{16}$ g.s$^{-1}$) & $h/r$\\\\  % \\hline % 1 & OY Car & 1.51 & 0.6 & 0.21 & 0.1 & 0.085 &2500 & 0.5 & 16&0.06\\\\ 2 & U Gem & 4.24 & 1.48 & 0.29 &0.36 &  0.7& 3500 & 5,10 & 100&0.09\\\\ 3 & Z Cam  & 6.96 & 2.17 & 0.35 &0.6 &0.7 &4200& 30,60  & 250&0.1\\\\ \\hline 4 & A 0620  & 7.8 & 3.9 & 0.21 & 0.1 & 0.7 & 3000 & 0.5 & $10^3$&0.1\\\\ \\hline % \\end{tabular} \\end{table*}  The thermal/viscous instability alone is not sufficient to explain the rich and diverse features of DNs and SXTs lightcurves. For example, \\citet{1989PASJ...41.1005O} suggested that a tidal instability driving the disc eccentric \\citep[see][]{1988MNRAS.232...35W} would account for the superoutburst phenomenon of SU UMa stars (a subclass of DNe) when coupled to the standard thermal/viscous instability. Tidal effects have also been proposed to account for the secondary maximum often present in SXTs lightcurves \\citep[see][]{2002MNRAS.337.1329T}.  In its most common formulation, the disc instability model (DIM) assumes that the mass transfer rate from the secondary is fixed, and given by the secular losses of angular momentum from the binary system. This assumption is most likely to be wrong and significant deviation of the mass transfer rate from its secular mean on time scales as short as the outburst duration could be an essential ingredient missing in the standard model. Numerical investigations showed that variations of the mass transfer rate could explain the outburst bimodality of DNe \\citep{1999AcA....49..383S} and that it could lead to long outbursts similar to superoutbursts \\citep{2000NewAR..44...15H}. \\cite{1993A&A...279L..13A} and \\cite{1993ApJ...408L...5C} sought to explain the rebrightenings of SXTs lightcurves by an increase of the mass transfer rate. However, in these works either an empirical relationship between the mass transfer rate and the mass accretion rate was assumed or mass transfer burst episodes were imposed without physical grounds. It has been often argued that a such mass transfer enhancement could result from irradiation of the secondary during an outburst. \\cite{2007A&A...475..597V} investigated this problem and found that, because the $L_1$ point is shielded from direct irradiation disc by the secondary star, the increase of the mass transfer rate is modest.  In this paper, we investigate two different mechanisms that could be responsible for a mass transfer enhancement, resulting in both cases from an increase of the temperature at the $L_1$ point. Since $L_1$ lies in the shadow cast by the accretion disc, any heating must be indirect. The first effect that we consider is heating by radiation emitted by the hot edge of the disc during an outburst, which produces a flux at $L_1$ comparable to the intrinsic stellar flux. The second effect that we consider is the heating of $L_1$ by the accretion luminosity scattered by a wind and/or a corona extending above the accretion disc. The scattered fraction is small, but the scattered flux is still significant in SXTs due to their high luminosity in outburst. We then turn to the determination of the mass transfer enhancement. First, we compute the vertical structure of the secondary envelope in DNe and SXTs both in quiescence and outbursts to determine the dependence of the mass transfer rate versus a given incoming heating flux. Our results show that $\\dot{M}$ can potentially be raised by a factor of up to $\\sim 100$ in SXTs, and up to $\\sim 10$ in DNe. We then discuss if the effects investigated here could lead to an efficient heating of the $L_1$ point. This is done by computing the fraction of the incoming radiation that penetrates the photosphere. This depends on the nature of the radiation, hence on its origin. We show that the thermal flux emitted by the edge of the disc can penetrate below the photosphere of the secondary. The resulting heating is probably too small to have any major effect. For soft X-ray transients, the backscattered radiation is in the soft X-ray band and is therefore strongly absorbed by the initially neutral hydrogen in the very upper part of the atmosphere. However, this large X-ray flux significantly affects the atmosphere. A full analysis of the radiative transfer should be undertaken to determine the exact structure of the atmosphere.   The paper is organized as follow: in Section 2, we discuss the two mechanisms mentioned above. Section 3 is devoted to the computation of the vertical structure of secondary stars in DNe and SXTs, and to the relationship between the mass transfer enhancement and the heating of $L_1$. Concluding remarks are given in section 4. ",
        "conclusions": "We have considered two mechanisms that could be responsible for heating of the $L_1$ point during outbursts of dwarf novae and soft X-ray transients. First, the outer edge of the disc is heated by tidal dissipation and reaches an effective temperature of the order of $10^4$ K (see Sect. 1). We have shown in Sect. 3 that a significant part of the resulting thermal radiation can penetrate below the photosphere of the secondary. The disc edge is able to heat the upper layer of the secondary with an incoming flux comparable to the intrinsic stellar flux. This has probably no significant effect in dwarf novae. In soft X-ray transients, the dominant effect could be a heating by scattered radiation. A disc corona, or a wind driven by the X-ray heating of the disc, could act as a scattering medium, allowing the radiation to overcome the screening by the accretion disc. It is hard however to infer a quantitative estimate of this effect. However, the luminosity of SXTs are so high that even a low efficiency ($10^{-3}-10^{-4}$) could lead to a significant heating. However, the bulk of this radiation is absorbed by neutral hydrogen, and initially nothing reaches the photospheric level of the secondary star. Nevertheless, the flux is so strong that the incident radiation probably strongly modifies the state of the atmosphere. Yet it is hard to assess whether in this case an efficient heating takes place or not. The resolution of the corresponding radiative transfer should be undertaken.  In Sect. 3 we determined the relation between the mass transfer enhancement and a given heating flux impinging the $L_1$ point (independently of the origin of this heating). We computed the vertical structure of the secondary with the assumption of hydrostatic equilibrium. The atmosphere is not hydrostatic near the $L_1$ point due to the mass outflow in this region. However, departure from hydrostatic equilibrium is likely to be important only in the near vicinity of the $L_1$ point. We found that the $L_1$ point lies higher in the atmosphere in soft X-ray transients than in dwarf novae. This is a consequence of a lower value of the mass transfer rate in soft X-ray transients, where the $L_1$ point lies in a nearly isothermal, radiative region of the atmosphere, whereas in dwarf novae the $L_1$ point lies in the convective part of the atmosphere, where temperature increases more rapidly with depth. As a consequence, an incident flux can heat a significant layer in soft X-ray transients, which yields a large enhancement of the mass transfer rate. On the other hand, in dwarf novae only a narrower region of the atmosphere is affected. For a dwarf nova below the period gap, the envelope is too massive to be fully affected on the time scale of an outburst. Our results suggest that the mass transfer enhancement could be potentially very important for soft X-ray transients, up to a factor of $\\sim 100$ for a heating flux equal to $10 F_\\star$. For dwarf novae, the mass transfer rate enhancement is more moderate, of a factor of $\\sim 10$ for a heating flux equal to $10 F_\\star$. In dwarf novae, none of the effects investigated here is able to produce such fluxes. The soft X-ray transient case needs further investigation, on one hand to determine the fraction of radiation that is back-scattered toward $L_1$, and on the other hand the efficiency of the heating of the secondary has to be determined properly."
    },
    "0807/0807.1328_arXiv.txt": {
        "abstract": "The majority of gamma-ray emission from Galactic dark matter annihilation is likely to be detected as a contribution to the diffuse gamma-ray background.  I show that dark matter substructure in the halo of the Galaxy induces characteristic anisotropies in the diffuse background that could be used to determine the small-scale dark matter distribution.  I calculate the angular power spectrum of the emission from dark matter substructure for several models of the subhalo population, and show that features in the power spectrum can be used to infer the presence of substructure.  The shape of the power spectrum is largely unaffected by the subhalo radial distribution and mass function, and for many scenarios I find that a measurement of the angular power spectrum by Fermi will be able to constrain the abundance of substructure.  An anti-biased subhalo radial distribution is shown to produce emission that differs significantly in intensity and large-scale angular dependence from that of a subhalo distribution which traces the smooth dark matter halo, potentially impacting the detectability of the dark matter signal for a variety of targets and methods. ",
        "introduction": "\\label{sec:intro}  A wealth of evidence suggests that most of the matter in the universe is in the form of non-baryonic dark matter \\citep[see][for reviews]{bergstrom_00, bertone_hooper_silk_05}, but the fundamental nature of this constituent remains unknown. Theoretical work has produced several candidates, of which many current favorites fall under the category of weakly interacting massive particles (WIMPs).  These include the lightest supersymmetric particle (LSP, often termed the neutralino) in supersymmetric extensions to the standard model, and the lightest Kaluza-Klein particle (LKP) arising in theories of universal extra dimensions.  Both the neutralino and the LKP are stable, but can produce photons and other standard model particles through self-annihilation.  Dark matter could thus be detected indirectly by observing these annihilation products.  The prospects for detecting dark matter via gamma-rays from annihilation have been studied extensively.  The Galactic Center \\cite{berezinsky_bottino_mignola_94, dodelson_hooper_serpico_08, serpico_zaharijas_08}, dwarf galaxies \\cite{baltz_briot_salati_etal_00, strigari_koushiappas_bullock_etal_08}, dark matter subhalos \\cite{baltz_taylor_wai_07,  pieri_bertone_branchini_08, kuhlen_diemand_madau_08}, and intermediate mass black holes \\cite{bertone_zentner_silk_05}, along with Galactic \\cite{bergstrom_edsjo_gondolo_etal_99, calcaneo-roldan_moore_00, berezinsky_dokuchaev_eroshenko_03} and extragalactic \\cite{bergstrom_edsjo_ullio_01, ullio_bergstrom_edsjo_etal_02, elsasser_mannheim_05, oda_totani_nagashima_05} diffuse emission, have all been considered as targets for indirect searches. However, detection methods which rely in whole or in part on the amplitude of the signal are hindered by uncertainties in the properties of the dark matter particle and by limited knowledge of the distribution of dark matter, particularly its clustering properties on small scales.  Even methods which make use of spectral information to extract the signal must contend with the presence of uncertain astrophysical foregrounds.  Recent work has proposed that angular features in the diffuse extragalactic gamma-ray background (EGRB) could be used to identify emission from dark matter.  An initial calculation of the angular power spectrum of the EGRB from dark matter annihilation in large-scale structures was made by \\citet{ando_komatsu_06}, and subsequent studies have further investigated this approach \\cite{ando_komatsu_narumoto_etal_07a, cuoco_hannestad_haugbolle_etal_07, cuoco_brandbyge_hannestad_etal_07}.  Estimates of the angular power spectrum of the EGRB from source classes other than dark matter have also been derived \\cite{zhang_beacom_04, miniati_koushiappas_di-matteo_07, ando_komatsu_narumoto_etal_07b}, providing a basis for comparison. While the large-scale distribution of dark matter is fairly well understood, the structure of dark matter halos on subgalactic scales, which is the focus of this study, is only minimally constrained by current observations.  The intrinsic properties of the dark matter particle influence the formation and survival of substructure within galaxy-sized halos, and thus the small-scale distribution of dark matter remains an important test of its fundamental nature.  For a generic WIMP cold dark matter (CDM) candidate, free-streaming and collisional damping induce a lower bound on the minimum mass of a CDM structure at around $10^{-6}$ to $10^{-4}$ M$_{\\odot}$ \\cite{green_hofmann_schwarz_05,loeb_zaldarriaga_05, bertschinger_06}, suggesting that the Galactic halo may be populated by an enormous number of dark matter clumps of roughly Earth-mass.  However, it is unclear whether the smallest halos are resilient enough to survive until the present day.  This issue has been investigated in numerical simulations  \\cite{diemand_moore_stadel_05, diemand_kuhlen_madau_06} and analytically \\cite{berezinsky_dokuchaev_eroshenko_03, berezinsky_dokuchaev_eroshenko_06, zhao_hooper_angus_etal_07, berezinsky_dokuchaev_eroshenko_08}, but has not yet been tested observationally.  In addition, numerical simulations predict far more subhalos in a galaxy-sized halo than the number of known luminous satellites of the Milky Way, spawning several proposed mechanisms for reducing small-scale structure.  For example, the abundance of substructure may be decreased in the CDM framework by a suppression of small-scale power in the primordial power spectrum \\cite{kamionkowski_liddle_00, zentner_bullock_03}.  Modifications of the properties of the dark matter particle can also alter predictions for structure on subgalactic scales.  Self-interacting dark matter (SIDM) \\cite{spergel_steinhardt_00, yoshida_springel_white_etal_00, dave_spergel_steinhardt_etal_01, colin_avila-reese_valenzuela_etal_02}, proposed to bring predictions for halo structure better in line with observations, relaxes the assumption that dark matter is collisionless by introducing self-interaction to standard CDM candidates via a non-negligible scattering cross-section.  While inheriting the successes of standard CDM on large scales, SIDM leads to less substructure in galaxy-sized halos, and halos with shallower inner density profiles.  Since SIDM candidates are variations of standard CDM candidates, they could similarly be detected by gamma-ray emission from annihilation. Warm dark matter (WDM) \\cite{hogan_dalcanton_00, colin_avila-reese_valenzuela_00, bode_ostriker_turok_01} and MeV-scale dark matter \\cite{hooper_kaplinghat_strigari_etal_07} have also been invoked to reduce structure on small-scales.  In these scenarios candidates are typically too light to produce gamma-ray emission, but may be detectable at lower energies.  The work presented here investigates the prospects for using anisotropies in the diffuse gamma-ray background to gain new insights into the Galactic dark matter distribution.  In contrast with previous studies of large-scale angular features from Galactic dark matter \\cite{dixon_hartmann_kolaczyk_etal_98, hooper_serpico_07, berezinsky_dokuchaev_eroshenko_07}, this approach focuses on fluctuations at small angular scales as an observational probe of Galactic dark matter substructure. Previous work has predicted that most subhalos will not be detectable individually by the Fermi Gamma-ray Space Telescope; this study suggests how Fermi could overcome this difficulty by using angular correlations in diffuse emission to statistically infer the presence of a subhalo population.  In many scenarios the annihilation signal from Galactic dark matter is expected to dominate over that from extragalactic dark matter \\cite{hooper_serpico_07}, and although emission from a Galactic dark matter component has by no means been indisputably found, possible indications have already been identified: \\citet{dixon_hartmann_kolaczyk_etal_98} presented evidence of a halo feature in the EGRET data, and \\citet{hooper_finkbeiner_dobler_07} proposed that the observed excess microwave emission in the inner Galaxy \\cite{finkbeiner_04, dobler_finkbeiner_08}, dubbed `the WMAP haze', is in fact synchrotron emission from electrons and positrons produced by dark matter annihilation. Moreover, even if the Galactic dark matter contribution at high latitudes is modest, it is necessarily a foreground to any measurement of the EGRB, and so an understanding of its angular structure will be essential for making use of existing predictions for the angular power spectrum of the EGRB\\@.  In this paper I calculate the angular power spectrum of gamma-ray emission from unresolved Galactic dark matter substructure under various assumptions about the properties of the subhalo population.  The terms `substructure' and `subhalos' are used here interchangeably to refer to dark matter clumps within a larger halo. I introduce the formulae for the gamma-ray emission from dark matter annihilation in \\S\\ref{sec:dmgamma}, and present the models used to describe the subhalo population in \\S\\ref{sec:dmmodel}. The angular distribution of the emission from substructure relative to the smooth halo emission is discussed in \\S\\ref{sec:maps}.    In \\S\\ref{sec:powerspect} I present the angular power spectrum of gamma-ray emission from substructure and evaluate the prospects for detecting this signal with Fermi.  I discuss the implications of my results for indirect detection of Galactic dark matter in \\S\\ref{sec:discussion}, and summarize my conclusions in \\S\\ref{sec:summary}. ",
        "conclusions": "\\label{sec:discussion}  A confident identification of dark matter as the source of an observed signal is one of the greatest challenges for indirect detection.  The angular power spectrum of the diffuse gamma-ray background is an unique diagnostic of the properties of the source population, and as such could strengthen other approaches by providing complementary information.  The angular power spectrum is determined solely by the angular distribution of the measured emission.  This study took advantage of this characteristic to explore the potential for constraining the small-scale dark matter distribution with a measurement of the angular power spectrum by Fermi.  Since expectations for the dark matter distribution on subgalactic scales depend on assumptions about the nature of dark matter and the adopted cosmology, predictions for the angular power spectrum of the emission are subject to considerable uncertainties. The models for the subhalo population considered here sample a few plausible scenarios, and while they certainly do not represent the full range of possibilities, they serve to illustrate that predictions for the small-scale distribution of dark matter could be tested by Fermi.  The mass functions, density profiles, and radial distributions used in this work are motivated by simulations in a $\\Lambda$CDM cosmology.  To approximate the subhalo population for other scenarios, such as SIDM or a suppression of power at small-scales in the primordial power spectrum, I employed a simple cut-off in the subhalo mass function below roughly the mass of a dwarf galaxy.  Although for these models the abundance of subhalos would be drastically reduced compared to the standard CDM case, a more sophisticated modification of the mass function would more accurately describe the subhalo population in these scenarios.  The sensitivity of the angular power spectrum to the assumed mass function should be studied in more detail as predictions from simulations which adopt alternate cosmologies are further refined.  My treatment did not account for the modification of the internal structure of the subhalos expected in SIDM models or from tidal stripping or stellar encounters.  For SIDM, the softening of the inner density profile may lead to a significant reduction in the total flux of a subhalo, but since since the angular size of the emission region of a subhalo is typically far smaller than the resolution limit of Fermi, I expect that its primary effect would be to reduce the overall intensity of the signal without altering the measured angular distribution of the emission significantly.  Similarly, dynamical processing of the subhalos is also likely to reduce the intensity of the emission slightly.  If subhalos are more efficiently disrupted in some regions of the Galaxy than others \\cite{berezinsky_dokuchaev_eroshenko_07}, the measured angular power spectrum could vary between different patches of the sky.  I considered two very different radial distributions in an attempt to bracket the most extreme cases, and found little difference in the angular power spectrum calculated for the entire sky in the relevant multipole range.  Although the anti-biased distribution does not show any large-scale dependence, for the unbiased distribution the angular power spectrum would likely differ somewhat across the sky due to the variation in subhalo number density. In addition, I assumed that the radial distribution of the subhalos is spherically symmetric about the Galactic Center, but asymmetric distributions are often found in simulations \\citep[e.g.,][]{zentner_kravtsov_gnedin_etal_05}, and also could lead to a directional dependence of the angular power spectrum.  More detailed models of the subhalo spatial distribution incorporating the effects of tidal disruption \\citep[e.g.,][]{berezinsky_dokuchaev_eroshenko_07} and correlations between halo concentration and radial distribution \\citep[e.g.,][]{diemand_kuhlen_madau_07b} have been developed.  These models may be able to more accurately predict the directional dependence of the power spectrum.  Uncertainties in the flux and spectrum of the gamma-ray emission from both dark matter and other source classes could make the overall amplitude of the substructure power spectrum difficult to determine, so initially the shape of the power spectrum may be a better indicator of the abundance of substructure than the measured amplitude.  I note, however, that the energy dependence and absolute amplitude of the power spectrum could provide useful information, particularly once emission from other source classes is sufficiently constrained.  This work reaffirms the findings of \\cite{kuhlen_diemand_madau_08} that the spatial distribution and intensity of the gamma-ray emission from Galactic dark matter annihilation are strongly affected by the radial distribution of the subhalo population.  An anti-biased subhalo distribution results in considerably less total emission than an unbiased distribution due to the larger typical distances of the subhalos, and in the case of a small minimum subhalo mass as generically predicted for CDM, the observed emission is almost entirely diffuse and devoid of large-scale features.  In addition to impacting the detectability of individual subhalos \\cite{kuhlen_diemand_madau_08}, an anti-biased radial distribution may alter the prospects for detecting the annihilation signal by several other methods. The prominence of the dipole from the Galactic dark matter halo  would be diminished relative to the unbiased distribution.   Predictions for scenarios in which the substructure mass function extends to very small masses would be most strongly impacted.  The additional flux from the outer regions of the halo would affect the angular dependence of the emission by increasing the mean observed intensity from dark matter annihilation in all directions. A related consequence concerns constraints on extragalactic dark matter emission derived from the detectability of the Galactic Center.  \\citet{ando_05} presented an argument using self-similarity of halo structure to constrain the fraction of the EGRB due to dark matter annihilation using the signal from the Galactic Center.  However, if a substantial fraction of the luminosity of a galaxy due to dark matter annihilation originates from the outer regions of the halo, as would be the case for an anti-biased subhalo distribution, the flux from the Galactic Center as measured from our position would provide only a weak constraint. A better understanding of these issues will be needed to clarify the interpretation of a Fermi detection (or non-detection) of predicted dark matter signals.      With the Fermi era just around the corner, this study and many others have focused on the exciting possibility that dark matter may soon be robustly detected for the first time via gamma-ray emission. In this work I considered the angular power spectrum of diffuse gamma-ray emission as a means of constraining the abundance and properties of dark matter substructure.  Using models motivated by numerical simulations to describe the Galactic dark matter distribution, I calculated the angular power spectrum from dark matter annihilation in substructure for several scenarios and evaluated the prospects for measuring this signal and distinguishing between dark matter models with observations by Fermi.  I showed that the presence of Galactic dark matter substructure  induces distinct features in the angular power spectrum of diffuse gamma-ray emission.  A cut-off in the substructure mass function below roughly the mass of known dwarf galaxies results in declining power ($C_{\\ell}$) in the multipole range $100 \\lesssim \\ell \\lesssim 1000$.  In contrast, if the substructure mass function extends down to the minimum masses predicted for CDM halos, the angular power spectrum coefficients $C_{\\ell}$ of the diffuse emission will be almost constant over this multipole range.  For several models, I found that a measurement of the angular power spectrum by Fermi will be able to differentiate between a halo populated by a multitude of subhalos and one in which substructure is scarce, as long as the contamination of the dark matter signal by the EGRB is not overwhelming.  The detectability of the signal is determined largely by the subhalo abundance and radial distribution, which can change the mean intensity of the emission from dark matter annihilation by several orders of magnitude.  The gamma-ray emission produced by an anti-biased subhalo distribution differs greatly in intensity and spatial distribution from that produced by an unbiased distribution.  An anti-biased distribution generates emission of similar intensity in all directions, with fluctuations confined to small angular scales, and does not enhance the large-scale dipole feature from the smooth halo component as an unbiased distribution does.  The detectability of dark matter annihilation by a variety of methods may be altered relative to the case of an unbiased distribution if the subhalo distribution is anti-biased.  \\ack I am greatly indebted to Angela Olinto for insights and guidance over the course of this work, and to Andrey Kravtsov for valuable advice and suggestions.  I also thank Maximo Ave, John Beacom, Brian Fields, Stephan Meyer, Vasiliki Pavlidou, Luis Reyes, Simon Swordy, and Carlos Wagner for helpful discussions.  This work was supported by NSF grant PHY-0457069 and by the Kavli Institute for Cosmological Physics at the University of Chicago through grants NSF PHY-0114422 and NSF PHY-0551142 and an endowment from the Kavli Foundation and its founder Fred Kavli.  \\appendix"
    },
    "0807/0807.2601_arXiv.txt": {
        "abstract": "The LHC will probe the nature of the vacuum that determines the properties of particles and the forces between them. Of particular importance is the fact that our current theories allow the Universe to be trapped in a metastable vacuum, which may decay in the distant future, changing the nature of matter. This could be the case in the Standard Model if the LHC finds the Higgs boson to be light. Supersymmetry is one favoured extension of the Standard Model which one might invoke to try to avoid such instability. However, many supersymmetric models are also condemned to vacuum decay for different reasons. The LHC will be able to distinguish between different supersymmetric models, thereby testing the stability of the vacuum, and foretelling  the fate of the Universe. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0807/0807.2437_arXiv.txt": {
        "abstract": "Star clusters that form in nuclear rings appear to be at slightly larger radii than the gas. We argue that the star clusters move out from the gas in which they are formed because of satellite-disk tidal interactions. In calculating the dynamics of this star cluster and gas ring system, we include the effects of dynamical friction of the background stars in the host galaxy on the star cluster, and inflowing gas along the bar onto the nuclear ring at the two contact points. We show that the final separation is of the order of the Hill radius of the nuclear ring, which is typically $20-30\\%$ of its radius. Massive star clusters can reach half of this separation very quickly and produce a factor of a few enhancement in the gas surface density. If this leads to star formation in addition to the (ongoing) formation of star clusters near the contact points, a possible (initial) azimuthal age gradient may become diluted or even disappear. Finally, if the star cluster are massive and/or numerous enough, we expect the nuclear ring to migrate inward, away from the (possibly) associated (inner) Lindblad resonance. We discuss how these predictions may be tested observationally. ",
        "introduction": "\\label{sec:intro}  Nuclear rings in barred spiral galaxies are regions of large gas surface densities and high star formation. Their existence is intimately related to the stellar bars with which they are associated. The stellar bar drives the interstellar medium (ISM) into the nuclear ring.  How the radius of the ring is determined is unclear. The radius may be set by the location of the inner Lindblad resonance (ILR) of the stellar bar \\citep[see review by][and references therein]{Buta1996}. The reasoning behind this idea is that the sign of the torque changes on either side of the Lindblad resonance. Alternatively, the radius may be at the location where the non-axisymmetric bar weakens, which is typically at $\\sim 10\\%$ of the size of the bar \\citep{Shlosman1990}. Finally, it may also be determined by where the transition from phase space domination of $X_1$ orbits (parallel to the major axis of the large-scale bar) to the phase space domination of $X_2$ orbits (perpendicular to the bar) occurs \\citep{Regan1997}.  Regardless of how they are formed they are interesting as a class because they are perhaps the largest population of nearby star bursting regions.  They form star clusters in prodigious amounts and are the only environment where super star clusters (SSCs) might be found in abundance in normal galaxies \\citep{Maoz2001}.  These SSCs can have masses in excess of $10^6$\\,\\Msun\\ and may be the progenitors of globular clusters if they can survive for $10^{10}$ years.  In some of these systems the location of these star clusters is curious. They appear to be at larger radii than the gas ring from which they are presumably formed. For instance, the morphology of NGC\\,1512 clearly show the star cluster(s) exterior (radially) to the gas as opposed to being embedded in the gas \\citep{Maoz2001}. Taking another example, in NGC\\,4314 the star clusters are at larger radii than the gas in the nuclear ring \\citep{Benedict2002}. Finally, \\cite{Martini2003a} find that in their sample of 123 galaxies, eight have strong nuclear rings and in each of the eight, star formation occurs outside of the dust ring. Two of their other galaxies have dust lanes immediately interior to what appears to be an older stellar ring.  There are two possible explanations for why these star clusters are external to the gas. First, while the gas can lose angular momentum due to dissipation, the stars and star clusters are collisionless and hence remain on larger orbits \\citep{Regan2003}. Second, the star clusters migrated outward after their formation. We explore the latter possibility in this paper.  We argue that tidal interactions between the star cluster and the nuclear ring from which it was formed leads to their separation. We apply the physics of ``gas shepherding'', which was studied by one of us \\citep[][hereafter C08]{Chang2008}, to this cluster-ring system. We derive a natural length scale for this separation, which we call the ``Hill radius'' of the ring. At the same time, we show that the outer edge of the ring can experience a surface density enhancement of a few, possibly leading to additional formation of stars and star clusters along the edge of the ring.  Finally, we find that if the star clusters are massive and/or numerous enough the cluster-ring systems as a whole is expected to migrate inward. We show that these predictions can be tested observationally.  We present our model in \\S~\\ref{sec:tidal}. After summarizing the observed properties of nuclear rings in \\S~\\ref{sec:properties}, we give the basic picture and introduce the Hill radius of the ring in \\S~\\ref{sec:basic_picture}, followed by the basic equations of our model in \\S~\\ref{sec:basic_equations}. The solutions for various scenarios are outlined in \\S~\\ref{sec:results}. A few observationally testable predictions that arises from our model are given and compared to observations in \\S~\\ref{sec:comparison}. We discuss some aspects of our model in \\S~\\ref{sec:discussion} and summarize our conclusions in \\S~\\ref{sec:conclusions}. ",
        "conclusions": "\\label{sec:conclusions}  We have studied the dynamics between a star cluster and the gas in the nuclear ring from which it was formed. The star cluster is subject to dynamical friction with the background bulge stars in the host galaxy. The resulting dynamical friction torque induces a tidal (and viscous) torque on the gas in the nuclear ring. In addition, we investigated a torque due to in-falling gas along the bar onto the nuclear ring at the contact points. We arrived at the following, observationally testable predictions based on our model.  \\begin{enumerate} \\item[(i)] The star clusters which are formed near the outer edge of the ring migrate outward. The final separation is of the order of the Hill radius of the ring, which is $20-30\\%$ of its radius if the ring's gas mass is the observed $1-3\\%$ of the enclosed mass. Similarly, the time to reach this final separation is about $20-30\\%$ of the dynamical friction time. The initial separation is very fast, reaching half the Hill radius within only $1\\%$ of the dynamical friction time. The latter can be as short as a few million years for a massive enough ($\\gtrsim 10^6$\\,\\Msun) star cluster. \\item[(ii)] While such a massive cluster moves out the surface mass density of the gas at the edge of the nuclear ring is pushed inward and gets enhanced by a factor of a few. If this enhancement reaches the critical value for star formation, new star clusters are expected to form along the ring. This is in addition to the (ongoing) formation of star cluster from shocking gas near the contact points.  The (initial) azimuthal age gradient that is expected from the latter gets diluted or even washed out by the former (secondary) star cluster formation. \\item[(iii)] If the (formed) star cluster are massive and/or numerous enough, and the gas mass inflow rate low enough (or even shut down), we expect the inward migration of the cluster-ring system as a whole.  This mean that, even though initially the nuclear ring could have formed as often suggested at an (inner) Lindblad resonance, it can migrate closer to the center of the host galaxy. \\end{enumerate}  The latter implies that the most massive star cluster, after an initial outward separation, turn around and migrate inward along with the nuclear ring. Consequently, any initial radial age gradient due to the dynamical friction time being inverse proportional to the mass of the star cluster, is diluted or even inverted. Even so, both a radial and azimuthal age gradient are observationally difficult to measure. When the most massive stars are still alive, the bright emission lines can be used for (relative) age dating \\citep[e.g.][]{Boker2008}, but after a few tens of million years it becomes already very challenging to distinguish the fading star clusters from the background stellar population \\citep[but see e.g.][]{Benedict2002}.  The outward migration of a least the most massive star clusters is easier to establish, and using the dust to trace the gas, observations indeed show the young star clusters exterior to the nuclear ring. The measured separation should be of the order of the Hill radius of the nuclear ring, which indeed is the case in NGC\\,4314. This well-studied nuclear ring also seems to have migrated inward, leaving a tail of older star clusters behind and being currently located between the inner and outer ILR. To firmly confirm such a ``shepherding of gas'' towards smaller radii, the dynamics of the host galaxy, including the (bar) pattern speed, needs to be measured accurately."
    },
    "0807/0807.0268_arXiv.txt": {
        "abstract": "We report on the detection of four extrasolar planets orbiting evolved intermediate-mass stars from a precise Doppler survey of G--K giants at Okayama Astrophysical Observatory. All of the host stars are considered to be formerly early F-type or A-type dwarfs when they were on the main sequence. 14 And (K0 III) is a clump giant with a mass of 2.2 $M_{\\odot}$ and has a planet of minimum mass $m_2\\sin i=4.8 M_{\\rm J}$ in a nearly circular orbit with a 186 day period. This is one of the innermost planets around evolved intermediate-mass stars and such planets have only been discovered in clump giants. 81 Cet (G5 III) is a clump giant with 2.4 $M_{\\odot}$ hosting a planet of $m_2\\sin i=5.3 M_{\\rm J}$ in a 953 day orbit with an eccentricity of $e=0.21$. 6 Lyn (K0 IV) is a less evolved subgiant with 1.7 $M_{\\odot}$ and has a planet of $m_2\\sin i=2.4 M_{\\rm J}$ in a 899 day orbit with $e=0.13$. HD 167042 (K1 IV) is also a less evolved star with 1.5 $M_{\\odot}$ hosting a planet of $m_2\\sin i=1.6 M_{\\rm J}$ in a 418 day orbit with $e=0.10$. This planet was independently announced by Johnson et al. (2008, \\apj, 675, 784). All of the host stars have solar or sub-solar metallicity, which supports the lack of metal-rich tendency in planet-harboring giants in contrast to the case of dwarfs. ",
        "introduction": "Precise Doppler surveys of evolved stars have opened a new frontier in extrasolar planet searches during the past several years. Since the successive discoveries of planets around K-type giants, $\\iota$ Dra (Frink et al. 2002) and HD 47536 (Setiawan et al. 2003), and a G-type giant, HD 104985 (Sato et al. 2003), about 20 substellar companions orbiting evolved stars have been identified so far (Setiawan 2003; Setiawan et al. 2005; Sato et al. 2007, 2008; Hatzes et al. 2003, 2005, 2006; Reffert et al. 2006; Johnson et al. 2007, 2008; Lovis \\& Mayor 2007; Niedzielski et al. 2007; D$\\ddot{\\rm{o}}$llinger et al. 2007; Liu et al. 2008). Planets in evolved stars is now one of the major subjects in the field of extrasolar planets.  In the past, planets around evolved stars were primarily studied from the interests in the fate of our solar system, that is, whether the Earth and the other planets will be engulfed by the Sun in the future (Sackmann et al. 1993; Duncan and Lissauer 1998). On the other hand, the current Doppler surveys of evolved stars have been mainly carried out in the context of planet searches around intermediate-mass (1.5--5 $M_{\\odot}$) stars. Intermediate-mass stars on the main sequence, that is early-type dwarfs (B--A dwarfs), are more difficult for Doppler planet searches because they have fewer absorption lines in their spectra than later type ones, which are often broadened due to their rapid rotation, and thus it is more difficult to achieve a high measurement precision in radial velocity that is enough to detect orbiting planets (but see eg. Galland et al. 2005, 2006, which have developed a technique to extract Doppler information from spectra of A dwarfs sufficient for detection of substellar companions). This is actually one of the major reasons why targets for planet searches had been limited to low-mass ($<1.5M_{\\odot}$) F--M dwarfs. On the contrary, late G to early K giants and subgiants, which are intermediate-mass stars in evolved stages, have many sharp absorption lines in their spectra appropriate for precise radial velocity measurements and their surface activity such as pulsation is quiet enough not to prevent us from detecting planets. Therefore, these types of stars have been identified as promising targets for Doppler planet searches around intermediate-mass stars.  The ongoing planet searches have already revealed that the properties of planets around evolved intermediate-mass stars are probably different from those around low-mass dwarfs (see Bulter et al. 2006 or Udry and Santos 2007, and references therein for properties of planets around low-mass stars): 1) high frequency of massive planets (Lovis \\& Mayor 2007; Johnson et al. 2007), which may be supported by the detection of a planet in an intermediate-mass giant in the Hyades open cluster (Sato et al. 2007) despite the absence of planets in low-mass hundred dwarfs in the cluster (Paulson et al. 2004), 2) lack of inner planets with orbital semimajor axes of $\\lesssim0.7$ AU (Johnson et al. 2007; Sato et al. 2008), and 3) lack of metal-rich tendency in the host stars of planets (Pasquini et al. 2007; Takeda et al. 2008). Although the properties must reflect history of formation and evolution of planetary systems dependent on their host stars' mass, they are still in need of confirmation by a larger number of samples. Combined with these observational progresses, planet formation theories applicable to more massive stars than the Sun have begun to develop rapidly (Ida and Lin 2005; Burkert and Ida 2007; Kennedy and Kenyon 2007) as extension of the standard theory for solar-mass ones (e.g., Hayashi et al. 1985; Ida and Lin 2004) for the first time since Nakano (1988) explored the idea nearly 20 years ago.  In this paper, we report on the detection of four planets orbiting intermediate-mass G--K giants and subgiants (14 And, 81 Cet, 6 Lyn, and HD 167042) from the Okayama Planet Search Program (Sato et al. 2005), one of which (HD 167042) was independently announced by Johnson et al. (2008, \\apj, 675, 784). We describe our observations in Section 2. Properties of the host stars, radial velocity data, and orbital parameters of their planets are presented in Section 3. Section 4 is devoted to an analysis of spectral line shape for the host stars. We summarize our results in Section 5 with a discussion about the properties of planets and host stars. ",
        "conclusions": "We here reported four new planets around evolved intermediate-mass stars from the Okayama Planet Search Program; two of them orbit clump giants and the other two orbit subgiants. In total, we discovered 7 planets and 1 brown dwarf around intermediate-mass clump giants (2.1--2.7$M_{\\odot}$) and 2 planets around subgiants (1.5 and 1.7$M_{\\odot}$) so far from the program. The host stars are considered to be formerly early F-type or A-type dwarfs when they were on the main sequence.  Like all of the other planets found around these types of stars, the four planets presented in this paper reside beyond $\\sim0.7$ AU from the central stars. Since many planets are known to exist within 0.7 AU around low-mass dwarfs, the lack of inner planets around evolved intermediate-mass stars must reflect different history of formation and evolution of planetary systems. Two scenarios can account for the orbital distribution; one is that inner planets are primordially deficient around intermediate-mass stars and the other is that they have been engulfed by the expanding central stars due to stellar evolution. Mass loss of the central star makes planets move outward because of their weakened gravitational pull on the planets, but orbital change due to mass loss is negligible in RGB phase for planets around intermediate-mass stars because the mass loss of those stars in RGB phase is negligible (see discussion in Sato et al. 2008). The lack of inner planets around less evolved subgiants (Johnson et al. 2007, 2008) may favor the former scenario. Since there is no observational bias against detecting planets with small semimajor axes around subgiants, whose radii are $\\sim5R_{\\odot}$ ($=0.025$ AU) at most and their intrinsic variability in radial velocity is adequately small ($\\lesssim10$ m s$^{-1}$), the lack of inner planets around them may be primordial. On the other hand, as for planets around intermediate-mass clump giants, evolutionary effect of the central stars can not be ignored. When we assume that many of clump giants are post-RGB stars (core-helium-burning stars), planets in orbits with $\\lesssim0.5$ AU could have been engulfed by the central stars at the tip of RGB ($R_{\\star}\\sim25$--$40R_{\\odot}$ for a 2--3 $M_{\\odot}$ star) due to tidal force from the central stars (Sato et al. 2008). Thus we can not reject the possibility that planets had originally existed in short orbital distances around progenitors of clump giants. It should be noted that all of the four innermost planets around evolved intermediate-mass stars with semimajor axes of 0.7--1 AU have been discovered around clump giants. Excluding the planets around clump giants, all of the six planets ever discovered around intermediate-mass subgiants reside beyond 1 AU, which expands the lack of inner planets around them. Although the number is still small, properties of planets might be different between the two populations when we consider that clump giants include $\\gtrsim2M_{\\odot}$ stars while subgiants are limited up to $\\lesssim2M_{\\odot}$. Larger number of planets, which will be provided by ongoing Doppler surveys of evolved stars among the globe, enables us to derive their statistical properties more clearly.  All of the host stars of the four planets presented here have solar or sub-solar metallicity. The results support the lack of metal-rich tendency in planet-harboring giants recently claimed by several authors (e.g., Pasquini et al. 2007; Takeda et al. 2008), which makes marked contrast to the case of dwarf stars, where planet-harboring stars tend to be generally metal-rich (see e.g., Gonzalez 2003 or Udry and Santos 2007, and references therein). Interpretation of the absence of metal-rich trend in planet-harboring giants still remains to be cleared. High efficiency in core-accretion in massive proto-planetary disks around massive stars combined with high efficiency in inward orbital migration of planets (finally falling into the central stars) in metal-rich disks might reproduce the lack of metal-rich trend. Alternatively, metallicity-independent planet formation scenario such as disk instability model (e.g., Boss 2002) might be possible. Pasquini et al. (2007) pointed out a possibility that metallicity excess in planet-harboring dwarfs originates from accretion of metal-rich material and the excess is diluted by deep convective envelope at the stage of giants. It should be noted, however, that there are no super-metal-rich stars with [Fe/H]$>+0.2$ in our sample regardless of the existence of orbiting planets (Takeda et al. 2008). Since the frequency distribution of planets steeply rises in [Fe/H]$>+0.2$ for solar-type dwarfs (Fischer \\& Valenti 2005), we might see metal-poor ([Fe/H]$<+0.2$) tail of the same distribution, where frequency of planets is less sensitive to metallicity, in the case of giants. Detailed investigation of metallicity distribution for a population of giants is required to derive a firm conclusion on planet-metallicity correlation in giants and its origin.  Host stars' metallicity could control the fate of orbiting planets as well as their birth. From the view point of the engulfment scenario, inner planets could be deficient around metal-rich clump giants of $\\sim2M_{\\odot}$ compared to around metal-poor ones. A metal-rich star with $\\sim2M_{\\odot}$ can have more than two times larger maximum radius ($>70R_{\\odot}$) than a metal-poor one at the tip of RGB (cf. evolutionary tracks by Girardi et al. 2000; Lejeune and Schaerer 2001; Claret 2004, 2006, 2007), which may increase a chance for the central star to engulf inner planets even up to $\\sim1$ AU. However, the current knowledge about evolutionary tracks of giants is not sufficient to estimate this effect precisely. Stellar evolutionary track at RGB phase for a $\\sim2M_{\\odot}$ star is sensitive to slight difference in stellar mass and metallicity because this mass range is a border where helium burning starts in a degenerated core or a non-degenerated one and thus the tracks are largely dependent on adopted models. To investigate orbital distribution of planets around post-RGB stars together with precise determination of stellar mass (by using technique of asteroseismology; e.g., Frandsen et al. 2002; Hatzes \\& Zechmeister 2007; Ando et al. 2008) and metallicity would conversely give constraints on stellar evolution for this problematic mass range from observational side.  As described above, the recent discoveries of planets around evolved stars have made stellar evolution and fate of planets a renewed area of study. This subject has been studied by several authors (Sackmann et al. 1993; Duncan and Lissauer 1998; Rasio et al. 1996; Siess and Livio 1999ab; Villanver and Livio 2007) mainly focusing on a solar-mass star based on the interests in the future of the solar system, and/or an asymptotic giant branch star with the scope of planets in planetary nebula and white dwarfs. Siess and Livio (1999b) proposed that the IR excess and high Li abundance observed in a few percent of G--K giants originate from the accretion of a substellar companion. Now is the time to promote such kinds of studies more extensively not only for solar-mass stars but also for intermediate-mass ones especially in RGB phase by both theoretical and observational approaches. \\\\  This research is based on data collected at Okayama Astrophysical Observatory (OAO), which is operated by National Astronomical Observatory of Japan (NAOJ). We are grateful to all the staff members of OAO for their support during the observations. Data analysis was in part carried out on ``sb'' computer system operated by the Astronomy Data Center (ADC) and Subaru Telescope of NAOJ. We thank the National Institute of Information and Communications Technology for their support on high-speed network connection for data transfer and analysis. B.S. is supported by MEXT's program \"Promotion of Environmental Improvement for Independence of Young Researchers\" under the Special Coordination Funds for Promoting Science and Technology. H.I. and M.Y. are supported by Grant-in-Aid for Scientific Research (C) No.13640247 and (B) No.18340055, respectively, from the Japan Society for the Promotion of Science (JSPS). This research has made use of the SIMBAD database, operated at CDS, Strasbourg, France."
    },
    "0807/0807.0083_arXiv.txt": {
        "abstract": "We investigate the influence of ram pressure on the star-formation rate and the distribution of gas and stellar matter in interacting model galaxies in clusters. To simulate the baryonic and non-baryonic components of interacting disc galaxies moving through a hot, thin medium we use a combined N-body/hydrodynamic code GADGET2 with a description for star formation based on density thresholds. Two identical model spiral galaxies on a collision trajectory with three different configurations were investigated in detail. In the first configuration the galaxies collide without the presence of an ambient medium, in the second configurations the ram pressure acts face on on the interacting galaxies and in the third configuration the ram pressure acts edge on. The ambient medium is thin ($10^{-28}$ g/cm$^3$), hot (3 keV $\\approx 3.6\\times10^7$K) and has a relative velocity of 1000 km/s, to mimic an average low ram pressure in the outskirts of galaxy clusters. The interaction velocities are comparable to galaxy interactions in groups, falling along filaments into galaxy clusters. The global star formation rate of the interacting system is enhanced in the presence of ram pressure by a factor of three in comparison to the same interaction without the presence of an ambient medium. The tidal tails and the gaseous bridge of the interacting system are almost completely destroyed by the ram pressure. The amount of gas in the wake of the interacting system is $\\sim50$\\% of the total gas of the colliding galaxies after 500 Myr the galaxies start to feel the ram pressure. Nearly $\\sim10-15\\%$ in mass of all newly formed stars are formed in the wake of the interacting system at distances larger than 20 kpc behind the stellar discs. As the tidal tails and the gaseous bridge between the interacting system feel the ram pressure, knots of cold gas ($T<1\\times10^5$K) start to form. These irregular structures contain several $10^6$ M$_{\\sun}$ of cold gas and newly formed stars and, as the ram pressure acts on them, they move far away (several 100 kpc) from the stellar discs. They can be classified as 'stripped baryonic dwarf' galaxies. These 'stripped baryonic dwarfs' are strongly affected by turbulence, e.g. Kelvin-Helmholtz instabilities, which are not resolvable within the presented SPH simulations. Heat conduction, which is not included, would affect these small structures as well. Therefore we give some estimate on the lifetime of these objects. ",
        "introduction": "Multiwavelength observations have shown that the star formation rate in interacting galaxies is enhanced in comparison to isolated galaxies (Bushouse 1987, Sulentic 1976, Stocke 1978, Solomon \\& Sage, 1988, Combes et al. 1994, Kewley et al. 2006, and references therein). To understand the physical processes involved in interacting galaxies numerical simulations are an ideal tool. Since the first publications in this field (Pfleiderer 1963, Toomre \\& Toomre 1972) it is evident that interactions are the sources of tidal tails, bridges and other signs of massive perturbations in and around the discs of galaxies. First calculations including star formation and gas depletion in interacting systems (e.g. Noguchi \\& Ishibashi 1986, Noguchi 1991, Olson \\& Kwan 1990 a,b, Mihos 1992, Mihos \\& Hernquist 1996) indicated that galaxy mergers are able to increase the total star formation of the system up to an order of magnitude and that these events of strong starbursts are able to deplete the cold gas reservoir of the system significantly. Many numerical investigations put special emphasis on modelling observed interacting systems, like NGC7252 (Mihos et al. 1998) or on the dependence of the star formation rates on interaction parameters like spatial alignment and minimum separation (di Matteo et al. 2007, Kapferer et al. 2005). Bournaud et al. (2005) investigated the remnants of galaxy mergers with different mass ratios; Mayer et al. (2006) showed via numerical simulations that dwarf galaxies lose their gas by ram pressure and tidal stripping during passages through a disc. Cox et al. (2004) investigated galaxy mergers with special emphasis on the heating process of gas due to shocks. The formation of dwarf galaxies in the debris of interacting galaxies was investigated by Duc et al. (2004). The accretion onto supermassive black holes in merging galaxies and the resulting suppression of star formation and the morphology of the elliptical remnants were investigated by Springel et al. (2005).\\\\ From observations of the Virgo cluster of galaxies it is reported that galaxies moving through a galaxy cluster suffer mass loss gas by the ram pressure of the intra-cluster medium (ICM) onto the inter-stellar medium (ISM) (Gunn \\& Gott 1972, Cayatte et al. 1990, Kenney et al. 2004, Vollmer et al. 2004). To study the effects of ram-pressure stripping in detail several numerical approaches were carried out, from very simple descriptions for the gas phase, like sticky particles (e.g. Vollmer et al. 2001), to advanced Eulerian grid techniques (e.g. Roediger et al. 2006) or Lagrangian descriptions of hydrodynamics (e.g. Schulz \\& Struck 2001, J{\\'a}chym et al. 2007), so called smoothed particle hydrodynamics (SPH). All these methods were applied to single model disc galaxies interacting with a time-dependent or a constant pressure from a hot gas phase outside the disc. Recently the investigation of the dependence of ram pressure on the star formation of a single model galaxy moving through an ambient medium was done by Kronberger et al. (2007).\\\\ On larger scales the effect of ram-pressure stripping on the chemical enrichment was studied in cosmological galaxy-cluster simulations. To model the physics involved below the resolution of such simulations an analytical approach for the mass loss by ram-pressure stripping was introduced, based on the Gunn \\& Gott criterion (Schindler et al. 2005, Domainko et al. 2006). In combination with the mass losses by galactic winds and starbursts, ram-pressure stripping is able to enrich the ICM to observed levels (Kapferer et al. 2007).\\\\ In this paper we concentrate on the distribution of baryonic, stellar and gaseous matter, in interacting disc galaxies moving through an ambient hot medium. By applying a density threshold technique to model star formation (Hernquist \\& Springel 2003) we present the impact of ram-pressure stripping on the star formation rate and the positions of star formation in comparison to galaxy interactions without the presence of ram pressure. ",
        "conclusions": "In order to investigate the influence of the ram pressure on interacting galaxies, we compared the results of a simulation of an interaction with and without a constant wind. We focus on the evolution of the star formation rate of the interacting system and on the distribution of the different baryonic components (i.e. gas and stars). The results can be summarised as follows:  \\begin{itemize} \\item The star formation rate of the interaction is enhanced in the presence of an ambient hot (3keV) and rare medium ($10^{-28} $g/cm$^3$) together with ram pressure by a factor of three in comparison to the same interaction without the ambient medium.  \\item The morphology of the interacting system is strongly distorted. The tidal tails and the bridge between the interacting system are completely destroyed by the ram pressure. The resulting gas and stellar mass distributions of the two galaxies would not be characterised by observers as interacting system after the first close encounter.  \\item The amount of gas in the wake of the interacting system is 50\\% of the total gas mass of the interacting system after 500 Myr of ram pressure acting on it. Approximately 10-15\\% of the gas is heated up to temperatures above $1\\times10^7$K, which would be observable in X-rays.  \\item After 500 Myr of ram pressure $\\sim10\\%$ of all newly formed stars are formed in the wake of the interacting system at distances larger than 20 kpc behind the stellar disc in the case of the face on ram pressure. The same behaviour can be observed in the case of the edge on ram pressure.  \\item As the tidal tails and the gaseous bridge between the interacting system feel the ram pressure, knots of cold gas ($T<10^5$K) start to form and these irregular structures start to form new stars. These knots are found to contain several $10^6$ M$_{\\sun}$ of cold gas and newly formed stars. They can be classified as 'stripped baryonic dwarf' galaxies. The lifetime of these 'stripped baryonic dwarfs' is limited by turbulence and heat conduction. If the objects are in the slipstream of the disc galaxies, they can survive for a several hundred Myrs up to a Gyr. As the ram pressure in the gaseous wake is decreasing due to the decreasing relative velocities, the timescales for Kelvin-Helmholtz instabilities is increasing as well, therefore extending the lifetime of these objects. Thermal conduction is affecting 'stripped baryonic dwarfs' as well, but the low densities of the surrounding ICM seems to result in evaporation timescales at least larger than 500 Myrs. \\end{itemize}  \\noindent Concluding we found that interacting galaxies affected by a moderate ram pressure show a completely different behaviour compared to the case of no ram pressure acting on them. The simulations presented in this work are idealised in order to be able to distinguish between different effects. As galaxies move through a real cluster, the ram pressure changes and complex interactions with the cluster potential and other galaxies are present. These factors will be investigated in upcoming papers."
    },
    "0807/0807.1616_arXiv.txt": {
        "abstract": "We present the results of a study with the \\emph{Swift} Burst Alert Telescope in the 14 -- 195 keV range of the long-term variability of 5 low mass X-ray binaries with reported or suspected super-orbital periods --- 4U 1636-536, 4U 1820-303, 4U 1916-053, Cyg X-2 and Sco X-1. No significant periodic modulation was detected around the previously reported values in the 4U 1916-053, Cyg X-2 or Sco X-1 light curves. The $\\sim$170 d period of 4U 1820-303 was detected up to 24 keV, consistent with the proposed triple system model. The $\\sim$46 d period in 4U 1636-536 was detected up to 100 keV, clearly inconsistent with variable photoelectric absorption via a warped precessing disc. We speculate that the appearance of this modulation after 4U 1636-536 entered the low/hard state indicates that this variability could be linked to jet precession such as observed in SS 433. ",
        "introduction": "Long-term ``super-orbital'' periodic variability has been observed in soft X-rays from over 30 X-ray binaries with the All Sky Monitor (ASM) on the \\emph{RXTE} satellite \\citep[e.g.][]{Farrell_so07}. Also known as `long' or `third' periods, these modulations are defined simply as any periodic variability greater than the orbital period. The precession of a warped accretion disc modulating the X-ray flux from the compact object is currently the favoured model \\citep[e.g][]{Farrell_wij99}. \\citet{Farrell_og01} showed that while radiation driven warping gives a coherent picture of super-orbital periods in some systems, this model cannot explain the observed phenomena in all cases.  The limited spectral coverage of the \\emph{RXTE} ASM has so far precluded long-term studies of super-orbital variability at high energies. The launch of \\emph{Swift} in 2004 saw the arrival of a wide field monitoring instrument capable of complementing the \\emph{RXTE} ASM in hard X-rays -- the Burst Alert Telescope \\citep[BAT;][]{Farrell_bar05}. In this paper we present the results of a search for the previously reported super-orbital modulation of 5 low mass X-ray binaries in the BAT 14 -- 195 keV light curves. ",
        "conclusions": "No evidence for the reported super-orbital periods in 4U 1916-053 \\citep[199 d;][]{Farrell_gr92}, Cyg X-2 \\citep[77 d;][]{Farrell_sm92} or Sco X-1 \\citep[37 d/62 d;][]{Farrell_pe96,Farrell_og01} was found in the ASM or any of the BAT light curves. The power spectra were instead dominated by significant low-frequency noise, with no peaks appearing to indicate the presence of real long-term periodic modulation. The super-orbital periods in 4U 1916-053 and Cyg X-2 were originally detected using \\emph{Vela 5B} data, while the detection of super-orbital modulation in Sco X-1 was made using only the first 9 months of ASM data. It is thus apparent that either the reported detections were spurious or the mechanisms behind the super-orbital variability in each of these three systems has since ceased. In comparison, the $\\sim$46 d and $\\sim$170 d variability previously reported for 4U 1636-536 \\citep{Farrell_sh05} and 4U 1820-303 \\citep{Farrell_pr84b} respectively were detected in both the ASM and BAT light curves (Figure \\ref{fig1}).  The $\\sim$170 d period in 4U 1820-303 has been attributed to variable accretion linked to the presence of a third body in this system, whereby the eccentricity of the inner orbit is modulated by tidal effects produced by the outer body \\citep{Farrell_cho01}. Transitions between the high/soft and low/hard states over this timescale are linked to variability in the accretion rate \\citep{Farrell_bl00b} and the movement in and out of the inner region of the accretion disc \\citep[see for example][]{Farrell_mal07}, thus modulating the thermal Comptonisation region of the spectrum. This thermal Comptonisation component has been seen to dominate the spectrum up to $\\sim$50 keV, above which a high energy non-thermal tail is dominant \\citep{Farrell_ta07}. The non-detection of the modulation above 24 keV is therefore due to a lack of sensitivity of the BAT (as the modulation should be present in the 24 -- 50 keV band), precluding us from testing the idea that the modulation also affects the non-thermal tail and thus extends to higher energies.  The detection of the $\\sim$46 d period in 4U 1636-536 up to 100 keV is clearly inconsistent with variable photoelectric absorption by a warped precessing disc, as absorption should effect only the low energy (ASM) data. \\citet{Farrell_fi06} linked the recent reduction in X-ray flux and the appearance of a high energy tail detected by \\emph{INTEGRAL} in the spectrum of 4U 1636-536 to the onset of jet formation. The detection of the modulation at energies where the high energy tail is dominant (i.e. $>$ 30 keV) indicates that both the thermal Comptonisation and high energy tail components are varying over this timescale. Coincidentally, the $\\sim$46 d modulation first appears in the ASM data around the same time as this transition to the low/hard state. The appearance of a high energy power law tail in the spectrum in conjunction with the first detection of the $\\sim$46 d period and the reduction in flux is intriguing, implying a link between the three phenomena. If the assertion made by \\citet{Farrell_fi06} is correct --- that the high energy tail is produced in compact jets --- then the $\\sim$46 d modulation is likely also associated in some way with the jets. The modulation could thus be tied to either transient jet formation or the precession of the jets (linked in turn to the precession of a warped disc) in a similar fashion as seen in SS 433 \\citep{Farrell_ma84}. Further observations at radio wavelengths are being sought in order to confirm this hypothesis.\\\\  \\begin{figure*} \\begin{center} \\scalebox{0.8}{\\includegraphics{farrell_s_fig1.eps}} \\caption{Normalised power spectra of the ASM and BAT light curves of 4U 1636-536 (left panel) and 4U 1820-303 (right panel). The red horizontal dashed lines indicate the 99$\\%$ white noise significance levels. The blue vertical dotted line indicates the location of the reported super-orbital periods at $\\sim$46 d (4U 1636-536) and $\\sim$170 d (4U 1820-303) respectively.}\\label{fig1} \\end{center} \\end{figure*}"
    },
    "0807/0807.3339_arXiv.txt": {
        "abstract": "{We present a study of the distribution and kinematics of the neutral gas in the low-inclination Scd galaxy NGC\\,6946. The galaxy has been observed for 192 hours at  21-cm with the Westerbork Synthesis Radio Telescope. These are among the deepest observations ever obtained for a nearby galaxy. We detect widespread high-velocity \\hi\\ (up to about 100 \\kms) and find 121 \\hi\\ holes, most of which are located in the inner regions where the gas density and the star formation rate are higher. Much of the high-velocity gas appears to be related to star formation and to be, in some cases, associated with \\hi\\ holes. The overall kinematics of the high-velocity gas is characterized by a slower rotation as compared with the regular disk rotation.  We conclude that the high-velocity gas in NGC\\,6946 is extra-planar and has the same properties as the gaseous halos observed in other spiral galaxies including the Milky Way. Stellar feedback (galactic fountain) is probably at the origin of most of the high-velocity gas and of the \\hi\\ holes. There are also indications, especially in the outer regions, --an extended \\hi\\ plume, velocity anomalies, sharp edges, and large-scale asymmmetries--  pointing to tidal encounters and recent gas accretion. } ",
        "introduction": "\\label{introduction}  Deep \\hi\\ surveys of nearby spiral galaxies carried out in recent years have revealed the presence of cold extra-planar gas and have opened a new chapter in the study of the disk-halo connection. An extended, lagging \\hi\\ halo  has been discovered in the edge-on galaxy NGC\\,891 \\citep{1997ApJ...491..140S,2007AJ....134.1019O}. Extra-planar gas with similar properties has also been detected in a number of nearby spirals \\citep[see a review by][]{2008arXiv0803.0109S}. Some of these \\citep[e.g.\\ NGC\\,2403,][]{2000A&A...356L..49S,2001ApJ...562L..47F} are not edge-on, and the presence  of extra-planar \\hi\\ has been inferred from its  anomalous velocities. In the Milky Way, halo gas components are the well-known Intermediate-Velocity Clouds (IVCs) and, as conclusively shown by recent distance determinations \\citep{2007ApJ...670L.113W,2008ApJ...672..298W}, also the High-Velocity Clouds (HVCs)  \\citep{1997ARA&A..35..217W}.  The origin of the halo gas is still an open question. One of the mechanisms proposed is the so called Galactic Fountain  \\citep{1976ApJ...205..762S,1980ApJ...236..577B} according to which  gas, heated by stars and supernovae, leaves the disk and moves via chimneys in the vertical direction to fall back, eventually, onto the disk. There is also evidence, however, that part of the extra-planar gas must be infall from intergalactic space \\citep{2008arXiv0803.0109S}. It is likely, for example, that the Galactic HVCs, because of their low metallicity, are such accreted gas \\citep{2004hvc..conf..195V}.  Most of the observations which have revealed the presence of extra-planar gas have been carried out for spiral galaxies viewed edge-on or highly inclined. Naturally, they have been used to study the vertical extent of the extra-planar \\HI and  its rotation. A more ``face-on'' view is required to investigate the gas motion perpendicular to the plane, to detect high-velocity gas and to study its possible connections with structures in the disk, such as spiral arms, \\hi\\ holes, star clusters, and \\hii\\ regions.  For this purpose we have carried out a very sensitive \\hi\\ survey with the Westerbork Synthesis Radio Telescope (WSRT) of the nearby spiral galaxy NGC\\,6946. This is a Scd galaxy seen at low inclination showing high-level star formation activity \\citep{1984ApJ...278..564D}.  The bright optical disk ($R<$ \\r25) shows many large \\HII complexes. Some \\HII regions are also present in the outer arms  \\citep{1998ApJ...506L..19F}. Furthermore, there is an extended disk of diffuse \\Ha  and X-ray emission \\citep{2003ApJ...598..982S}.  Over the last 30 years, NGC\\,6946 has been observed several times in the 21-cm line. \\citet{1973A&A....22..111R} made the first synthesis observation at the Owens Valley Radio Observatory. Their \\HI map already showed a gas disk that extends well beyond the optical image and deviations from circular rotation could be seen in their velocity field despite the low resolution of 2$^{\\prime}$. \\citet{1986ApJ...308..600T} confirmed these results with their 40$^{\\prime\\prime}$ data from the Very Large Array. \\citet{1990A&A...234...43C} and \\citet{1992A&A...266...37B} made a more detailed study of the kinematics and the distribution of the \\HI in NGC\\,6946 using the WSRT. They were the first to report the presence of holes in the \\HI distribution. In addition, they reported the detection of a diffuse fast gas component and of a few isolated high-velocity \\HI clouds. \\citet{1993A&A...273L..31K} studied the widespread high-velocity \\HI and found that this is predominantly located in the direction of the bright optical disk  suggesting a link with stellar winds and supernovae (i.e. a galactic fountain). \\citet{1993PhDT.......185K} also made the first detailed study of \\HI holes in this galaxy.  Catalogues of \\HI shells have been compiled for the gas-rich galaxies in the Local Group, such as M\\,31 \\citep{1986A&A...169...14B}, M\\,33 \\citep{1990A&A...229..362D}, the LMC \\citep{1999AJ....118.2797K} and SMC \\citep{1999MNRAS.302..417S}. Many \\HI holes have been detected in galaxies in neighbouring groups: Ho~II \\citep{1992AJ....103.1841P}, NGC\\,2403 \\citep{1998A&A...332..429T}, IC~2574 \\citep{1999AJ....118..273W}, M~101 \\citep{1988AJ.....95.1354V,1993PhDT.......185K}, and IC 10 \\citep{1998AJ....116.2363W}. The shells and holes are commonly thought to have been produced by clustered supernova explosions and stellar winds \\citep{1986PASJ...38..697T,1987ApJ...317..190M,1988ApJ...324..776M,1988ARA&A..26..145T}. Simulations show that these are energetic enough to form kpc-size bubbles and chimneys \\citep[see e.g.][]{2001MNRAS.328..708D}, which would appear as holes in the \\HI distribution, when observed in an external galaxy.  Here, we report the main results of a new, deep \\hi\\ study of NGC 6946 with the WSRT. In particular, we draw attention to the presence of a large number of \\HI holes (Section \\ref{sec:holes}) and of gas complexes with anomalous velocities (Section \\ref{sec:hvgas}). We discuss the role of star formation and environment in the formation of these features (Section \\ref{sec:discussion}).  \\begin{figure*}[!ht] \\centering \\includegraphics[width=0.95\\textwidth]{0120fg01.eps} \\caption{The left panel is a colour composite of the Digitized Sky Survey plates. The right panel is the deep (192 hours integration) \\HI map on the same scale as the optical. Column densities range from $6\\times10^{19}$ cm$^{-2}$ to $3.7\\times10^{21}$ cm$^{-2}$. The \\HI in the centre is seen in absorption against the bright core, producing a black spot on the \\HI map.\\label{2_optical_totalHI}} \\end{figure*} ",
        "conclusions": "\\label{sec:conclusions}  We have reported 21-cm line observations, among the deepest for a spiral to date, of the low-inclination galaxy NGC\\,6946. We have found:  1) A large number of holes in the HI disk. Sizes are up to 3 kpc diametre. They are mostly located within the bright optical disk and preferentially in regions of high \\hi\\ column density and star formation. Their ages are in the range of 1 to $6 \\times 10^7$ yr. Some of them have \\HII regions at their rims.  2) Widespread, clumpy high-velocity \\hi. The velocities are up to about 100 \\kms\\ (for the faintest features at the detection limit). Most of the high-velocity gas is seen in the direction of the inner bright optical disk. It has a rotational component following the disk rotation but lagging behind it. There is also some \\hi\\ at forbidden velocities. The total mass of the gas with \\vdevabs\\,$>50$\\,\\kms\\ is $2.9\\times10^8$ \\msun, ~4\\% of the total \\HI mass.  3) A large \\hi\\ outer plume, sharp outer edges, and a lopsided structure and kinematics. In the outer parts there are also pronounced velocity wiggles and associated \\hi\\ cavities.  4) \\hi\\ disk velocity dispersion reaching maximum values of 12-13  \\kms\\ inside $R_{\\rm 25}$ and decreasing (linearly) down to 6 \\kms in the outer parts.  We conclude that:  1) The high-velocity \\hi\\ is extra-planar gas which rotates more slowly than the disk. This has the same properties as the \\hi\\ halos found in nearby spiral galaxies such as NGC\\,891, NGC\\,2403 and also the Milky Way (IVCs and HVCs).  2) Most of this extra-planar \\hi\\ is fountain gas which originates from the regions of star formation in the bright inner disk and is probably related to the presence of the numerous \\hi\\ holes.  3) There is evidence, mainly in the outer parts of NGC\\,6946, pointing to  recent tidal encounters and minor mergers."
    },
    "0807/0807.1420_arXiv.txt": {
        "abstract": "Current cosmological constraints from Cosmic Microwave Background (CMB) anisotropies are typically derived assuming a standard recombination scheme, however additional resonance and ionizing radiation sources can delay recombination, altering the cosmic ionization history and the cosmological inferences drawn from CMB data. We show that for recent observations of CMB anisotropy, from the Wilkinson Microwave Anisotropy Probe satellite mission 5-year survey (WMAP5) and from the ACBAR experiment,  additional resonance radiation is nearly degenerate with variations in the spectral index, $n_s$, and has a marked effect on uncertainties in constraints on the Hubble constant, age of the universe, curvature and the upper bound on the neutrino mass. When a modified recombination scheme is considered, the redshift of recombination is constrained to $z_*=1078\\pm11$, with uncertainties in the measurement weaker by one order of magnitude than those obtained under the assumption of standard recombination while constraints on the shift parameter are shifted by 1$\\sigma$ to ${\\cal R}=1.734\\pm0.028$. Although delayed recombination limits the precision of parameter estimation from the WMAP satellite, we demonstrate that this should not be the case for future, smaller angular scales measurements, such as those by the Planck satellite mission. ",
        "introduction": "The recent measurements of the Cosmic Microwave Background (CMB) flux provided by the five year Wilkinson Microwave Anisotropy Probe (WMAP) mission (see \\cite{wmap5cosm,wmap5komatsu} and the ACBAR collaboration (see \\cite{acbar}) have confirmed several aspects of the cosmological standard model and improved the constraints on several key parameters. The constraints on the neutrino mass, for example, coming just from CMB data are drastically improved and are now competitive with current direct laboratory measurements (see e.g. \\cite{wmap5komatsu,fogli}). Moreover, the presence of a neutrino background is now inferred at more than $95 \\%$ c.l. by CMB data alone (see e.g., \\cite{wmap5komatsu,fdebe08}). Finally, inflationary parameters, curvature, baryonic and dark matter densities are also now better determined due to improved treatment of systematics (see \\cite{wmap5komatsu,kinney08}).  These spectacular results, apart from the experimental improvements, have been possible  due to the high precision of the CMB theoretical predictions that have now reached an accuracy close to $0.1\\%$ over a wide range of scales. A key ingredient in the CMB precision cosmology is the accurate computation of the recombination process. Since the seminal papers by Peebles and Z'eldovich (see \\cite{Peebles68,Zeldovich68}) detailing the recombination process, further refinements to the standard scheme were developed \\cite{Seager99} allowing predictions at the accuracy level found in data from the WMAP satellite and predicted for the future Planck satellite \\cite{Hu95,Seljak03,hirata}.  While the attained accuracy on the recombination process is impressive, it should be noticed that these computations rely on the assumption of standard physics. Non-standard mechanisms such as (just to name a few) high redshift stars or active galactic nuclei, topological defects and dark matter decays could produce extra sources of radiation and modify the recombination process. With the WMAP results and the future Planck data, it therefore becomes conceivable that deviations from standard recombination may be be detected.  Several papers in recent years have indeed investigated this possibility. For example, one could use a phenomenological approach such as in \\cite{Hannestad01} or \\cite{Lewis:2006ym}. Other works have adopted a more physically motivated (but model-dependent) approach and considered modified recombination by allowing time variations in the fine-structure (\\cite{alpha}) or gravitational constants (\\cite{zalzahn}). Here we instead focus on delayed recombination mechanisms based on the hypothesis of extra sources of ionizing and resonance radiation at recombination (see e.g. \\cite{Seager00,Naselsky02,Dorosh02}).  Extra photon sources can be generated by a variety of mechanisms. A widely considered process is the decay or annihilation of massive particles \\cite{Sarkar:1983,Scott:1991,Ellis:1992,Adams:1998,Doroshkevich:2002ff,Naselsky:2003zj,Zhang:2006fr}. The decay channel depends on the nature of the particles, and could, for example, include charged and neutral leptons, quarks or gauge bosons.  These particles may then decay further, leading to a shower cascade that could,  among other products, generate a bath of lower energy photons that would interact with the primordial gas and cosmic microwave background. Interestingly these models, as well as injecting energy at recombination, boost the ionization fraction after recombination and can distort the ionization history of the universe at even later times, during galaxy formation and reionization \\cite{Pierpaoli:2003rz,Chen:2003gz,Padmanabhan:2005es,Mapelli:2006ej,Lewis:2006ma}. Other mechanisms include evaporation of black holes \\cite{Naselsky87,Naselsky02} or inhomogenities in baryonic matter \\cite{Naselsky02}.  Several authors have already compared delayed recombination with cosmological data (see for example \\cite{bms,bms2,naso,petruta,cinesi}). With respect to these previous analyses, we assess below the improvements given by more recent data from the WMAP five year survey and from the ACBAR experiment. Moreover, we will study in detail the impact of delayed recombination on the current constraints on cosmic parameters. As showed in \\cite{Seager00} delayed recombination has two main effects: damping of the CMB anisotropy and polarization at small angular scales and a shift of the acoustic peaks in their angular spectra. As we will see, these two effects change in a significant way the current constraints on such parameters as the spectral index $n_S$ and the Hubble constant $H_0$. We then study the impact on the current determination of the redshift of recombination $z_*$ and on the shift parameter $\\cal R$. Both parameters are used to provide complementary geometric constraints on the background expansion history; $R$ as a measure of the angular diameter distance to last scattering, and $z_*$ in interpreting the scale of baryon acoustic oscillations \\cite{Eisenstein:2005su,Percival:2007yw}, which in combination with supernova type Ia data, can impose constraints on the dark energy component (see e.g. \\cite{efstathiou,corasaniti,Elgaroy:2007bv}. Here we show that delayed recombination has decisive effects on the determination of these parameters. Finally, we  discuss the implications of a modified recombination scheme for the Planck satellite.  The paper proceeds as follows: in section \\ref{sec2} we briefly describe the delayed recombination scheme. In \\ref{sec3} we analyze the latest CMB data and place new constraints on delayed recombination. Moreover, we forecast the constraints prospectively attainable by the Planck satellite mission, and study the impact that a modified recombination scheme can have on several key cosmological and astrophysical parameters. In \\ref{sec4} we draw our conclusions and review their future implications. ",
        "conclusions": "\\label{sec4}  In this paper, we have updated the upper bounds that can be placed on the contribution of extra Ly-$\\alpha$ and ionizing photon-producing sources in light of the new WMAP and ACBAR data. We have found that, adopting a simple parametrization using constant, effective values for $\\epsilon_{\\alpha}$ and $\\epsilon_{i}$, the WMAP5 data constraints $\\epsilon_{\\alpha} < 0.31$ and $\\epsilon_{i}<0.053$ at the $95\\%$ level.  We have studied the implications of delayed recombination on several current cosmological constraints derived from CMB anisotropies. While we found no relevant changes on current CMB estimates of curvature, optical depth, baryon and cold dark matter densities, delayed recombination proves to have a significant impact on current constraints for inflationary parameters. Current conclusions for theoretical inflationary models, motivated by evidence for a deviation from scale invariance from recent CMB data in the standard recombination scheme, would have to be relaxed when delayed recombination is considered, since the inflationary spectral index of scalar perturbations is in complete agreement with a scale invariant HZ spectrum.  The delayed recombination also has an impact on estimates of the shift factor, $\\cal R$, and recombination redshift, $z_*$, currently used in estimating angular diameter distance and baryon acoustic oscillation measures of the cosmic expansion history. Moreover, constraints on particle physics parameters like the neutrino mass can also be relaxed when non-standard recombination is considered.  Physically motivated models for non-standard recombination, like those based on primordial black hole or super heavy dark matter decay, are possible and provide a good fit to the current data. Future observations in both temperature and polarization, as those expected from the Planck satellite will be needed if we are to more stringently test and reduce the dependency of other cosmological parameters on delayed recombination models.  \\medskip \\textbf"
    },
    "0807/0807.4669_arXiv.txt": {
        "abstract": "In this paper, we provide new 1 Ms $Chandra$ ACIS spectra of two quasars in the Chandra Deep Field South (CDF-S), which were previously reported to show strong and extremely blueshifted X-ray emission/absorption line features in previous 1 Ms spectra, with outflowing bulk velocity $v\\sim$0.65-0.84c. In the new 1 Ms spectra, the relativistic blueshifted line feature is solidly confirmed in CXO CDFS J033225.3-274219 (CDFS 46, $z$ = 1.617), and marginally visible in CXO CDFS J033260.0-274748 (CDFS 11, $z$ = 2.579), probably due to the increased Chandra ACIS background in the new 1 Ms exposure. The new data rule out the possibility (though very tiny already based on the old 1Ms data) that the two sources were selected to be unusual due to noise spikes in the spectra. The only likely interpretation is extremely blueshifted iron absorption/emission line or absorption edge due to relativistic outflow. We find that the rest frame emission line center in CDFS 46 marginally decreased from 16.2 keV to 15.2 keV after 7 years. The line shift can be due to either decreasing outflowing velocity or lower ionization level. Including the two quasars reported in this paper, we collect from literature a total of 7 quasars showing blueshifted emission or absorption line feature with $v\\geq0.4c$ in X-ray spectra, and discuss its connection to jet and/or BAL (broad absorption line) outflow. ",
        "introduction": "It's well known that quasars interact with their environments through collimated jets and outflowing winds, which is required as the feedback (from the central supermassive black hole to the host galaxy) for the the M-$\\sigma_*$ relation (Gebhardt et al. 2000). The collimated jets and outflowing winds are also natural diagnostics to study the gas flow patterns in the innermost regions of black holes, the source geometry and energy generation mechanism (e.g., see Cappi 2006).  Intrinsic outflowing systems may be categorized by their outflowing velocities. The jets are often seen in radio-loud quasars with velocities very close to the speed of light (except for shocks and lobe-dominated sources, which are interacted with the surrounding medium and appear to be sub-relativistic, e.g., Kellermann et al. 2004). The broad absorption lines (BALs) in UV/optical band are blueshifted with velocities up to 0.2c, presenting in the spectra of % 10$\\sim$20\\% quasars. The BAL outflow can be seen in the X-ray spectra at even higher velocity. Chartas et al. (2002, 2003, 2007a, 2007b) discovered very broad absorption lines with outflow velocity of 0.4-0.67c in three gravitational lensed radio quiet BAL quasars (PG 1115$+$080, H 1413$+$117, and APM 08279$+$5255; see table 1 for details), which are the only BALs reported with relativistic outflows ($v\\geq0.4c$) in their X-ray spectrum.  Meanwhile, relativistic outflows with $v\\sim0.6-0.8c$ are also seen in the X-ray spectra of four non-BAL quasars. Yaqoob et al. (1998) reported blueshifted O$_{VII}$ emission line in the ASCA spectrum of PKS 0637-752 with outflow velocity of 0.77c. The line is however invisible in the Chandra exposure obtained $\\sim$ 2 years later, which instead detected an extended X-ray jet in this core-dominated radio-loud quasar. Yaqoob et al. (1999) reported the detection of blueshifted iron K emission line in radio-loud quasar PKS 2149-306 with velocity of 0.71c (see Table 1) in the ASCA spectrum. Such feature is also visible in following up XMM spectrum (see Fig. 2 of Page et al. 2005), and our independent analysis gives a confidence level of 99\\% in XMM PN spectrum. With the 1 Ms Chandra exposure on Chandra Deep Field South (CDF-S), Wang et al. (2003; 2005) reported the detection of strong blueshifted iron K line features with $v\\sim$0.65-0.84c, in CXO CDFS J033225.3-274219 (CDFS 46, $z$ = 1.617, emission line, radio-loud) and CXO CDFS J033260.0-274748 (CDFS 11, $z$ = 2.579, absorption line, radio-quiet) respectively. The blueshift velocities of the three emission lines reported are far larger than that could be produced by the rotational velocities of the accretion disk, and can only be due to outflow (see Wang et al. 2003).    In this paper, we report new 1 Chandra Ms observation\\footnote{The $Chandra$ Deep Field South (CDF-S, e.g., Giacconi et al. 2002; Rosati et al. 2002) now has a newly 1 Ms exposure (from 2007 September to 2007 November), and the data has been released. See http://cxc.harvard.edu/cda/whatsnew.html\\#CDFS2000-2007}  of the two CDF-S sources with relativistic outflows. Their multi-band properties are described in details in Wang et al. (2003; 2005). The X-ray data and spectral fitting are reported in \\S2 and the nature of the blueshifted features are discussed in \\S3. Throughout the paper we assume a cosmological model with $H_0 = 70$ km s$^{-1}$ Mpc$^{-1}$,  $\\Omega_m$ = 0.3, and $\\Omega_{\\Lambda}$ = 0.7. ",
        "conclusions": "\\label{lineorigin}  Wang et al. (2003; 2005) have ruled out the possibilities that the unusual X-ray features in CDFS 46 and CDFS 11 are due to any instrumental artifact, or due to improper background subtraction. The conclusions hold for the new 1 Ms exposure, and get stronger since it's even more unlikely that any artifact mentioned above, can affect both exposures obtained 7 years apart. Statistical fitting results show that the confidence level of the unusual features in both CDFS 46 and CDFS 11 is lower in the new 1 Ms exposure, than early reported with old 1 Ms exposure. This is actually expected due to the increasing \\emph{Chandra} ACIS background with time.\\footnote{The ACIS background rates remained relatively constant from the year 2000 until the end of 2003, and started to increase in 2004 with 10\\% per year, independent of variability induced by the solar cycle. (See Section 6.16.2 of the Proposers' Observatory Guide v. 10 $-->$ http://cxc.harvard.edu/proposer/POG/) } Note the blueshifted absorption feature in CDFS 11 is only marginally visible in the new 1 Ms spectrum, and we are unable to tell whether it varied based on current data. Except for providing marginally support to the reality of the unusual absorption feature in CDFS 11 with new 1 Ms spectrum, we are not discussing on its nature more than Wang et al. (2005) did. In this paper we will focus on CDFS 46.   Wang et al. (2003) has interpreted the X-ray emission line in CDFS 46 as relativistic blueshifted iron K line intrinsic to the quasar at $z$ = 1.617 (see below for further discussion). We noticed that Basu (2006) proposed a different scenario to interpret the optical and X-ray spectrum of CDFS 46. He suggested that all the optical and X-ray emission lines are blueshifted, resulting from an ejection. However, we find this scenario is unacceptable due to several reasons. First, all the emission lines were identified to have negative redshift, i.e., the observed optical emission lines were identified as infrared lines in the rest frame. However, such extremely unusual galaxies with high blueshifts have never been reported. Furthermore, the identification of three optical emission lines yields a redshift/blueshift scatter of 2.8\\%, significantly larger than 0.8\\% we obtained. The scatter (corresponding to 8400 km/s) is even larger than the full width of the broadest emission line identified in CDFS46 (see Fig. 6 of Szokoly et al. 2004). Finally, the strong X-ray emission line was classified as S K$\\alpha$, which is supposed to be rather weak, and has never been detected before.   It's interesting to note that in the GOODS-S ACS images, a nearby fainter galaxy ($i$ = 24.7) is detected, 1.1\\arcsec\\ to the southwest of the quasar CDFS 46. We note that the possibility to have a random galaxy brighter than $i$ = 24.7 within 1.1\\arcsec\\ to certain known galaxy in GOODS-S is only 4\\%, suggesting that the two galaxies are likely at the same redshift. VLT/FORS spectroscopy observations have classified the nearby fainter galaxy as an emission line galaxy at redshift 1.609 based on the detection of [OII] 3727, and photometric measurements (Vanzella et al. 2008). This is within 0.5\\% to the redshift of CDFS 46 ($z$ = 1.617)\\footnote{The redshift of CDFS 46 was identified based on the detection of CIV, CIII] and MgII emission lines, see Wang et al. 2003; Szokoly et al. 2004.}. This provides further evidence showing the redshift measurement of CDFS 46 is reliable. CDFS 11 is not covered by GOODS ACS images.   \\subsection{The Nature of the Outflow}  Wang et al. (2003) has listed several possible explanations to the unusual strong emission line in CDFS 46, including iron emission line from a relativistic outflow intrinsic to the quasar, strong iron absorption edge due to a cold relativistic outflow, or intervening low-redshift ($z$ $\\sim$ 0.034) type 2 AGN. In the high spatial resolution GOODS ACS images, only CDFS 46 at $z$ = 1.617 and a nearby galaxy at $z$ = 1.609 (1.1\\arcsec\\ apart) are visible within 3.5\\arcsec\\ radius. Also, although strong Fe K$\\alpha$ line from neutral iron at 6.4 keV is not unusual in type 2 AGN, variation in the line central energy is not expected, contrary to what find in CDFS 46. These facts firmly rule out an intervening low-redshift type 2 AGN.  As Wang et al. (2003) stated, the emission line feature could also be statistically fitted by a blueshifted strong Fe absorption edge due to cold outflow. Spectral fitting to the 2 Ms spectrum shows that a heavy absorption (with $N_H$ = 3.2$^{+0.4}_{-0.3}$ $\\times$ 10$^{24}$ cm$^{-2}$), covering 99.92$\\pm0.04$\\% of the direct continuum, is required. Statistically we can not rule out this model based on the 2Ms exposure, however, we point out that, in addition to a Compton-thick cold outflow with a velocity of 0.65c, such a model can only work under more extreme conditions. These include that, the intrinsic X-ray luminosity of CDFS 46 needs to be $\\gtrsim$ 500 times larger (L$_{2-10keV}$ $\\sim$ 3.0  $\\times$ 10$^{46}$ erg/s), which is way too luminous for its radio and optical emission\\footnote{The intrinsic X-ray to optical ratio $\\alpha_{ox}$ obtained under this model is 0.2.}; the leaking X-ray emission (or scattered X-ray emission) can not exceed 0.12\\%; the cold outflow needs to slow down from 0.69c\\footnote{The velocity of 0.74c (page L90, right column, 5th line) in Wang et al. (2003) should be corrected to 0.69c. There is a mistake in the calculation from Doppler factor to velocity in Wang et al. (2003).} to 0.65c after 2.7 years (in the quasar rest frame), and have to be locate at a large distance from the unusual strong central continuum, to avoid to be ionized, otherwise we would see strong soft X-ray emission which is free from absorption by an ionized absorber. Assuming a density of 10$^{11}$ cm$^{-3}$, which is typical of the Broad Line Region in AGN, we obtained a distance  $r$ $>$ $2\\times10^{16}$ cm from the central emission. Assuming the outflow cloud has a round or square shape, the outflow rate is then $\\dot{M}=4\\pi f_c *r^2n*m_p*v_{wind}\\simeq$ $(4\\pi f_c)*2.1\\times 10^4 M_{\\sun}$ yr$^{-1}$, which is higher than the Eddington accretion rate $\\dot{M_E}=2.2$ M$_{\\sun}$ yr$^{-1}$ of a 10$^8$ solar mass black hole, even for a very small covering factor of the outflow $f_c$ $\\sim$ 10$^{-4}$.  Therefor, the line feature we detected is most plausibly blueshifted iron line due to a relativistic outflow intrinsic to the quasar. The Doppler factor of 2.2 (for H- like Fe$_{XXVI}$ K$_\\alpha$, 6.97 keV) implies bulk velocity of $\\sim$0.65c (head-on) that must be responsible for the blueshift and relativistic boosting intensity. Assuming a lower ionization state would yields higher velocity up-to 0.7c (for neutral Fe K$_{\\alpha}$ line). If the line is due to fluorescent emission or recombination emission from plasma ionized by X-ray continuum, the continuum must illuminate the outflow from sides to produce strong line with equivalent width above several keV. In this case, the equivalent width of the line can be boosted by a factor of DF$^{3+\\Gamma}$, where  DF means the Doppler factor for the bulk motion, and $\\Gamma$ the photon index of the X-ray continuum.  If the outflow is not ionized by the central X-ray emission, considering the recombination timescale for Fe$_{XXV}$ is $t_{recomb}\\sim 3\\times 10^4 Z^{-2}T_5^{1/2}n_9^{-1}$ s (Allen 1973, here T$_5$ is the temperature in unit of 10$^5$ K, and n$_9$ the electron density in unit of $10^9$cm$^{-3}$.), typically in the range between 4.4$\\times 10^3$ and 4.4 s (Chartas, et al. 2002), there must be some mechanics (such as magnetic driving) which keep heating the relativistic outflow. Migliari, Fender, \\& Mendez (2002) had reported the discovery of blueshifted and very strong iron emission lines (with velocity of 0.26c and equivalent width of 13 keV, Migliari, private communication) from extended X-ray emission in the X-ray binary system SS 433. Their discovery implied the existence of large-scale reheating of baryon, which may work for the extreme blueshifted iron line in CDFS 46. We note that the new HETG observation of SS 433 also shows relativistic red- and blue- shifted lines from the central part indicating two side relativistic jets (Marshall et al. 2002, Lopez et al. 2006). Taking the velocity of the outflow as 0.70c, the distance which the outflow travelled during the 7 years' observation gap is $\\sim0.6$ pc = 1.8$\\times 10^{18}$ cm (in the quasar rest frame), which is $\\sim$5 times the dust sublimation radius ($R_d=0.4L_{45}^{1/2}\\simeq$0.13 pc, here L$_{45}$ is the intrinsic 0.5-10 keV luminosity in unit of $10^{45}$ erg s$^{-1}$.). The distance is too large for an adiabatic cooling process except for the existence of reheating (Migliari, Fender, \\& Mendez 2002).  The marginal change in the observed line central energy can either be due to a slow down of the outflow (from 0.69c to 0.65c for H-like Fe K$\\alpha$ line), or a change of the ionization state, for example, from H-like Fe K$\\alpha$ dominated to He-line Fe K$\\alpha$ dominated (see table 3). In the former case, we expect weaker emission line due to the weaker boosting effect, which is consistent with observation (see Fig 1 and table 3). In the later case, the line could either goes stronger or weaker, depending on the exact ionization state. We note that Chartas et al. (2007a) discovered 0.92-yr and 5.9-day (rest frame) variability of the outflowing absorbing gas, with velocity of $\\sim$0.1c (0.05c) and $\\sim$0.4c (0.36c) due to He-like Fe K$_\\alpha$ (H-like Fe K$_\\alpha$) resonant absorption, in the large gravitational lensed BAL quasar PG 1115$+$80.   \\subsection{Why So Rare?} \\label{rare} Our two quasars are among the several AGNs which show X-ray outflows with velocity $\\geq$ 0.4 c (see Table 1). Considering the large number of X-ray spectra obtained for quasars with advanced X-ray telescopes such as XMM and Chandra, why there are only several sources show relativistic outflowing features in their X-ray spectra? We contribute this to the orientation and observational selection effect. The broad absorption line outflow in quasars, with velocity up-to 0.1c, is believed to have a sky coverage of 20\\%. We propose that relativistic outflow with velocity $\\geq 0.4$c has similar or smaller sky coverage, thus can only be detected in a small fraction of quasars.  The instrumental bandpass, such as \\emph{Chandra} (0.3--10 keV) and \\emph{XMM-Newton} (0.3--12 keV), also limits the detection of relativistic outflowing features to certain redshift range. Assuming a Doppler factor $\\geq$2 (or a bulk velocity of $\\geq$0.6 c), only H- like Fe K$_\\alpha$ line at redshift above 1.0 can be detected below 7 keV in the observed frame (where \\emph{Chandra} \\& \\emph{XMM} are sensitive to detect enough counts). Furthermore, to confirm an emission or absorption line feature in an X-ray spectrum requires high number of detected photons (more than several hundreds), thus they can only be detected in luminous sources or by long enough exposure.  \\subsection{Connection with jet or BAL Outflows in AGN} We have stated above that the strong blueshifted emission line could be possibly due to line emission from collimated jets. It's interesting to note that all the three sources presenting blueshifted emission lines in Table 1 are radio-loud, suggesting the blueshifted emission lines and jets are somehow connected.  Historically, in additional to jet, the most dominant examples of AGN outflows have been found in Broad Absorption Line Quasars (BAL QSOs), which show absorption troughs in UV and optical lines with velocities up to tens of thousands of km s$^{-1}$ (as large as 0.2 c, e.g., Hidalgo et al. 2007). Disk wind model is one of most popular model to BAL outflows (see Proga 2007), and the presence of X-ray shielding gas is required to prevent the gas from over ionization and attenuate the X-ray continuum as observed (with $N_H$ $\\geq10^{23}$ cm$^{-2}$, Gallagher \\& Everett, 2007). We propose that the innermost region of the X-ray shielding gas could be outflowing at a higher velocity (such as 0.6 $\\sim$ 0.7 c), and more ionized, which produces no photo-electric but He-like and H-like iron absorption to central X-ray continuum. Actually, Chartas et al. (2002; 2003; 2007a; 2007b) have identified blueshifted X-ray absorption features in BAL QSOs at higher velocities and higher ionization level ( see table 1). An iron absorption line is produced when our line of sight is covered by the innermost shielding gas, and an emission line is observed when our line of sight is not covered but close to the outflowing direction (see Fig. \\ref{diag}). Note the most recent simulation work of Sim et al. (2008) shows a similar and clear pattern that outflow can produce both iron absorption and emission features, depending on the viewing angle. Taking into account the relativistic boosting effect, the blueshifted Fe K emission line can be seen at a significant level. The innermost region of the X-ray shielding gas might have sky coverage different from the broad absorption line region, thus the presence of X-ray absorption or emission line features does not necessarily link to the presence of broad absorption line in optical/UV band. However, it's worth to look for possible blueshifted optical/UV features in CDFS 46 and CDFS 11, which requires optical spectra with much higher quality. On the other hand, the search for blueshifted X-ray features in high S/N X-ray spectra of BAL QSOs would also be helpful.  Finally, we note that BAL outflow could be aligned with jets, at least in some BAL QSOs, and they might be physically connected (e.g. see Wang et al. 2008). This suggest that both mechanisms (jet or BAL outflow) could probably work simultaneously to produce blueshifted emission/absorption lines with $v$ $>$ 0.4c in the X-ray spectra of quasars."
    },
    "0807/0807.3049_arXiv.txt": {
        "abstract": "{In the course of the life of a massive star, wind-wind interaction can give rise to the formation of circumstellar nebulae which are both predicted and observed in the nature.} {We present generic model calculations to predict the properties of such nebulae for blue supergiants.} {From stellar evolution calculations including rotation, we obtain the time dependence of the stellar wind properties and of the stellar radiation field. These are used as input for hydro-calculations of the circumstellar medium throughout the star's life.} {Here, we present the results for a rapidly rotating 12$\\mso$ single star. This star undergoes a blue loop during its post main sequence evolution, at the onset of which its contraction spins it up close to critical rotation. Due to the consequent anisotropic mass loss, the blue supergiant wind sweeps up the preceding slow wind into an hour glass structure. Its collision with the previously formed spherical red supergiant wind shell forms a short-lived luminous nebula consisting of two polar caps and a central inner ring. With time, the polar caps evolve into mid-latitude rings which gradually move toward the equatorial plane while the central ring is fading. These structures are reminiscent to the observed nebulae around the blue supergiant Sher~25 and the progenitor of SN~1987A. } {The simple model of an hour glass colliding with a spherical shell retrieves most of the intriguing nebula geometries discovered around blue supergiants, and suggests them to form an evolutionary sequence. Our results indicate that binarity is not required to obtain them.} ",
        "introduction": "During the course of their evolution, massive stars have strong winds which eject matter into their surroundings. During their post-main sequence evolution, these stars can move back and forth from the blue to the red side of the Hertzsprung-Russell (HR)~diagram and back to the red, with little time spent at intermediate effective temperatures (e.g., Langer~{\\cite{L91b}}). Hydrodynamic considerations imply that each such transition does produce a circumstellar shell: When the star moves from the blue to the red side of the HR~diagram, the slow red supergiant (RSG) wind will be stalled by the high pressure of the previously created hot wind bubble, and will accumulate into a shell at the location where this pressure equals the RSG wind ram pressure (Garc\\'{i}a-Segura et al.~{\\cite{GLM96b}}). We call such a more or less stationary shell the RSG~shell. When the star moves from the red to the blue side of the HR~diagram, the wind speed increases and the blue supergiant (BSG) wind plows up the preceding RSG wind into a rapidly expanding shell, which we call the BSG~shell.  Consequently, we expect a spectacular circumstellar phenomenon for stars undergoing so called blue loops, namely that it triggers the formation of an expanding BSG~shell, which will at some point smash into the previously formed stationary RSG~shell. While both, the RSG and the BSG~shell by itself, may be difficult to observe, their violent interaction may release enough energy to provide an observable nebula.  Despite this simple and intriguing expectation, there are so far only few attempts to obtain quantitative prediction for the outcome of the described shell interaction (see Blondin et al.~{\\cite{BL93}}, Martin et al.~{\\cite{MA95}}, Podsiadlowski et al.~{\\cite{PM05}}). Within an effort to describe this phenomenon through generic calculations, which use detailed stellar evolution models as input for the circumstellar hydrodynamic modeling (Chi\\c{t}\\v{a} et. al., in preparation), we focus here on the results for a rotating 12$\\mso$ single star. ",
        "conclusions": "\\begin{figure}[!h] \\centering \\resizebox{\\hsize}{!}{\\includegraphics[width=0.95 \\textwidth]{cartoonsab.eps}} \\caption{ Schematic representation of the interaction of the hour glass shaped BSG shell (blue) with the RSG shell (orange). The collision regions are marked with red color; they form the brightest parts of the nebula.} \\label{FigCartoon} \\end{figure}   Figure~\\ref{FigCartoon} illustrates a simplified picture of the formation of multiple ring nebulae, according to our model. It contains two kinematic components: a stationary, spherical RSG shell and an expanding hour glass structure. The strongly emitting parts of the structure are the collision surfaces, which are marked in red in the figure. We believe that both kinematic components are realized by nature in the circumstellar medium of massive stars. RSG shells are unambiguously predicted (Garc\\'{i}a-Segura~{\\cite{G07}}) and while they are not yet observationally confirmed, there seems to be no way to avoid their formation. Expanding hour glass structures, on the other hand, are a well documented feature in circumstellar nebulae of low and high mass stars (see Nota et al.~{\\cite{N95}}; Langer et al.~{\\cite{LGM99}}) and are thought to be confined by a circumstellar disk in the equatorial plane of the central star.  A number of predictions emerge from this simple model. First, the collision starts about 10$^4$\\,yr after the onset of the blue loop. This timing is set by the expansion speed of the BSG shell and the radius of the RSG shell. Second, the life time of the nebula is determined by the duration of the collision phase, as the emission time scale is shorter than that. In our example, this is about $10^4$\\,yr, or about 1\\% of the core helium burning life time. This provides an upper limit to the expected number of triple ring nebulae. Third, the rotation rate of the central star during the collision is high for a BSG, since it just about recovers from critical rotation. At the time of maximum brightness of the nebula, the equatorial rotation rate of our central star is about 80$\\kms$ (Fig.~\\ref{FigMvvo}). Fourth, as all the material in the nebula is ejected after the first dredge-up phase of the central star, the nebula material is nitrogen-rich, here enhanced by a factor of~6.5, and carbon and oxygen depleted by factors of~6.5 and~1.5, respectively. We note that the level of N-enrichment predicted by current stellar evolution models is quite uncertain (see Hunter et al. {\\cite{Hu08}}), but a RSG phase is still expected to produce some nitrogen enhancement. Due to the assumptions of efficient rotational mixing (Heger \\& Langer~{\\cite{HL00}), the star and nebula in our model are more enriched than expected from non-rotating stellar models. Fifth, one ingredient of our simple model, namely the RSG shell, is expected for massive stars, but not so for low mass star which produce planetary nebulae. Therefore, while quite analogous expanding hour glass structures are observed for both cases (Langer~{\\cite{L00}}), multiple ring nebulae formed through the collision process shown in Fig.~\\ref{FigCartoon} are expected around massive stars, but not as planetary nebulae. In this sense, the polar caps observed around the blue supergiant Sher~25 might be considered as the first indirect empirical confirmation of a RSG shell.  Previous models of multiple ring nebulae were mostly constructed in the context of the triple-ring structure observed around SN\\,1987A (Burrows et al.~{\\cite{B95}}; Crotts \\& Heathcote~{\\cite{CH00}}). While single star models often fail to explain important features (e.g., Martin \\& Arnett~{\\cite{MA95}}; Meyer~{\\cite{Me97}}; Woosley et al.~{\\cite{W97}}), many invoke rather complex binary phenomena (e.g., Podsiadlowski et al.~{\\cite{P91}}; Blondin \\& Lundqvist~{\\cite{BL93}}; Llyod et al.~{\\cite{L95}}; Morris \\& Podsiadlowski~{\\cite{PM05}}). Whereas we do not attempt to reproduce the circumstellar medium of SN~1987A, a single star approach with suitable choices for the major parameters in our model (initial mass, initial rotation rate, metallicity) appears promising and will be pursued in the near future. The current failure of single star models to produce suitable blue loops and blue supergiant pre-supernova models may have to do more with missing physics in stellar evolution models rather than supporting the evidence for a binary progenitor of SN\\,1987A (Woosley et al.~{\\cite{W97}})  Various multiple ring nebulae around blue supergiants have been observed in the last 20~years (Smith et al.~{\\cite{SBW07}}). While our generic numerical model was not designed to correspond to any of them, many of the general properties of these nebulae are well reproduced. Most striking is the agreement of the emission geometries. While the nebula around the B1.5~Ia supergiant Sher~25 shows two polar caps and one equatorial ring (Brandner et al.~{\\cite{BGCW97}}) and the other objects discussed by Smith et al.~{\\cite{SBW07}} rather show narrow rings, including the ``twin'' of the SN~1987A nebula around HD~168625 (Smith~{\\cite{Sm07}}) all these structure occur as an evolutionary sequence in our model. Expansion velocities of the inner ring ($\\sim 18\\kms$) and the outer collision products ($\\sim 50\\kms$), the spatial scale of about 1\\,pc, and the kinematic nebula age agree rather well with empirical values. The rotation velocity of our stellar model fits well to the derived value of $\\sim 70\\kms$ for Sher~25 (Hendry et al.~{\\cite{He08}}). Central star and nebula of our model are nitrogen enriched, as are most of the observed nebulae.  We note that the emission in our model is caused by compressional heating, which may be in conflict with evidence for photoionization being the dominant process in some observed multiple ring nebulae (see Smith et al.~{\\cite{SBW07}}). And indeed, looking at the density distributions shown in Fig.~\\ref{FigDenstemp}, which might resemble emission geometries in the pure photo-ionization case, the situation appears more complex. In our simulation, the thick RSG shell ($\\Delta r \\simeq r \\simeq 1\\,$pc) collapses in two parts (at $r\\simeq 1.2\\,$ pc and $r\\simeq 1.7\\,$) due to a cooling instability. However, this collapse is questionable since it requires a long timescale --- our shell has an age of close to $10^6\\,$yr, while in many cases the shell will be much younger at the time of collision ---, and since the employed cooling function is uncertain for temperatures below $10\\,000\\,$K. Without this collapse, its density would be only about $2\\times10^{-25}\\,$g\\,cm$^{-3}$ (or 0.1 particles/cm$^3$), which may render it unobservable even if it were photoionized. However, even in the case of the collapsed RSG shell as in our simulation, the collision leads to a clear density enhancement. In panel~2 of Fig.~\\ref{FigDenstemp}, we see that in our model, the density enhancement in the polar caps is about a factor~5. This is to be considered a lower limit, as higher resolution models might approach the theoretically expected enhancement factor of about~100, which follows from the (well realized) isothermal shock approximation and a Mach number of about~10. The lower panels of Fig.~\\ref{FigDenstemp} show that in order to represent the rings of SN~1987A, further refinements are required, which, as we think, might be achieved by altering the properties of the RSG shell. For this particular case, this may indeed be justified, as the life time of the final RSG stage of the progenitor of SN~1987A might have been quite short (Woosley~{\\cite{W88}}, Langer~{\\cite{L91a}}).  Despite that our model does not fit any of the observed cases in detail, the approximate agreement with most general properties of this class of objects encourages to produce tailored models for individual nebulae as next step. Our results indicate that stars with multiple ring nebulae might just have left the RSG branch --- as stellar evolution models argued for the case of SN~1987A (Woosley~{\\cite{W88}}, Langer~{\\cite{L91a}}) and furthermore, that binarity may not be required to obtain multiple ring emission geometries."
    },
    "0807/0807.4175_arXiv.txt": {
        "abstract": "We present two-dimensional line-of-sight stellar kinematics of the lens galaxy in the Einstein Cross, obtained with the \\gemini\\ 8m telescope, using the \\gmos\\ integral-field spectrograph. The stellar kinematics extent to a radius of $4$\\arcsec\\ (with $0.2$\\arcsec\\ spaxels), covering about two-thirds of the effective (or half-light) radius $\\Reff \\simeq 6$\\arcsec\\ of this early-type spiral galaxy at redshift $z_l \\simeq 0.04$, of which the bulge is lensing a background quasar at redshift $z_s \\simeq 1.7$.  The velocity map shows regular rotation up to $\\sim 100$\\,\\kms\\ around the minor axis of the bulge, consistent with axisymmetry. The velocity dispersion map shows a weak gradient increasing towards a central value of $\\sigma_0 = 166 \\pm 2$\\,\\kms. We deproject the observed surface brightness from \\hst\\ imaging to obtain a realistic luminosity density of the lens galaxy, which in turn is used to build axisymmetric dynamical models that fit the observed kinematic maps. We also construct a gravitational lens model that accurately fits the positions and relative fluxes of the four quasar images. We combine these independent constraints from stellar dynamics and gravitational lensing to study the total mass distribution in the inner parts of the lens galaxy.  We find that the resulting luminous and total mass distribution are nearly identical around the Einstein radius $\\Rein = 0.89$\\arcsec, with a slope that is close to isothermal, but which becomes shallower towards the center if indeed mass follows light. The dynamical model fits to the observed kinematic maps result in a total mass-to-light ratio $\\MLdyn = 3.7 \\pm 0.5$\\,\\MLsunI\\ (in the $I$-band). This is consistent with the Einstein mass $\\Mein = 1.54 \\times 10^{10}$\\,\\Msun\\ divided by the (projected) luminosity within $\\Rein$, which yields a total mass-to-light ratio of $\\MLein = 3.4$\\,\\MLsunI, with an error of at most a few per cent. We estimate from stellar populations model fits to colors of the lens galaxy a stellar mass-to-light ratio $\\MLstar$ from $2.8$ to $4.1$\\,\\MLsunI. Although a \\emph{constant} dark matter fraction of 20 per cent is not excluded, dark matter may play no significant role in the bulge of this $\\sim$\\Ls\\ early-type spiral galaxy. ",
        "introduction": "\\label{sec:intro}  In the cold dark matter (CDM) paradigm for galaxy formation \\citep[e.g.][]{2002SciAm.286..36K}, galaxies are embedded in extended dark matter distributions with a specific and (nearly) universal radial profile \\citep[e.g.][]{1997ApJ...490..493N, 2005ApJ...624L..85M}. Measurements of rotation curves from neutral hydrogen (\\HI) observations in the outer parts of late-type galaxies have provided evidence for the presence of dark matter in these systems already more than two decades ago \\citep[e.g.][]{1985ApJ...295..305V}. In the outer parts of early-type galaxies, however, cold gas is scarce \\citep[but see e.g.][]{1994ApJ...436..642F, 1997AJ....113..937M, 2008MNRAS.383.1343W}, so evidence for a dark matter halo has to come from other tracers, such as kinematics of stars, planetary nebulae and globular clusters \\citep[e.g.][]{1995ApJ...441L..25C, 2001AJ....121.1936G, 2003Sci...301.1696R, 2003ApJ...591..850C} or hot X-ray gas \\citep[e.g.][]{2006ApJ...636..698F}.  However, these tracers are not always (sufficiently) available, the observations are often challenging, and the interpretation modeling dependent. Also in the inner parts of galaxies the amount, shape and profile of dark matter is still poorly known \\citep[e.g.][]{2004IAUS..220...53P}, even though stellar kinematics are readily available.  A fundamental problem in using stellar kinematics (as well as other collisionless kinematic tracers) for this purpose is the mass-anisotropy degeneracy: a change in the measured line-of-sight velocity dispersion can be due to a change in total mass, but also due to a change in velocity anisotropy. Both effects can be disentangled by measuring also the higher-order velocity moments \\citep{1987MNRAS.224...13D, 1993ApJ...407..525V, 1993MNRAS.265..213G}, but only the inner parts of nearby galaxies are bright enough to obtain the required high-quality kinematic measurements \\citep[e.g.][]{1991MNRAS.253..710V, 2001AJ....121.1936G, 2006MNRAS.366.1126C}. A unique alternative to break the mass-anisotropy degeneracy is provided by gravitational lensing. In case of strong lensing, the mass of a foreground galaxy bends the light of a distant bright object behind it, resulting in multiple images. From the separation and fluxes of the images the total mass distribution of the lens galaxy can be inferred directly. The velocity anisotropy then can be determined from the observed velocity dispersion, without the need for higher-order velocity moments.  \\citeauthor{2004ApJ...611..739T} (\\citeyear{2004ApJ...611..739T}, and references therein) have applied this approach to several strong lensing systems, of which 0047-281 \\citep{2003ApJ...583..606K} is the best constrained case, with three radially separated velocity dispersion measurements extending to about the effective radius of the lens galaxy. They measure the mass within the Einstein radius by fitting a singular isothermal ellipsoid to the positions of the quasar images.  This Einstein mass is used to set the amplitude of the total (stellar and dark) matter distribution, which they assume to be spherical. The constant stellar mass-to-light ratio $\\MLstar$ determines the contribution of the stars with a fixed radial profile, with the remainder due to dark matter with a single power-law profile with slope $\\gamma$. They then compare the dispersion profile predicted by the spherical Jeans equations, for an ad-hoc assumption of the velocity anisotropy $\\beta$, with the observed dispersion measurements. Based on a reasonably constrained $\\MLstar$ and an upper limit on $\\gamma$, they conclude that a significant amount of dark matter is present in the inner parts of the lens galaxy, with a slope shallower than the nearly isothermal total mass distribution. An additional (external) constraint on $\\MLstar$ (or $\\gamma$) is needed to go beyond this limit on the dark matter distribution. Even so, their results are limited by too few kinematic constraints (which leaves the anisotropy degenerate), and by the use of a simple spherical dynamical model.  Clearly, most lens galaxies are significantly flattened and so cannot be well-described by spherical models. Non-spherical models provide a more realistic description of the lens galaxy, but the increase in freedom requires also (significantly) more spatially resolved kinematic measurements to constrain them. Only a few of the known strong gravitational lens systems are close enough to obtain such kinematic data. Even then, one can make (ad-hoc) assumptions on the velocity distribution, such as velocity isotropy in the meridional plane of an axisymmetric model. The latter restriction to a distribution function of two (instead of three) integrals of motion allows for the recovery of the intrinsic shape and mass distribution in the presence of two-dimensional kinematic data \\citep{2007ApJ...666..726B, 2008MNRAS.384..987C}, possibly even without the additional constraint provided by gravitational lensing \\citep[e.g.][]{2006MNRAS.366.1126C}.  In this paper, we relax the two-integral assumption and find that one can still derive independent determinations of the intrinsic mass distribution from both gravitational lensing and stellar dynamics with two-dimensional kinematic data. We have observed the gravitational lens system QSO\\,2237+0305, well-known as the Einstein Cross, with the integral-field spectrograph \\gmos\\ on the \\gemini-North Telescope.  We fit the extracted stellar velocity and velocity dispersion maps of the inner parts of the lens galaxy at a redshift $z_l \\simeq 0.04$ with axisymmetric dynamical models. The resulting intrinsic mass distribution agrees well with that inferred from the lens model that fits the four quasar image positions and relative fluxes. Furthermore, we compare this total mass distribution with the stellar mass distribution, based on an independent estimate of the stellar mass-to-light ratio $\\MLstar$, to place constraints on the dark matter distribution in the inner parts of the lens galaxy.  In Section~\\ref{sec:observations}, we briefly describe the Einstein Cross, and we present the photometric and spectroscopic observations. In Section~\\ref{sec:analysis}, we extract the two-dimensional stellar kinematics, describe the dynamical modeling method, and construct lens models. From the latter we infer the total mass distribution, which we compare in Section~\\ref{sec:results} with the luminous mass distribution inferred from the observed surface brightness. We then build axisymmetric dynamical models, and compare the resulting dynamical mass-to-light ratio with that inferred from the lens models, and in turn with an estimate of the stellar mass-to-light ratio. In Section~\\ref{sec:discconcl} we discuss our findings and we summarize our conclusions.  Throughout we adopt the \\wmap\\ cosmological parameters for the Hubble constant, the matter density and the cosmological constant, of respectively $H_0=73$\\,\\kmsM, $\\Omega_M=0.24$ and $\\Omega_L=0.76$ \\citep{2007ApJS..170..377S}. ",
        "conclusions": "\\label{sec:discconcl}  We used the \\gmos-North integral-field spectrograph to obtain two-dimensional stellar kinematics of the lens galaxy in the Einstein Cross. In addition to the four bright quasar images and the distance of the lens galaxy ($D_l = 155$\\,Mpc), in particular the presence of sky lines in the observed Ca II triplet region made the extraction of the absorption line kinematics challenging. Even so, we were able to derive high-quality line-of-sight velocity $V$ and dispersion $\\sigma$ maps of the bulge-dominated inner region $R \\lesssim 4$\\arcsec, reaching about two-thirds of the effective radius $\\Reff \\simeq 6$\\arcsec\\ of this early-type spiral galaxy. The $V$ map shows regular rotation up to $\\sim 100$\\,\\kms\\ around the minor axis of the bulge, consistent with axisymmetry. The $\\sigma$ map shows a weak gradient increasing towards a central value of $166 \\pm 2$\\,\\kms.  The only other direct (single) measurement of the velocity dispersion is $215 \\pm 30$\\,\\kms\\ by \\cite{1992ApJ...386L..43F}. For a singular isothermal sphere lens model, we can use the relation $\\Delta\\theta=8\\pi(\\sigma_\\mathrm{SIS}/c)^2D_{ls}/D_s$ \\citep[e.g.][]{2000ApJ...543..131K} with a separation $\\Delta\\theta = 1.8$\\arcsec\\ of the four quasar images, to obtain a simple estimate for the dispersion of $\\sigma_\\mathrm{SIS} \\simeq 180$\\,\\kms.  Taking into account aperture correction and a range in velocity anisotropy, \\cite{2003MNRAS.344..924V} convert this to a central velocity dispersion of $168 \\pm 17$\\,\\kms. The King and de Vaucouleurs models of \\cite{1988AJ.....96.1570K} predict a similar value of $\\sim 166$\\,\\kms, and also \\cite{1999MNRAS.309..641B} find a value of $165 \\pm 23$\\,\\kms\\ based on their two \\HI\\ rotation curve measurements. It is likely that the long-slit measurement by \\cite{1992ApJ...386L..43F} is affected by the bright quasar images, whereas our spatially resolved measurements allow for a clean separation of the quasar contribution.  A large variety of different lens models have been constructed for the Einstein Cross, most of which fit the positions of the quasar images but not their relative flux ratios. Adopting the scale-free lens model of \\cite{2003MNRAS.345.1351E}, we found that fitting at the same time also the (radio) flux ratios constrained the slope of the total mass surface density $\\Sigma \\propto R^{\\alpha-2}$ to be $\\alpha = 1.0 \\pm 0.1$. The total mass within the Einstein radius $\\Rein = 0.89$\\arcsec, i.e., the Einstein mass, is $\\Mein = 1.54 \\times 10^{10}$\\,\\Msun, nearly independent of the slope $\\alpha$, and consistent with previous measurements in the literature. Dividing by the projected luminosity within $\\Rein$, as measured from the observed $I$-band surface brightness, we obtained a mass-to-light ratio of $\\MLein = 3.4$\\,\\MLsunI, with an error of at most a few per cent.  We determined the mass-to-light ratio in an additional, independent way by fitting dynamical models to the (combined) observed $V$ and $\\sigma$ maps. We used the solution of the axisymmetric Jeans equations to predict this second-order velocity moment, with the gravitational potential inferred from the deprojected $I$-band surface brightness, assuming that mass follows light. We expect the resulting best-fit constant mass-to-light ratio to be an overestimation due to possible residual contribution from the quasar light. Even so, restricting the fit to the unaffected and accurately measured central second-order velocity moment, we found $\\MLdyn = 3.7 \\pm 0.5$\\,\\MLsunI. When we used instead the gravitational potential that follows directly, i.e., without mass-to-light ratio conversion, from the surface mass density of the best-fit lens models, we arrived at a similar prediction of the second-order velocity moment. We showed that the reason is that the luminous and total mass distribution, as inferred from respectively the surface brightness and lens model, are very similar. This implies that the inner region of the lens galaxy mass (closely) follows light, with a total mass-to-light ratio $\\MLtot \\simeq 3.4$\\,\\MLsunI.  By fitting single stellar population models to measured colors of the center of the lens galaxy, we estimated an $I$-band stellar mass-to-light ratio $\\MLstar$ from 2.8 to 4.1\\,\\MLsunI. Although a \\emph{constant} dark matter fraction of 20 per cent is thus not excluded, it is likely that dark matter does not play a significant role in the inner region of this early-type spiral galaxy of luminosity $\\sim L\\star$. This is consistent with indications that less-luminous early-type galaxies \\citep[e.g.][]{2001AJ....121.1936G, 2006MNRAS.366.1126C, 2006ApJ...649..599K, 2007MNRAS.382..657T} and late-type galaxies \\citep[e.g.][]{2000AJ....120.2884P, 2001ApJ...550..212B, 2006ApJ...643..804K} lack the dark matter that is ubiquitous in dwarf galaxies as well as in (at least the outer parts of) giant elliptical and spiral galaxies.  The constraint $\\alpha = 1.0 \\pm 0.1$ on the slope of the lens model implies that the intrinsic total mass density is close to isothermal, consistent with previous studies of lens galaxies \\citep[e.g.][]{2006ApJ...649..599K}. However, one has to be very careful not to over-interpret this result, not only due to assumptions in the modeling, but most of all because the slope is only (well) constrained around the Einstein radius $\\Rein$. If indeed mass follows light in the inner region of this lens galaxy, the (deprojected) surface brightness indicates deviations from isothermal towards the center where the slope becomes shallower as well as a possible steepening at larger radii. Interestingly, the construction of realistic galaxy density profiles with a stellar and dark matter component that are both non-isothermal, shows that the combined slope and corresponding lensing properties are nevertheless consistent with isothermal around $\\Rein$ \\citep{vandeVenetal2008,Mandelbaumetal2008}. With more extended images, e.g. in the case of galaxy-galaxy lensing, one can place a stronger constraint on the total mass distribution of the lens galaxy \\citep[e.g.][]{2007ApJ...666..726B, 2008arXiv0804.2827S}. Since the extension in these cases is still mostly tangential, constraining a large radial range is in particular possible when lensing occurs in multiple (redshift) planes, resulting in images at different Einstein radii \\citep[e.g.][]{2008ApJ...677.1046G}.  Nevertheless, already the quasar images provide an accurate constraint on the total mass within $\\Rein$, nearly independent of the details of the lens model \\citep[e.g.][]{1991ApJ...373..354K, 2001MNRAS.327.1260E}.  Therefore, at least around this radius no higher-order velocity moments are needed to break the mass-anisotropy degeneracy. Towards the center the surface brightness increases steeply, so that high enough S/N to measure velocity moments beyond $V$ and $\\sigma$ might be achievable even at higher redshift. In the outer parts, the degeneracy might be (partially) broken by combining the kinematics of stars and/or discrete tracers, such as globular clusters and planetary nebulae, with total mass estimates from hot X-ray gas and/or weak lensing measurement \\citep[e.g.][]{2007ApJ...667..176G}. In the current and even more in the upcoming extensive and deep photometric surveys, numerous (strong) gravitational lensing systems will be discovered.  They provide important and independent constraints on the total mass distribution in galaxies, especially in combination with (resolved) stellar kinematics, as we showed in this paper \\citep[see also][]{2007ApJ...666..726B, 2008MNRAS.384..987C}."
    },
    "0807/0807.3878_arXiv.txt": {
        "abstract": "In this paper, we combine the latest observational data, including the WMAP five-year data (WMAP5), BOOMERanG, CBI, VSA, ACBAR, as well as the Baryon Acoustic Oscillations (BAO) and Type Ia Supernoave (SN) ``Union\" compilation (307 sample), and use the Markov Chain Monte Carlo method to determine the cosmological parameters, such as the equation-of-state (EoS) of dark energy, the curvature of universe, the total neutrino mass and the parameters associated with the power spectrum of primordial fluctuations. Our results show that the $\\Lambda$CDM model remains a good fit to the current data. In a flat universe, we obtain the tight limit on the constant EoS of dark energy as, $w=-0.977\\pm0.056$ ($1~\\sigma$). For the dynamical dark energy models with time evolving EoS parameterized as $w_{\\DE}(a) = w_0 + w_1 (1-a)$, we find that the best-fit values are $w_0=-1.08$ and $w_1=0.368$, implying the preference of Quintom model whose EoS gets across the cosmological constant boundary during evolution. For the curvature of universe $\\Omega_k$, our results give $-0.012<\\Omega_k<0.009$ ($95\\%$ C.L.) when fixing $w_{\\DE}=-1$. When considering the dynamics of dark energy, the flat universe is still a good fit to the current data, $-0.015<\\Omega_k<0.018$ ($95\\%$ C.L.). Regarding the neutrino mass limit, we obtain the upper limits, $\\sum m_{\\nu}<0.533$ eV ($95\\%$ C.L.) within the framework of the flat $\\Lambda$CDM model. When adding the SDSS Lyman-$\\alpha$ forest power spectrum data, the constraint on $\\sum m_{\\nu}$ can be significantly improved, $\\sum m_{\\nu}<0.161$ eV ($95\\%$ C.L.). However, these limits can be weakened by a factor of $2$ in the framework of dynamical dark energy models, due to the obvious degeneracy between neutrino mass and the EoS of dark energy model. Assuming that the primordial fluctuations are adiabatic with a power law spectrum, within the $\\Lambda$CDM model, we find that the upper limit on the ratio of the tensor to scalar is $r<0.200$ ($95\\%$ C.L.) and the inflationary models with the slope $n_s\\geq1$ are excluded at more than $2~\\sigma$ confidence level. However, in the framework of dynamical dark energy models, the allowed region in the parameter space of ($n_s$,$r$) is enlarged significantly. Finally, we find no strong evidence for the large running of the spectral index, $\\alpha_s=-0.019\\pm0.017$ ($1~\\sigma$) for the $\\Lambda$CDM model and $\\alpha_s=-0.023\\pm0.019$ ($1~\\sigma$) for the dynamical dark energy model, respectively. ",
        "introduction": "\\label{Int}  With the accumulation of observational data from cosmic microwave background measurements (CMB), large scale structure surveys (LSS) and supernovae observations and the improvements of the data quality, the cosmological observations play a crucial role in our understanding of the universe, and also in constraining the cosmological parameters, such as the EoS of dark energy models, the curvature of universe, the total neutrino masses and those associated with the running of the spectral index and gravitational waves. In our previous work \\cite{Xiaplanck}, we have used the Markov Chain Monte Carlo (MCMC) method to constrain cosmological models from the astronomical observational data, including the WMAP three-year data (WMAP3) \\cite{WMAP3GF,WMAP3Other}, small-scale CMB data, LSS data \\cite{SDSS,2df} and SNIa ESSENCE sample \\cite{Essence}. We found that the cosmological constant is consistent with the data, however the dynamical dark energy models are still allowed and interestingly the model with its EoS getting across $w = -1$ the {\\it Quintom} model \\cite{Quintom} is the best fit model. And we found no strong significant evidence for the non-flat universe and massive neutrino. Within the $\\Lambda$CDM model, the scale-invariant spectrum and the spectra with $n_s>1$ are disfavored by more than $2~\\sigma$ confidence level. Due to the degeneracy between the EoS of dark energy and tensor fluctuation, those inflationary models excluded within the $\\Lambda$CDM model, will revive in the framework of dynamical dark energy models. Furthermore, we did not find any significant evidence for the tensor fluctuations and the large running of the spectral index.  Given the precision of current observations, these results are not conclusive for the time being. Recently, the WMAP experiment has published its five-year data of temperature and polarization power spectra \\cite{WMAP5GF1,WMAP5GF2,WMAP5Other}. The Arcminute Cosmology Bolometer Array (ACBAR) experiment has also published its new CMB temperature power spectrum \\cite{ACBAR}. These new CMB data can strengthen the constraints on the cosmological parameters, especially for the inflationary models \\cite{KinneyWMAP5,WMAP5GF1,WMAP5GF2}. Furthermore, the Supernova Cosmology Project has made an unified analysis of the world's supernovae datasets and presented a new compilation ``Union\" (307 sample) \\cite{Union} which includes the recent samples of SNIa from SNLS and ESSENCE Survey, as well as some older datasets, \\emph{etc}. In the literature \\cite{Union,LinderUnion}, this Union compilation combining with the shift parameter derived from CMB and the BAO information has been used to constrain cosmological models. However, in these studies the CMB information considered is just the shift parameter instead of the full CMB data, which will lose some information of CMB and lead to a biased result \\cite{Compare}. Thus, in this paper, we revisit the issue on the determination of these cosmological parameters and update our previous results with the latest observational data.  Our paper is organized as follows: In Section II we describe the method and the latest observational datasets we used; Section III contains our main global fitting results on the cosmological parameters and the last section is the summary. ",
        "conclusions": "\\label{Sum}  Recently many experimental groups have published their new observational data, such as temperature and polarization power spectra of WMAP5 \\cite{WMAP5GF1,WMAP5GF2,WMAP5Other}, temperature power spectrum of ACBAR \\cite{ACBAR} and supernovae dataset of Union compilation \\cite{Union}. In this paper we report the updated constraints on the cosmological parameters from these latest observational data, such as the EoS of dark energy, curvature of universe, neutrino mass and inflation parameters.  For dark energy, we explore the constraints on two kinds of dynamical models. Assuming a flat universe, within the WCDM model, we find that the latest observational data yield the limit on the constant EoS of dark energy, $w=-0.977\\pm0.056$ ($1~\\sigma$). For RunW model with a flat universe, we find that the best fit model is $w_0=-1.08$ and $w_1=0.368$ described by the Quintom theory with EoS across the cosmological constant boundary. And because the precision of current observations are not good enough to determine the dark energy EoS conclusively, the dynamical dark energy models are not excluded and the $\\Lambda$CDM model remains a good fit.  The Quintom scenario, with the particular feature that its EoS can cross the cosmological constant boundary smoothly, has been applied to many aspects of cosmology theoretically. Firstly, the Quintom dark energy models can also give rise to interesting prediction on the fates of the universe, different from the Quintessence or Phantom models, such as the cyclic universe \\cite{FengOsc,Oscillating}, an expanding universe with oscillating EoS. Secondly, applying a Quintom matter for the early universe can provide a scenario of bouncing cosmology, which can avoid the notorious issue of initial singularity \\cite{Bounce}.  Our results also show that the universe is very close to flatness and the upper limit on the total neutrino mass is $\\sum m_{\\nu}<0.533$ eV ($95\\%$ C.L.), from the combination of CMB, BAO and SN data. Given the efficiency of Ly$\\alpha$ forest data on constraining the total neutrino mass, we also perform a calculation with the inclusion of the SDSS Lyman-$\\alpha$ forest power spectrum and find that $\\sum m_{\\nu}<0.161$ eV ($95\\%$ C.L.). This result might lead to the exclusion of degenerate pattern of neutrino mass, when combining the results of neutrino oscillation experiments. Due to the degeneracy between the neutrino mass and EoS of dark energy, however in the presence of dynamics of dark energy, the constraints on $\\sum m_{\\nu}$ can be relaxed by a factor of $2$.  Finally, for the inflationary models, within the $\\Lambda$CDM framework, we find that the latest observational data prefer the inflation models with ``red\" tilt, namely, $n_s=0.961\\pm0.012$ ($1~\\sigma$) and small tensor fluctuations, $r<0.200$ ($95\\%$ C.L.). Because of the degeneracy between $r$ and EoS of dark energy, the upper limit on $r$ is relaxed to $r<0.268$ ($95\\%$ C.L.) and the parameter space in ($n_s$,$r$) panel are enlarged in the framework of dynamical dark energy models. Therefore, the inflationary model with HZ primordial spectrum ($n_s=1$, $r=0$), some hybrid models and some models with a ``blue\" tilt ($n_s>1$), which are excluded more than $2~\\sigma$ confidence level in the $\\Lambda$CDM model, will be consistent with the current observations now. Furthermore, in our analysis we do not find any significant evidence for the running of spectrum index."
    },
    "0807/0807.1085_arXiv.txt": {
        "abstract": "{The solar analogues are fundamental targets for a better understanding of our Sun and Solar System. Notwithstanding the efforts, this research is usually limited to field stars. The open cluster M67 offers a unique opportunity to search for solar analogues because its chemical composition and age are very similar to those of our star. In this work, we analyze FLAMES@VLT spectra of about one hundred of M67 main sequence stars with the aim to identify solar analogues. We first determine cluster members which are likely not binaries, by combining both proper motions and radial velocity measurements. Then, we concentrate our analysis on the determination of stellar effective temperature, using the analyzes of line-depth ratios and H$\\alpha$ wings. Finally, we also compute lithium abundance for all the stars. Thanks to the our analysis, we find ten solar analogues, which allow us to derive a solar color $(B-V)=0.649\\pm0.016$ and a cluster distance modulus of $9.63\\pm0.08$, very close to values found by previous authors. Among them, five are the best solar twins with temperature determinations within 60 K from the solar values. Our results lead us to do further spectroscopic investigations because the solar analogues candidates are suitable for planet search. ",
        "introduction": "The search for exo-planets has been mostly developed around field stars. Such stars present several advantages, for instance a wide range of stellar characteristics (mass, age, effective temperature, chemical composition, etc.), which allows us to study the dependence of planet formation on stellar parameters.  Another line of research is the specificity of our Sun and the opportunity to find solar stars hosting exo-planets. This search is best performed in open clusters, which, showing homogeneous age and chemical composition, common birth and early dynamical environment (\\citealt{Randich2005}), provide us an excellent laboratory for investigating the physics of planetary system formation. The old open cluster M67 is to this purpose a perfect target, having many solar-type stars and showing an age encompassing that of the Sun ($3.5-4.8$ Gyrs; \\citealt{Yadav2008}), a solar metallicity ([Fe/H]=$0.03\\pm0.02$, \\citealt{Randich2006}) and lithium depleted G stars (\\citealt{Pasquini1997}).  The present paper is the culmination  of a work, which involved the chemical determination of this cluster (\\citealt{Randich2006}), photometry and astrometry (\\citealt{Yadav2008}) to obtain membership, and FLAMES/GIRAFFE high resolution spectroscopy to clean this sample from binaries, and to look for the best solar analogues using the line-depth ratios method and the wings of the H$\\alpha$ line to determine accurate temperatures with respect to the Sun (\\citealt{Pasquini2008}). ",
        "conclusions": "By using spectroscopic observations performed with FLAMES/GIRAFFE at the VLT, we have found the best solar twins in M67 thanks to accurate determinations of radial velocity, lithium abundance, and effective temperature.  Thanks to our promising results, we plan to apply our method to other open clusters younger and older than the Sun in order to get informations about ``progenitors'' and ``descendants'' of our unique star.  Our method has proved to be suitable for the determination of both the solar color and the cluster distance modulus. It could therefore be applied for the same purpose to other open clusters."
    },
    "0807/0807.0495_arXiv.txt": {
        "abstract": "Time resolved spectra of $\\beta$\\,Cep yield the following average results: T$_{\\rm eff}$\\,=24000\\,$\\pm$\\,250~K, $\\log g$\\,=\\,3.91\\,$\\pm$\\,0.05 and $\\xi$\\,=\\,8.1\\,$\\pm$\\,0.9~km~s$^{-1}$. N, O, Ne, Al, Si and S abundances are solar while C, Mg and Fe are slightly under-abundant. Pulsational amplitudes of $\\Delta$T$_{\\rm eff}$\\,$\\sim$\\,700~K and $\\Delta$$\\log g$\\,$\\sim$\\,0.2 dex are found from H$\\beta$. The metal lines give similar amplitudes but centred on T$_{\\rm eff}$\\,$\\sim$\\,25000~K. An upper limit of 1.0 km s$^{-1}$ to the variability of the microturbulence is derived from the Si {\\sc iii} triplet at 455~nm. The radial velocity amplitude derived from the core of H$\\beta$ is $\\sim$\\,15$\\%$ greater than that from the metal lines. ",
        "introduction": "The capability of oscillations to probe the internal structure of stars is fully testified by solar studies. For this reason, pulsating stars, that are spread in the HR diagram, are expected to largely improve our knowledge of stellar evolution and physical assumptions in models.  \\citet{osaki71} has shown that both radial and non-radial modes of pulsating stars  are observationally distinguishable via the analysis of line profile variation (LPV). From this pioneering work, there has been an increasing interest in LPV modeling driven by the improvement in resolution and efficiency of modern spectrographs: \\citet{lee90} included the influence of stellar rotation, \\citet{briquet01} considered photospheric spots, \\citet{aerts94} performed a mode identification via the moments of the line profile or \\citet{telting97} via the Doppler imaging technique.  \\citet{dupret02} noted that a pure kinematic approach is not enough for LPV modeling and that a realistic representation of the strength and broadening of spectral lines were necessary. They considered the atmosphere of a pulsating star as a perturbed status of the hydrostatic equilibrium and represented the temperature perturbation of atmospheric layers with a series expansion in terms of effective temperature and gravity.  This operative method does not clearly represent the physics of $\\beta$\\,Cep stars, pulsation modes do not change the energy production and the very small (1\\%) stellar radius variation \\citep{aerts94} cannot modify neither the effective temperature nor the surface gravity in an observable way. However along a pulsation cycle, the compression and expansion of photospheric layers change the dependence of temperature and pressure on optical depth and let the effective temperature appear variable, with the pulsation period, if determined in a limited spectral interval. A phenomenon indeed observed in some $\\beta$ Cephei stars by \\citet{deridder02} on the basis of a photometric approach. These authors derived a T$_{\\rm eff}$ variability up to 1000 K with the pulsation period, from Geneva colours.  With the density and temperature, the variation of atomic level populations should also result in apparent surface gravity variations. In fact, there is an equivalence between stellar surface gravity and electron pressure in the framework of spectral line formation \\citep{Gray92}.  A further contribution to the strength and profile of spectral lines comes from the microturbulence ($\\xi$), that in principle could be modified along compression and expansion phases.  In this paper we analyse a set of time resolved spectra of $\\beta$ Cep (=~HD\\,205021~=~HR\\,8238), which is generally accepted to be the proto-type of a class of hot pulsating stars, with the aim of ascertaining the apparent variability of the effective temperature, gravity and microturbulence with the pulsation period. These variations would strongly constrain any LPV method, as the one proposed by Dupret and collaborators.   \\begin{table} \\caption{For each spectrum we report the Heliocentric Julian Date, the radial velocity measured from the Si {\\sc iii} triplet at 4550\\AA, the T$_{\\rm eff}$ and $\\log g$ determined through the  H$\\beta$ line.} \\label{vr} \\centering \\begin{tabular}{c r c c } \\hline \\hline HJD         & V$_{\\rm rad}$~~~~~ & T$_{\\rm eff}$ & $\\log g$ \\\\ (2554041.0+) & (km s$^{-1}$)~~ &     (K)       &                 \\\\ \\hline 0.2510 & $-$19.70\\,$\\pm$\\,0.97 & 24500\\,$\\pm$\\,310 & 3.95\\,$\\pm$\\,0.04  \\\\ 0.2587 & $-$17.86\\,$\\pm$\\,0.85 & 24150\\,$\\pm$\\,240 & 3.87\\,$\\pm$\\,0.05  \\\\ 0.2674 & $-$16.68\\,$\\pm$\\,0.40 & 24300\\,$\\pm$\\,240 & 3.89\\,$\\pm$\\,0.05  \\\\ 0.2746 & $-$13.67\\,$\\pm$\\,0.49 & 24250\\,$\\pm$\\,220 & 3.90\\,$\\pm$\\,0.05  \\\\ 0.2819 & $-$10.01\\,$\\pm$\\,0.35 & 24150\\,$\\pm$\\,220 & 3.88\\,$\\pm$\\,0.05  \\\\ 0.2892 &  $-$6.86\\,$\\pm$\\,0.28 & 24050\\,$\\pm$\\,210 & 3.88\\,$\\pm$\\,0.05  \\\\ 0.2964 &  $-$3.95\\,$\\pm$\\,0.40 & 23800\\,$\\pm$\\,200 & 3.81\\,$\\pm$\\,0.05  \\\\ 0.3040 &     0.76\\,$\\pm$\\,0.19 & 23700\\,$\\pm$\\,160 & 3.78\\,$\\pm$\\,0.04  \\\\ 0.3113 &     4.89\\,$\\pm$\\,0.30 & 23750\\,$\\pm$\\,170 & 3.78\\,$\\pm$\\,0.04  \\\\ 0.3186 &     8.18\\,$\\pm$\\,0.82 & 23700\\,$\\pm$\\,160 & 3.77\\,$\\pm$\\,0.04  \\\\ 0.3261 &    10.63\\,$\\pm$\\,0.55 & 23850\\,$\\pm$\\,210 & 3.81\\,$\\pm$\\,0.05  \\\\ 0.3334 &    11.97\\,$\\pm$\\,0.64 & 23850\\,$\\pm$\\,210 & 3.85\\,$\\pm$\\,0.06  \\\\ 0.3406 &    12.18\\,$\\pm$\\,0.26 & 23850\\,$\\pm$\\,220 & 3.84\\,$\\pm$\\,0.06  \\\\ 0.3515 &    10.70\\,$\\pm$\\,0.22 & 24150\\,$\\pm$\\,210 & 3.88\\,$\\pm$\\,0.05  \\\\ 0.3588 &     8.24\\,$\\pm$\\,0.50 & 24400\\,$\\pm$\\,220 & 3.92\\,$\\pm$\\,0.04  \\\\ 0.3660 &     4.95\\,$\\pm$\\,0.28 & 24200\\,$\\pm$\\,210 & 3.89\\,$\\pm$\\,0.05  \\\\ 0.3733 &     1.75\\,$\\pm$\\,1.01 & 24300\\,$\\pm$\\,240 & 3.90\\,$\\pm$\\,0.05  \\\\ 0.3806 &  $-$1.00\\,$\\pm$\\,0.99 & 24400\\,$\\pm$\\,260 & 3.92\\,$\\pm$\\,0.04  \\\\ 0.3878 &  $-$5.13\\,$\\pm$\\,1.42 & 24250\\,$\\pm$\\,250 & 3.89\\,$\\pm$\\,0.06  \\\\ 0.3955 &  $-$8.74\\,$\\pm$\\,2.13 & 24200\\,$\\pm$\\,250 & 3.86\\,$\\pm$\\,0.06  \\\\ 0.4028 & $-$12.14\\,$\\pm$\\,1.04 & 24200\\,$\\pm$\\,250 & 3.89\\,$\\pm$\\,0.06  \\\\ 0.4100 & $-$15.52\\,$\\pm$\\,0.71 & 24350\\,$\\pm$\\,250 & 3.91\\,$\\pm$\\,0.06  \\\\ 0.4173 & $-$18.43\\,$\\pm$\\,0.29 & 24350\\,$\\pm$\\,240 & 3.90\\,$\\pm$\\,0.05  \\\\ 0.4246 & $-$20.11\\,$\\pm$\\,0.94 & 24350\\,$\\pm$\\,210 & 3.92\\,$\\pm$\\,0.04  \\\\ 0.4318 & $-$20.13\\,$\\pm$\\,0.88 & 24300\\,$\\pm$\\,210 & 3.90\\,$\\pm$\\,0.05  \\\\ 0.4475 & $-$19.34\\,$\\pm$\\,1.26 & 24050\\,$\\pm$\\,210 & 3.87\\,$\\pm$\\,0.05  \\\\ 0.4548 & $-$15.45\\,$\\pm$\\,2.05 & 24050\\,$\\pm$\\,230 & 3.87\\,$\\pm$\\,0.05  \\\\ 0.4620 & $-$14.02\\,$\\pm$\\,1.02 & 23850\\,$\\pm$\\,210 & 3.83\\,$\\pm$\\,0.05  \\\\ 0.4693 & $-$11.86\\,$\\pm$\\,1.57 & 23900\\,$\\pm$\\,210 & 3.84\\,$\\pm$\\,0.05  \\\\ 0.4766 &  $-$8.07\\,$\\pm$\\,0.89 & 23800\\,$\\pm$\\,190 & 3.81\\,$\\pm$\\,0.05  \\\\ 0.4838 &  $-$4.53\\,$\\pm$\\,0.92 & 23800\\,$\\pm$\\,190 & 3.81\\,$\\pm$\\,0.05  \\\\ \\hline \\end{tabular} \\end{table} ",
        "conclusions": "With the aim of determining the variations of the effective temperature, surface gravity and microturbulence velocity that occur in hot pulsating stars during the pulsational cycle, we performed time resolved R\\,=\\,20000 spectroscopic observations of the prototype $\\beta$~Cep along a complete period coverage.  As a first step, we confirm the T$_{\\rm eff}$ and $\\log g$ found in a previous paper by \\citet{catanzaro08} and we measure $\\xi$\\,=\\,8.1\\,$\\pm$\\,0.9~km~s$^{-1}$, from a sample of single lines of O{\\sc ii}. With these values, the abundances resulted to be solar for nitrogen, oxygen, neon, aluminium, silicon and sulfur, while slight under-abundances have been inferred for carbon, magnesium and iron.  We find that along a pulsation cycle, after a correction for the pulsational Doppler shift, the H$\\beta$ line profile changes are matched assuming a variability of the effective temperature in the range of 23700 to 24400 K and a variability of gravity in the range of 3.76 to 3.93. The amplitude of $\\approx$~0.08~dex measured in the gravity variation could lead the reader to the erroneous belief that it corresponds to a radius variation larger than the one computed by \\citet{aerts94}, via cinematic considerations. However, the variability of $\\log g$ here determined for $\\beta$~Cep is not a consequence of changes in the surface gravity because of expansions and compressions of the star as a whole. We ascribed the $\\log g$ variability to a change of the electron pressure \\citep{Gray92} induced by temperature, that modulates the ionisation state of metals. Thus, it should not be linked directly to the radius of the star.  Metal lines are characterised by equivalent widths variable with the pulsation cycle. These are explained by a variable effective temperature between 25000 and 26000 K as a consequence of the changed ionisation state. It could be that the most extended formation region of the H$\\beta$ line is at the origin of the lower temperature, in agreement with the 15\\% larger amplitude presented by the radial velocity measured through the H$\\beta$ core.  \\begin{figure} \\centering \\includegraphics[width=8cm]{ew_4lines.ps} \\caption{Equivalent widths as a function of effective temperature.} \\label{ew_T} \\end{figure}  As to the possibility that the microturbulence depends on the compression and expansion of the atmospheric layers, we find that it is constant within a 1.0 km~s$^{-1}$.  We conclude that pulsations modify the dependence of the temperature and pressure on the optical depth in the atmosphere of $\\beta$~Cep. This mimics a variability of the effective temperature and gravity when a small range of wavelengths is selected.  We hope that this information can help in modeling the variability of spectral lines that is not only due to the pulsation velocity field \\citep{dupret02}.  Along the expansion phase, the negative radial velocity is expected to increase up to a null value just in coincidence with the maximum of the stellar radius. At this phase of minimum compression, the photosphere is expected at the temperature minimum. Differently, we observe a delay in the temperature minimum about one tenth of the pulsation period. According to \\citet{deridder02}, this extra-phase shift is due to the non-adiabatic condition that occurs in the  pulsating layers."
    },
    "0807/0807.0176_arXiv.txt": {
        "abstract": "We present  new results for the molecular gas, dust emission and the ionized gas in J1023+1952, an \\hi\\ rich intergalactic star-forming tidal dwarf galaxy candidate. It is  located at the projected intersection of two faint stellar tidal streams wrapped around the interacting pair of galaxies NGC~3227/6 (Arp~94). Using the IRAM 30m telescope, emission from \\co\\  and \\CO\\  was detected across the entire extent of the neutral hydrogen cloud associated with J1023+1952, a region of the size of 8.9$\\times$5.9~kpc,  as well as in the nuclear region and outer disk of NGC~3227. The molecular gas is found to be abundant over the entire \\hi\\ cloud, with H$_2$-to-\\hi\\ gas mass ratios between 0.5 and 1.7. New \\spitzer\\ mid-infrared observations at 3.6, 4.5, 5.8, 8.0, 15 and 24 \\mi\\  show that young SF  is restricted to the southern part of the cloud. Despite the relatively uniform H$_2$ and \\hi\\ column density across the cloud, young SF occurs only in those regions where the velocity dispersion in the CO and \\hi\\ is a factor of $\\sim$two lower (FWHM of 30 - 70 \\kms) than elsewhere in the cloud (FWHM of 80 - 120 \\kms). Thus the kinematics of the gas, in addition to its column density, seems to be a crucial factor in triggering SF. Optical/infrared spectral energy distributions (SEDs) and \\halpha\\ photometry confirm that all the knots  are young, % with a tentative age sequence from the south-west (oldest knots) to the north-east (youngest). Optical spectroscopy of the brightest SF region allowed us to determine the metallicity (12+log(O/H) =  8.6 $\\pm$ 0.2) and the extinction ($A_B$=2.4). This shows that J1023+1952 is made from metal-enriched gas which is inconsistent with the hypothesis that it represents a pre-existing dwarf galaxy. Instead, it must be formed from recycled, metal-enrichd gas, expelled from NGC~3227 or NGC~3226 in a previous phase of the interaction. ",
        "introduction": "Merger-driven galaxy evolution is  important on a wide range of scales at different cosmic epochs and in different environments. The observed decline in the global star formation rate (SFR) at redshifts z$<$2 and the increased dominance of small galaxies in the local Universe has been suggested to imply interaction-driven 'downsizing' of galactic structure and star formation (Cowie et al. 1996, Madau et al. 1996, Brinchmann et al.  1998, Treu et al.  2005, Thomas et al. 2005) in contrast to hierarchical galaxy formation. The recently-discovered free-floating intergalactic and extraplanar \\hii regions around nearby gas-rich galaxies (Gerhard et al. 2002, Ryan-Weber et al. 2004) and H$\\alpha$-emitting star-forming knots in galaxy clusters (Sakai et al. 2002) are the smallest intergalactic structures, and may represent an important evolutionary link or even a new mode of star formation (SF)  in the low-density environment of the far outskirts of the galaxy halo. Their young star-formation ages suggest they are tidal in origin, forming from gaseous tidal debris in intergalactic space.   The largest \\hii regions are typically found at the apparent tip of tidal tails, at more than 100 kpc from the parent galaxies,  in maxima of the \\hi\\ column density where we also detect large quantities of molecular gas (Braine et al. 2000, 2001, Lisenfeld et al. 2002) and where enough tidal material is often available to build a dwarf galaxy (Tidal Dwarf Galaxies, hereafter TDGs). Indeed, tidal interactions between approximately equal-mass gas-rich galaxies have long been known to result in dramatic tidal damage, producing bridges and tails of debris extending out to several galaxy diameters (e.g. Toomre \\& Toomre 1972, Haynes et al. 1984, Barnes \\& Hernquist 1992, Mihos 2001, Mundell et al. 1995).  Tidal tails often comprise primarily neutral hydrogen since \\hi\\ is usually the most spatially extended component of a spiral galaxy disk, and is therefore most easily stripped during the earliest stages of a tidal encounter (e.g. Hibbard \\& van Gorkom 1996). Molecular gas is generally  observed to be more centrally concentrated than \\hi\\ in disks of galaxies outside clusters (Nishiyama, Nakai \\& Kuno 2001) so it is less likley to be stripped on first passage.  Here we present a study of an extragalactic region with active SF  embedded in an extensive gas cloud named J1023+1952. The region is associated with the interacting Seyfert system Arp~94 (Mundell et al. 1995; Mundell et al. 2004). Optically, the system is dominated by two disturbed Seyfert galaxies: NGC~3227, an SAB(s) pec barred spiral Seyfert galaxy and its elliptical companion, NGC3226 (E2 pec) (Rubin \\& Ford 1968, Mundell et al. 2004).  Neutral hydrogen (\\hi) imaging of the system (Mundell et al. 1995) revealed two gaseous tidal tails extending $\\sim$100~kpc north and south of NGC~3227, well-ordered gas in the disk of NGC~3227 and a massive \\hi\\ cloud (M$_{\\rm \\hi}$~=~3.8~$\\times$~10$^8$~M$_{\\odot}$), that lies at the base of the northern tail and is close to, but physically and kinematically distinct from the disk of NGC~3227. Mundell et al. (1995) suggested that this cloud (hereafter J1023+1952) might be a dwarf galaxy that is either pre-existing and being accreted by Arp~94, or a newly-created tidal dwarf galaxy (TDG) forming from the tidal debris.  The subsequent discovery of a region of very blue star-forming knots with luminosities between $10^6$ and $5\\times 10^6$ \\lsol, embedded in a high \\hi\\ column density ridge in the southern half of J1023+1952 (Mundell et al. 2004) confirmed its classification as a dwarf galaxy; the inferred SFR, UV luminosities and H$\\alpha$ equivalent widths of the knots suggested a starburst age less than 10~Myr. Near infrared imaging provided further evidence for the youth of J1023+1952; no additional embedded star formation or old stellar population were found (Mundell et al. 2004). These results suggest that the observed SF is currently being triggered in the interstellar medium (ISM) of J1023+1952 and does not merely represent young stars tidally stripped from the disk of NGC3227 (Mundell et al. 2004).   The Arp 94 interacting system offers an interesting nearby laboratory in which to study the interplay between SF and the physics of the multiphase ISM in a tidally disrupted system with young, active intergalactic SF regions. Here we present a study of the optical, infrared and millimetre properties of the Arp 94 system that aims to determine the nature and origin of J1023+1952, its role in the observed interaction and the underlying reasons for its observed star formation properties. Our dataset comprises (a) \\co\\ and \\CO\\ mapping of the full extent of J1023+1923, the nucleus of NGC3227 and the outer edges of the galactic disk of NGC3227; (b) multi-band infrared imaging with the {\\it Spitzer Space Telescope}  of J1023+1952 with particular emphasis on investigating the known SF regions and searching for further embedded SF; (c) optical spectroscopy of the brightest \\hii region to determine the gas-phase metallicity.  The paper is structured as follows: the observations and data reduction are described in Section 2, molecular and atomic gas and star formation properties are derived in Section 3 and the implications for the triggering of SF and  the origin of J1023+1952 as well as its role in the system are discussed in Section 4. We conclude in Section 5. Assuming a distance of 20.4 Mpc to the Arp~94 system (Tully 1988), 1\\arcsec corresponds to  $\\sim$100~pc. ",
        "conclusions": "We presented a comprehensive study of the ISM of J1023+1952, based on new millimeter  and infrared data, with the goal of determining its nature and origin. The main results and conclusions are:  1) \\co\\ and \\CO\\ emission was detected for the first time and shown to be extended across the entire \\hi-emitting  cloud. The intensity-weighted  velocities and the line widths of the CO and \\hi\\ emission agree well.  2) The optical spectroscopy provided measures of the extinction and metallicity of the brightest SF knot (knot 1). We derived an extinction of 1.4 mag for the \\halpha\\ line (corresponding to $A_B =2.4$) and a metallicity of 12+log(O/H) = 8.6 $\\pm$ 0.2 showing that the gas is pre-enriched.  3) We identified seven SF knots in optical and {\\it Spitzer} images, and presented the results of aperture photometry. The SEDs of these objects were indicative of young star formation (ages between 1 and 10-20 Myr), although they showed surprising differences. We identified a tentative age gradient from the south-east, where the two knots (knot 4 and 5) had the reddest color, to the north-west (knots 1 and 1a), where a very young stellar population was present. One  knot (knot 5) had \\halpha\\ emission with no associated 8\\mi\\ emission. Knot 2 has  the strongest blue emission but showed no \\halpha\\  and no dust emission in spite of its blue color. A possible reason might be the lack of dust which would lead to an overestimate of the extinction and thus to a too blue color.   4) We studied the properties of the 8\\mi\\ and 24\\mi\\ emission in comparison with the extinction corrected \\halpha\\ emission, in order to investigate the dust properties and the use of these two emissions as SF tracers. Neither the 8-to-\\halpha\\ ratio nor the 24-to-\\halpha\\ ratio of most knots showed a  significant difference in comparsion to \\hii\\ regions in M81 and M51. Only for the combined emission of knots 1+1a we found a slight excess of \\halpha\\ with respect to the 24\\mi\\ luminosity.  5) We discussed various possible scenarios for the origin of J1023+1952. A strong constraint is the high metallicity, which is supported by the abundant CO emission, much higher than expected from its luminosity if J1023+1952 were a preexisting dwarf galaxy. Based on this fact, we can exclude a chance superposition of a dwarf galaxy or an infalling, tidally disrupted  small galaxy  as the origin of J1023+1952. Instead, we conclude that J1023+1952 must have   formed from recycled gas from the parent galaxies."
    },
    "0807/0807.2459_arXiv.txt": {
        "abstract": "The next-to-leading order (NLO) QCD corrections are calculated for the pair-annihilation of spin-1 dark matter (DM) by dimensionally regularizing both ultraviolet and infrared singularities in non-relativistic limit ($v\\ll 1$). The complete $O(\\alpha_{s})$ correction is about $8\\%$ due to the massless gluon contribution. An extra $5\\%$ will be added if there is a new interaction from a massive gluon of approximately same mass as the DM particle. The NLO QCD correction could give the sizable shift to the DM mass constrained by relic density measurements. ",
        "introduction": "The spin-1 dark matter candidate $B^{\\prime}$, which is the $Z_{2}$-odd partner of the hypercharge gauge boson $B$, appears in interesting models like universal extra dimensions (UED) \\cite{tappel} and Little Higgs (LH) \\cite{narkani}, motivated by solving the gauge hierarchy problem. The phenomenology of spin-1 DM ($B^{\\prime}$) has been well studied at the tree level. With advances in precision measurements in cosmology and astrophysics as well as the advent of LHC era, the NLO calculation of the $B^{\\prime }B^{\\prime}$ annihilation becomes timely for accurate analysis. We present the next-to-leading order (NLO) QCD corrections for the pair-annihilation of spin-1 bosonic dark matter by dimensionally regularizing both ultraviolet (UV) and infrared (IR) singularities in the non-relativistic limit ($v\\ll1$). Our analysis of the process ($B^{\\prime}B^{\\prime}\\rightarrow q\\overline{q}$) is applicable to general spin-1 DM, as in UED, LH, etc. We assume that the quark field $q$ interacts with its $Z_{2}$-odd partner $\\widetilde{q}$ in the form.%  \\begin{equation} \\mathcal{L}\\supset-g_{Y}\\overline{\\widetilde{q}}\\gamma^{\\mu}(\\widetilde{Y}% _{L}P_{L}+\\widetilde{Y}_{R}P_{R})qB_{\\mu}^{\\prime}+h.c. \\end{equation}   The coupling $g_{Y}$ and the number $\\widetilde{Y}_{L}$ in the UED would be the usual hypercharge coupling and the corresponding quantum number,%  \\[ g_{Y}=\\frac{e}{\\cos\\theta_{W}},\\text{ \\ }\\widetilde{Y}_{L}(u,d)=\\frac{1}% {6},\\text{ \\ }\\widetilde{Y}_{R}(u)=\\frac{2}{3},\\text{ \\ }\\widetilde{Y}% _{R}(d)=-\\frac{1}{3}% \\]   However in the Littlest Higgs model (L$^{2}$H) \\cite{narkani}, $\\widetilde {Y}_{L}(u,d)=\\frac{1}{10}$ and otherwise $\\widetilde{Y}_{R}=0$.  We also assume that the mass $\\widetilde{M}$ of $\\widetilde{q}$ is not too far above the mass $M$ of the spin-1 DM particle $B^{\\prime}$. The approximate relation $\\widetilde{M}\\approx M$ is valid in UED, but a choice of parameters in other models is required. However, such a choice allows us to carry through analytical calculation and give simple results.  Fig.1 shows the Born diagrams for $B^{\\prime}(p_{1})B^{\\prime}(p_{1}^{\\prime })\\rightarrow q(p_{2})\\overline{q}(p_{2}^{\\prime})$. The exchange of the $Z_{2}$-odd quark line is bold-faced. The amplitude for the Born diagrams is given by%  \\begin{equation} \\mathcal{M}_{B}=-g_{Y}^{2}\\widetilde{Y}_{L}^{2}\\overline{u}(p_{2})\\left( \\gamma^{\\mu}\\frac{{\\not p  }_{2}-{\\not p  }_{1}+\\widetilde{M}}{t-\\widetilde {M}^{2}}\\gamma^{\\nu}+\\gamma^{\\nu}\\frac{{\\not p  }_{2}-{\\not p  }_{1}^{\\prime }+\\widetilde{M}}{u-\\widetilde{M}^{2}}\\gamma^{\\mu}\\right)  P_{L}v(p_{2}% ^{\\prime})\\epsilon_{\\mu}(p_{1})\\epsilon_{\\nu}(p_{1}^{\\prime})\\text{\\ }% +\\text{RH}% \\end{equation}  \\begin{figure} [ptb] \\begin{center} \\includegraphics[ trim=0.000000in 0.000000in -0.080171in 0.000000in, height=2.179cm, width=5.0215cm ]% {fig1.eps}% \\caption{The Born diagrams for $B^{\\prime}B^{\\prime}\\rightarrow q\\overline{q}% $. The bold wavy lines represent the DM candidate gauge bosons ($B^{\\prime}$), the bold solid lines represent the heavy quarks (the $Z_{2}$-odd partners of the outgoing quarks) and the light solid lines correspond to the quarks.}% \\label{fig1}% \\end{center} \\end{figure}   In the extremely non-relativistic limit ($v\\simeq0$), $p_{1}\\simeq p_{1}^{\\prime}\\simeq(p_{2}+p_{2}^{\\prime})/2\\simeq(M,\\mathbf{0})$. \\ The invariant Mandelstam parameters are%  \\begin{subequations} \\begin{align} s  &  =(p_{1}+p_{1}^{\\prime})^{2}=(p_{2}+p_{2}^{\\prime})^{2}\\simeq4M^{2}\\\\ t  &  =(-p_{1}+p_{2})^{2}=(p_{1}^{\\prime}-p_{2}^{\\prime})^{2}\\simeq-M^{2}\\\\ u  &  =(-p_{1}^{\\prime}+p_{2})^{2}=(p_{1}-p_{2}^{\\prime})^{2}\\simeq-M^{2}% \\end{align}   The formula reduces to%  \\end{subequations} \\begin{equation} \\mathcal{M}_{B}=g_{Y}^{2}\\widetilde{Y}_{L}^{2}\\overline{u}(p_{2})\\frac {({p}_{2}-{p}_{2}^{\\prime})^{\\mu}\\gamma^{\\nu}+({p}_{2}-{p}_{2}^{\\prime})^{\\nu }\\gamma^{\\mu}}{M^{2}+\\widetilde{M}^{2}}P_{L}v(p_{2}^{\\prime})\\epsilon_{\\mu }(p_{1})\\epsilon_{\\nu}(p_{1}^{\\prime})\\ +\\text{RH}% \\end{equation}   The consistent annihilation rate in $d=4-2\\epsilon$ dimensions is%  \\begin{equation} \\sigma_{B}v=\\frac{(1-\\epsilon)\\Gamma(1-\\epsilon)}{\\Gamma(2-2\\epsilon)}\\left( \\frac{4\\pi\\mu^{2}}{4M^{2}}\\right)  ^{\\epsilon}\\frac{2g_{Y}^{4}(\\widetilde {Y}_{L}^{4}+\\widetilde{Y}_{R}^{4})N_{c}}{9\\pi}\\frac{M^{2}}{\\left( M^{2}+\\widetilde{M}^{2}\\right)  ^{2}}% \\end{equation}   where $N_{c}$ is number of colors, $v$ denotes the relative velocity between spin-1 dark matter candidate pair, $\\mu$ is an arbitrary mass scale and $\\sigma_{B}$ is the Born cross section$.$ ",
        "conclusions": "The next-to-leading order (NLO) QCD corrections are calculated for pair-annihilation of spin-1 bosonic dark matter by dimensionally regularizing both ultraviolet and infrared singularities in the non-relativistic limit. The order $\\alpha_{s}$ correction amounts to about $8\\%$ and can enhance to $13\\%,$ when including heavy gluon. The NLO QCD correction could give the sizable shift to the DM mass constrained by relic density measurements."
    },
    "0807/0807.2035_arXiv.txt": {
        "abstract": "As part of a broader effort to characterize the population of star-forming galaxies in the local universe, we have carried out an H$\\alpha$+[NII] imaging survey for an essentially volume-limited sample of galaxies within 11 Mpc of the Milky Way. This first paper describes the design of the survey, the observation, data processing, and calibration procedures, and the characteristics of the galaxy sample.  The main product of the paper is a catalog of integrated H$\\alpha$ fluxes, luminosities, and equivalent widths for the galaxies in the sample.  We briefly discuss the completeness properties of the survey and compare the distribution of the sample and its star formation properties to other large H$\\alpha$ imaging surveys. These data form the foundation for a series of follow-up studies of the star formation properties of the local volume, and the properties and duty cycles of star formation bursts in dwarf galaxies. ",
        "introduction": "Much of our accumulated knowledge of the star formation properties of nearby galaxies has been derived from the imaging, photometry, and spectroscopy of galaxies in the H$\\alpha$ nebular emission line (e.g., Kennicutt 1998a and references therein).  CCD imaging observations have proven to be especially valuable because they provide not only integrated measurements of star formation rates (SFRs), but also information on the spatial distribution and the luminosity functions of individual star forming regions, with spatial resolutions that are not easily matched at other wavelengths such as the ultraviolet or far-infrared.  The power of this technique is evidenced by the large number of published H$\\alpha$ observations of galaxies, now comprising more than 100 papers and nearly 3000 galaxies measured (Kennicutt et al. 2008, in preparation).  However this combined body of H$\\alpha$ observations is heterogeneous, with strong biases favoring luminous, face-on, and high-surface brightness galaxies.  These provide representative sampling of a diverse galaxy population, but do not constitute a statistically complete star formation inventory of the local universe.   The most complete inventories of local star-forming galaxies have come from spectroscopic surveys. Objective prism surveys such as the Universidad Complutense de Madrid (UCM) survey (Zamorano et al. 1994, 1996; Alonso et al. 1999), the Kitt Peak International Spectroscopic Survey (KISS; Salzer et al. 2000, 2001; Gronwall et al. 2004, Jangren et al. 2005), and the CIDA--UCM--Yale survey (Bongiovanni et al. 2005) provide H$\\alpha$-based SFRs for redshift-limited samples over areas of order $10^2 - 10^3$ deg$^2$ and survey volumes of order $10^5 - 10^6$ Mpc$^3$.  These offer emission line selected complete sampling of the most luminous and highest equivalent width (EW) H$\\alpha$ emitters, and have yielded some of the most widely-applied measurements of the local comological SFR density (e.g., Gallego et al. 1995).  Although such prism surveys have well-defined sensitivity and detection properties across their survey volumes, they become severely incomplete for galaxies with spatially extended low surface brightness H$\\alpha$ emission, and for galaxies with strong superimposed continuua (i.e., low H$\\alpha$ EWs).  As a result as much as half of the aggregate star formation can be missed in these surveys (Sakai et al. 2004).  These problems can be reduced by undertaking a comprehensive spectroscopic survey that uniformly targets galaxies within a given magnitude and redshift range, such as the Sloan Digital Sky Survey (SDSS), and SDSS data have produced the most comprehensive characterization of the low-redshift star forming galaxy population to date (Brinchmann et al. 2004).  This approach has unmatched statistical power, with spectroscopically measured H$\\alpha$ fluxes now available for hundreds of thousands of galaxies, but the data still suffer from incompleteness introduced by aperture undersampling, and the catalogs are heavily weighted to luminous spirals.  Another, complementary approach to this general problem is to obtain deep H$\\alpha$ imaging for large and complete (or at least robustly defined) samples of galaxies.  A number of ongoing surveys are exploiting this approach.  These include two surveys of rich galaxy clusters, the GOLDmine survey of the Virgo cluster and Coma superclusters (Gavazzi et al. 2003), and the MOSAIC H$\\alpha$ survey, which has produced an H$\\alpha$-selected sample of star-forming galaxies in 8 $z < 0.03$ Abell clusters (Sakai et al. 2008, in preparation). Other surveys of (predominantly) field galaxies include the H$\\alpha$ Galaxy Survey (H$\\alpha$GS; James et al. 2004), which is based on a diameter and radial velocity selected sample of 334 $z \\le 0.01$ galaxies taken from the Uppsala Galaxy Catalog (Nilson 1973), the Survey for Ionization in Neutral-Gas Galaxies (SINGG; Meurer et al. 2006), which is based on an HI-selected subsample of 468 $z \\le 0.0423$ galaxies from the HI Parkes All Sky Survey (HIPASS; Barnes et al. 2001), the Analysis of the Interstellar Medium of Isolated Galaxies (AMIGA; Verdes-Montenegro et al. 2005), which is observing a sample of 206 members selected from the Catalog of Isolated Galaxies (Karachentseva 1973), and a series of surveys of local irregular galaxies by Hunter and collaborators, culminating in a survey of 140 galaxies by Hunter and Elmegreen (2004, hereafter denoted as HE04).  The most direct approach to assembling a truly complete inventory of local star-forming galaxies would be to image {\\it every} galaxy within a fixed distance of the Milky Way.  Such a perfect census would be free of selection biases and would provide a full characterization of the star-forming population, missing only those objects that are too rare to appear in a small local volume.  A survey of this kind would provide especially powerful constraints on the star formation and evolutionary properties of low-mass dwarf galaxies, which will dominate the numbers in any complete volume-limited survey.  Of course such an idealized survey cannot yet be undertaken, because existing galaxy catalogs are incomplete for dwarf galaxies even at small distances (e.g., Belokurov et al. 2007, Irwin et al. 2007, and references therein). Within the Galactic zone of avoidance this incompleteness increases and extends to brighter luminosities.  Nevertheless one can achieve many of the scientific advantages of a volume-complete survey by imaging the known galaxies in the local volume.  In this spirit we have undertaken to obtain and compile H$\\alpha$ fluxes (and in most cases imaging) for most of the 261 known spiral and irregular galaxies within 11 Mpc of the Milky Way, above a Galactic latitude $|b| = 20\\arcdeg$, and with $B \\le 15$, and for another 175 galaxies at lower latitudes and/or fainter magnitudes within this volume.  Within these limits the resulting sample is virtually complete for M$_B \\sim$ $-$15.5 and M(HI) $\\sim 10^8$~M$_\\odot$ (see \\S5), and is statistically correctable for incompleteness to luminosities and masses roughly an order of magnitude lower. Another H$\\alpha$ imaging survey of galaxies within the local 10 Mpc volume is being carried out by Karachentsev and collaborators (Kaisin \\& Karachentsev 2007 and references therein), and published data from that project have been used to supplement our own observations.  Our resulting dataset provides a statistically rigorous sample for characterizing the distributions of absolute and mass-normalized star formation rates and HII region populations in the local volume.  The sample also allows us to study the temporal variation of star formation, because in a volume-limited sample the relative numbers of galaxies observed in various phases of a starbust cycle (for example), will scale with the relative durations of those phases in the cycle. The sample also forms the foundation for two subsequent Legacy surveys on the {\\sl Galaxy Evolution Explorer (GALEX)} (Martin et al. 2005) and the {\\it Spitzer Space Telescope} (Werner et al. 2004). The 11 Mpc Ultraviolet and H$\\alpha$ Survey (11HUGS) is aimed at obtaining deep ultraviolet imaging for most of the galaxies in the complete ($|b| = 20\\arcdeg$, $B \\le 15$) H$\\alpha$ sample.  The {\\it Spitzer} Local Volume Legacy project (LVL) is aimed at obtaining infrared imaging at 3.6 -- 160 $\\mu$m for the sample.  The combination of H$\\alpha$, ultraviolet, and infrared imaging of the galaxies will provide extinction-corrected SFRs, and information on the time dependence and spatial migration of the massive star formation over timescales of $\\sim10^7 - 10^8$ yr.  This is the data paper for a series of papers that will use the H$\\alpha$ observations of this sample to characterize the star formation properties of the local galaxy population (Lee et al. 2007), constrain the temporal behavior of star formation in dwarf galaxies (Lee 2006; Lee \\& Kennicutt 2008, in preparation), and develop quantititative metrics of galactic star formation that can be applied to more distant galaxy samples. In this paper we describe the design, construction, and execution of the survey, and present the complete H$\\alpha$ integrated flux and equivalent width (EW) catalog.  We also discuss the statistical properties of the sample, in the context of other local narrowband imaging surveys.  The remainder of the paper is organized as follows.  We first describe the design and construction of the sample (\\S 2).  In \\S 3 we describe the observations, reduction, and calibration procedures.  The H$\\alpha$ flux and emission-line equivalent width catalog is presented in \\S4, and we briefly describe the properties of the galaxy sample and the dataset in \\S5. We conclude the main paper in \\S6 with a brief description of the ancillary ultraviolet and infrared observations now being obtained for the sample, and the multi-wavelength public archive that will include the imaging data presented in this paper. Appendices provide details on the emission-line flux calibrations and corrections for [NII] emission. ",
        "conclusions": ""
    },
    "0807/0807.1200_arXiv.txt": {
        "abstract": "s{After reviewing the modified Newtonian dynamics (MOND) proposal, we advocate that the associated phenomenology may actually not result from a modification of Newtonian gravity, but from a mechanism of ``gravitational polarization'' of some dipolar medium playing the role of dark matter. We then build a relativistic model within standard general relativity to describe (at some phenomenological level) the dipolar dark matter polarizable in a gravitational field. The model naturally involves a cosmological constant, and is shown to reduce to the concordance cosmological scenario ($\\Lambda$-CDM) at early cosmological times. From the mechanism of gravitational polarization, we recover the phenomenology of MOND in a typical galaxy at low redshift. Furthermore, we show that the cosmological constant $\\Lambda$ scales like $a_0^2$, where $a_0$ is the constant MOND acceleration scale, in good agreement with observations.} ",
        "introduction": "The mysteries of the nature of dark matter and dark energy are perhaps the most important ones of contemporary cosmology. Dark matter, which accounts for the observed discrepancy between the dynamical and luminous masses of bounded astrophysical systems, is usually formulated within the so-called \\textit{particle dark matter} approach, in which the dark matter consists of unknown non-baryonic particles, e.g. neutralinos as predicted by super-symmetric extensions of the standard model of particle physics (see~\\cite{Be.al.05} for a review). Furthermore, the dark matter triggers the formation of large-scale structures by gravitational collapse and explains the distribution of baryonic matter from galaxy cluster scales up to cosmological scales by the non-linear growth of initial perturbations. Simulations suggest some universal dark matter density profile around distributions of ordinary baryonic matter~\\cite{Na.al.97}. An important characteristic of dark matter, required by the necessity of clustering matter on small scales, is that it should be cold (or non-relativistic) at the epoch of galaxy formation. Together with the observational evidence of dark energy (presumably a cosmological constant $\\Lambda$) measured from the Hubble diagram of supernovas, the particle dark matter hypothesis has yielded the successful concordance model of cosmology called $\\Lambda$-CDM, which reproduces extremely well the observed cosmic microwave background spectrum~\\cite{HuD02}.  However, despite these successes at cosmological scales, the particle dark matter approach has some difficulties~\\cite{McBl.98} at explaining in a natural way the flat rotation curves of galaxies, one of the most persuasive evidence for the existence of dark matter, and the Tully-Fisher empirical relation between the observed luminosity and the asymptotic rotation velocity of spiral galaxies. In order to deal with these difficulties, all linked with the properties of dark matter at galactic scales, an alternative paradigm has emerged in the name of the modified Newtonian dynamics (MOND)~\\cite{Mi1.83,Mi2.83,Mi3.83}. Although MOND in its original formulation cannot be considered as a viable physical model, it is now generally admitted that it does capture in a very simple and powerful ``phenomenological recipe'' a large number of observational facts, that any pertinent model of dark matter should explain.  It is frustrating that the two alternatives $\\Lambda$-CDM and MOND, which are successful in complementary domains of validity (say the cosmological scale for $\\Lambda$-CDM and the galactic scale for MOND), seem to be fundamentally incompatible. In the present paper, we shall propose a different approach, together with a new interpretation of the phenomenology of MOND, which has the potential of bringing together $\\Lambda$-CDM and MOND into a single unifying relativistic model for dark matter and dark energy. This relativistic model will be shown to benefit from both the successes of $\\Lambda$-CDM at cosmological scales, and MOND at galactic scales. ",
        "conclusions": ""
    },
    "0807/0807.3802_arXiv.txt": {
        "abstract": "We calculate the reionization history for different models of the stellar population and explore the effects of primordial magnetic fields, dark matter decay and dark matter annihilation on reionization. We find that stellar populations based on a Scalo-type initial mass function for Population II stars can be ruled out as sole sources for reionization, unless star formation efficiencies of more than $10\\%$ or very high photon escape fractions from the parental halo are adopted. When considering primordial magnetic fields, we find that the additional heat injection from ambipolar diffusion and decaying MHD turbulence has significant impact on the thermal evolution and the ionization history of the post-recombination universe and on structure formation. The magnetic Jeans mass changes the typical mass scale of the star forming halos,  and depending on the adopted stellar model we derive upper limits to the magnetic field strength between  $0.7$ and $5\\,$nG (comoving). For dark matter annihilation, we find an upper limit to the thermally averaged mass-weighted cross section of $\\langle\\sigma v \\rangle/m_{DM}\\le10^{-33}\\ \\mathrm{cm}^3\\mathrm{/s/eV}$. For dark matter decay, our calculations yield a lower limit to the lifetime of dark matter particles of $\\tau\\ge3\\times10^{23}\\,$s. These limits are in agreement with constraints from recombination and provide an independent confirmation at a much later epoch. ",
        "introduction": "Over the last decades, our understanding of cosmological reionization has improved considerably. Observations of high-redshift quasars clearly indicate that reionization must end around $z\\sim6$ \\citep{Becker}, and the WMAP 5-year-measurement finds a reionization optical depth of $\\tau=0.087\\pm0.017$ \\citep{Komatsu, WMAPAngular}, yielding clear evidence that reionization started significantly earlier and is thus a continuous process. Reionization through miniquasars would produce a soft-X-ray background that is significantly higher than the observed background, and can thus be ruled out \\citep{Dijkstra, Salvaterra}. Simulations of stellar reionization further show that the process is highly inhomogeneous and based on the growth and merging of ionized bubbles around the first stellar sources \\citep{Gnedin, Ciardi2, Kohler}. Understanding the essential physical processes, it is possible to build semi-analytic models that describe stellar reionization, which can be tested over a wide parameter space \\citep{Shapiro, Haiman, Barkana, Loeb, Choudhury, Schneider}. Additional constraints on the cosmic star formation rate are now available from gamma-ray burst observations \\citep{Yuksel}. The reionization framework thus provides an increasingly reliable test for the stellar population during reionization and can be used to constrain global physical conditions in the early universe.\\\\ \\\\  Regarding the stellar population, the works \\cite{Abel} and \\cite{Bromm} suggest that the first stars were top-heavy with a peak in the IMF at around $100\\ M_\\odot$. On the contrary, \\cite{Clark} and \\cite{Omukai} indicate that gas especially in more massive systems can fragment because of dips in the equation of state, which may lead to the formation of a stellar cluster. It was shown that cooling in previously ionized gas is enhanced and leads to typical stellar masses of $\\sim10\\ M_\\odot$ \\cite{Yoshida1, Yoshida2}. It was further suggested that the presence of weak magnetic fields is sufficient to lead to a more present-day like mode of star formation, resulting in considerably lower masses \\citep{Silk}. Recently, a new phase of stellar evolution was suggested in which the stars would be powered by dark matter annihilation instead of nuclear fusion \\cite{Spolyar}. Various follow-up works have explored such a scenarios in more detail. Studies by \\citet{Iocco, Freese} explored the main-sequence and pre-main-sequence phase of dark stars, and other works calculated the effect of dark matter capture by off-scattering from baryons \\citep{IoccoBressan,Yoon,FreeseSpolyar,Taoso}. While many predictions are still model-dependent, it has often been suggested that such stars may have typical masses of $800\\ M_\\odot$, giving rise to a very bright main-sequence phase, or may have much longer lifetimes due to modifications in the stellar evolution. In the end, all these suggestions must face the constraint that the stellar population must be able to provide the correct reionization optical depth. \\\\ \\\\  Complications may arise through the presence of additional physics that are often not considered in standard reionization calculations and simulations on the first stars. Such possibilities include the presence of primordial magnetic fields, dark matter decay and dark matter annihilation. Magnetic fields have been observed on all scales in the universe, and recently, it has been demonstrated that they were present already in high-redshift galaxies \\citep{Bernet}. They are found in the interstellar gas as well as in the intergalactic medium \\citep[][]{Beck96, Beck01, Carilli02}, but their origin is still unclear. There is a viable possibility that these fields have a primordial origin \\citep{Grasso01, Widrow02}. So far, the most stringent constraints on the strength of these putative fields come from the measurements of the cosmic microwave background radiation (CMBR) and from big-bang nucleosynthesis (BBN) calculations.  A homogeneous magnetic field would produce temperature anisotropies in the CMBR~\\citep{Zeldovich83} whose maximal amplitudes are limited by the COBE satellite measurements which in turn limit the field strength, as measured today, to $B_0 \\simle 3.5\\times 10^{-9} \\, \\rm{G}$~\\citep{Barrow97}. The presence of any primordial field would also alter the CMBR power spectrum by changing the characteristic velocities. With the sensitivity of the PLANCK satellite \\footnote{http://sci.esa.int/science-e/www/area/index.cfm?fareaid=17} one should be able to detect fields with present day strength of $B_0 > 5\\times 10^{-8} \\, \\rm{G}$~\\citep{Adams96}.  So far, measurements by the WMAP satellite \\footnote{http://map.gsfc.nasa.gov/} are compatible with the absence of primordial magnetic fields.  Strong magnetic fields in the early universe can also change the abundance of relic $^4$He and other light elements during the big bang nucleosynthesis~\\citep{Matese69, Greenstein69}. To comply with observational limits on light element abundances these primordial fields must not exceed $10^{12} \\, \\rm{G}$ at the time when the universe was $T = 5\\times 10^9 \\, \\rm{K}$ which corresponds to a present day field $B_0 \\simle 3\\times 10^{-7} \\, \\rm{G}$\\cite{Greenstein69}.  Effects of primordial magnetic fields have already been considered in \\cite{Kim}, finding that density perturbations can be enhanced by the Lorentz force from tangled magnetic fields. Both the evolution of perturbations in the presence of magnetic fields as well as their effect on the thermodynamics via ambipolar diffusion heating and decaying MHD turbulence was considered in \\cite{Sethi05}. Recent calculations of \\cite{Tashiro} show that the enhancement of structure due to magnetic fields is pronounced at about $5\\times10^6\\ M_\\odot$, but becomes less effective on larger mass scales and appears as a subdominant contribution on the scale of the magnetic Jeans mass. Consequences for $21$ cm observations have been explored as well \\cite{Tashiro21,Schleicher}. In this work, we examine the consequences for reionization in more detail and calculate the backreaction on structure formation according to the work of \\citep{Gnedin}. The WMAP 5 year data \\citep{WMAPAngular} have measured the Thomson scattering optical depth from reionization and allow to constrain different reionization scenarios in the early universe.  The nature of dark matter is still unclear, and consequences of various particle physics like massive neutrinos or axion decay have been explored early \\citep[e. g.][]{Doroshkevich,Berezhiani, Berezhiani2}. Observational progress allowed to refine these studies and to explore such scenarios in the framework of $\\Lambda$CDM cosmology \\citep{Chen, Hansen, Avelino, Kasuya, Pierpaoli, Bean, Padmanabhan}. Recently, the possibility was discussed to detect such effects using future $21$ cm telescopes\\cite{FurlanettoDM}, and it was shown that the fraction of the energy absorbed into the IGM can be calculated in detail for specific models of dark matter decay and annihilation \\cite{Ripamonti}. Indeed, secondary ionization through the decay / annihilation products can provide a way to ionize the IGM which is independent of the stellar contribution \\citep{Mapelli}, whereas the additional heat input increases the Jeans mass and delays the formation of the first structures.  In this work, we show how reionization constrains the properties of the stellar population as well as some additional heat sources like primordial magnetic fields, dark matter annihilation and decay. In \\S\\ref{reionization}, we present our model for stellar reionization, which considers the IGM as a two-phase-medium consisting of ionized bubbles and overall neutral gas. Based on the thermal evolution, we self-consistently determine the minimal mass scale of halos which can collapse. In \\S \\ref{bfields}, we explain our treatment of primordial magnetic fields and show how they modify the thermal evolution. The treatment of dark matter annihilation and decay, as well as some implications for dark stars, are discussed in \\S \\ref{darkness}. The optical depth for different models is given in \\S \\ref{cmbfast}. Further discussion and outlook is given in \\S \\ref{outlook}. ",
        "conclusions": "\\label{outlook} We have calculated the reionization optical depth for different stellar models in the presence of primordial magnetic fields as well as dark matter annihilation and decay. The results indicate which star formation efficiencies are required in the presence of some additional heating mechanism for a given stellar model. Considering different stellar models and primordial magnetic fields, we find the following results: \\begin{enumerate} \\item Independent of the model for the stellar population, we can securely exclude primordial magnetic fields larger than $5$ nG. \\item For the most realistic case with a mixed stellar population (model C), we even find an upper limit of $0.7$ nG, as higher magnetic fields would require star formation efficiencies larger than $1\\%$, which is unrealistic. Similar results are found for model B, assuming reionization completely due to Pop. III stars. \\item Reionization only due to population II stars (model D) is ruled out completely. \\end{enumerate} For dark matter, we found the following results: \\begin{enumerate} \\item Dark matter annihilation provides noticeable contributions to the reionization optical depth only for thermally averaged mass-weighted cross sections $\\langle\\sigma v \\rangle/m_{DM}\\ge10^{-33} \\mathrm{cm}^3/\\mathrm{s}/\\mathrm{eV}$. \\item Parameters $\\langle\\sigma v \\rangle/m_{DM}\\ge3\\times10^{-33}\\ \\mathrm{cm}^3/\\mathrm{s}/\\mathrm{eV}$ can be ruled out by $3\\sigma$ on the basis of WMAP 5 year data. \\item Dark matter decay becomes important for the reionization optical depth for lifetimes below $3\\times10^{24}$ s. \\item Dark matter lifetimes below $3\\times10^{23}$ s are ruled out by $3\\sigma$. \\end{enumerate} These results are in agreement with conservative constraints obtained from the gamma-ray background \\citep{Mack}. We further showed that reionization can not be due to $\\sim1000\\ M_\\odot$ dark stars alone, as the corresponding optical depth would be significantly too high. One might wonder whether heating from dark matter annihilation might help to significantly delay stellar reionization, in order to reconcile this model with observations. However, as can be seen in Fig. \\ref{img:annihilation}, this is more than compensated by the effects of secondary ionization, once that dark matter annihilation starts to have a significant influence on the IGM. In case collider experiments like the LHC \\footnote{http://lhc.web.cern.ch/lhc/} or other dark matter detection experiments \\footnote{A list of dark matter detection experiments is given at http://cdms.physics.ucsb.edu/others/others.html.} find evidence for a self-annihilating dark matter candidate, this can be seen as evidence for a rapid transition towards a different mode of star formation, or a problem in our understanding of dark stars.  To reconcile dark stars with observations, the following scenarios seem feasible: \\begin{enumerate} \\item Dark stars with $1000\\ M_\\odot$ may be extremely rare objects, and their actual mass scale is closer to the mass scale for typical Pop. III stars. In such a case, reionization could not distinguish between dark stars and conventional Pop. III stars. \\item The transition to lower-mass stars occurs very rapidly, such that dark stars only contribute to the very early phase of reionization. Reasons for that might be chemical, radiative as well as mechanical feedback (see also the discussion in \\S \\ref{dark star}). In fact, even a double-reionization scenario might be conceivable, in which very massive dark stars ionize the universe at high redshifts. Due to chemical and radiative feedback, formation of such stars might be suppressed and the universe might become neutral again, until the formation of less massive stars becomes efficient enough to reionize the universe. \\item Dark stars do not reach a main-sequence phase, but are disrupted earlier by some non-linear instability. Such an instability would be constrained by the fact that it neither appears in a linear stability analysis nore in 1D simulations. Violent explosions of dark stars might even be considered as a source for Gamma-Ray bursts. \\end{enumerate} While too definite conclusions on the existence of dark stars are not yet possible, it is at least indicated that the possibilities mentioned above should be explored in more detailed, and a better understanding of the properties of the dark stars. A better understanding of their evolution after the dark phase will certainly help to better understand their possible role during reionization.  The constraints derived here are independent of other works that essentially rely on the physics of recombination to derive upper limits on additional physics \\citep[e. g.][]{Barrow97, ZhangAnn, ZhangDec}. In particular, magnetic fields can evolve dynamically and their field strength may thus change between these epochs. It has recently been suggested that the Biermann Battery effect creates magnetic fields in the presence of an electron pressure gradient\\cite{Xu}. We thus need to probe magnetic fields at different epochs. As we will show in a separate work \\citep{Schleicher}, upcoming $21$ cm measurements will allow to probe the thermal history before and during reionization in great detail, and may allow to detect primordial magnetic fields of the order of $0.1$ nG.  Additional ways of probing the reionization history and the dark ages exist as well: Scattering of CMB-photons in fine-structure lines of heavy elements may lead to a frequency-dependent CMB power spectrum and may allow to measure metal-abundances as a function of redshift \\citep{Basu04}. Before reionization, molecules may form in the IGM and introduce further frequency-dependent features in the CMB \\citep{Schleicher2}. Such features are likely enhanced in the presence of either primordial magnetic fields or dark matter decay / annihilation, as the increased electron fraction catalyses the formation of molecules, and the additional heat input leads to a departure of the level populations from the radiation temperature.  Further improvements are expected from the upcoming measurement of Planck, that will measure the reionization optical depth with unprecedent accuracy and thus allow to strengthen the constraints obtained here on the stellar populations and additional physics. In addition, a more accurate determination of the cosmological parameters will remove further uncertainties in the present models of reionization."
    },
    "0807/0807.1346_arXiv.txt": {
        "abstract": "Scalar field dark energy evolving from a long radiation- or matter-dominated epoch has characteristic dynamics.  While slow-roll approximations are invalid, a well defined field expansion captures the key aspects of the dark energy evolution during much of the matter-dominated epoch.  Since this behavior is determined, it is not faithfully represented if priors for dynamical quantities are chosen at random.  We demonstrate these features for both thawing and freezing fields, and for some modified gravity models, and unify several special cases in the literature. ",
        "introduction": "}  The acceleration of the cosmic expansion can be explained by a multitude of models suggested in the literature but few or none derived from first principles. We look, therefore, for common characteristics among classes of models and for physically motivated generic behaviors.  One example of effective characterization is the designation of the thawing and freezing patterns of dark energy evolution \\cite{caldlin}.  These involve scalar fields leaving or approaching the behavior of a cosmological constant in the dark energy equation of state $w$ and its time variation $w'=dw/d\\ln a$.  Here we seek further unifying features in the evolution of the dark energy field.  For inflation a slow-roll approach, neglecting higher time derivatives in the field evolution, can be used.  This is not valid for dark energy, even at early times when $w$ may be near $-1$ and the dark energy is a small contribution to the total energy density, and we must develop a different formalism. Several suggestions for specific approximations exist in the literature and we will see that these can be unified into a single approach.  This can also be extended in part to the freezing class of fields, traditionally difficult to characterize generically.  One key conclusion is that one must take into account that our universe is old: the scalar field has been evolving for many Hubble times in a background that was initially radiation-dominated then matter-dominated. This defines particular initial conditions and determines the dynamical behavior.  Employing random, Monte Carlo initial conditions may lead to underrepresentation of thawing or freezing behavior, due to neglecting the physics of a long evolution.  In \\S\\ref{sec:dyn} we show why slow roll conditions are invalid for dark energy and discuss new methods for evaluating the dynamics.  We provide a specific example of a complete solution in \\S\\ref{sec:difq}. Identifying a particular characteristic combination of parameters in \\S\\ref{sec:flow}, we show how the physics constrains the dark energy evolution within the classes of dark energy.  We investigate extending this relation to some modified gravity scenarios in \\S\\ref{sec:Halpha}. ",
        "conclusions": "}  Our universe is old, having expanded by a factor of perhaps $10^{28}$ since the last period of acceleration during the inflationary epoch. This strongly affects the evolution of a scalar field that may give rise to the current epoch of acceleration and determines some key properties of the dark energy.  Although conventional slow roll approximations are invalid, we show that analysis in terms of an integral relation between the Hubble friction and potential driving terms and a well characterized field expansion can give insights into the evolutionary behavior.  For the case of thawing fields, the field expansion provides a clear initial track in the $w$-$w'$ phase space, a unification of a number of interesting special cases, and a rapid convergence in the evolution of the dark energy density.  For tracking fields, the ratio between the Klein-Gordon terms reaches a constant value.  In both cases, the evolution of the deviation of $\\op$ from zero and of the tilt $1+w$ scale in a manner constrained by the long matter (and radiation) dominated era.   Phrased in terms of a flow parameter combining the scalings, this ratio is nearly conserved until quite recent times, $z\\approx1-2$, when the dark energy finally begins to take over.  This physical behavior means that dark energy dynamics is not random, or equally probable, e.g.\\ in the sense of a uniform prior over $w$-$w'$, but is focused -- ``flows'' -- in specific ways. We have also tested this for two modified gravity theories and found some similar behavior but also some deviations that could offer clues to the nature of the acceleration.   While one can always arrange initial conditions such that the dark energy comes out of the matter-dominated era with arbitrary behavior, this involves fine tuning.  The oldness of our universe does provide a natural path for dark energy dynamics."
    },
    "0807/0807.3343_arXiv.txt": {
        "abstract": " ",
        "introduction": "Current observations of Type Ia supernovae luminosity distances indicate that our Universe is in a phase of accelerated expansion \\cite{Astier:2005qq}. Various proposals have been put forward to explain the present acceleration of the Universe.  One  can roughly distinguish two classes. On  the one hand, the acceleration might be caused by  the presence of  dark energy, a fluid with negative equation of state $w$. This may be provided by a tiny  cosmological constant which is characterized by \\mbox{$w=-1$} or by  some ultralight scalar field whose potential is presently dominating the energy density of the Universe. This is usually dubbed quintessence \\cite{ArmendarizPicon:2000dh} (see \\cite{Copeland:2006wr} for a comprehensive review). On the other hand, the acceleration might be due to a modification of standard gravity at large distances. This happens in $f(R)$ theories \\cite{Carroll:2003wy} and in extra-dimension inspired models, like DGP \\cite{Dvali:2000hr}. Understanding which class of models Nature has chosen will represent not only a  breakthrough in cosmology, but also in the field of high energy physics.  Mapping the expansion of cosmic scales and the growth of large scale structure in tandem can provide insights to distinguish between the two possible origins of the present acceleration. For such reason, there has been increasing interest in analysing the time evolution of the dark matter perturbation. Several recent works deal with characterizing the growth of dark matter perturbations in different frameworks \\cite{Knox:2005rg, Koyama:2005kd, Koyama:2006ef, Kunz:2006ca,   Uzan:2006mf, Carroll:2006jn, Bertschinger:2006aw, Linder:2007hg, Tsujikawa:2007gd, Acquaviva:2007mm, Jain:2007yk, Wei:2008vw, Wei:2008ig, Zhang:2008kx, Bertschinger:2008zb}.  The evolution of the growth function of dark matter perturbations \\mbox{$g=\\delta_c/a\\,$,} which is the ratio between the perturbation $\\delta_c$ and the scale factor of the Universe $a$, can be parameterized in a useful way using the growth index $\\gamma$ \\cite{Linder:2005in}, defined in Eq. \\ref{gammadef}. In a pure matter-dominated Universe, $g$ does not evolve in time (remains equal to one) and $\\gamma$ is zero. However, in the presence of a dark energy background, $g$ changes in time, $\\gamma$ is different from zero and its value can be approximated by \\be \\label{eqlind} \\gamma=0.55+0.05\\left[1+w(z=1)\\right], \\ee which provides a fit to the evolution of $g$ to better than 0.2\\% for $-1\\lesssim w$ and a broad range of initial conditions for the dark matter abundance \\cite{Linder:2005in}. Typically, the growth index in modified gravity models turns out to be  significantly different (for instance $\\gamma\\simeq 0.68$ for DGP \\cite{Linder:2005in}) and therefore it is in principle distinguishable from the one predicted for dark energy models\\symbolfootnote[2]{However, see \\cite{Kunz:2006ca} and \\cite{Bertschinger:2008zb}.}. The available data on the growth of structures are still poor and there is a long way to go before we can talk about precision cosmology in this respect. The methods developed to study the growth of structure involve baryon acoustic oscillations, weak lensing, observations of \\mbox{X-ray} luminous clusters, large scale galaxy surveys, {\\mbox Lyman-$\\alpha$} power spectra and the integrated \\mbox{Sachs-Wolfe effect} on the Cosmic Microwave Background. There are however various works that use these kind of techniques to place constraints on the growth index (and some also on the equation of state of dark energy) as well as  forecasts for its determination based on future observations \\cite{Huterer:2006mva, Amendola:2007rr, Sapone:2007jn, Nesseris:2007pa, Dore:2007jh, Mantz:2007qh, DiPorto:2007ym, Yamamoto:2007gd,  Heavens:2007ka, Acquaviva:2008qp}. In particular, it is found in \\cite{Heavens:2007ka} using Bayesian methods that a next generation weak lensing survey like DUNE \\cite{Refregier:2008js} can strongly distinguish between two values of $\\gamma$ that differ by approximately $0.05\\,$. The authors of \\cite{Amendola:2007rr} made a forecast for the same kind of satellite proposal and concluded that it will be possible to measure the growth index with an absolute error of about 0.04 at 68\\% confidence level. In \\cite{Sapone:2007jn} a slightly bigger error of 0.06 at the same confidence level is given for a forecast based on baryon acoustic oscillations. Finally, for a combination of weak lensing, supernovae and Cosmic Microwave Background data an error of about 0.04\\, is estimated in \\cite{Huterer:2006mva} after marginalizing over the other cosmological parameters. Since the growth index  is approximately equal to 0.55, the nearest future observations should be able to determine it with a relative error of around 8\\%.  While much effort has been put into determining the value of the growth index in dark energy and in modified gravity models, less attention has been devoted to the possible effect on $\\gamma$ of  non--vanishing dark energy perturbations. The latter do not affect the background evolution,  but are fundamental in determining the dark energy clustering properties. They will have an effect on the evolution of fluctuations in the matter distribution and, consequently, on $\\gamma$. While minimally coupled scalar field (quintessence) models commonly have a non-adiabatic speed of sound close or equal to unity, and therefore dark energy perturbations can be neglected for them; other non-minimal models, for instance the adiabatic Chaplygin gas model, motivated by a rolling tachyon \\cite{Gibbons:2002md},  have  a speed of sound which is approximately zero. Observational implications of dark energy perturbations with a small  speed of sound in a variety of dark energy models have been  recently discussed in  k-essence \\cite{Erickson:2001bq,DeDeo:2003te}, condensation of dark matter \\cite{Bassett:2002fe} and the Chaplygin gas, in terms of the matter power spectrum \\cite{Sandvik:2002jz,Beca:2003an} and combined full CMB and large scale structure measurements \\cite{Bean:2003fb,Amendola:2003bz}. Let us also emphasize that dark energy perturbations may not be consistently set to zero in perturbation theory \\cite{Kunz:2006wc} even if $w= -1$. Indeed, it is unavoidable that dark energy perturbations are generated, even if set to zero on some initial hypersurface, due to the presence of a non--vanishing gravitational potential. Therefore, the expression (\\ref{eqlind}) rigorously holds only in the physical limit in which the speed of sound is very close to unity \\mbox{(if $w\\neq -1$)} so that dark energy perturbations are sufficiently suppressed.  In this Letter we study the effect of dark energy perturbations on the growth index $\\gamma$. Our main motivation is to understand if the introduction of the new degrees of freedom introduced by dark energy perturbations imply changes in $\\gamma$ large compared to the forecasted errors $\\Delta\\gamma\\simeq {\\cal O}(0.04)$ (at 68\\% confidence level). Following the common lore, see for instance \\cite{Bean:2003fb}, and to simplify the analysis, we will assume that the speed of sound associated with the dark energy perturbations and the equation of state do not change appreciably in the proper time range and that the dark energy perturbations have no shear. This is a good approximation in linear perturbation theory for dark energy models with a scalar field. Under these assumptions, we  provide an analytical formula for the growth index $\\gamma$ as a function of the speed of sound, the equation of state $w$, the dark matter abundance and the comoving scale. As we will see, in the presence of dark energy perturbations, the growth index differs from the corresponding value without dark energy perturbations by an amount which is comparable to the  realistic forecasted errors, especially for small speed of sound and $w$ significantly different from $-1$. This opens up the possibility that the presence of dark energy perturbations may leave a significant imprint on the growth function of dark matter perturbations.  The Letter is organized as follows. In Section 2 we summarize our framework and provide the necessary equations for the perturbations at the linear level. In Section 3 we discuss the growth index and in Section 4 we give our results and summarize of our work. ",
        "conclusions": "\\noindent In this section we present a combination of numerical results and an analytical formula for the growth index $\\gamma$ as a function of the relevant cosmological parameters.  In Figure~\\ref{functionwf} we plot the growth index at $z=1$ versus $w$ for several values of the speed of sound of dark energy and two different scales. Notice that the curves for the two different values of the comoving momenta coincide for $\\hat{c}_s=1$ and in the limit of very small speed of sound. The figure indicates that the dark energy speed of sound and the scale determine whether $\\gamma$ grows or decreases as a function $w$ at a given redshift. This is one of the reasons why having a more complete parameterization than (\\ref{eqlind}) is important. Choosing another redshift would have the effect of an overall shift of the merging point at $w=-1$ together with modifications in the curvatures of the lines.  To gain some insight on the change of the value of $\\gamma$ from $\\cshat=1$ to $\\cshat\\ll 1$, we observe that, in the limit $\\cshat\\simeq 0$ and from Equation (\\ref{depvel}),  the dark energy velocity perturbation promptly decays in time. One is left with the following solution for $\\delta_x$ \\begin{equation} \\label{aa} \\delta_x(a)=\\delta_{x(mr)}\\left(\\frac{a}{a_{(mr)}}\\right)^{3w}+ (1+w)a^{3w}\\int \\, \\tilde{a}^{-3w-1}\\,\\dot{\\delta}_c\\left(\\cshat=1\\right)\\,{\\rm d}\\tilde{a}\\,, \\end{equation} where the dot  stands for differentiation with respect to $\\ln a$. As a first approximation, we can solve Equation (\\ref{aa}) plugging in the dark matter perturbation $\\delta_c\\left(\\cshat=1\\right)$ obtained taking $\\cshat=1$, which for this purpose corresponds to the case in which no dark energy perturbations are present. From Equation (\\ref{first}), it is clear  that the dark energy perturbations provide an extra source for the dark matter pertrubation growth. We then solve numerically Equation (\\ref{first}) with this new known source and $\\theta_x=0\\,$. The difference between the true value of $\\gamma$ and the one obtained with such an approximation is plotted in Figure~\\ref{diff}.  \\begin{figure} \\begin{center} \\psfrag{a}[t][t][1.3][0]{$w$} \\psfrag{b}[b][t][1.3][0]{$\\left(\\gamma(z=1)/\\gamma_{ap}(z=1)-1\\right)\\times100$}              % \\includegraphics[width=0.9\\linewidth]{diff} \\caption{\\small{Relative error as a function of $w$ between the exact numerical result for $\\gamma(z=1)$ with very small dark energy speed of sound and the approximation $\\gamma_{ap}$ at the same redshift based on Equations (\\ref{aa}) and (\\ref{first}) with zero $\\theta_x$. The figure has been done for $\\cshat=10^{-6}$, $\\Omega_c^0=0.30$ and $k=0.050\\, h {\\rm Mpc}^{-1}\\,$.}} \\label{diff} \\end{center} \\end{figure}  In Figure~\\ref{functioncs} we show the growth index at $z=1$ versus $\\log_{10}\\hat{c}_s$ for different values of the equation of state of dark energy and two scales $k$. From this plot it is clear that the effect of changing the scale is an overall shift along the $\\log_{10}\\hat{c}_s$ axis. Notice that the intersecting points for the two sets of lines have the same value of the growth index, $\\gamma\\simeq0.547$, which corresponds to the merging point in Figure~\\ref{functionwf}.  The redshift dependence of the growth index has already been studied without taking into account dark energy perturbations \\cite{Polarski:2007rr} concluding that $\\rmd \\gamma /\\rmd z\\sim-0.02$ at $z=0\\,$; being this value nearly independent of $z$ for a given $\\Omega_c^0\\,$. However, including dark energy perturbations, we find that it is actually the derivative of $\\gamma$ with respect to the scale factor $a$ which is constant. Therefore the redshift dependence of the growth index can be better modeled with a $1/z$ term plus a constant term. We will later see that the growth index actually has an almost constant slope as a function of the scale factor when dark energy perturbations are taken into account.  \\begin{figure} \\begin{center} \\psfrag{z}[b][t][1.3][0]{$\\gamma(z=1)$} \\psfrag{w}[t][t][1.3][0]{$\\log_{10}\\hat{c}_s$} \\psfrag{A}[c][c][0.8][0]{$w=-0.99$} \\psfrag{B}[c][c][0.8][0]{$w=-0.9$} \\psfrag{C}[c][c][0.8][0]{$w=-0.8$} \\psfrag{D}[c][c][0.8][0]{$w=-0.7$} \\includegraphics[width=0.9\\linewidth]{varycs} \\caption{\\small{$\\gamma(z=1)$ as a function of $\\log_{10}\\hat{c}_s$ is shown for four values of $w$. Red curves correspond to $k=0.03\\, h {\\rm Mpc}^{-1}$ and blue dashed ones to $k=0.08\\, h {\\rm Mpc}^{-1}$.}} \\label{functioncs} \\end{center} \\end{figure}  Our next step is to obtain an analytical parameterization of the growth index as a function of the cosmological parameters. We start with the following generic ansatz: \\be \\label{linans} \\gamma\\left(\\Omega_c^0,\\cshatn,k,w,a\\right)=\\gamma_{eq}\\left(\\Omega_c^0,\\cshatn,k,w\\right)+\\zeta\\left(\\Omega_c^0,\\cshatn,k,w\\right)\\left[a-a_{eq}\\left(\\Omega_c^0,w\\right)\\right]\\;, \\ee where $a_{eq}$ is the value of the scale factor at which ``dark equality'' ($\\Omega_c=\\Omega_x=1/2$) takes place: \\be \\label{equality} a_{eq}=\\left(\\frac{1}{\\Omega_c^0}-1\\right)^{\\frac{1}{3w}}\\;. \\ee We want to fit the growth index for $a$ in the interval $[a_{eq},1]$ which approximately corresponds to a redshift $z\\in[0,0.55]$ for the ranges of the equation of state of dark energy and its relative energy density that we consider. Ideally one would wish to be able to use (\\ref{eqgamma}) and the equations for the perturbations to infer completely the analytical dependence of $\\gamma_{eq}\\left(\\Omega_c^0,\\cshatn,k,w\\right)$ and $\\zeta\\left(\\Omega_c^0,\\cshatn,k,w\\right)$ in their variables. This turns out to be difficult and we find it efficient to make a numerical fit directly. The generic form (\\ref{linans}), which can be viewed as a first order Taylor expansion in the scale factor, is motivated by the nearly zero variation of $\\rmd \\gamma /\\rmd a$. The choice of $a_{eq}$ as the point around which we make the expansion is a convenient one,  but the fit could in principle be done taking a model independent value of $a$ as the fiducial point. We use the same ansatz to fit $\\gamma_{eq}$ and $\\gamma_0\\,$, which is the growth index at $a_0=1$, and doing so we directly obtain the slope $\\zeta\\,$ from (\\ref{linans}): \\be\\label{slope} \\zeta\\left(\\Omega_c^0,\\cshatn,k,w\\right)=\\frac{\\gamma_0-\\gamma_{eq}}{1-a_{eq}}\\;. \\ee In particular, we assume the following parameterization for $\\gamma_{eq}$ and $\\gamma_0$: \\be\\label{ordinate} \\gamma_{j}\\left(\\cshatn,k,w\\right)=h_{j}(w)\\tanh\\left[\\left(\\log_{10} \\cshatn - g_{j}(k)\\right)\\frac{r_{j}(w)}{h_{j}(w)}\\right]+f_{j}(w)\\;,\\, j=\\{eq,0\\}\\,. \\ee Notice that we have taken $\\gamma_{eq}$ and $\\gamma_0$ to be independent of $\\Omega_c^0\\,$ and we incorporate this assumption in our notation, so we will refer to $\\gamma_{j}\\left(\\cshatn,k,w\\right)$ from now on. The functions $f_{j}(w)\\,$, $g_{j}(k)\\,$, $h_{j}(w)\\,$ and $r_{j}(w)\\,$ are polynomials in their variables. It turns out that the fit obtained with this procedure can be importantly improved with the addition of a polynomial correction to $\\zeta$ that depends on $\\Omega_c^0$, so finally: \\bea \\nonumber \\gamma\\left(\\Omega_c^0,\\cshatn,k,w,a\\right)&=&\\gamma_{eq}\\left(\\cshatn,k,w\\right)\\\\ \\label{gammacorr} &&+\\left[\\zeta\\left(\\Omega_c^0,\\cshatn,k,w\\right)+\\eta\\left(\\Omega_c^0\\right)\\right]\\left[a-a_{eq}\\left(\\Omega_c^0,w\\right)\\right]\\,. \\eea  \\begin{figure} \\begin{center} \\psfrag{z}[t][t][1.3][0]{$\\log_{10}\\hat{c}_s$} \\psfrag{g}[b][t][1.3][0]{$\\gamma_{eq}\\left(\\cshatn,k,w\\right)$} \\psfrag{A}[c][c][0.8][0]{A} \\psfrag{B}[c][c][0.8][0]{B} \\psfrag{C}[c][c][0.8][0]{C} \\psfrag{D}[c][c][0.8][0]{D}      % \\includegraphics[width=0.9\\linewidth]{figureeq} \\caption{\\small{$\\gamma_{eq}\\left(\\cshatn,k,w\\right)$ versus $\\log_{10}\\hat{c}_s$ for different combinations of the pair $\\{k,w\\} \\,$: \\newline ${\\rm A}=\\{0.08\\, h\\, {\\rm Mpc^{-1}},-0.95\\}\\,$, ${\\rm B}=\\{0.02\\, h\\, {\\rm Mpc^{-1}},-0.7\\}\\,$, ${\\rm C}=\\{0.04\\, h\\, {\\rm Mpc^{-1}},-0.87\\}\\,$ and ${\\rm D}=\\{0.06\\, h\\, {\\rm Mpc^{-1}},-0.75\\}$. Red lines are the exact numerical result and blue dashed ones the corresponding fits.}} \\label{figeq} \\end{center} \\end{figure}  The set of equations (\\ref{slope}), (\\ref{ordinate}) and (\\ref{gammacorr}) constitute the full fitting formula for the growth index. The resulting nine polynomials through which the fit can be expressed are the following: {\\small \\bea f_{eq}(w)&=&4.498\\cdot10^{-1}- 2.176\\cdot10^{-1}\\, w - 1.041\\cdot10^{-1}\\, w^2 + 5.287\\cdot10^{-2}\\,w^3\\nonumber\\\\&+&4.030\\cdot10^{-2}\\,w^4\\,,\\\\ f_{0}(w)&=&4.264\\cdot10^{-1}- 3.217\\cdot10^{-1}\\, w - 2.581\\cdot10^{-1}\\, w^2- 5.512\\cdot10^{-2}\\,w^3\\nonumber\\\\&+&1.054\\cdot10^{-2}\\,w^4\\,, \\eea \\bea g_{eq}(k)&=&-5.879\\cdot10^{-1} - 2.296\\cdot10^{-2}\\, k + 2.125\\cdot10^{-4}\\, k^2 - 1.177\\cdot10^{-6}\\, k^3\\nonumber\\\\  &+& 3.357\\cdot10^{-10}\\, k^4-3.801\\cdot10^{-12}k^5\\,,\\\\ g_{0}(k)&=&-6.401\\cdot10^{-1} - 2.291\\cdot10^{-2}\\, k + 2.119\\cdot10^{-4}\\, k^2 - 1.173\\cdot10^{-6}\\, k^3\\nonumber\\\\ &+& 3.344\\cdot10^{-10}\\, k^4-3.787\\cdot10^{-12}k^5\\,, \\eea \\bea h_{eq}(w)&=&1.759\\cdot10^{-1} + 4.066\\cdot10^{-1}\\, w + 3.254\\cdot10^{-1}\\, w^2+9.470\\cdot10^{-2}\\, w^3\\,,\\\\ h_{eq}(w)&=&2.008\\cdot10^{-1} + 4.644\\cdot10^{-1}\\, w + 3.713\\cdot10^{-1}\\, w^2+1.076\\cdot10^{-1}\\, w^3\\,, \\eea \\bea r_{eq}(w)&=& 5.158\\cdot10^{-1} + 1.203 w + 9.697\\cdot10^{-1}\\, w^2+2.827\\cdot10^{-1}\\, w^3\\,,\\\\ r_{0}(w)&=& 6.093\\cdot10^{-1} + 1.435 w +  1.1668 w^2+3.412\\cdot10^{-1}\\, w^3\\,, \\eea \\be \\eta(\\Omega_c^0)= 8.037\\cdot10^{-3} + 4.676\\cdot10^{-2}\\,\\Omega_c^0  - 2.829\\cdot10^{-1} \\left(\\Omega_c^0\\right)^2. \\ee } The truncation of the coefficients above has been done in such a way that the figures in the Letter can be reproduced and that the maximum relative error between the numerical value of $\\gamma$ and the fitting formula does not exceed 0.2\\% for any combination of the parameters. In fact, this error turns out to be much smaller for generic choices of the parameters.  In Figure~\\ref{figeq} we show $\\gamma_{eq}\\left(\\cshatn,k,w\\right)$ versus the decimal logarithm of $\\cshatn$ for several combinations of $k$ and $w$. The red curves represent the exact numerical growth index and the blue dashed lines are the corresponding fits. In Figures~\\ref{figuraw}, \\ref{figuraOmega} and \\ref{figurak} we show $\\gamma\\left(\\Omega_c^0,\\cshatn,k,w,a\\right)$ versus the scale factor for several values of $w$, $\\Omega_c^0$ and $k$ respectively, as explained in the captions. The other parameters are kept fixed. The colour code, as in Figure~\\ref{figeq}, is that the red curves represent the exact numerical growth index and the blue dashed lines are the corresponding fits. These figures are meant to illustrate the goodness of fit for several choices of the parameters.  \\begin{figure} \\begin{center} \\psfrag{a}[t][t][1.2][0]{$a$} \\psfrag{b}[b][t][1.2][0]{$\\gamma\\left(\\Omega_c^0,\\cshatn,k,w,a\\right)$} \\psfrag{r1}[c][c][0.8][0]{$w=-0.70$} \\psfrag{r2}[c][c][0.8][0]{$w=-0.75$} \\psfrag{r3}[c][c][0.8][0]{$w=-0.80$} \\psfrag{r4}[c][c][0.8][0]{$w=-0.90$} \\psfrag{r5}[c][c][0.8][0]{$w=-0.99$}            % \\includegraphics[width=0.9\\linewidth]{functionw} \\caption{\\small{$\\gamma\\left(\\Omega_c^0,\\cshatn,k,w,a\\right)$ versus $a$ for $k=0.033\\, h\\, {\\rm Mpc^{-1}}\\,$, $\\Omega_c^0=0.27\\,$ and $\\cshat=0.01\\,$. Different values of $w$ are chosen as shown in the figure. The red lines are the numerical results from the differential equations and the blue dashed ones are the fits to them.}} \\label{figuraw} \\end{center} \\end{figure}  \\begin{figure} \\begin{center} \\psfrag{a}[t][t][1.3][0]{$a$} \\psfrag{b}[b][t][1.3][0]{$\\gamma\\left(\\Omega_c^0,\\cshatn,k,w,a\\right)$}              % \\includegraphics[width=0.9\\linewidth]{varyOmega} \\caption{\\small{$\\gamma\\left(\\Omega_c^0,\\cshatn,k,w,a\\right)$ versus $a$ for $k=0.03\\, h\\, {\\rm Mpc^{-1}}\\,$, $w=-0.92$ and $\\cshat=0.0036\\,$. The value of $\\Omega_c^0\\,$ runs between 0.25 and 0.30 in steps of 0.01 from top to bottom of the figure. The red lines are the numerical results from the differential equations and the blue dashed ones are the fits to them.}} \\label{figuraOmega} \\end{center} \\end{figure}  \\begin{figure} \\begin{center} \\psfrag{a}[t][t][1.3][0]{$a$} \\psfrag{b}[b][t][1.3][0]{$\\gamma\\left(\\Omega_c^0,\\cshatn,k,w,a\\right)$}              % \\includegraphics[width=0.9\\linewidth]{varyk} \\caption{\\small{$\\gamma\\left(\\Omega_c^0,\\cshatn,k,w,a\\right)$ versus $a$ for $w=-0.80\\,$, $\\cshat=0.01$ and $\\Omega_c^0=0.27\\,$. The scale $k$ in units of $h\\, {\\rm Mpc^{-1}}$ takes the values $\\{0.023, 0.027, 0.037, 0.067 \\}\\,$ from bottom to top of the figure. The red lines are the numerical results from the differential equations and the blue dashed ones are the fits to them.}} \\label{figurak} \\end{center} \\end{figure}   Equations (25)-(36) offer an analytical expression for the growth index in terms of the relevant cosmological parameters in the case in which dark energy perturbations are present. The case without dark energy perturbations is reproduced by assuming $\\cshat=1\\,$. The analytical parameterization fits the numerical results  in the assumed range of parameters to a precision of 0.2\\% (in the worst cases) or better for the growth index. Our findings show that $\\gamma$ can vary from 0.55 by an amount $\\Delta\\gamma$ as large as $\\sim0.03\\,$. We have checked that this result holds for any redshift between $z_{eq}$ (at the time of dark equality) and $z=1$. This difference is of the same order of magnitude of  the 68\\% c.l. forecasted error band. The predicted value of $\\gamma$ may differ by this amount from the value without dark energy perturbations if the speed of sound is tiny and if the equation of state substantially deviates from $-1$. This opens up the possibility that a detailed future measurement of the growth factor might help in revealing the presence of dark energy perturbations. Finally, let us reiterate that our results have been obtained under the assumption that $\\cshat$ and $w$ do not evolve in time, at least for mild values of redshift. Furthermore, we have assumed that  the dark energy perturbations have no shear."
    },
    "0807/0807.1988_arXiv.txt": {
        "abstract": "{ The fractal structure of the interstellar medium suggests that the interaction of UV radiation with the ISM as described in the context of photon-dominated regions (PDR) dominates most of the physical and chemical conditions, and hence the far-infrared and submm emission from the ISM in the Milky Way. } { We investigate to what extent the Galactic FIR line emission of the important species \\co, \\czero, \\cplus, and \\ozero, as observed by the Cosmic Background Explorer (COBE) satellite can be modeled in the framework of a clumpy, UV-penetrated cloud scenario. } { The far-infrared line emission of the Milky Way is modeled as the emission from an ensemble of clumps with a power law clump mass spectrum and mass-size relation with power-law indices consistent with the observed ISM structure. The individual clump line intensities are calculated using the \\kosma{} for spherical clumps. The model parameters for the cylindrically symmetric Galactic distribution of the mass density and volume filling factor are determined by the observed radial distributions. A constant FUV intensity, in which the clumps are embedded, is assumed. } { We show that this scenario can explain, without any further assumptions and within a factor of about 2, the absolute FIR-line intensities and their distribution with Galactic longitude as observed by COBE. } {} ",
        "introduction": " ",
        "conclusions": ""
    },
    "0807/0807.2291_arXiv.txt": {
        "abstract": "Recently, analysis of near-infrared broad-band photometry and {\\it Spitzer} IRS spectra has led to the identification of a new ``pre-transitional disk\" class whose members have an inner optically thick disk separated from an outer optically thick disk by an optically thin gap.  This is in contrast to the ``transitional disks\" which have inner disk holes (i.e. large reductions of small dust from the star out to an outer optically thick wall). In LkCa 15, one of these proposed pre-transitional disks, detailed modeling showed that although the near-infrared fluxes could be understood in terms of optically thick material at the dust sublimation radius, an alternative model of emission from optically thin dust over a wide range of radii could explain the observations as well.  To unveil the true nature of LkCa 15's inner disk we obtained a medium-resolution near-infrared spectrum spanning the wavelength range 2-5 {\\micron} using SpeX at the NASA Infrared Telescope Facility. We report that the excess near-infrared emission above the photosphere of LkCa 15 is a black-body continuum which can only be due to optically thick material in an inner disk around the star. When this confirmation of a primordial inner disk is combined with earlier observations of an inner edge to LkCa 15's outer disk it reveals a gapped structure. Forming planets emerge as the most likely mechanism for clearing the gap we detect in this evolving disk. ",
        "introduction": "The origin of a star and its planets is intricately tied to the evolution of the system's primordial accretion disk.  These disks are composed of gas and dust and are formed in the collapse of the star's natal, molecular cloud \\citep{terebey84}.  As time passes, the dust grains in these primordial disks evolve:  they collide and stick, eventually growing in size and perhaps forming planetary systems much like our own \\citep{weiden97}.  The finer details of how the disk material evolves from an initially well-mixed distribution of gas and dust to a system composed mostly of large solids like our own solar system is not well understood.  In recent years a growing number of primordial disks with signatures of dust evolution hinting to the early stages of planet formation have been identified.  Most of these cases, dubbed transitional disks, \\citep{strom89} consist of stars with inner holes in their disks that are mostly devoid of material. Within the past few years, observations at mid-infrared wavelengths by the {\\it Spitzer Space Telescope} have led us to define transitional disks as those objects with small or negligible near-infrared flux excesses over photospheric fluxes but with a substantial excess in the mid-infrared and beyond.  This flux deficit at near-infrared wavelengths relative to full disks accreting material onto the star has been explained by modeling transitional disks as optically thick disks with inner cleared regions; the mid-infrared emission originates in the inner edge or ``wall\" of the truncated disk which is frontally illuminated by the star \\citep{calvet05}.   A small number of these transitional disks with detailed {\\it Spitzer} IRS spectra have been analyzed to date.  The estimated truncation radii of these disks cover a wide range, from 4 AU in DM Tau to 24 AU in GM Aur \\citep{calvet05}.  Transitional disks have been found in all ages where protoplanetary disks have been identified, from the $\\sim$1--2 Myr old Taurus population \\citep{calvet05} to the 10 Myr old TW Hya association \\citep{calvet02} and 25 Ori (Espaillat et al, in preparation).  In each case, the disk is accreting mass onto the star so we can conclude that gas is being transported through the inner cleared disk.  In addition, in some cases a small amount of micron or sub-micron dust coexists with the gas in this region, giving rise to an excess over photospheric fluxes detected in the near infrared.  Details of the distribution of this optically thin dust are largely unknown, although near-infrared interferometric observations suggest that this material is highly structured \\citep{ratzka07}.   A new class of evolving disks has been identified very recently. Disks in this class have inwardly truncated outer disks, as the transitional disks.  However, their significantly larger near-infrared excess, comparable to that of full disks, points to the existence of a remaining optically thick disk separated by a gap from the outer disk. A handful of these disks have been analyzed to date, including four around intermediate mass stars \\citep{brown07} and two around classical T Tauri stars \\citep{espaillat07b}.  In two of these cases, LkCa 15 and LkH${\\alpha}$ 330, millimeter interferometry has confirmed the truncation of the outer disk \\citep{pietu06,brown08} in agreement with the SED modeling \\citep{espaillat07b,brown07}. However, the substantial near-infrared excess of LkCa 15 could be explained by either optically thick material or by 5$\\times$10$^{-11}$ ${\\msun}$ of optically thin dust mixed with the gas in the inner disk \\citep{espaillat07b}. These contradictory hypotheses cannot not be properly tested with the existing near-infrared data, which consists of 2MASS and Spitzer broad-band photometry.  Here, we present detailed spectroscopic data that allows us to unambiguously search for optically thick material in the inner disk of LkCa 15, a CTTS in the Taurus cloud, in which an outer disk truncated at 46 AU has been imaged in the millimeter \\citep{pietu06}. We confirm that LkCa 15 belongs to the new class of disks around young stellar objects, the ``pre-transitional disks,\" \\citep{espaillat07b} where a gap is opening within the disk around a low mass pre-main sequence star. ",
        "conclusions": "Unseen planets \\citep{quillen04,rice03} have been proposed to clear out the large cavities within the primordial disks around GM Aur \\citep{calvet05}, DM Tau \\citep{calvet05}, TW Hya \\citep{calvet02,uchida04,hughes07}, and CoKu Tau/4 \\citep{dalessio05}. However, several other explanations are also possible.  The inner disk holes in most of these ``transitional disks\" can be explained by inside-out evacuation mechanisms like the magneto-rotational instability (MRI; Chiang \\& Murray-Clay 2007) and, in the case of CoKu Tau/4, by photoevaporation \\citep{alexander07}.  In addition, a stellar companion can inwardly truncate the outer disk as is most likely the cause of the inner hole of CoKu Tau/4 \\citep{ireland08}.  \\citet{alexander07} found that the accretion rates and disk masses of GM Aur, DM Tau, and TW Hya suggest planet formation in these systems, but perhaps the best observational evidence that links planets with the holes seen in transitional disks is the detection of a planet around TW Hya \\citep{setiawan08}.  Nonetheless, this still does not discount that the MRI is the main clearing agent in the inner disk given that this mechanism is still viable in the presence of a planet.  Consequently, inner disk holes are not conclusive signatures of planet formation.  The origin of the gap in LkCa 15 can be evaluated against different mechanisms that have been proposed to clear the inner disk, namely the MRI, photoevaporation, stellar companions, and planet formation.  The MRI operates on the ionized, frontally illuminated wall of the inner disk and allows material to accrete onto the star leading to inside-out clearing \\citep{chiang07}; the MRI cannot account for a remnant optically thick inner disk.  In the photoevaporation model, a stellar wind creates a small gap and halts the inward accretion of mass.  Without replenishment the inner disk quickly accretes onto the central star and only then can the hole grow outward \\citep{clarke01}. When the hole is about 46 AU, as is seen in LkCa 15, no inner disk remains \\citep{alexander07} so LkCa 15's gap cannot be due to photoevaporation.  A companion star would have to be located at about 18 -- 26 AU \\citep{artymowicz94} in order to truncate the outer disk at 46 AU and studies of LkCa 15 have revealed no stellar mass companion down to about 4 AU \\citep{leinert93, ireland08}.  Planet formation emerges as the most likely explanation since a planet can create a gap about its orbit \\citep{paardekooper04, varniere06}. The large gap of LkCa 15, which encompasses the orbits of Mercury (.4 AU) and Neptune (30 AU), raises the interesting possibility that we are seeing clearing due to multiple planets which would suggest that LkCa 15 may be an early analog of our own Solar system.  In conclusion, our observations confirm the presence of an inner optically thick disk in LkCa 15. The existence of optically thick material inside a truncated disk provides significant insight into the models presented to date to explain the transitional disks and calls for more detailed studies of this new class of disk.      \\vskip -0.1in"
    },
    "0807/0807.0893_arXiv.txt": {
        "abstract": "The study of Gamma-ray bursts (GRBs) is a key field to expand our understanding of several astrophysical and cosmological phenomena. SVOM is a Chinese-French Mission which will permit to detect and rapidly locate   GRBs, in particular those at high redshift, and to study their multiwavelength emission. The SVOM satellite, to be launched in 2013,  will carry wide field instruments operating in the X-/$\\gamma$-ray band and narrow field optical and soft X-ray telescopes. Here we describe a small soft X-ray telescope (XIAO) proposed as an Italian contribution to the SVOM mission. Thanks to a grazing incidence X--ray telescope with effective area of $\\sim$120 cm$^{2}$ and a short focal length, coupled to a very compact, low noise, fast read out CCD camera, XIAO  can substantially contribute to the overall SVOM capabilities for both GRB and non-GRB science. ",
        "introduction": "The study of Gamma-ray bursts (GRBs), extremely luminous transient sources appearing in the sky when black holes are born in the explosions of massive stars or in the merging of compact stellar objects, has become a key field to expand our understanding of several astrophysical and cosmological phenomena. These  include, e.g., the evolution of the  young universe, the history of star formation, the metal enrichment of galaxies, the mechanisms driving supernova explosions, the physics of ultra-relativistic shocks.   SVOM (Space-based multi-band astronomical Variable Objects Monitor) is a mission, developed in collaboration by the Chinese and French space agencies (CNSA and CNES), for the study of GRBs and other high-energy transients. SVOM is designed  to detect and rapidly locate all kinds of GRBs, in particular those at high redshift, and to study their emission on a broad spectral range, from the visible to the MeV region. This will be possible thanks to a satellite payload composed of  wide field X-/$\\gamma$-ray instruments and narrow field optical and soft X-ray telescopes, complemented by dedicated ground telescopes and a system for rapid distribution of the GRB positions.   Building on the successful experience of Swift\\cite{geh04}, the SVOM  operations are based on the following steps: \\textsl{(i)} GRB detection with a wide field gamma-ray imaging instrument able to derive on-board its localization with a few arcmin precision\\cite{sch07} ; \\textsl{(ii)} the GRB position is immediately transmitted to ground through a network of ground stations and at the same time the satellite slews rapidly to position the GRB in the narrow fields of view of its  X-ray and optical telescopes, which will study the afterglow and provide refined coordinates.  SVOM will adopt an optimized observation  strategy \\cite{cor08}, based on antisolar pointing and avoidance of the galactic plane, in order to permit follow-up observations with large telescopes and maximize the number of redshift measurements. The knowledge of redshift and the determination of the spectral shape over  an extended energy range are in fact essential to derive the bursts energetics and to study the empirical correlations used for cosmological studies.  The SVOM satellite will carry the following instruments:  \\begin{itemize}  \\item \\textit{Camera X and Gamma (CXG)}: a wide field instrument operating in the 4-300 keV energy range, with a field of view of 2 sr and a location accuracy of several arcmin.  This instrument, providing the GRB triggers and initial localizations, is based on an array of CdTe pixels with a sensitive area of 1000 cm$^{2}$ coupled to a coded mask aperture. It is based on the instrument originally proposed for the Eclairs satellite \\cite{sch06}.  \\item \\textit{XIAO}: a narrow field soft X--ray telescope to locate and study the GRB afterglows (see next Section).   \\item  \\textit{Visual Telescope (VT)}: an optical telescope operating in the 400-950 nm range, with a field of view of $\\sim$21$'$, reaching a sensitivity of V$\\sim$23 magnitudes in 300 s exposure times.  \\item \\textit{Gamma-ray Burst Monitor (GRM)}: a non-imaging spectrometer to measure the GRB spectra in the 50 keV-100 MeV energy range over a wide field of view  \\end{itemize}  SVOM will be placed in a near earth orbit, of $\\sim$600 km altitude and \\ltsima30$^{\\circ}$   inclination. The  GRB coordinates, and a small set of relevant information, will be transmitted on ground in real time by means of a network of VHF stations, as successfully done in the past by the HETE-2 satellite. The bulk of the data will be downloaded (a few times per day) when the satellite is in contact with the main ground station(s). The SVOM satellite will be launched in 2013, with the goal to  detect and precisely locate about 200 GRBs in a nominal mission duration  of 2.5 years. ",
        "conclusions": ""
    },
    "0807/0807.2919_arXiv.txt": {
        "abstract": "We present the Submillimeter Array observation of the CO $J$=2$-$1 transition towards the northern galaxy, ARP~302N, of the early merging system, ARP~302. Our high angular resolution observation reveals the extended spatial distribution of the molecular gas in ARP~302N. We find that the molecular gas has a very asymmetric distribution with two strong concentrations on either side of the center together with a weaker one offset by about 8 kpc to the north. The molecular gas distribution is also found to be consistent with that from the hot dust as traced by the 24 $\\mu$m continuum emission observed by the Spitzer. The line ratio of CO $J$=2$-$1/1$-$0 is found to vary strongly from about 0.7 near the galaxy center to 0.4 in the outer part of the galaxy. Excitation analysis suggests that the gas density is low, less than 10$^3$ cm$^{-3}$, over the entire galaxy. By fitting the SED of ARP~302N in the far infrared we obtain a dust temperature of $T\\rm _d$=26--36 K and a dust mass of M$\\rm _{dust}$=2.0--3.6$\\times$10$^8$ M$\\rm _\\odot$. The spectral index of the radio continuum is around 0.9. The spatial distribution and spectral index of the radio continuum emission suggests that most of the radio continuum emission is synchrotron emission from the star forming regions at the nucleus and ARP302N-cm. The good spatial correspondance between the 3.6 cm radio continuum emission, the Spitzer 8 \\& 24 $\\mu$m data and the high resolution CO $J$=2$-$1 observation from the SMA shows that there is the asymmetrical star forming activities in ARP~302N. ",
        "introduction": "Most luminous infrared (IR) galaxies (LIRGs; 10$^{11}$ L$\\rm _\\odot$ $<$ L$\\rm _{IR}$ $\\leq$ 10$^{12}$ L$\\rm _\\odot$) are found in interacting or merging systems with disturbed morphologies. In the standard scenario proposed by Sanders et al. (1988a,b), the formation of LIRGs, or ultra-luminous infrared galaxies (ULIRGs) is due to the interaction and merging of two gas-rich galaxies. Numerical simulations of interacting galaxies \\citep{barnes91,mihos96} indicate that a stable bar is formed during the first encounter. Under the gravitational torque induced by the bar, the molecular gas is channelled into the central regions of the galaxies forming a huge reservoir to fuel intense starbursts. Recent high angular resolution observations \\citep{downes98} seem to support this picture by revealing rotating disks of molecular gas that has been driven into the centers of the mergers.  A lot of (U)LIRGs have been studied at high angular resolution, for example ARP~299 \\citep{casoli99}, NGC5394/95 \\citep{kaufman02}, ARP~220 \\citep{wiedner02}, NGC4038/39 \\citep{zhu03}, NGC5194/95 \\citep{matsushita04}, NGC~6090 \\citep{wang04}, NGC5218 \\citep{olsson07}. A few samples of the sources along the merging sequence are also been investigated. \\citet{xilouris04} found evidence for an increase in 100- to 850-$\\mu$m flux density ratio along the merging sequence and interpreted this trend as an increase of the mass fraction of the warm dust in comparison to the cold dust, inferring the possible enhanced star formation rate (SFR). \\citet{bergvall03}, however, question the ubiquitous connection between interactions and starburst activity. \\citet{gao99} find the enhanced star formation efficiency (SFE) for the later merging systems, but the effect is cancelled by the decrease in the amount of molecular gas. That might result in no change in the infrared luminosity, i.e. the indicator of the starburst activity. Thus more sources in high resolution are needed and multiple CO transtions are essential to investigate the gas properties for the mergers at different stages.  Here we present another case study of an early merger system, ARP~302. This galaxy is also named as VV 340 or UGC 9618, which consists of two galaxies (ARP~302N and ARP~302S) with a separation of about 40$\\arcsec$ between the two nuclei, which is $\\sim$27 kpc at the angular distance of 138 Mpc (the luminosity distance is 147 Mpc), and has an infrared luminosity of 4.5$\\times$10$^{11}$ L$_\\odot$. ARP~302 is the most CO luminous system known in the local universe, and more luminous in CO than those ultraluminous galaxies out to redshifts of 0.27 (Solomon et al. 1997). The detected molecular gas mass is 9.1$\\times$10$^{10}$ M$_\\odot$ \\citep{lo97} (corrected by the adopted comsmology here). HI data gives a mass of atomic hydrogen 2.7$\\times$10$^{10}$ M$_\\odot$ \\citep{van01}. In all observations it was found that ARP~302N contributes most of the IR emission and molecular gas mass.  Previous high resolution CO $J$=1$-$0 observation on ARP~302 by \\cite{lo97} (the angular resolution is 7$\\arcsec$$\\times$5$\\arcsec$) reveals two extended CO disks ($>$ 10 kpc), and ARP~302N shows roughly an exponential disk with a scale length of 7 kpc extending out to 23 kpc in diameter. Based on the unusually extended distribution of the molecular gas, \\cite{lo97} concluded that the high IR luminosity in ARP~302 arises simply because an extremely large amount of molecular gas is forming high-mass stars at a rate similar to the Galaxy, instead of the starburst activities in the galactic nucleus.  In this study, we report the high angular resolution CO $J$=2$-$1 observation of ARP~302N and the analysis of archival data in the radio and mid/far infrared to study the physical properties of the gas as well as the star formation activities in this early merger. In \\S\\ref{observations}, we show the details of our observations together with the information about the archive data used in this paper. In \\S\\ref{spatial}, \\S\\ref{ratio} and \\S\\ref{kinematics}, we present the results of spatial distribution and the kinematics of the molecular gas of our observations. In \\S\\ref{vla} and \\S\\ref{spitzer}, the results of the radio emission and the mid/far infrared emission from the archive data are presented. Finally we present the dust properties using a dust model in \\S\\ref{dust}, the excitation condition of the gas in \\S\\ref{excitation}, the star formation activities in \\S\\ref{formation} and a comparison with other mergers in \\S\\ref{compar}. A summary of this study is given in \\S\\ref{conclusion}. In this paper, we use a $\\Lambda$ cosmology with $\\rm H_{\\rm0} = 71$~km~s$^{-1}$~Mpc$^{-1}$, $\\rm \\Omega_\\Lambda=0.73$ and $\\rm \\Omega_{m}=0.27$ \\citep{spergel03}. ",
        "conclusions": "In this paper we present the SMA observations of the CO $J$=2$-$1 transition towards the northern galaxy, ARP~302N, of the early merging system, ARP~302, to study the physical properties of the gas. To well understand the star formation activities, the VLA data at 3.6~cm, 6.0~cm, 20~cm and the Spitzer (both IRAC and MIPS data) archive data are also shown here. Our main results are summarized as the following:  \\begin{enumerate}  \\item Our high angular resolution CO $J$=2$-$1 observation reveals the extended spatial distribution of the molecular gas in ARP~302N. We find that the molecular gas has a very asymmetric distribution with two strong concentrations on either side of the center together with a weaker one offset by about 8 kpc to the north. This spatial distribution is very similar to the morphology seen previously in the CO J=2--1 line.  \\item The distribution of the molecular gas is also found to be consistent with that for the PAH and the hot dust as traced by the 8 $\\mu$m and 24 $\\mu$m continuum emission observed by Spitzer.  \\item The line ratio CO J=2--1/1-0 shows that the line ratio varies strongly from about 0.7 near the galaxy center to 0.4 in the outer part of the galaxy. Excitation analysis suggests that the gas density is low, less than 10$^3$ cm$^{-3}$, over the entire galaxy.  \\item By fitting the far infrared part of the SED of ARP~302N, we obtain a dust temperature $T\\rm _d$ in the range 26--36 K and a dust mass of M$\\rm _{dust}$=2.0--3.6$\\times$10$^8$ M$\\rm _\\odot$.  \\item We determine the spectral index of the radio continuum to be approximately 0.9. The spatial distribution and spectral index of the radio continuum emission suggests that most of the radio continuum emission is synchrotron emission from the star forming regions at the nucleus and ARP302N-cm. The good spatial correspondance between the 3.6 cm radio continuum emission, the Spitzer 8 and 24 $\\mu$m data and the high resolution CO $J$=2$-$1 observation from the SMA shows that there is the asymmetrical star forming activities in ARP~302N and ARP302N-cm is the strongest active star forming region in the galaxy.   \\end{enumerate}"
    },
    "0807/0807.4823_arXiv.txt": {
        "abstract": "{Plaskett's star (\\object{HD~47\\,129}) is a very massive O~+~O binary that belongs to the Mon OB2 association. Previous work suggested that this system displays the Struve-Sahade effect although the measurements of the secondary radial velocities are very difficult and give controversial results. Both components have powerful stellar winds that collide, and produce a strong X-ray emission.} {Our aim is to study in detail the physical parameters of this system, and to investigate the relation between its spectral properties and its evolutionary status.} {We present here the analysis of an extensive set of high resolution optical spectra of \\object{HD~47\\,129}. We use a disentangling method to separate the individual spectra of each star. We derive a new orbital solution and discuss the spectral classification of both components. A Doppler tomography technique applied to the emission lines H$\\alpha$ and \\ion{He}{ii}~$\\lambda$~4686 yields a Doppler map that illustrates the wind interactions in the system. Finally, an atmosphere code is used to determine the different chemical abundances of the system components and the wind parameters.} {\\object{HD~47\\,129} appears to be an O8~III/I~+~O7.5~III binary system in a post RLOF evolutionary stage, where matter has been transferred from the primary to the secondary star. The He overabundance of the secondary supports this scenario. In addition, the N overabundance and C underabundance of the primary component confirm previous results based on X-ray spectroscopy and indicate that the primary is an evolved massive star. We also determine a new orbital solution, with $M_{\\mathrm{P}}\\,\\sin^3i = 45.4 \\pm 2.4$~M$_{\\odot}$ and $M_{\\mathrm{S}}\\,\\sin^3i = 47.3 \\pm 0.3$~M$_{\\odot}$. Furthermore, the secondary star has a large rotational velocity ($v\\,\\sin\\,i \\sim 300$~km~s$^{-1}$) that deforms its surface, leading to a non-uniform distribution in effective temperature. This could explain the variations in the equivalent widths of the secondary lines with phase. We suggest that the wind of the secondary star is confined near the equatorial plane because of its high rotational velocity, affecting the ram pressure equilibrium in the wind interaction zone.}{} ",
        "introduction": "Plaskett's star (\\object{HD~47\\,129}) is a non-eclipsing early type binary which was analysed for the first time by \\citet{Plaskett22}. He found doubled spectral lines and described the system as consisting of a bright O5e primary component and a fainter secondary component with weak and diffuse lines. When he determined the elements of the spectroscopic orbit, the large mass of this system attracted his attention ($(M_{\\mathrm{P}} + M_{\\mathrm{S}}) \\sin^3 i = 138.9$~M$_{\\odot}$). This result motivated \\citet{Struve48} to re-observe this object in 1947$-$1948 over two complete orbital cycles. He detected that some lines in the secondary spectrum (H$\\gamma$, \\ion{He}{i}~$\\lambda\\lambda$~4026,~4471) were weaker when this component was receding. This effect is nowadays called the Struve-Sahade effect (S-S effect hereafter) and is currently defined as the apparent strengthening of the secondary spectrum of a binary system when the secondary is approaching, and the corresponding weakening of its lines when it is receding. \\citet{Struve48} also noted differences in the intensity of the secondary spectrum from one cycle to another, at a given orbital phase. He explained these variations with the presence of streams of gas in the system. A third observational campaign, conducted between 1956 and 1958 by \\citet{SSH58}, allowed these authors to conclude that the behavior of the radial velocities of the secondary component suggests that it ejects matter in an irregular manner, and that these ejections are larger over the hemisphere facing the primary star. \\citet{HC76} confirmed the presence of the S-S effect in the system and suggested that the primary star has a stellar wind which is stronger on the rear side of the star with respect to its orbital motion. IUE data allowed \\citet{Stickland87} to determine a radial velocity curve and an orbital solution for the primary component, but not to identify the secondary spectral lines. \\citet{BGW92} used a tomographic technique on the same IUE data to separate the two components of the system. They found a mass ratio of $1.18 \\pm 0.12$ (the secondary being more massive than the primary) and spectral types of O7.5~I and O6~I for the primary and the secondary, respectively. Given that they estimate an inclination around 71\\degr, they found masses equal to 42.5~M$_{\\odot}$ for the primary star and 51.0~M$_{\\odot}$ for the secondary star. The secondary lines appeared rotationally broadened ($v\\,\\sin\\,i = 310 \\pm 20$~km~s$^{-1}$) but are visible throughout the orbit. Finally they found that the secondary contribution to the cross-correlation functions is relatively constant with orbital phase. \\citet{Stickland97} also used IUE data and cross-correlation functions to study specifically the S-S effect in Plaskett's star. The conclusion was that the secondary radial velocities ''make no sense dynamically'', and that they are probably due to the presence of gaseous streams or envelopes in the system. However, \\citet{BGR99} used the same method and concluded that there is no evidence of the S-S effect in \\object{HD~47\\,129}.  \\begin{table}[t] \\caption{Journal of the observations of \\object{HD~47\\,129}. The first column indicates the instrument used to obtain the data, the second lists the Heliocentric Julian Date (HJD), the third provides the orbital phase ($\\Phi$) as computed from our orbital solution (see Table \\ref{tab:orbit:47129}) and the fourth and fifth ones give the radial velocities used to compute the orbital solution, respectively for the primary and secondary stars. The dispersion on the radial velocities are on average less than 2.0\\,km\\,s$^{-1}$ for the primary, and of the order of 20\\,km\\,s$^{-1}$ for the secondary component. } \\label{tab:journal:plaskett} \\centering \\begin{tabular}{c c c r r} \\hline \\hline Instrument&HJD &$\\Phi$ &\\multicolumn{2}{c}{RV (km s$^{-1}$)}\\\\ &$-$ 2 450 000 & &Primary&Secondary\\\\ \\hline \\multicolumn{5}{c}{}\\\\ FEROS&2335.565& 0.28&$230.8$&$-149.2$\\\\ FEROS&2336.532&  0.35&$197.7$&$-141.0$\\\\ FEROS&2336.551& 0.35&$192.9$&$-119.7$\\\\ FEROS&2337.527& 0.42&$120.6$&$-117.8$\\\\ FEROS&2337.534& 0.42&$118.3$&$-118.2$\\\\ FEROS&2338.528& 0.49&$60.8$&$-29.2$\\\\ FEROS&2339.535& 0.56&$-28.4$&-\\\\ AURELIE&3648.675&0.49&-&-\\\\ AURELIE&3650.662&0.63&-&-\\\\ AURELIE&3652.673&0.77&-&-\\\\ AURELIE&3654.663&0.91&-&-\\\\ FEROS&3738.692 &0.74&$-169.0$&$134.5$\\\\ FEROS&3772.629 &0.10&$153.9$&$-151.4$\\\\ FEROS&3772.640 &0.10&$158.5$&$-129.7$\\\\ AURELIE&3775.486&0.30&-&-\\\\ AURELIE&3776.486&0.37&-&-\\\\ AURELIE&3777.478&0.44&-&-\\\\ AURELIE&3778.462&0.51&-&-\\\\ AURELIE&3779.318&0.57&-&-\\\\ AURELIE&3780.288&0.63&-&-\\\\ FEROS&3796.515& 0.76&$-171.7$&$124.9$\\\\ FEROS&3797.521& 0.83&$-152.0$&$145.2$\\\\ FEROS&3798.510& 0.90&$-80.9$&$147.1$\\\\ FEROS&3799.511& 0.97&$-23.2$&$74.6$\\\\ FEROS&3800.507& 0.04&$74.5$&$-$\\\\ AURELIE&3843.329&0.01&-&-\\\\ AURELIE&4033.700&0.24&-&-\\\\ AURELIE&4034.692&0.30&-&-\\\\ AURELIE&4035.672&0.37&-&-\\\\ SOPHIE&4162.432& 0.18&$205.9$&$-181.6$\\\\ SOPHIE&4164.307& 0.31&$219.7$&$-190.5$\\\\ SOPHIE&4168.468& 0.60&$-78.6$&$155.3$\\\\ SOPHIE&4172.305& 0.86&$-123.5$&$187.6$\\\\ SOPHIE&4174.333& 0.00&$31.1$&-\\\\ SOPHIE&4176.327& 0.14&$208.5$&$ -135.2$\\\\ SOPHIE&4182.321& 0.56&$-45.3$&$100.2$\\\\ SOPHIE&4184.350& 0.70&$-164.5$&$160.0$\\\\ \\hline \\end{tabular} \\end{table}  \\begin{figure}[t] \\resizebox{\\hsize}{!}{\\includegraphics{disspec_new.ps}} \\caption{Normalized disentangled spectra of \\object{HD~47\\,129}, over 3 distinct wavelength ranges. The secondary spectrum was vertically shifted by 0.4 units for clarity.} \\label{fig:plaskett_disspec} \\end{figure}  Both components of Plaskett's star are very massive and luminous O-stars, and they thus both possess strong stellar winds that collide, producing X-ray emission in the interaction zone. An X-ray study of this object has been made by \\citet{LRP06}, using new \\textit{XMM-Newton} data and \\textit{ROSAT} archival data. The RGS (Reflection Grating Spectrometers of the \\textit{XMM-Newton} satellite) data showed a nitrogen overabondance by a factor of 7.44$^{+4.08}_{-2.59}$ in 2002 and 5.49$^{+4.51}_{-2.71}$ in 2003. We detected a modulation of the X-ray flux phased with the orbital motion, but could not find short term variability (from a few minutes to about one hour) that might have been expected from hydrodynamical instabilities of the wind-wind collision.  In this paper, we discuss a set of high resolution optical spectra of Plaskett's star, in order to constrain the evolutionary status of this peculiar system, and to determine the role played by the Struve-Sahade effect. Indeed, a better knowledge of the processes that take place in HD~47129 would be useful to understand the massive O+O binaries in general. The data have been obtained between 2002 and 2007 at the European Southern Observatory (ESO, Chile) and at the Observatoire de Haute-Provence (OHP, France).  The paper is organized as follows. Section~\\ref{observations} describes the observations and data reduction. In Section~\\ref{rvcurve}, a radial velocity curve, computed with a disentangling method, and the corresponding orbital solution are presented. We discuss the spectral type of \\object{HD~47\\,129} in Section~\\ref{spectraltype}, the behavior of the emission lines using a Doppler tomography technique in Section~\\ref{emission_lines} and the behavior of the absorption lines in Section~\\ref{sseffect}. We use a model atmosphere code in Section~\\ref{martins} in order to fit theoretical spectra to the disentangled spectra of the components of Plaskett's star and discuss the evolutionary status of the system in Section~\\ref{evolution}. Finally, Section~\\ref{conclusions} summarizes the results and gives the conclusions. ",
        "conclusions": "}  We have presented a spectroscopic analysis of the massive O-type binary \\object{HD~47\\,129}.  The radial velocities of both components have been inferred using a disentangling method. The orbital solution gave two massive components of 54~M$_{\\odot}$ and 56~M$_{\\odot}$ for the primary and the secondary star respectively (assuming an inclination of 71\\degr). The projected rotational velocities, determined using a Fourier technique, are very different from each other ($\\sim$~75~km~s$^{-1}$ for the primary, and~$\\sim$~300~km~s$^{-1}$ for the secondary).  The spectral types, calculated from the two disentangled spectra, gave O8~III/I~+~O7.5~III, which are too late spectral types for the dynamical masses. This discrepancy cannot be solved by taking into account the large rotational velocity of the secondary star. Whilst this effect indeed acts on the secondary luminosity, it is not sufficient to explain such a difference and it explains nothing concerning the primary star.  The optical spectra display a phase-locked variability in the strong emission lines \\ion{He}{ii}~$\\lambda$~4686 and H$\\alpha$. A Doppler tomography technique allowed us to obtain the Doppler maps of both lines which present their highest peak near the primary component velocity. Its velocity coordinates indicate a motion from the secondary toward the primary, which can be interpreted by a colliding wind interaction, where the shock is displaced toward the primary star. The Doppler map also shows an annular feature, almost centered on the system's center of mass, that seems to be another consequence of the secondary's rotation. Indeed, it likely represents the secondary's wind compressed into its equatorial plane by the fast rotation.  Furthermore, we used the high resolution spectra in order to study the Struve-Sahade effect in this system. Actually, while Plaskett's star presents strong variations in its absorption lines, it is difficult to connect them to the 'standard' definition of the Struve-Sahade effect, which is: ''the apparent strengthening of the secondary spectrum of the binary when it is approaching, and the corresponding weakening of its lines when it is receding'' \\citep{BGR99}. The variations are visible in both stars, even though they are stronger in the secondary's spectrum, but it is not clear that they are phase-locked. They appear dominated by a modulation that is either stochastic or at least not tied to the orbital phase. If correct, the interpretation of a flattened secondary wind would imply that these apparent variations result from the effect of the turbulence and other velocity perturbations on the emissions formed in the disk around the secondary star. The deformation of the star itself due to the rotation can also be in this case an additional, but not determining, perturbation. In the case of the main sequence stars of \\citet{LRS07}, the Struve-Sahade effect was already found to be due to an inhomogeneous distribution of the line formation, because of the deformation of the star and/or mutual heating, in at least three out of four studied objects (\\object{HD~100\\,213}, \\object{HD~159\\,176} and \\object{DH~Cep}).  Finally, we compared our observations with theoretical spectra computed with the model atmosphere code CMFGEN. We then determined effective temperature and chemical abundances for both components, and wind properties for the primary star. The nitrogen overabundance in the primary star, already detected in the XMM-Newton data, is confirmed in the optical spectra. This comes with an helium overabundance of the secondary star, which suggests a past RLOF in the system, from the primary star to its companion.  In conclusion, we can say that most of the peculiarities in the spectra of \\object{HD~47\\,129} can be explained today by the evolutionary status of the system (a post-case A RLOF) and by the fact that the projected rotational velocity of the secondary star is quite important. However, the large dynamical masses of both components are still not compatible with their spectral types and luminosities."
    },
    "0807/0807.1547_arXiv.txt": {
        "abstract": "To test the idea that ultraluminous X-ray sources (ULXs) in external galaxies represent a class of accreting intermediate-mass black holes (IMBHs), we have undertaken a program to identify ULXs and a lower luminosity X-ray comparison sample with the highest quality data in the {\\it Chandra} archive. We establish as a general property of ULXs that the most X-ray-luminous objects possess the flattest X-ray spectra (in the {\\it Chandra} bandpass). No prior sample studies have established the general hardening of ULX spectra with luminosity. This hardening occurs at the highest luminosities (absorbed luminosity $\\geq$5$\\times$10$^{39}$~erg~s$^{-1}$) and is in line with recent models arguing that ULXs are actually stellar-mass black holes. From spectral modeling, we show that the evidence originally taken to mean that ULXs are IMBHs - i.e., the ``simple IMBH model'' - is nowhere near as compelling when a large sample of ULXs is looked at properly. During the last couple of years, {\\it XMM-Newton} spectroscopy of ULXs has to a large extent begun to negate the simple IMBH model based on fewer objects. We confirm and expand these results, which validates the {\\it XMM-Newton} work in a broader sense with independent X-ray data. We find that (1) cool-disk components are present with roughly equal probability and total flux fraction for any given ULX, regardless of luminosity, and (2) cool-disk components extend below the standard ULX luminosity cutoff of 10$^{39}$~erg~s$^{-1}$, down to our sample limit of 10$^{38.3}$~erg~s$^{-1}$. The fact that cool disk components are not correlated with luminosity damages the argument that cool disks indicate IMBHs in ULXs, for which strong statistical support was never found. ",
        "introduction": "Ultraluminous X-ray sources (ULXs) have long been hailed as direct observational evidence for the existence of accreting intermediate-mass black holes \\citep[IMBHs;][]{col99}. The X-ray spectral model that has emerged as a central pillar for this argument (the ``simple IMBH model'') is that which is commonly applied as the canonical X-ray spectral fit to Galactic black-hole binaries with stellar mass black holes \\citep{mcc06}.  This model consists of a thermal accretion disk component plus a power-law (PL) continuum component. When applied to ULX spectra, the derived disk temperatures are 0.1$-$0.3~keV \\citep[e.g.][]{mill04}, much lower than Galactic black holes (at 0.6$-$1~keV). A cooler disk implies a bigger disk; so assuming that the disk approximately extends inward to the last stable orbit around the black hole, this would imply bigger, and more massive black holes. Such cool disks were indeed found in a few ULXs \\citep[e.g.][]{mill04}.  To counter this argument, many recent papers have pointed out both theoretical and observational problems with the simple IMBH model as a global explanation for all ULXs \\citep[e.g.][]{gon06, wil06, rob07}. The observed accretion disk components can be fairly weak; thus they do not provide a reliable measure of black hole mass. Also, the simple IMBH model does not necessarily approximate well the X-ray spectra of many ULXs. Attention has switched to perhaps less exotic models to explain some of the ULXs, such as beaming \\citep{king01} or super-Eddington accretion \\citep{beg02}, both of which explain ULX X-ray properties without the need for an IMBH. Galactic super-Eddington sources are known, such as stellar-mass black hole binaries like GRS~1915+105 \\citep{fen04}, V4641~Sgr \\citep{rev02} and possibly SS~433. The latter could be an example of both beaming {\\it and} super-Eddington emission, the combination of which could easily explain even the most luminous ULXs \\citep{beg06, pou07}. Cool accretion disks can also be physically explained by ``coupled disk-corona'' models \\citet{don06}, blurred emission and absorption lines from surrounding (outflowing) gas \\citep{gon06}, or a microblazar with magnetized jets \\citep{fre06}, that can transfer disk energy into the jet (thus making the disk fainter and cooling it the same time).  Recent detailed X-ray spectral modeling has revealed properties that further complicate any simple global interpretation, suggesting that multiple classes of ULXs exist.  Some very bright ULXs have been found by several authors \\citep{zez02, sor07, sorw06, soc06} to have relatively flat spectra, not usually expected in high states for accretion states of black holes \\citep{mcc06}. A flat spectrum suggests an inverse correlation between the slope of the spectrum and source luminosity \\citep[see also NGC~5204 X-1;][]{rob06}. Such an inverse correlation is hard to explain with current IMBH models, because in the typical high state the spectrum is soft, dominated by the disk component and with a steep PL \\citep{mcc06}. Specifically, {\\it XMM-Newton} spectroscopy of ULXs has to a large extent already begun to directly negate the simple IMBH model \\citep[see, e.g.][]{wil06, gon06}.  Two {\\it Chandra} surveys suggest that ULXs may in fact be an extension of normal lower-luminosity galaxy populations to higher luminosities. Using simple PL models applied to spectra with typically ~50 counts each, \\citet{swa04} compared ULXs to a lower luminosity sample of X-ray sources with L$_{X}=$~10$^{38}-$10$^{39}$~erg~s$^{-1}$ and found both distributions of photon indices to be well fitted by Gaussians centered at about 1.9. The samples also have similar X-ray colors, time series and positions within their host galaxies. In another {\\it Chandra} study, \\citet{col04} found no discernible difference between the X-ray colors of ULXs and lower luminosity sources in spiral galaxies. Both analyses were done with data of fairly poor spectral quality in terms of fitting detailed models and the latter relied on a color-color analysis rather than spectral fitting. These works also did not include two-component model fits that would identify spectral states and directly test the simple IMBH model  To search carefully for spectral properties that can differentiate ULXs, we need the best data available that will allow us to distinguish between simple spectral models. In this paper we use the highest quality X-ray spectra for a large, complete sample of ULXs from the {\\it Chandra} archives to test various ULX models.  We are able to provide a statistically strong comparison of the results with lower luminosity X-ray sources of equal data quality.  We pay special attention to properties usually associated with ULXs, such as the signature of a cool disk, which has never been tested for a uniform and large sample of good quality {\\it Chandra} spectra. Is the cool disk preferentially found in ULXs? If yes, does the disk dominate the total emission? We also search for other spectral behaviors found more recently in individual ULXs, such as a correlation between hardness and luminosity, and what this might mean.  For the first time, strong statistical tests of various ideas of ULX models can be provided to the ULX community.  With its unmatched spatial resolution, {\\it Chandra} is better suited than {\\it XMM-Newton} for studying point sources in crowded regions or resolving point sources in distant galaxies.  This is particularly true for the starburst galaxies that host populations of ULXs (e.g NGC 3256, Lira et al. 2002; Cartwheel galaxy, Gao et al. 2003), where only {\\it Chandra}'s unparalleled X-ray optics can spatially and spectrally resolve the emission of ULXs from that of the underlying galaxy. We have searched all public data available in the {\\it Chandra} archive for ULXs and lower-luminosity comparison objects with at least 1000 counts. In section 2 we present our source selection process, methods for identifying rejected objects, and an estimate of contamination from background objects. In section 3 we discuss the spectral fitting procedures and compare the spectral properties of the two samples. Our goal is to determine whether ULXs as a class have different spectral properties than the less luminous, ``normal'' X-ray sources and to offer an improved diagnosis by using the high-quality spectral data available in the {\\it Chandra} archive. In section 4 we present results from the variability analysis. Finally, in Section 5 we interpret our results and discuss the insight provided into the nature of ULXs. ",
        "conclusions": "From our X-ray spectral comparison between ultraluminous X-ray sources (ULXs) and other X-ray point sources in nearby galaxies, we find an interesting subclass of nine ULXs that have unique properties compared to the other sources that are classified as ULXs. This subclass of ULXs also differs from the lower-luminosity sample of X-ray point sources. While most of the ULXs we analyzed can plausibly be explained as scaled-up versions of Galactic black hole binaries, this particular subclass cannot. We discuss this subclass of ULXs followed by our general results, especially how our results relate to current evidence that supports the idea that ULXs host intermediate-mass black holes (IMBHs).  \\subsection{Luminous, Hard (flat-spectrum) X-ray ULXs}  Our analysis has identified nine ULXs with very high luminosities and hard (flat) X-ray spectra. This sample (U4, U5, U10, U11, U14, U18, U19, U20, and U43) is shown in Figure~3a (upper left corner). These ULXs are all well fitted by a power-law (PL) model with a photon index of $<$1.7. More complex spectral fitting using two component models reveals that when trying to add an accretion disk component (the MCD model described above), the contribution of this component to the total X-ray flux is very small, indicating that the relative contribution of emission from the accretion disk to the total spectrum is small. This result is shown in Figure~5, where the ratio of the absorbed MCD component flux to the total flux is plotted against the PL index for the PLMCD model. By using absorbed fluxes we do not specifically show the absolute physical strength of the accretion disk component; however, the ratios measured in this way better indicate the significance of detecting a soft excess and are also less dependent on the modeling.  The nine members of the subclass of flat-spectrum ULXs are also very luminous, so they would seem to be the best candidates for hosting an IMBH based on a simple Eddington limit argument, which predicates that higher mass black holes are required to explain the most luminous accreting sources. The spectra of these ULXs resemble Galactic black holes in a hard state, but such spectral shapes are usually associated with a low-luminosity state in the case of Galactic black holes. If these ULXs are indeed accreting IMBHs in a low state (i.e. low/hard state), our result begs the question as to why we do not also see ULXs in a high state (high/soft state) with even higher luminosities. Moreover, if these are IMBHs in a low state, such a scenario implies very high mass black holes ($>$10$^4$~M$_{\\odot}$). The formation of such black holes is not easy to explain.  It seems more plausible that this subclass of hard (flat) and luminous ULXs are accreting sources in the PL dominated very high state \\citep[VHS;][]{mcc06}, with an unusually weak soft X-ray component. A model that describes the properties of a hard PL with very little flux from the disk, at least in the {\\it Chandra} band, is the ``coupled disk-corona'' model proposed by \\citet{don06}. In this model, the underlying accretion disk emission is distorted by a process that drains energy from the disk into the corona. In an extreme case, the inner disk emission could be almost completely Comptonized, and thus only the visible outer disk would contribute to the accretion disk component. As this obviously only appears at low temperatures it could be easily absorbed in some galaxies and also hard to detect. However, if our nine luminous ULXs are interpreted as stellar-mass black hole systems in a high state, we would need to explain why their X-ray spectra are much harder compared to those observed in our Galaxy \\citep{mcc06}, which have typical photon indices of ~2.5 in the VHS. We would also need to explain how such low-mass black hole systems could reach such high luminosities. A model that could explain both the flat non-thermal component and the weak soft disk component in luminous ULXs has been proposed by \\citet{soc06}. Their ULX model shows that at super-Eddington accretion rates, in the inner region of the disk, magnetic fields in the corona can prevent strong winds, thus the radiative efficiency is not reduced by photon trapping. The resulting spectrum is dominated by the coronal emission from the inner region and the soft thermal component is generated only in the outer disk.   There is some direct evidence that high-luminosity states in ULXs correlate with the hardness of the PL tail in their spectra. \\citet{rob06} have shown this for a long {\\it Chandra} observing campaign of NGC~5204 X-1 (our U36). These data were not available when we searched the archive. \\citet{rob06} found that the spectrum becomes harder as the flux increases, over time-scales of days to weeks. The model used was a Comptonized disk model, and the results showed a cool disk ($\\sim$0.1~keV) and an optically thick corona. This model demonstrates that flux variations correlate with the corona temperature. \\citet{rob06} favor a stellar mass black hole interpretation for this ULX and suggest an unusual VHS, probably produced by extreme mass transfer from a massive star. We should note, however, that the PL slope is much steeper for NGC~5204 X-1 compared to the ULXs we discuss; therefore it is much easier to interpret as a VHS in comparison to what is observed for Galactic black holes.  It is possible that our nine ULX spectra appear to be harder than they actually are due to the limited energy band covered by {\\it Chandra}. If these objects intrinsically possess a break or curvature in their spectra and the break occurs at an energy above the {\\it Chandra} bandpass or where the sensitivity of {\\it Chandra} falls off significantly, this might bias our modeling to measuring these sources as ``hard''. This appears to be the case for one ULX in our sample, NGC 1313 X-1 (U5). Using {\\it XMM-Newton} spectra, \\citet{sto06} found evidence for a break (or curvature) at 4.9~keV, with a photon index for the high-energy PL of 2.16 (much closer to Galactic black holes values). The authors show, however, that such breaks are easier to explain if ULXs contain stellar mass black holes rather than IMBHs. The curvature would be likely to originate in optically thick coronae. This theory would need to be tested for the remaining ULXs in our subclass by obtaining better quality spectra.  In conclusion, the subclass of nine ULXs with very high luminosities and hard (flat) X-ray spectra suggests a PL-dominated VHS, in line with recent models of stellar mass black hole systems in very high accretion states. The fact that the highest luminosity ULXs are explained more easily with such models argues strongly against IMBHs as the only explanation of ULXs.    \\subsection{Cool disks and the IMBH interpretation}  Using our sample of ULXs and lower luminosity X-ray sources, we have found that the spectral signature of a cool accretion disc is not specific to ULXs. The results of the widely used PLMCD model (see Fig.~6 and Table~\\ref{table5}) show that many sources in both samples have MCD components with low inner-disk temperatures. Cool disks have been used until recently as support for the IMBH interpretation. Our results show that this evidence is nowhere near as compelling when a large sample of ULXs is looked at properly.  In the standard accretion disk model, cool disks are not expected for stellar-mass black holes accreting near their Eddington limits. The disk temperature scales with the black hole mass as T$_{in}\\propto$~M$^{-1/4}$, and is $\\sim$~1~keV for stellar mass black holes. However, cool disks can be seen in a low (hard) state, because the temperature dependence on the accretion rate for standard disks is T$_{in}\\propto~\\dot{M}^{-1/4}$ \\citep{mcc06,mill06}. Cool disks have indeed been found in some non-ULX sources \\citep[e.g.][]{sto06}. The authors note the similarity of these spectral fits with those typically used for ULXs. They also suggest that the soft excess in some cases could be otherwise explained by contamination from the host galaxy.  Most sources in Figure~6b possess low disk temperatures, within both the ULX and comparison samples. Indeed, it is surprising that we do not see many states that are typical (high) states for stellar mass black hole binaries, with a prominent $\\sim$1~keV blackbody component (although a hard PL tail would be difficult to discern in the limited {\\it Chandra} bandpass). Only U2, U37, C3 and C9 show such high temperatures. We find a similar result from our cool-disk model (PLMCD0.25). A large fraction of our ULX spectra (70\\%) are well fitted by this model, but a significant number of lower luminosity objects (59\\%) are as well.  There are theoretical models that do not require the presence of an IMBH to explain a cool disk at high accretion rates and high luminosities (i.e., for ULXs). We already mentioned the model proposed by \\citet{don06}, that explains cool disks by a process of draining energy from the disk to launch an optically-thick corona that obscures the hot inner regions of the disk. \\citet{fre06} developed microblazar models with magnetized jets that cause a transference of disk energy into the jet, thus making the disk fainter and cooling it at the same time. Other phenomenological models include the ``dual thermal'' model of \\citet{wil06}, in which the soft excess comes from an optically thick outflow produced at high accretion rates \\citep[see][]{king03}, which is seen in addition to a disk component with a temperature similar to those seen in stellar-mass black hole binaries. This latter model was proposed to explain the alternate model of \\citet{wil06}.  \\citet{don06} found evidence of a Galactic black hole that supports the interpretation of ULXs as stellar mass black holes in a VHS. The microquasar XTE J1550-564 has a ``strong VHS'' \\citep[see also][]{kub04}, where the disc temperature decreases with luminosity, reaching values of 0.3$-$0.4~keV. A similar behavior has been found for NGC~1313 X-2 \\citep{fen07}. This suggests a new type of VHS, a so-called ultraluminous branch, which is very similar to the ULX spectra \\citep{rob07,sor07}. In this interpretation, ULXs represent the high end of such an accretion state, with black hole masses up to 100~M$_{\\odot}$ and accretion rates up to 20 times the Eddington limit. Forming black holes with such masses is much easier to explain than forming IMBHs. For example \\citet{bel04} showed that black holes with masses of 80~M$_{\\odot}$ or more can form through binary mergers. \\citet{sorw06} suggested that black holes of up to 200~M$_{\\odot}$ could form by large-scale dynamical collapse of protoclusters in active regions in galaxies. These formation mechanisms are supported by the association between ULXs, star-forming regions and colliding galaxies.    \\subsection{Conclusion}  We have found that the highest luminosity ULXs tend to have the hardest X-ray spectra in the {\\it Chandra} bandpass and are well fitted by a simple power-law model, without evidence for thermal accretion disc components.  Such spectra are not consistent with current IMBH models, but are more in line with current models of extreme very high states, or perhaps a new ``ultraluminous state'' \\citep{rob07}, in stellar mass black holes.  Our work shows that cool accretion disks are not exclusive to the ULX class, suggesting that low-temperature IMBHs are not the only explanation for this phenomenon.  In general, our results show that ULXs are likely to be composed of several distinct types of objects, and that these types may extend into lower X-ray luminosity classes, such as classical Galactic black hole candidates and other objects in our comparison sample.  Our conclusions provide another ``nail in the coffin'' for assumptions that ULXs are simply a class of accreting IMBHs.  No other specific properties have been found for the ULX group, except for spectral hardening at the highest luminosities. All these results suggest that ULXs are the highest luminosity end of stellar mass black hole binaries, with the largest black holes permitted by current formation mechanisms and/or accreting at super-Eddington rates."
    },
    "0807/0807.3542_arXiv.txt": {
        "abstract": "We present optical and near-IR (OIR) observations of the major outbursts of the \\ns\\ soft \\xrt\\ binary system \\aql, from summer 1998 -- fall 2007. The major outbursts of the source over the observed timespan seem to exhibit two main types of light curve morphologies, (a) the classical {\\em Fast-Rise and Exponential-Decay} (FRED) type outburst seen in many soft X-ray transients and (b) the {\\em Low-Intensity State} (LIS) where the optical-to-soft-X-ray flux ratio is much higher than that seen during a FRED. Thus there is no single correlation between the optical (R-band) and soft X-ray (1.5-12 keV, as seen by the ASM onboard RXTE) fluxes even within the hard state for \\aql, suggesting that LISs and FREDs have fundamentally different accretion flow properties. Time evolution of the OIR fluxes during the major LIS and FRED outbursts is compatible with thermal heating of the irradiated outer accretion disk. No signature of X-ray spectral state changes or any compact jet are seen in the OIR, showing that the OIR \\cmd\\ (CMD) can be used as a diagnostic tool to separate thermal and non-thermal radiation from X-ray binaries where orbital and physical parameters of the system are reasonably well known. We suggest that the LIS may be caused by truncation of the inner disk in a relatively high \\mdot\\ state, possibly due to matter being diverted into a weak outflow. ",
        "introduction": "}  \\aql\\ is one of the most frequently recurring soft X-ray transients in the sky. It goes into outburst approximately once every year although the periodicity is not fixed \\citep{simon2002}. The frequent outburst recurrence makes it an ideal source for studying accretion flow during an outburst as well as changes in accretion flow from outburst to outburst. The physical and orbital parameters of this system are relatively well known. The orbital period of the binary system was determined to be $18.71\\pm0.06$ hours by \\citet{wry2000} from optical photometry during quiescence. \\citet{c1999} estimated the quiescent V-band magnitude to be $21.6\\pm0.10$ magnitude. However, there is a contaminating star about 0.48\\arcsec\\ East of \\aql\\ which has a V magnitude of $19.42\\pm0.06$ \\citep{c1999}. The distance to the system is between 4-6.5 kpc as inferred both from (a) limits on optical brightness of the Roche lobe filling donor star of spectral type between K6-M0 \\citep{c1999,r2001} and (b) peak X-ray fluxes during type-1 bursts \\citep{ccg1987,r2001}.  In systems like \\aql\\ and other Roche lobe overflowing \\sxt\\ (SXT) binary systems, the thermal viscous disk instability model \\citep[DIM;][]{l2001} is usually invoked to explain the typical {\\em fast-rise and exponential-decay} (FRED) type outburst light curve morphology \\citep{csl1997}. The DIM alone however cannot produce realistic FRED outbursts. Another mechanism that plays a key role during outbursts of SXTs is X-ray irradiation. The irradiation is caused by the intense X-rays from the inner regions of the accretion disk and \\ns\\ surface (if the system harbors a \\ns) impinging on the outer disk/donor, and by re-emission of this impinged energy in lower energies \\citep{c1976,v1990,vPM1994,m1995,d1999,obrien2002,h2002,r2006}. Another catagory of modifications to the DIM is generated by the possible truncation of the inner region of the accretion flow, where the disk inflow can be replaced by a radiatively inefficient accretion flow (RIAF). There is an increasing body of work on this topic \\citep{i1977,r1982,ny1994,bb1999,m2001}, but the overall impact of these efforts on the outburst morphology is currently still unclear.  Collimated relativistic outflows, loosely termed as ``jets'', have been observed in several \\sxt s \\citep[see][for a review]{f2006} and their broad band spectral energy distribution (SED) from radio to X-rays have been fitted with jet+disk models comprising of a freely expanding ``Blandford - K\\\"onigl'' jet \\citep{bk1979,fb1999,mnw2005} and a thermal-viscous ``Shakura-Sunyaev'' accretion disk \\citep{ss1973}. Observations suggest that compared to black hole (BH) systems, the neutron star (NS) SXTs tend to have intrinsically weaker jets and also the transition from optically-thick to optically-thin jet spectrum in NS systems probably occurs in the mid-IR \\citep[][this work]{mf2006,migliari2006}, compared to that in near-IR for BH systems \\citep{cf2002,bb2004,h2005,r2006}.  During outburst, light curve morphology of many sources are not FREDs. The X-ray spectral and temporal features also show an enormous diversity during outbursts of SXT systems, and are categorized into various {\\em states}. See e.g. \\citet{mr2006} and \\citet{hb2005} for extensive discussions on definition of states and \\citet{vdk2006} for their timing properties. Despite the multitude of states observed in the different systems and different outbursts, two ``canonical'' states can be identified during the outburst of most SXTs, characterised by presence or absence of a thermal and nonthermal component in the energy spectrum. In the temporal domain, presence or absence of a thermal component usually shows up respectively as low or high rms variability in the light curve. A host of other intermediate and extreme states are also seen, which could be a superposition of the thermal and nonthermal components with varying contributions. Precise determination of X-ray state requires modeling of spectral energy distribution and Fourier power density spectra. Since this work primarily focusses mainly on OIR observations, we have used dates of X-ray state transitions of \\aql\\ from published studies of X-ray spectral and temporal properties of the source \\citep{mc2003a,mc2003b,rks2004,mb2004,rodriguez2006,y2007} to study correlation between X-ray state and OIR flux-density.  We have observed \\aql\\ for $\\sim9$ years, through many outbursts, to explore the nature of outburst morphologies in a single object. Even though the OIR (as well as X-ray) light curve morphologies from outburst to outburst differ, we find systematic and repeated tracks traced by the source in OIR \\cmd s. Since we have long-term data in two bandpasses (R and J) only, we have used a simple model to show that the color evolution and the resulting track are consistent with heating of a constant area blackbody and can be caused by an irradiated outer disc.  In \\S\\ref{sec:data} we describe the general reduction procedures adopted in this work. The light curve morphologies in J, R and ASM bands (of the RXTE satellite) and optical--soft X-ray correlations during the major outbursts are described in \\S\\ref{sec:ltc}. Evolution of \\aql\\ on the OIR \\cmd\\ is presented in \\S\\ref{sec:colors}. Finally we present our discussions in \\S\\ref{sec:discuss}. ",
        "conclusions": "}  Multi-wavelength monitoring of the neutron star soft X-ray transient \\aql\\ in near-IR, optical and soft X-ray data from RXTE/ASM over the past nine years has allowed us to study the evolution of the source with unprecedented coverage. The OIR and X-ray light curves show a wide range of outburst light curve patterns. However the major outbursts show two distinct light curve morphologies. The {\\em fast-rise and exponential-decay} \\citep[FRED;][]{csl1997} type outbursts characteristic of many soft X-ray transients is seen in \\aql\\ too. However, only three of the major outbursts seen since 1998 June show just a FRED type light curve profile. More often a prolonged, somewhat variable ``Low-Intensity state'' \\citep[LIS;][]{w2002}, uniquely characterized by an extremely high optical-to-soft X-ray flux-ratio, is seen either before or after a FRED. It has been observed that for some of these ``soft X-ray faint'' LIS outbursts, \\aql\\ never comes out of the X-ray hard state during the entire outburst \\citep{rodriguez2006}. Other than these major outbursts, small-scale flaring activities were also observed. While the origin of such small-scale short activity remains unknown, such activity has been observed in other sources like XTE J1550-564 by \\citet{ss2005} who suggested that these small outbursts could be due to discrete accretion events. We suggest the following framework for understanding these results:  {\\em Irradiation dominates OIR flux for both FRED and LIS outbursts of \\aql.} During the course of FRED and LIS outbursts of \\aql\\, the evolution of the source on the OIR \\cmd\\ is consistent with heating/cooling of an irradiated outer accretion disk as a constant area \\bb. A classical viscous accretion disk (e.g. a Shakura-Sunyaev disk) spectrum has a frequency dependence of $F_{\\nu}\\sim \\nu^{1/3}$ between $kT(R_{out})/h \\ll \\nu \\ll kT(R_{in})/h$ where  $T(R_{in})$ and $T(R_{out})$ are the disc temperatures at the inner and outer edges of the disk. At frequencies much smaller than $kT(R_{out})/h$ the disk spectrum is simply the Rayleigh-Jeans region of a \\bb\\ with $F_{\\nu}\\sim \\nu^2$. For soft X-ray transients $kT(R_{in})/h$ is typically in far-UV or soft X-rays. Therefore in our case the flux ratios (a.k.a. colors) of two bands, from a purely viscous disc are expected to be constant as long as {\\em both bands} are either on the viscous ($\\nu^{1/3}$) or Rayleigh-Jeans ($\\nu^2$) regime. In the special case when the break frequency between the viscous dissipation and the Rayleigh-Jeans regime given by $\\nu_{break}=kT(R_{out})/h$ lies between R and J bands, the spectral index should lie between $1/3$ and $2$. The observed systematic and continuous change in color (and the corresponding range in the spectral index from $-0.95$ to $0.27$; \\S\\ref{sec:colors}) seen in Fig.~\\ref{r-j_1} and Fig.~\\ref{r-j_2} is strongly suggestive of an irradiative origin for the observed OIR emission and not of a purely viscous origin.  Neither the OIR color nor its brightness change sharply during an X-ray spectral state transition. This confirms that for \\aql\\, the outer accretion disk, which emits most of the thermal OIR radiation, is not affected by X-ray spectral state transitions which change the physical environment of the inner regions of the accretion disk. This is in sharp contrast with \\bh\\ binaries like 4U 1543-47, GX 339-4 \\citep{bb2004,h2005} where near IR colors are strongly correlated with X-ray state changes and the correlation interpreted as evidence for non-thermal emission. In FREDs, a slow increase in brightness above that predicted by our constant area \\bb\\ model, as the source becomes bluer and brighter, is most likely due to increase in the surface area of the disk (e.g most prominent for outburst \\#2 and \\#5 in Fig.~\\ref{aql-cmd1}). It is possible that at highest mass accretion rates the high X-ray flux from the central inner disk causes the flaring of the outer disk and hence increase the effective surface area. In contrast, during the LISs, when the source probably spends the whole period of activity in hard X-ray state, there does not seem to be any significant change in the disk area (e.g. outburst \\#7 in Fig.~\\ref{aql-cmd2} ). Therefore it is likely that the effect of X-ray state transition on the outer accretion disk is both an increase of the outer disk temperature due to enhanced irradiation from the inner disk and an increase in the emitting area of the outer disk due to flaring. Total irradiation (hence OIR flux) is a bit smaller in LIS than in peak FRED. This could be due either to lower \\mdot, or to radiative inefficiency lowering $L/\\mdot$ ratio.  {\\em Broad band quasi-simultaneous observations in LIS suggest outflow.} Soft X-ray flux is very low in LIS, but hard X-rays are high \\citep[][Linares et al. in preparation,]{atel259}. While absence of soft, thermal X-ray photons suggests that the inner edge of the accretion disk is cold and probably receded during the LIS, the high OIR brightness (see Fig.~\\ref{allfluxes} for the dramatic difference in $f_{opt}/f_{X-ray}$ between LIS and FRED outbursts ) as well as type I X-ray bursts suggests that a significant amount of mass is being accreted by the outer accretion disk and fed to the inner region.  Radio detection of the source during LIS \\citep{atel286} with flat spectral slope hints toward a jet outflow. However, since the evolution of the source on OIR \\cmd\\ is consistent with irradiation, it can be concluded that the jet, if present, is either quite weak in terms of the input power, or the transition from optically thick to optically thin jet emission occurs at a wavelength much longer than J-band. Current models for jets from compact systems suggest that this might imply a large jet launching area or the jet base, or much weaker particle acceleration \\citep{mff2001}. Thus our observations support the current hypothesis that jet-luminosity is not significant in most neutron star X-ray binary systems \\citep{mf2006,f2006}.  {\\em A possible scenario for the LIS might be that an \\mdot\\ that would ordinarily be high enough to move the disk inwards, and thus trigger a soft state, somehow doesn't do that, and launches an outflow instead.} Thus the standard processes associated with the DIM are disrupted by the absence of an inner disk in this high \\mdot\\ state. If the accretion flow in the inner region in a LIS is radiatively inefficient, it can drive the observed outflow, as suggested by recent theoretical works \\citep[see e.g.][]{bb1999,m2001,n2005}."
    },
    "0807/0807.4362_arXiv.txt": {
        "abstract": "We present the science database produced by the Formation and Evolution of Planetary Systems (FEPS) \\Spitzer\\ Legacy program. Data reduction and validation procedures for the IRAC, MIPS, and IRS instruments are described in detail. We also derive stellar properties for the FEPS sample from available broad-band photometry and spectral types, and present an algorithm to normalize Kurucz synthetic spectra to optical and near-infrared photometry. The final FEPS data products include IRAC and MIPS photometry for each star in the FEPS sample and calibrated IRS spectra. ",
        "introduction": "The Formation and Evolution of Planetary Systems (FEPS) \\Spitzer\\ Legacy program \\citep{Meyer06} was designed to characterize the evolution of circumstellar gas and dust around solar-type stars between ages of 3~Myr and 3~Gyr. To achieve these goals, FEPS obtained spectro-photometric observations with the \\Spitzer\\ Space Telescope \\citep{Werner04} for a sample of 328 stars \\citep[see][for a description of the sample]{Meyer06}. The observing strategy was to measure the spectral energy distribution (SED) between wavelengths of 3.6\\micron\\ and 70\\micron\\ with IRAC \\citep[InfraRed Array Camera;][]{Fazio04} and MIPS \\citep[Multiband Imaging Photometer for Spitzer;][]{Rieke04} photometry, and between 8 and 35\\micron\\ with low-resolution IRS \\citep[Infrared Spectrograph;][]{Houck04} spectra. In addition, the FEPS program obtained MIPS~160\\micron\\ photometry for 80 stars to search for colder dust, and high-resolution IRS spectra for 33 sources to probe for circumstellar gas.  The FEPS team has produced several studies on the incidence of dusty debris disks around solar type stars, including the discovery of a debris system in the initial \\Spitzer\\ observations \\citep{Meyer04}, a census of warm debris \\citep{Stauffer05,Silverstone06,Hines06,Meyer08}, the identification of Kuiper-belt analogs \\citep{Kim05,Hillenbrand08}, and an investigation of debris disks around stars with known planets \\citep{Moro07}. The FEPS team has also analyzed the processing of dust in optically thick, primordial disks \\citep{Bouwman08}, and has produced a series of papers on the evolution of gas in solar-type stars \\citep{Hollenbach05,Pascucci06,Pascucci07}.  This paper describes the data reduction procedures for IRAC (3.6, 4.5, and 8\\micron) and MIPS (24 and 70\\micron) images and IRS low resolution spectra obtained by the FEPS program. Data reduction methods for the MIPS~160\\micron\\ images and IRS high resolution spectra are discussed in Kim \\etal~(in preparation) and \\citet{Pascucci06}, respectively. The adopted reduction procedures for the IRAC, MIPS~24\\micron, MIPS~70\\micron, and IRS observations are presented in \\S\\S\\ref{irac}-\\ref{irs}. We also investigate the effects of source confusion on the 24\\micron\\ and 70\\micron\\ photometry (\\S\\ref{confusion}) and the relative calibration accuracy between \\Spitzer\\ instruments (\\S\\ref{crosscal}). The series of FEPS papers frequently utilized synthetic spectra derived from Kurucz model atmospheres to infer the presence of infrared excesses diagnostic of circumstellar dust. In the Appendices, we describe the data and algorithm used to obtain normalized synthetic spectra for individual stars. The primary data products from the FEPS program are a tabulation of IRAC and MIPS photometry presented in Table~\\ref{tbl:phot}, and extracted, calibrated spectra which are available electronically. ",
        "conclusions": "\\label{summary}  The FEPS \\Spitzer\\ Legacy program was designed to obtain infrared photometry from 3.6 to 160\\micron\\ and low resolution spectra from 5 to 35\\micron\\ for 328 solar-type stars spanning ages from 3~Myr to 3~Gyr. The broad goal of FEPS was to determine the incidence of circumstellar disks and place the results in context with the expected evolution of our Solar System. An essential component of this study was to construct carefully calibrated spectral energy distributions. Here, we outline the data reduction procedures adopted by the FEPS team to obtain accurate and well-characterized \\Spitzer\\ photometry and spectra.  The adopted image processing steps for the IRAC, MIPS, and IRS data closely follow the recommended procedures by the \\Spitzer\\ Science Center and \\Spitzer\\ Instrument Teams. We describe in detail the data reduction methods for each instrument and the procedures used to validate the data products. We present in Table~\\ref{tbl:phot} the measured IRAC (3.6, 4.5, and 8\\micron\\ bands) and MIPS (24 and 70\\micron) flux densities and uncertainties. The extracted, calibrated IRS spectra are available electronically."
    },
    "0807/0807.2687_arXiv.txt": {
        "abstract": "It is often assumed that primordial perturbations are statistically isotropic, which implies, among other properties, that their power spectrum is invariant under rotations.  In this article, we test this assumption by placing model-independent bounds on deviations from rotational invariance of the primordial spectrum. Using five-year WMAP cosmic microwave anisotropy maps, we set limits on the overall norm and the amplitude of individual components of the  primordial spectrum quadrupole.  We find that there is no significant evidence for primordial isotropy breaking, and that an eventually non-vanishing quadrupole  has to be subdominant. ",
        "introduction": "Observations of the cosmic microwave background show that the early universe contained tiny density perturbations, from which structures developed as the universe expanded.  In recent times, we have gained a wealth of information about  these primordial perturbations. We have precisely measured their spectrum, and we have placed quite stringent limits on their properties \\cite{Komatsu:2008hk}. Prompted by these advances, the nature of the primordial perturbations has entered the standard cosmological model, a set of a few parameters and assumptions that summarizes what we know about our universe.  There is however an assumption in the standard cosmological model that has not been subject to much observational or theoretical scrutiny: the statistical isotropy of the primordial perturbations. Cosmological perturbations are statistically isotropic if their probability distribution functionals are invariant under rotations, which implies in particular that the power spectrum of statistically isotropic perturbations only depends on the magnitude of the wave vector, $\\mathcal{P}(\\mathbf{k})=\\mathcal{P}(k)$. Though some papers have analyzed the impact of statistically anisotropic perturbations on structure \\cite{ArmendarizPicon:2005jh,Pullen:2007tu, Ando:2008zz}, while others have proposed mechanisms for their generation \\cite{Ackerman:2007nb,ArmendarizPicon:2007nr,Pereira:2007yy,Gumrukcuoglu:2007bx,Akofor:2007fv,Erickcek:2008sm,Yokoyama:2008xw},  to date no model-independent limits on the deviations from statistical isotropy of the primordial perturbations exist.\\footnote{On the other hand, the statistical isotropy of the CMB itself has been extensively investigated: see \\cite{Souradeep:2006dz} and references therein.}  In this article we set  precise bounds, and thus verify one of the key ingredients in our understanding of the origin of structure. ",
        "conclusions": "Inspection of Table \\ref{tab:limits} quickly reveals that the amplitude of the quadrupole components  is consistent with statistical isotropy. In particular, if there is a non-vanishing quadrupole in the primordial spectrum, it clearly has to be subdominant. The results in Table \\ref{tab:criteria} also imply that there is no evidence for primordial statistical anisotropy. Although a non-zero primordial quadrupole significantly increases the likelihood, the information criteria that penalize the introduction of additional parameters strongly favor isotropy.  Because any deviation from statistical isotropy can be cast as in equation (\\ref{eq:P}), the limits that we have found are mode-independent, and can thus be directly applied to any of the models for the generation of (adiabatic) statistically anisotropic perturbations discussed in the literature \\cite{Ackerman:2007nb,ArmendarizPicon:2007nr,Gumrukcuoglu:2007bx,Akofor:2007fv,Erickcek:2008sm,Yokoyama:2008xw}. We have not studied how our bounds constrain the parameters of these models, but it should be straight-forward to do so. On the other hand, our null results confirm again the predictions of the simplest inflationary models.  The study of the statistical isotropy of the primordial perturbations is still in its infancy, and our analysis is just a first step toward  preciser measurements of the primordial spectrum.  With more data, improved analysis techniques, and better control of systematics, it should be possible in principle to obtain much tighter constraints \\cite{Pullen:2007tu}."
    },
    "0807/0807.0227_arXiv.txt": {
        "abstract": "We present a catalogue of 99 candidate clusters and groups of galaxies in the redshift range 0.1 $< z_{phot} <$ 1.3 discovered in the $Spitzer$ First Look Survey (FLS).  The clusters are selected by their R$_{c}$ - 3.6$\\micron$ galaxy color-magnitude relation using the cluster red sequence algorithm. Spectroscopic redshifts from numerous FLS followup projects confirm the photometric redshifts of 29 clusters and demonstrate that the R$_{c}$ - 3.6$\\micron$ red sequence color provides photometric redshifts with an accuracy of $\\Delta$z = 0.04 in the redshift range 0.1 $< z <$ 1.0.  Using this cluster sample we compute the 3.6$\\micron$, 4.5$\\micron$, 5.8$\\micron$, \\& 8.0$\\micron$ cluster luminosity functions (LFs).  Similar to previous studies, we find that for the bands that trace stellar mass at these redshifts (3.6$\\micron$, 4.5$\\micron$) the evolution in M$^*$ is consistent with a passively evolving population of galaxies with a high formation redshift ($z_{f}$ $>$ 1.5). Using the 3.6$\\micron$ LF as a proxy for stellar luminosity we remove this component from the MIR (5.8$\\micron$ \\& 8.0$\\micron$) cluster LFs and measure the LF of dusty star formation/AGN in clusters. We find that at $z <$ 0.4 the bright end of the cluster 8.0$\\micron$ LF is well-described by a composite population of quiescent galaxies and regular star forming galaxies with a mix consistent with typical cluster blue fractions; however, at $z >$ 0.4, regular star forming galaxies are insufficient to account for the excess of 8.0$\\micron$ galaxies, and an additional population of dusty starburst galaxies is required to properly model the 8.0$\\micron$ LFs. Comparison to field studies at similar redshifts shows a strong differential evolution in the field and cluster 8.0$\\micron$ LFs with redshift. At $z \\sim$ 0.65 8.0$\\micron$-detected galaxies are more abundant in clusters compared to the field, but thereafter the number of 8.0$\\micron$ sources in clusters declines with decreasing redshift and by  $z\\sim$ 0.15, clusters are underdense relative to the field by a factor of $\\sim$ 5.  The rapid differential evolution between the cluster and field LFs is qualitatively consistent with recent field galaxy studies that show the star formation rates of galaxies in high density environments are larger than those in low density environments at higher redshift. ",
        "introduction": "Since the compilation of the first large samples of galaxy clusters almost 50 years ago (Zwicky 1961; Abell 1958), clusters have been used as fundamental probes of the effect of environment on the evolution of galaxies. Over this time, our understanding of this phenomenon has grown significantly, and a basic picture of the formation and evolution of cluster galaxies between 0 $< z <$ 1 has emerged.  Studies of the stellar populations of cluster galaxies via the fundamental plane (e.g., van Dokkum et al. 1998; van Dokkum \\& Stanford 2003; Holden et al. 2005) and the evolution of the cluster color-magnitude relation (e.g., Ellis et al. 1997; Stanford et al. 1998; Gladders et al. 1998; Blakeslee et al. 2003; Holden et al. 2004; Mei et al. 2006; Homeier et al. 2006; Tran et al. 2007) have shown that the majority of stars in cluster galaxies are formed at high-redshift ($z >$ 2) and that most of the evolution thereafter is the passive aging of these stellar populations.  Studies of the evolution of the near-infrared (NIR) luminosity functions (LFs) of clusters have shown that not only are the stellar populations old, but that the bulk of the stellar mass is already assembled into massive galaxies at high-redshift (e.g., De Propris et al. 1999; Toft et al. 2003; Strazullo et al. 2006; Lin et al. 2006; Muzzin et al. 2007a).  Furthermore, it appears that the cluster scaling relations seen locally ($z <$ 0.1, e.g., Lin et al. 2004; Rines et al. 2004; Lin et al. 2003), such as the Halo Occupation Distribution, Mass-to-Light ratio, and the galaxy number/luminosity density profile are already in place by at least $z \\sim$ 0.5 (e.g., Muzzin et al. 2007b; Lin et al. 2006). \\newline\\indent These studies suggest a picture where the formation of the stars in cluster galaxies, as well as the assembly of the galaxies themselves occurs at a higher redshift than has yet been studied in detail; and that, other than the passive aging of the stellar populations, clusters and cluster galaxies have changed relatively little since $z \\sim$ 1.  This picture appears to be a reasonable zeroth-order description of the evolution of cluster galaxies; however, there are still properties of the cluster population which cannot be explained within this context. In particular, there are significant changes in the morphology (Dressler et al. 1997; Postman et al. 2005, Smith et al., 2005), color (e.g., Butcher \\& Oemler 1984; Rakos \\& Schombert 1995; Smail et al. 1998; Ellingson et al. 2001; Margoniner et al. 2001; Loh et al. 2008) and star-formation properties (e.g., Balogh et al. 1999; Dressler et al. 1999; Poggianti et al. 1999; Dressler et al. 2004; Tran et al. 2005a; Poggianti et al. 2006, although see Kodama et al. 2004) of cluster galaxies since $z \\sim$ 1.  The fraction of blue, star forming galaxies increases from almost zero at $z =$ 0 to as much as 50\\% at $z \\sim 0.5$ (the so-called Butcher-Oemler Effect), and correspondingly, the fraction of S0 galaxies in clusters drops by a factor of 2-3, with similar increase in the number of spiral/irregular galaxies over the same redshift range (Dressler et al. 1997).  Naively, these results suggest that gas-rich, star-forming galaxies at high-redshift have their star-formation truncated by the cluster environment at moderate redshift and become the dominate S0 population seen locally.  How such a transformation occurs, and how it avoids leaving a notable imprint on the stellar populations, is still not well-understood. \\newline\\indent Citing an abundance of post-starburst (k+a) galaxies in clusters at $z \\sim$ 0.4, Poggianti et al. (1999) and Dressler et al. (2004) suggested that there may be an abundance of dusty starburst galaxies in clusters at moderate redshift, and that the dusty starburst and k+a galaxies may represent the intermediate stages between regular star forming late-type galaxies and S0 galaxies (e.g., Shioya et al. 2004; Bekki \\& Couch 2003).  In particular, they suggested that the cluster e(a)\\footnote{An e(a) galaxy is defined as a galaxy with EW([OII]) $<$ -5\\AA$ $ and EW(H$\\delta$) $>$ 4\\AA$ $ by Dressler et al. (1999). These are emission line galaxies with a strong A star component to their spectrum suggesting a recent, possibly obscured, burst of star formation.} galaxies would be the best candidates for dusty starburst galaxies because their inferred star formation rates appear larger from H$\\alpha$ emission than from [OII] emission.  If the cluster environment excites a dusty starburst from harassment, tidal interaction, or ram-pressure stripping, then this may quickly deplete a star forming galaxy of its gas, transforming it first into a k+a galaxy, and then leaving it an S0.  More detailed work on two $z \\sim$ 0.5 clusters by Moran et al. (2005) also showed an abundance of starbursting galaxies conspicuously near the cluster virial radius, suggesting a environmental origin to their ``rejuvenation''. $ISO$ observations of relatively nearby clusters have detected significant amounts of dust-obscured star formation (e.g., Fadda et al. 2000; Duc et al. 2002; Biviano et al. 2004; Coia et al. 2005), and this has recently been confirmed at even higher redshift ($z =$ 0.2 - 0.8) by $Spitzer$ observations (Geach et al. 2006; Marcillac et al. 2007; Bai et al. 2007; Fadda et al. 2007; Saintonge et al. 2008; Dressler et al. 2008).  Despite this, it is currently unclear whether there is a  population of dusty starbursts which is sufficiently abundant to be the progenitors of the large number of cluster k+a galaxies. \\newline\\indent Alternatively, there is evidence from other cluster samples that the S0 population may simply be the result of the truncation of star formation in infalling late-type galaxies via gas strangulation (e.g., Abraham et al. 1996; Balogh et al. 1999; Treu et al. 2003; Moran et al. 2006) and that no accompanying starburst occurs.  Most likely, the star formation and morphology of galaxies are transformed both ``actively'' (as in a starburst triggered from merging/harassment/tidal forces) and ``passively'' (from gas strangulation or ram-pressure stripping), and that the magnitude of each effect varies significantly from cluster-to-cluster and possibly by epoch, which may explain why studies of small numbers of clusters have found discrepant results.  Interestingly, both processes can be active within massive clusters as was demonstrated by Cortese et al. (2007), who found two interesting galaxies in Abell 1689 and Abell 2667, one of which seems to be undergoing gas strangulation and ram-pressure stripping, while the other is experiencing an induced starburst. There is evidence that galaxies in clusters that are less dynamically relaxed have larger star formation rates (e.g.,  Owen et al. 1999; Metevier et al. 2000; Moss \\& Whittle 2000; Owen et al. 2005; Moran et al. 2005; Coia et al. 2005, and numerous others) and that the accretion of large substructures induces starbursts from harassment and tidal forces. \\newline\\indent The most obvious way to understand whether dusty starbursts are important in the evolution of cluster galaxies is to observe their abundances directly in the mid-infrared (MIR).  In particular, differences in the MIR LFs of the cluster and field environments can be used to determine if dusty starbursts are more common in the cluster environment.  If so, it would suggest that environmental processes may be responsible for triggering these events. \\newline\\indent The InfraRed Array Camera (IRAC) onboard $Spitzer$ provides a unique tool for studying this problem.  IRAC images in 4 bands simultaneously (3.6$\\micron$, 4.5$\\micron$, 5.8$\\micron$, 8.0$\\micron$) and this is particularly advantageous because 3.6$\\micron$ and 4.5$\\micron$ observations are a good proxy for the stellar mass of cluster galaxies between $0 < z < 1$, and 5.8$\\micron$ and 8.0$\\micron$ are sensitive to emission from warm dust (i.e., from dusty star forming regions) over the same redshift range.  In particular, the Polycyclic Aromatic Hydrocarbons (PAHs) emit strong line emission at rest frame 3.3$\\micron$, 6.2$\\micron$, 7.7$\\micron$, 8.6$\\micron$, and 11.3$\\micron$ (e.g., Gillett et al. 1973; Willner et al. 1977).  These features, in addition to the warm dust continuum, are sensitive indicators of dusty star formation, and several studies have already shown a good correlation between 8.0$\\micron$ flux and star formation rate (SFR\\footnote{Although there is a direct correlation between 8$\\micron$ flux and SFR, the scatter in the correlation is approximately a factor of 2 for metal rich galaxies in the local universe, and metal poor galaxies can deviate by as much as a factor of 50 (e.g., Calzetti et al. 2007).  Because of the large scatter and metallicity dependence, throught this paper we do not use the 8$\\micron$ data to quantitatively measure SFRs. Instead, we use the presence of enhanced 8$\\micron$ flux as a qualitative indicator of increased dusty star formation.}; e.g., Calzetti et al. 2005; Wu et al. 2005; Calzetti et al. 2007).  Therefore, examining the suite of IRAC cluster LFs at redshifts 0 $< z <$ 1 shows both the evolution of the majority of stellar mass in cluster galaxies, as well as the evolution of dusty star formation in the same galaxies. \\newline\\indent The obvious approach to measuring the presence of dusty star formation in clusters is to observe a handful of ``canonical'' galaxy clusters with IRAC.  However, given that determining the LF from a single cluster suffers significantly from Poisson noise, and perhaps most importantly,  is not necessarily representative of the average cluster population at a given mass/epoch, a better approach would be to stack large numbers of clusters in order to improve the statistical errors, and avoid peculiarities associated with individual clusters.  This approach requires targeted observations of numerous clusters, which is time-consuming compared to other alternatives.  For example, large-area $Spitzer$ surveys such as the 50 deg$^2$ $Spitzer$ Wide-area Infrared Extragalactic Survey (SWIRE\\footnote{SWIRE data are publically available at http://swire.ipac.caltech.edu/swire/}, Lonsdale et al. 2003), the 8.5 deg$^2$ IRAC Shallow Survey (Eisenhardt et al. 2004), and the 3.8 deg$^2$ $Spitzer$ First Look Survey (FLS\\footnote{The FLS data are publically available at http://ssc.spitzer.caltech.edu/fls/}, Lacy et al. 2005) are now, or soon-to-be, publically available and these fields already contain significant amounts of optical photometry.  These wide optical-IRAC datasets can be employed to find clusters in the survey area itself using optical cluster detection methods such as the cluster red sequence (CRS) technique (Gladders \\& Yee 2000, hereafter GY00), or photometric redshifts (e.g., Eisenhardt et al. 2008; Brodwin et al. 2006).  Subsequently, the IRAC survey data can be used to study the LFs of clusters at a much larger range of masses and redshifts than could be reasonably followed up by $Spitzer$. Furthermore, these surveys also provide panoramic imaging of clusters out to many virial radii, something that has thus far rarely been attempted because it is time-consuming. \\newline\\indent Finding clusters with the CRS algorithm is relatively straightforward with the ancillary data available from these surveys.  The technique exploits the fact that the cluster population is dominated by early type galaxies, and that these galaxies form a tight red sequence in color-magnitude space.  If two filters which span the 4000\\AA$ $ break are used to construct color-magnitude diagrams, early types are always the brightest, reddest galaxies at any redshift (e.g., GY00) and therefore provide significant contrast from the field.  The CRS technique is well-tested and provides photometric redshifts accurate to $\\sim$ 5\\% (Gilbank et al. 2007a; Blindert et al. 2004) as well as a low false-positive rate ($<$5\\%, e.g., Gilbank et al. 2007a; Blindert et al. 2004; Gladders \\& Yee 2005).  The method has been used for the 100 deg$^2$ red sequence Cluster Survey (RCS-1, Gladders \\& Yee 2005) and is also being used for the next generation, 1000 deg$^2$ RCS-2 survey (Yee et al. 2007).  Variations of the red sequence method have also been used to detect clusters in the Sloan Digital Sky Survey (the ``BCGmax'' algorithm, Koester et al. 2007; Bahcall et al. 2003) as well as in the fields of X-ray surveys (e.g. Gilbank et al. 2004; Barkhouse et al. 2006). \\newline\\indent In this paper we combine the $Spitzer$ FLS R$_{c}$-band and 3.6$\\micron$ photometry and use it to detect clusters with the CRS algorithm. Given the depth of the data, and that the R$_{c}$ - 3.6$\\micron$ filter combination spans the rest-frame 4000\\AA$ $ break to $z >$ 1, we are capable of detecting a richness-limited sample of clusters out to $z \\sim$ 1. Using the sample of clusters discovered in the FLS we compute the 3.6$\\micron$, 4.5$\\micron$, 5.8$\\micron$, and 8.0$\\micron$ LFs of clusters 0.1 $< z <$ 1.0 and study the role of dusty star formation in cluster galaxy evolution.  A second paper on the abundance of dusty starburst galaxies detected at 24$\\micron$ in the same clusters using the FLS MIPS data is currently in preparation by Muzzin et al. (2008). \\newline\\indent The structure of this paper is as follows.  In \\S 2 we give a brief overview of the optical, IRAC, and spectroscopic data used in the paper.  Section 3 describes the cluster-finding algorithm used to detect clusters and \\S 4 contains the FLS cluster catalogue, and a basic description of its properties.  In \\S 5 we present the IRAC cluster LFs and \\S 6 contains a discussion of these results as well as a comparison of the cluster and field LFs.  We conclude with a summary in \\S 7.  Throughout this paper we assume an $\\Omega_{m}$ = 0.3, $\\Omega_{\\Lambda}$ = 0.7, H$_{0}$ = 70 km s$^{-1}$ Mpc $^{-1}$ cosmology.  All magnitudes are on the Vega system. ",
        "conclusions": "We have presented a catalogue of 99 candidate clusters and groups at 0.1 $< z_{phot} <$ 1.3 discovered in the $Spitzer$ First Look Survey using the cluster red sequence technique.  Using spectroscopic redshifts from FLS followup campaigns and our own spectroscopic followup of clusters we have shown that the R - 3.6$\\micron$ color of the cluster red sequence is an accurate photometric redshift estimator at the $\\Delta$z = 0.04 level at $z <$ 1.0.  Furthermore, we demonstrated that the properties of the FLS cluster catalogue are similar to previous cluster surveys such as the RCS-1.  Using this cluster sample we studied the evolution of the cluster 3.6$\\micron$, 4.5$\\micron$, 5.8$\\micron$, and 8.0$\\micron$ LFs.  The main results from these LFs can be summarized as follows: \\begin{itemize} \\item In agreement with previous work, the evolution of the 3.6$\\micron$ and 4.5$\\micron$ LFs between 0.1 $< z <$ 1.0 is consistent with a passively evolving population of galaxies formed in a single-burst at $z >$ 1.5. Given that the 3.6$\\micron$ and 4.5$\\micron$ bandpasses are reasonable proxies for stellar mass, this suggests that the majority of stellar mass in clusters is already assembled into massive galaxies by $z \\sim$ 1. \\item The MIR color cuts used to select AGN by Lacy et al. (2004) and Stern et al. (2005) suggest that the fraction of cluster galaxies that host MIR-bright AGN at $z <$ 0.7 is low.  We estimate that the AGN fraction of cluster galaxies detected at 3.6$\\micron$ is 1$^{+1}_{-1}$\\%.  AGN are a larger, but still modest component of the 5.8$\\micron$ and 8.0$\\micron$ cluster population,  approximately 3$^{+3}_{-3}$\\% of these galaxies. \\item The cluster 5.8$\\micron$ and 8.0$\\micron$ LFs do not look similar to the 3.6$\\micron$ and 4.5$\\micron$ LFs, and this is due to the presence of  the cluster star forming galaxies. Star forming galaxies are much brighter in these bandpasses than early type galaxies and their varying fractions with redshift cause deviations from the shape of the 3.6$\\micron$ and 4.5$\\micron$ LFs.  The 5.8$\\micron$ and 8.0$\\micron$ LFs are well-described using different fractions of three basic types of galaxies: quiescent, regular star forming, and dusty starburst by assuming that the fractions of the latter two are proportional to the cluster f$_{b}$. \\item The 8.0$\\micron$ cluster LFs suggest that both the frequency and SSFR of star forming cluster galaxies is increasing with increasing redshift.  In particular it appears that when compared to star forming galaxies in the local universe, the intensity of star formation in clusters evolves from ``weak'' to ``regular'' to ``starburst'' with increasing redshift.  Qualitatively, this evolution mimics the evolution in the universal star formation density with redshift suggesting that this evolution is at least in part caused by the accretion of star forming galaxies into the cluster environment. \\item Comparing the 3.6$\\micron$ and 4.5$\\micron$ cluster and field LFs with similar normalization shows that the LFs in these environments are similar, with evidence for a small excess in the brightest galaxies in clusters, likely caused by the cluster giant ellipticals. In agreement with previous K-band studies this suggests that the distribution of galaxies as a function of stellar mass in both environments is roughly equivalent. \\item There is a significant differential evolution in the cluster and field 5.8$\\micron$ and 8.0$\\micron$ LFs with redshift.  At $z =$ 0.65 the cluster is more abundant in 8.0$\\micron$ galaxies than the field; however, thereafter the relative number of 5.8$\\micron$ and 8.0$\\micron$ galaxies declines in clusters with decreasing redshift and by $z =$ 0.15 the cluster is underdense in these sources by roughly a factor of 5. This differential evolution could be explained if starbursts are preferentially triggered during the early formation stages of the cluster but then preferentially quenched thereafter by the high density environment. \\end{itemize} \\indent A well-sampled spectroscopic study of several high-redshift clusters with MIR data would be extremely valuable for verifying our interpretation of the IRAC cluster LFs because it is always difficult to draw incontrovertible conclusions from LFs alone.  Still, the cluster LFs do show a strong increase in the number of 5.8$\\micron$ and 8.0$\\micron$ sources in clusters with increasing redshift which must almost certainly be attributed to increased amounts of dusty star formation in higher redshift clusters. \\newline\\indent One of the strengths of this analysis is that it is based on a relatively large sample of galaxy clusters.  It has become clear from the handful of clusters studied thus far by $ISO$ and $Spitzer$ that the MIR properties of cluster galaxies can be quite different from cluster to cluster.  They may depend  on dynamical state, mass, f$_{b}$, or other parameters (e.g., Coia et al. 2005; Geach et al. 2006).  The advantage of using many clusters is that it provides a metric of how the ``average'' cluster is evolving as a function of redshift.  Detailed studies of individual clusters with significant ancillary data will pave the way to a better understanding of the physics behind the evolution of dusty star formation in cluster galaxies; however, large statistical studies such as this one will indicate whether the clusters studied in future work are representative of the cluster population as a whole, or are potentially rare, biased clusters with unusual properties caused by an ongoing merger or some other event. \\newline\\indent It is worth noting that although the quality of the LFs provided by the 99 clusters in the FLS is good, these LFs would still benefit from a larger statistical sample.  In particular, a larger sample would allow for the separation of clusters by other properties such as mass or morphology, and to understand if these properties play a role in shaping the MIR cluster galaxy population.  We are currently working on a survey to detect clusters in the much larger SWIRE survey: the Spitzer Adaptation of the Red sequence Cluster Survey (SpARCS).  This project has 13 times more area than the FLS and is a factor of 2 deeper in integration time in the IRAC bands.  The analysis of that sample should provide a significant improvement in the quality of the cluster LFs."
    },
    "0807/0807.3971_arXiv.txt": {
        "abstract": "We present a numerical model describing a circularly symmetric gaseous disk around the Be star $\\chi$~Ophiuchi.  The model is constrained by long-baseline interferometric observations that are sensitive to the H$\\alpha$ Balmer line emission from the disk.  For the first time our interferometric observations spatially resolve the inner region of the circumstellar disk around $\\chi$~Ophiuchi and we use these results to place a constraint on the physical extent of the H$\\alpha$-emitting region.  We demonstrate how this in turn results in very specific constraints on the parameters that describe the variation of the gas density as a function of radial distance from the central star. ",
        "introduction": "Models of the circumstellar gas around some B-type stars have been constructed for a number of years.  Because the recombination in the circumstellar gas produces a series of emission lines, most commonly in the hydrogen Balmer series, such systems are usually denoted B-emission or Be stars for short.  Models describing the geometry of this circumstellar material have varied from spherical shells~\\citep{Gehrz74}, to slabs with constant thickness~\\citep{Kastner89}, to disks with constant opening angles~\\citep{Waters86}, to disks with an exponential density fall-off in the direction perpendicular to the plane of the disk~\\citep{Marlborough69}.  Although the first models assumed a constant temperature throughout the disk, more recent models, such as the ones developed by \\citet{Millar98, Millar99}, \\citet{Carciofi06}, and \\citet{Sigut07}, obtain a self-consistent temperature throughout the disk by enforcing radiative equilibrium (namely, by balancing the microscopic heating and cooling rates operating within the disk).  While the temperature structure of Be star disks is now the subject of sophisticated modeling, the density structure largely remains {\\it ad-hoc\\/} with simple, parameterized models. Perhaps the most popular is a power-law density decrease with radius in the equatorial plane coupled with the aforementioned exponential drop perpendicular to the disk. This vertical density structure is assumed to result from the gravitational equilibrium set by the balance between the gradient in the gas pressure and the vertical component of the star's gravitational acceleration. Despite the simplicity of such models, they have been quite successful in matching a wide range of observational diagnostics that are sensitive to the geometric distribution of the circumstellar gas. These include the infrared excess exhibited by Be stars over normal B stars of the same spectral class \\citep{dou94}, the observed continuum linear polarization of Be stars~\\citep{woo97}, and the shapes of the emission lines observed in Be star spectra \\citep[see for example][]{hum00a}. All of these diagnostics rely only on spatially unresolved spectra or spectropolarimetry.  However, even more powerful constraints on the geometry of the emitting regions in Be stars can be obtained by combining these spatially {\\it unresolved\\/} diagnostics with new optical interferometric observations that directly {\\it resolve\\/} the structure of the emitting region on the sky~\\citep{Quirrenbach97, Tycner06b, Gies07}.   To demonstrate how specific constraints can be placed on a disk model using observations that spatially resolve the source, we chose to concentrate on the Be star $\\chi\\;$Oph (B2Ve; HR~6118; HD~148184). This relatively nearby star with a {\\it Hipparcos} distance of $150\\pm 17$~pc~\\citep{Perryman97} has a strong H$\\alpha$ emission and therefore is an ideal candidate for probing the disk structure at these wavelengths. ",
        "conclusions": "It is useful to compare our best-fit disk model for $\\chi$~Oph to previous results of \\citet{Waters86} who modeled the IRAS IR excesses of a large sample of Be stars with disk models with fixed opening angles of $15^{\\circ}$ and where the density varied only with the distance from the stellar center, $r$, in the form of \\begin{equation} \\label{eq:rhorens} \\rho(r) = \\rho_0' (r/R_{\\star})^{-n} . \\end{equation} Assuming an isothermal disk with a temperature of 18,000~K, a temperature that is considerably larger than the $\\sim$10,400~K density-weighted average temperature of our best-fit model, \\citet{Waters86} found a radial power-law index $n$ of 2.4 and a density at the base of the disk of $\\log \\rho_0=-11.4$.  These results compare well to our results of $n=2.5$ and $\\log\\rho_0=-10.7$, even though the disk-density models used in the two studies are not the same.  The smaller base disk density found by \\citet{Waters86} is likely the result of the fact that his density distribution is constant along radial arcs over the entire opening angle of the disk which places more material in the disk as compared to our models.   \\citet{Porter99} modeled the circumstellar regions of Be stars as isothermal viscous disks and predicted $n$ of $7/2$.  Viscous models that drop the assumption of an isothermal disk can find a range of $n$ values as the disk fills or empties \\citep{JSP08}.  \\citet{Porter99} also analyzed the IR excess of $\\chi$~Oph using the data of \\citet{Waters86}, but with a disk model density essentially equivalent to ours. In the case of an isothermal disk (also set to 18,000~K), \\citet{Porter99} found an index $n$ of 2.2 with $\\log \\rho_0 =-11.2$; in the case of a disk in which the density-averaged temperature was allowed to vary with radial distance as a power-law, \\citet{Porter99} obtained $n=1.9$ and $\\log(\\rho_0)=-11.2$.  Porter indicates that the non-isothermal disk formally fits the IR data best, but that the improvement is not large.  In addition to the comparison between model and observed H$\\alpha$ interferometric visibilities and the H$\\alpha$ line profile, it is also instructive to compare the predicted IR excess of our best-fit model to the IR observations reported in the literature. Figure~\\ref{fig:IR-fit} illustrates the comparison between the SED of our best-fit model and the visual and IR observations from \\citet{Waters86}.  The overall agreement is quite reasonable, especially considering that the model was not fit to the SED data, although the observations in the $1-10\\,\\mu$m region fall below the model.  We should also point out that variability between the observations used by \\citet{Waters86} and our 2006 H$\\alpha$ observations (interferometric and spectroscopic) cannot be ruled out. For example, \\citet{ban00} obtained the H$\\alpha$ profile for $\\chi$~Oph in 1998 and reported an equivalent width of approximately $-3.6$~nm, as compared to the $-7.1\\pm0.2$~nm found in the present work for the 2006 epoch~(recall Fig.~\\ref{fig:Halpha}). Hence, the H$\\alpha$ emission of $\\chi\\,$Oph has approximately doubled in 8 years, and assuming that the H$\\alpha$ emission was similarly weaker at the time of observations of \\citet{Waters86} we can then expect our best-fit model to over-predict IR excess as compared to Waters' results.  For this reason, we feel that our model cannot be improved using the IR observations that were acquired using the IRAS satellite in the early 1980's."
    },
    "0807/0807.3140_arXiv.txt": {
        "abstract": "In recent years a plethora of future surveys have been suggested to constrain the nature of dark energy. In this paper we adapt a binning approach to the equation of state factor ``w'' and discuss how future weak lensing, galaxy cluster counts, Supernovae and baryon acoustic oscillation surveys constrain the equation of state at different redshifts. We analyse a few representative future surveys, namely DES, PS1, WFMOS, PS4, EUCLID, SNAP and SKA, and perform a principal component analysis for the ``w'' bins. We also employ a prior from Planck cosmic microwave background measurements on the remaining cosmological parameters. We study at which redshifts a particular survey constrains the equation of state best and how many principal components are significantly determined. We then point out which surveys would be sufficiently complementary. We find that weak lensing surveys, like EUCLID, would constrain the equation of state best and would be able to constrain of the order of three significant modes. Baryon acoustic oscillation surveys on the other hand provide a unique opportunity to probe the equation of state at relatively high redshifts. ",
        "introduction": "One of the most challenging puzzles in modern physics is the cause for the observed accelerated expansion of the Universe. The renewed interest in accelerating cosmologies was born out of Type Ia Supernovae observations a decade ago \\citep{Perlmutter:97,Riess:98,Perlmutter98}. In the meantime the cosmic concordance model \\citep{Ostriker:95} has been confirmed by complementary observations of the Cosmic Microwave Background (CMB) \\citep{Spergel03}, the large scale distribution of galaxies \\citep{Efstathiou02,Tegmark:04} and clusters of galaxies \\citep{Allen:04,Rapetti:04}. In addition observations of large scale CMB anisotropies \\citep{Scranton:03, Padmanabhan:04} and Baryon Acoustic Oscillations (BAO) in the matter distribution \\citep{Cole:05,Eisenstein:05} also indicate the presence of late time acceleration in the Universe.  The simplest way to model an accelerated expansion in the Universe is by including a cosmological constant in Einstein's equations of gravity. However, the introduction of a cosmological constant requires an extreme fine tuning of the initial conditions of our Universe to about 120 orders of magnitude. This motivates the search for alternative explanations for the speeding up of the expansion rates, although there have been recent successes in obtaining universes with the observed cosmological constant in a range of the multi verses of the meta-stable vacua of the string theory landscape \\citep{Kachru:03}. This, in connection with the anthropic principle \\citep{Weinberg:87,Efstathiou:95,Bousso:06,Peacock:07,Cline:07}, could motivate the observed value of the cosmological constant. Nevertheless, it is still valid to seek other alternatives to a cosmological constant. One approach is to introduce a dynamical scalar field, usually dubbed Quintessence, which is either slowly rolling down a potential or is trapped in a false vacuum state \\citep{Wetterich:88,Peebles:88,Ratra:88,1995PhRvL..75.2077F,Ferreira:98,Zlatev:98,Albrecht:00}. These models can lead to cosmic acceleration with less fine tuning. Another possibility was that the Universe is permeated by a network of cosmic defects, which also acts like a source in the energy-momentum tensor and can lead to accelerated expansion \\citep{Battye:99,Battye:05b}. The dynamical evolution of the Universe in Quintessence or defect based models is governed by the equation of state of the dominant components. The equation of state is given by the ratio of the pressure to the density of these fluids. For a cosmological constant this is -1. However dynamical dark energy models can deviate from this value and the ratio can also vary with time. Typically the ratio is named ``$w=p_{\\rm de}/\\rho_{\\rm de}$''. The precise measurement of the value of ``$w$'' is hence one of the foremost tasks for observational cosmology. Any significant measurement of a deviation of ``$w$'' from -1, would be a major result. Currently conservative constraints on ``$w$'' are in the region of $w=-1\\pm 0.15$ \\citep{Tegmark:04,Astier:06}.  Up to now, it has been difficult to make any connection between suggested Quintessence fields and fundamental theories. Another possibility to obtain accelerated expansion is the modification of Einstein's gravity on large scales. There are two approaches along these lines. One is to modify a proposal by \\cite{Starobinsky:80}, which includes higher order curvature terms in the gravitational action to obtain accelerated expansion in the early Universe. This has been adapted to explain the observed accelerated expansion \\citep{Carroll:04,Carroll:05}. Another approach is to obtain effective four dimensional Friedman equations, from a five or higher dimensional theory, where our standard model is confined to a brane, with only gravity allowed to leak out into the extra dimension(s) \\citep{Dvali:00,Deffayet:02,Dvali:03}. Although these models might have structural problems such as ghosts \\citep{Koyama:05,DeFelice:06} it is a worthwhile exercise to explore their cosmological phenomenology. In particular they are different to Quintessence models in the way structures grow in the Universe \\citep{Song:05,Koyama:06}. Finally it is possible, at least in theory, that the observed acceleration is an effect, which results from an averaging procedure to describe the large scale dynamics of an intrinsically inhomogeneous Universe \\citep{Buchert:05}.  All the models described above have in common that they can, but don't have to, deviate from the evolution of a so called $\\Lambda$-cold dark matter ($\\Lambda$CDM) Universe consisting of a cosmological constant and dark matter. Effectively the background evolution for all these models can be mimicked by a suitable chosen additional component in the energy-momentum tensor. Hence the quest for ``$w$'' goes beyond its usual application to Quintessence models, although they are the only models where there is actually a physical meaning associated with ``$w$''. It is therefore an interesting question to ask if ``$w$'' is different from -1 the value for a cosmological constant. The first attempt to address this question is to look for different {\\em  constant} values of ``$w$''. However, this approach is very restrictive given that dark energy models allow for evolving equation of state factors ``w'' \\citep{Weller:00,Weller01,Maor:2001ku}. Hence it is important to allow for a time varying equation of state parameter ``$w$''. A simple Taylor expansion in redshift \\citep{Weller01} is too restrictive if we want to include data sets which are sensitive to high redshifts, like the CMB. A more successful approach is to parameterise a time varying equation of state with a Taylor expansion in the scale factor $a$, like $w=w_0+w_a(1-a)$ \\citep{Chevallier:01,Linder:03}. This was also the parameterisation of choice adapted by the Dark Energy Task Force (DETF) \\citep{DETF:06}. The DETF chose this parameterisation to be able to compare different proposed surveys on an equal footing. In addition this parameterisation allows to some extend the distinction of two subclasses of Quintessence models, namely freezing and thawing out models \\citep{Linder:06}.  In this paper we want to pursue a more general approach. In order to figure out the redshift sensitivity to the dark energy density, or better its time derivative, we will bin the equation of state in redshift \\citep{Huterer:03,Crittenden:05,Huterer:07a,Albrecht:07,Sullivan:07}. Although, there are some ambiguities how to exactly perform the binning \\citep{dePutter:07}, the binning approach is in general the most model independent approach to fit for the background evolution of the Universe. Note that this approach is inherently related to the weighting function method \\citep{Saini:03,Simpson:06} in the limit of an infinite number of bins.  We are now faced with the difficult task of selecting surveys to pursue our analysis. We emphasize that our selection is subjective and our aim is that we have at least one representative survey for each of the probes we are discussing. In this introduction we will mention briefly the surveys and what their technical specification is expected to be, while we describe how they exploit particular cosmological probes in the sections where we discuss constraint from the different probes.  We start with Stage II surveys according to the classification by the DETF \\citep{DETF:06}. Already online in its simplest configuration is the PANoramic Survey Telescope and Rapid Response System (PAN-Starrs)\\footnote{see at: {\\tt pan-starrs.ifa.hawaii.edu}}. PAN-Starrs in it's full configuration will consist of four 1.8 metre telescopes, equipped with a wide field camera with a field of view of 7 arcminutes and a total of 1.4 Gigapixels. Currently only one of the four planned telescopes is deployed. We refer to this setup as PAN-Starrs1 (PS1) in contrast to the full setup PAN-Starrs4 (PS4). The PAN-Starrs survey is planned  to take exposures of $\\sim 60$ seconds in each of the grizy filters in the PS1 configuration and should go much deeper in the PS4 configurations with grizy filters. One of the surveys with PS1 will be a nearly all sky, $20.000\\,{\\rm deg}^2$, survey up to a median redshift of $z\\approx 0.5$. The nearly all sky survey is hoped to be exploited for large scale structure and BAOs \\citep{Peebles:70,Hu98;3,Blake:03} experiments as well as weak gravitational lensing observations. The medium deep survey will be suitable for cosmological constraints with Type Ia SNe out to redshifts of $z\\sim 0.8$. PS1 in the DETF category is a Stage II survey, while PS4 is Stage III.  Another survey already deployed, which is relevant to dark energy science, is the South Pole Telescope (SPT) \\citep{Ruhl:04}. The SPT is a 10 metre telescope with a bolometer array in its focal plane, designed to detect thousands of clusters of galaxies via their Sunyaev-Zel'dovich \\citep{Sunyaev:72} decrement in the sub-millimeter domain over $4.000\\,{\\rm deg}^2$. Current plans are to operate at 150, 219 and 274 GHz. This survey in addition with a redshift survey will allow to constrain dark energy with the clustering and redshift distribution of galaxy clusters.  We describe now near-future surveys, which are currently not deployed. The Dark Energy Survey (DES) \\citep{DES} is an optical-near infrared survey under construction. The optical part will go down to 24$^{\\rm th}$ magnitude in the grizy bands on the 4 metre Blanco telescope at Cerro Tololo Inter-American Observatory (CTIO). The IR counterpart will come from part of the Vista Hemisphere Survey (VHS) survey performed by ESA's 4m Visible and Infrared Survey Telescope for Astronomy (VISTA) in the JHK bands to $\\sim$20{\\rm th} magnitudes. The DES camera has a 2.2 deg field of view, which allows to obtain imaging and photometric informations of millions of galaxies. It is planned to observe $5.000\\,{\\rm deg}^2$, which overlap with the area of the SPT survey. This will allow to constrain dark energy via weak lensing tomography, galaxy clustering and galaxy cluster redshift distribution. In addition there is a wide-area survey where 10\\% of the allocated time will be used to follow Type Ia SNe light-curves.  The Wide Field MultiObject Spectrograph (WFMOS) \\citep{Glazebrook:05} will consist of a new wide-field (1.5 deg) 4000 fibers spectrograph with a passband of 0.39-1.0 microns which would be deployed in a 8{\\rm m} class telescope. Current suggestions are that there will be 10 low dispersion spectrographs with R=1800 in the blue and R=3500 in the red. This will enable the measurements of BAOs in $0.5<z<1.3$ and $2.3<z<3.3$ using the redshifts of millions of galaxies over $\\sim 2,000\\,{\\rm deg}^2$ at low z and $\\sim 500\\,{\\rm deg}^2$ at high z.  We will now introduce two proposed satellite missions, which might come online in the long term. Both NASA and ESA in their ``Beyond Einstein'' and ``Cosmic Vision'' studies have proposed to pursue missions which are able to probe dark energy\\footnote{see at: {\\tt universe.nasa.gov/home.html} and {\\tt www.esa.int/esaSC/SEMA7J2IU7E\\_index\\_0.html}}. One of the contenders for NASA is the SuperNovae Acceleration Probe (SNAP) \\citep{SNAP:05a,SNAP:05b}. SNAP is a 2 metre telescope in space with a 0.7 square-degree wide field imager and a $R\\sim 100$ spectrograph. Both are sensitive in the 0.4-1.7 $\\mu$m wave-band. SNAP is designed to probe dark energy with a SuperNovae and weak lensing survey, where the weak lensing takes advantage of the 9 filters with a depth of 26.6AB magnitude. We will discuss the particulars of the two surveys in the relevant sections later in the paper. The EUCLID survey is a combination of the former SPACE \\citep{space} and DUNE proposals. The Dark UNiverse Explorer (DUNE) \\citep{Refregier:06} is a proposed wide field space imager on a 1.2 metre telescope with a 1 deg$^2$ visible/near-IR field of view. It is designed to measure cosmic lensing over 20.000 deg$^2$ of the sky and will exploit this as its main cosmological probe. DUNE is designed to use one broad visible band (R+I+Z) for accurate shape measurement for weak lensing and Y,J,H in the near infrared to complement optical photometry which should be available by 2017 for accurate photometric redshifts in the range $0<z<2$.  The Square Kilometer Array (SKA) is a radio interferometer planned to operate in a large range of frequencies (60{\\rm MHz} - 35{\\rm GHz}) with a sensitivity of 20,000${\\rm m^2/K}$ in the range 0.5-1.4{\\rm GHz} with a field of view of $\\sim$ 200{\\rm deg$^2$} at 500{\\rm MHz} \\citep{2004astro.ph..9274C}. This interferometer would allow to locate galaxies in the Universe given their 21{\\rm cm} line emission and would allow us to perform large surveys for galaxy evolution and Cosmology \\citep[][Abdalla et al. in preparation] {2004NewAR..48.1013R,2004NewAR..48.1063B,2005MNRAS.360...27A, 2007MNRAS.381.1313A}. The project is planned to be an on-going project which should build up from $1\\%$ demonstrators in the following years to a $10-15\\%$ core in $\\sim 5-10$ year time to its full completion in $\\sim 2020$.  This concludes our preview of the surveys being discussed in this paper. However we will exploit one more survey, not for its ability to constrain dark energy, rather for its ability to constrain other cosmological parameters to high precision. The forthcoming Planck satellite mission will observe the sky in 9 radio wave bands in order to measure the anisotropies to in the CMB \\citep{Planck}. Given that the primary anisotropies are mainly a probe of the angular diameter distance to the surface of last scattering, Planck alone is not a strong probe of dark energy. However it puts strong constraints on other cosmological parameters\\citep{Planck}. We will hence use forecasts for the Planck surveyor to put prior constraints on the remaining cosmological parameters.  We start the paper by discussing the method of principal components analysis (PCA) in the context of the binning of the equation of state in Section \\ref{sec:pc}. In Section \\ref{sec:cmb} we discuss constraints arising from the Planck surveyor and introduce the covariance matrix we will use in the later sections of the paper. In Section \\ref{sec:sne} we study the principal components of Type Ia Supernovae as cosmological probes. In this section we will also analyse the impact of marginalization of the remaining cosmological parameters. In Section \\ref{sec:wl} we will study the PCA in the context of cosmic lensing. Section \\ref{sec:cc} will analyse cluster counts as cosmological probes. In the Section \\ref{sec:bao} we will examine the ability of BAOs to constrain dark energy. Before concluding in Section \\ref{sec:conc} we will discuss how to perform a joint principal component analysis between complementary surveys. ",
        "conclusions": "\\label{sec:conc} In this paper we have studied future constraints on $w$-bins for four dark energy probes, Type Ia Supernovae, weak lensing, cluster counts and baryon acoustic oscillations.  Instead of assuming that the equation of state parameter $w$ can be modeled by a simple function with a few parameters, we bin $w$ given by Eqn.~(\\ref{eq:w_z}) and treat each $w_{\\rm i}$ as independent parameter. Throughout this article, we choose fixed $\\Delta z=0.05$ to make sure that we can compare between different surveys. For each probe, we choose a few representative future surveys. We also use a prior from the forthcoming Planck CMB experiment as prior on all remaining cosmological parameters. We find that in future mainly weak lensing and BAO surveys are complementary, in a sense that weak lensing gives tight constrain at low redshifts but BAO allows one to push to higher redshifts. However the role of Supernovae is also complementary for low redshifts. The high redshift sensitivity of the BAO would allow one to study also early dark energy models \\citep{2003ApJ...591L..75C} in an efficient way. However at lower redshift weak lensing and Supernovae are the more efficient probes, with the exception of SKA BAO measurements. Typically forthcoming surveys constrain 2-3 modes. Since it is likely that there will be a multiple of stage III probes it is sensible to look at the combined power of these probes and we find that in this case five modes would be constraint. We would advocate that the number of well constraint modes is actually a better figure of merit than the one put forward by the DETF, since it encompasses the ability of future surveys not just to measure the equation of state at two redshifts, today and at the pivot point. We address the question of the number of modes in a Bayesian model selection way \\citep{Saini:04,Liddle:04}. Although this method has recently attracted some criticism \\citep{2008arXiv0802.3185E} it provides a well defined framework to establish the significant number of modes for a particular survey. It takes into account the goodness of fit, the bias and Occam's razor. However it is not only important to establish the significant number of modes. In addition the redshift distribution of the significant modes has to be taken into consideration to start talking about the complementarity of surveys. The complementarity can than be exploited to constrain eight modes for the joint stage III probes and twelve for stage IV if we include SKA. One question we have not addressed in this paper is how useful surveys are to probe ``dark energy'' beyond the background evolution, for example by probing structure formation. It is likely that in this context future weak lensing, and to some extend galaxy cluster surveys, will play a pivotal role. To conclude we would argue that all four discussed probes have their merits, but weak lensing and baryon acoustic oscillation surveys seem to bear the largest promise for revealing the nature of the accelerated expansion of the Universe."
    },
    "0807/0807.1974_arXiv.txt": {
        "abstract": "InfraRed Dark Clouds (IRDCs) are extinction features against the Galactic infrared background, mainly in mid-infrared band. Recently they were proposed to be potential sites of massive star formation. In this work we have made $^{12}$CO, $^{13}$CO, and C$^{18}$O (J=1$\\rightarrow$0)\\ survey of 61 IRDCs, 52 of which are in the first Galactic quadrant, selected from a catalog given by Simon et al. (2006), while the others are in the outer Galaxy, selected by visually inspecting the MSX images. Detection rates in the three CO lines are 90\\%, 71\\%, and 62\\% respectively. The distribution IRDCs in the first Galactic quadrant is consistent with the 5 kpc molecular ring picture, while slight trace of spiral pattern can also be noticed, which needs to be further examined. The IRDCs have typical excitation temperature of 10 K and typical column density of several $10^{22}$ cm$^{-2}$. Their typical physical size is estimated to be several $pc$s using angular sizes from the Simon catalog. Typical volume density and typical LTE mass are $\\sim$5000 cm$^{-3}$ and $\\sim$5000 M$_\\odot$ respectively. The IRDCs are in or near virial equilibrium. The properties of IRDCs are similar to those of star forming molecular clumps, and they seem to be intermediate between giant molecular clouds and Bok globules, thus they may represent early stages of massive star formation. ",
        "introduction": "The infrared dark clouds (IRDCs), revealed by the Midcourse Space Experiment (MSX), are dark extinction features against the Galactic mid-infrared background \\citep{Egan1998a}, usually with filamentary or compact shapes, mostly at 8.6 -- 10.6 ~$\\mu$m. High opacity suggests the IRDCs might have high gas column density. \\citet{Carey1998a, Carey2000a} showed that the IRDCs have column densities as high as $10^{23}$ cm$^{-2}$, gas density $\\sim$$10^5$ -- $10^6$ cm$^{-3}$, with distances to the sun $\\sim$1 -- 8 kpc, and diameters $\\sim$0.4 -- 15 pc. The masses of the clouds they observed are estimated to be from tens of to thousands of solar masses, while the temperatures are typically 10 -- 25 K. Based on millimeter-to-mid-IR continuum observation of an MSXDC, \\citet{Rathborne2005a} found very high luminosities, 9000 -- 32,000 L$_\\odot$, of the cores, and they concluded that it must be forming massive stars. \\citet{Ragan2006a} mapped 41 IRDCs in N$_2$H$^+$ 1$\\rightarrow$0, CS 2$\\rightarrow$1, and C$^{18}$O 1$\\rightarrow$0. They found that different species often show striking differences in morphologies, which were attributed to differences in evolutionary state and/or the presence of undetected, deeply embedded protostars. Average mass of the clouds is estimated to be $\\sim$2500 M$_\\odot$ using N$_2$H$^+$ observations, which is consistent with the previous studies of massive star-forming regions. The typical line width of the clouds they observed is 2 -- 3 km s$^{-1}$. \\citet{Rathborne2006a} found that the mass spectrum of IRDC cores derived from millimeter maps is consistent with the stellar IMF. Assuming each core will form a single star, they concluded that the majority of these cores will form OB stars. The IRDC cores are similar to the hot cores associated with individual, young high-mass stars, except that they are much colder, thus \\citet{Rathborne2006a} suggested that the IRDCs represent an earlier evolutionary phase in massive star formation, and they may be cold precursors to star clusters. A water maser survey of 140 IRDC compact cores of \\citet{Wang2006a} revealed that the detection rate of H$_2$O masers for higher mass cores is significantly higher than that of lower mass cores. They suggest that the most massive IRDC cores without H$_2$O maser may be at an earlier stage than the protostellar phases. A survey in\\ $^{13}$CO by \\citet{Simon2006b} established kinematic distances to 313 IRDCs. They derived typical sizes of $\\sim$5 pc, peak column densities of $\\sim$10$^{22}$ cm$^{-2}$, LTE masses of $\\sim$5$\\times10^3$ M$_\\odot$, and volume-averaged H$_2$ densities of $\\sim$2$\\times10^3$ cm$^{-3}$. \\citet{Beuther2007a} discovered a protostar in an IRDC, which will probably become a massive star. These previous studies suggest that IRDCs can be ideal candidates to study the initial conditions and early stages of massive star formation.  An IRDC catalog has been published by \\citet{Simon2006a}. Clouds in this catalog are in the first and fourth quadrant of the Galaxy. The authors used the MSX A band data to construct a background model, and then searched for regions with rapid decrease with respect to this background. The IRDC candidates are defined by contiguous regions bounded by closed contours of 2$\\sigma$ decremental contrast threshold. The catalog contains 10,931 IRDCs, and 12,774 cores, with reliability estimated to be 82\\%. The distribution of IRDCs is in good correspondence with the diffuse Galactic infrared background. There are more IRDCs toward the star-forming regions, the spiral arm tangents, and the 5 kpc molecular ring.  As part of our campaign to search for candidates of massive star forming clouds, we selected a sample of 61 clouds mainly from the IRDC catalog given by \\citet{Simon2006a} and observed them in $J=1\\rightarrow0$ lines of $\\rm ^{12}CO$, $\\rm ^{13}CO$, and $\\rm C^{18}O$ in single point mode. In this article we report the results of this survey. ",
        "conclusions": ""
    },
    "0807/0807.1235_arXiv.txt": {
        "abstract": "The five planets discovered around the main-sequence star 55 Cnc may represent a case of stable chaos. By using both the Frequency Map Analysis and MEGNO we find that about 15 \\% of the systems that can be build from the nominal orbital elements of the system are highly chaotic.  However, in spite of the fast diffusion rate in the phase space, the planetary system is not destabilized over 400 Myr and close encounters between the planets are avoided. ",
        "introduction": "\\label{sec:intro}  The most crowded extrasolar planetary system discovered so far is that around the main-sequence star 55 Cnc(= $\\rho^1$ Cancri). Doppler shift measurements strongly suggest the existence of five planets orbiting the star with semimajor axes ranging from 0.038 to 5.901 AU \\citep{fima} . The innermost and smallest planet, $e$, is a Neptune-mass object while the outer and most massive one, planet $d$,  has a minimum mass of about 4 Jupiter masses.  A self-consistent dynamical fit to the stellar wobble data performed by \\cite{fima} gives a set of orbital elements for the planets reported in Table 1. It is claimed \\citep{fima} that this nominal system of five planets described in Table 1 is dynamically stable at least over a timescale of 1 Myr. We would like to point out that planets $b$ and $c$ are no longer in a 3:1 mean motion resonance as in the previous solution given by \\cite{McA}. In addition, the eccentricity of planet $c$ in \\cite{fima} is significantly smaller compared to that derived by \\cite{McA}.  We have performed a detailed exploration of the stability of the nominal system by applying the Frequency Map Analysis (hereinafter FMA) method \\citep{las1,las2,las3,masch} on 400 varied systems. These systems are obtained by varying orbital elements of the nominal system. The varied systems are analysed by adding randomly some inclination (lower than 5 deg) to the Keplerian orbits to make the system more realistic. The masses of the planets are scaled accordingly. Among all the systems analysed with FMA we obtain cases with large diffusion speed in the phase space suggesting a chaotic evolution and possible instability. However, integrating a few of these chaotic systems over timescales of $10^8$ yrs, we never obtain destabilization. We, therefore, think that these chaotic system are examples for stable chaos. This term was introduced by \\cite{mino} to indicate the peculiar behaviour of asteroids with short Lyapunov times (of the order of some $10^3$ yrs) which however show a remarkable stability over the age of the solar system \\citep{mino}. For the 55Cnc planetary system we find that in about 15\\% of all cases the orbital elements of the four inner planets exhibit a chaotic evolution on a timescale of a few million of years. This stable chaos may be due to the topology of the phase space: Quasi--periodic solutions in the proximity of a chaotic island can act like ``quasi--barriers'' for the diffusion \\citep{tsiva}. The chaotic evolution appears to be confined and the system does not destabilize at least on a timescale of 400 Myr. This peculiar behaviour may be due to the closeness of the system to the planetary 3:1 mean motion resonance between planets $b$ and $c$ and secular resonances.  To confirm that stable chaos is not limited to the nominal case but can be retrieved also in a larger region of the phase space, we randomly varied the orbital elements of the nominal case and found a similar behaviour. For a few selected cases, we also applied the MEGNO method (the acronym of Mean Exponential Growth factor of Nearby Orbits) proposed by Cincotta \\& Sim\\`o (2000). The results obtained by MEGNO confirm the FMA results.   \\begin{table*} \\begin{center} \\caption{\\label{table}Orbital parameters for the self-consistent dynamical fit of the 55 Cancri planetary systems. These data are taken from Table 4 of Fischer et al. (2008).} \\begin{tabular}{ccccccc} \\hline Planet& $\\begin{array}{c} \\textrm{Period}\\\\ \\textrm{(days)} \\end{array}$& $\\begin{array}{c} T_p\\\\ \\textrm{(JD-2440000)} \\end{array}$& $\\begin{array}{c} e\\\\ \\textrm{} \\end{array}$& $\\begin{array}{c} \\omega \\\\ \\textrm{(deg)} \\end{array}$& $\\begin{array}{c} M\\,\\sin\\, i\\\\ (M_{Jup}) \\end{array}$& $\\begin{array}{c} a \\\\ \\textrm{(AU)} \\end{array}$ \\tabularnewline \\hline  b& $14.651262$& $7572.0307$& $0.0159$& $164.001$& $0.8358$& $0.115$ \\tabularnewline  c& $44.378710$& $7547.5250$& $0.0530$& $57.405$& $0.1691$& $0.241$ \\tabularnewline  d& $5371.8207$& $6862.3081$& $0.0633$& $162.658$& $3.9231$& $5.901$ \\tabularnewline  e& $2.796744$& $7578.2159$& $0.2637$& $156.500$& $0.0241$& $0.038$ \\tabularnewline  f& $260.6694$& $7488.0149$& $0.0002$& $205.566$& $0.1444$& $0.785$ \\tabularnewline \\hline  \\end{tabular} \\end{center} \\end{table*}    \\begin{figure} \\includegraphics[width=8. cm, height=5 cm]{histo.eps} \\caption[]{Histogram showing the dsitribution of the diffusion rate within our sample of 400 55Cnc planetary systems. } \\label{histo} \\end{figure}  \\newpage ",
        "conclusions": "\\label{sec:conclusions}  Among the possible dynamical configurations of the 55Cnc planetary system there is also stable chaos. About 15\\% of the systems built from the nominal dynamical fit given in \\cite{fima} have a high diffusion speed in the phase space. However, the chaotic island seems to be limited in extent and surrounded by quasi--periodic systems. In this way, the chaotic systems do not increase their eccentricity to planet crossing values and they are stable over a long timescale. Systems with intermediate values of diffusion speed $v_d$ as measured by FMA ($\\sim 4$) do not show this behaviour. They possibly populate the outer border of the chaotic region and their chaotic evolution is much slower. Systems with smaller values of  $v_d$ are possibly stable over timescale of $10^9$ yrs. Additional FMA and MEGNO computations show that this behaviour is not only peculiar for the nominal system given by Table 1 but it is extended over a wider range of orbital elements for the system. The weak resonances responsible for the stable chaos are present also for different combinations of the initial orbital elements."
    },
    "0807/0807.1529_arXiv.txt": {
        "abstract": "We present the discovery of optical/X-ray flux correlations on rapid timescales in the low/hard state of the Galactic black hole GX 339--4. The source had recently emerged from outburst and was associated with a relatively-faint counterpart with mag $V$$\\approx$17. The optical (\\vlt/ULTRACAM) and X-ray (\\rxte/PCA) data show a clear positive cross-correlation function (CCF) signal, with the optical peak lagging X-rays by $\\sim$ 150 ms, preceded by a shallow rise and followed by a steep decline along with broad anti-correlation dips. Examination of the light curves shows that the main CCF features are reproduced in superpositions of flares and dips. The CCF peak is narrow and the X-ray auto-correlation function (ACF) is broader than the optical ACF, arguing against reprocessing as the origin for the rapid optical emission. X-ray flaring is associated with spectral hardening, but no corresponding changes are detected around optical peaks and dips. The variability may be explained in the context of synchrotron emission with interaction between a jet and a corona. The complex CCF structure in GX~339--4 has similarities to that of another remarkable X-ray binary XTE~J1118+480, in spite of showing a weaker maximum strength. Such simultaneous multi-wavelength, rapid timing studies provide key constraints for modeling the inner regions of accreting stellar sources. ",
        "introduction": "\\gx339 is an X-ray binary system hosting one of the most promising Galactic black hole (BH) candidates, with a mass function of $\\sim 6$~\\Msun\\ \\citep{hynes03_gx339}. It is classified as a micro-quasar, based on tight radio/X-ray correlations found over a large dynamic range in source luminosity \\citep[e.g. ][]{gallo03}. Extensive X-ray studies have revealed a rich timing structure \\citep[e.g. ][]{dunn08, belloni05, nowak99, miyamoto91}. In its low state, the source is typically associated with an optically-bright counterpart, during which flickering on timescales as short as 10~ms has been observed. Rare transitions to the X-ray--off state have enabled placing some constraints on the much fainter optical companion star \\citep{shahbaz01}. And during X-ray outbursts, the optical and X-ray fluxes display complex correlations as well as anti-correlations over extended timescales \\citep[e.g. ][]{makishima86,russell06}.  Rapid timing observations simultaneous over a broad energy range remain as one of the lesser-explored aspects of \\gx339. The last published multi-wavelength studies to probe sub-second timescales were carried out over 20 years ago (but see Spruit et al. in prep.). Besides high state observations by \\citeauthor{makishima86}, \\citet[][]{motch82, motch83, motch85} studied the source during a low- to high-state transition, and found an anti-correlated cross-correlation function (CCF) signal with optical leading X-rays by a few seconds. Although there was an indication in the low-frequency light curves that the optical and 13--20 keV X-ray fluxes were correlated at a significance of about 98 per cent, no obvious high frequency ($\\ltsim 20$ s) correlation was uncovered at any of the selected energies.  The Galactic-halo X-ray binary \\xtej1118\\ has recently been a target of extensive, simultaneous multi-wavelength coverage. Among its interesting properties, a complex optical/X-ray CCF has been found \\citep[e.g. ][]{kanbach01, spruitkanbach02}, including positive and negative correlation components and an optical-vs.-X-ray peak lag of $\\sim 0.5$ s. Complex CCFs are now being found in other sources as well (\\citealt{durant08}). These provide important time domain constraints for physical models.  In this Letter, we present the first simultaneous optical/X-ray timing analysis of \\gx339\\ on rapid timescales of $\\sim 50-130$ ms in a low/hard state associated with a relatively-faint optical counterpart. A clear cross-correlation signal is detected with the optical peak lagging X-rays by $\\sim 150$ ms. The CCF has a complex pattern with some similarities as well as differences to the CCF found for \\xtej1118, and is likely to be result of distinct interacting accretion/ejection components. Full details of the timing and spectral analysis are presented in a forthcoming paper. ",
        "conclusions": "The main question that we wish to address is the origin of the rapidly variable optical power and the complex cross-correlation. Published models for the \\xtej1118\\ CCF include, among others: \\citet{merloni00}, who invoke a magnetically-dominated corona, \\citeauthor{esin01} (\\citeyear{esin01}, a dominant advection-dominated flow [ADAF] with additional synchrotron), \\citet[][ a pure jet]{markoff01}, \\citet[][ feedback in a common jet+corona reservoir]{malzac04} and \\citet[][ an ADAF and a jet dominating at different energies]{yuan05}. Most can explain the averaged broad-band energetics, and provide a qualitative description of the expected variability. But all agree that the origin and the details of the rapid variability patterns are likely to be complicated. The case of \\gx339\\ may well be similar, and a full investigation of the parameter space of the various models is beyond the scope of this Letter. Nevertheless, several important conclusions can be drawn from our simultaneous multi-wavelength data.  \\subsection{Reprocessing?}  The peak of the CCF time delay (150~ms) corresponds to a light-travel distance of $5000$ \\gravrad\\ [$\\equiv GM/c^2$] for $M$$\\gtsim$6~\\Msun, too small for reprocessing on the companion star (\\gx339\\ has a binary separation of $\\approx 25$ light-seconds). An auto-correlation analysis of the individual light curves can also be used to constrain the emission processes. Fig.~\\ref{fig:autocorr} shows the computed X-ray and optical auto-correlation functions (ACFs). Low count-rate Poisson noise dominating the X-ray ACF at zero lag has been corrected-for by subtracting white noise from an X-ray power spectrum, followed by an inverse Fourier transform. The final X-ray ACF is broader than the optical one, similar to the result found in the case of \\xtej1118. This also argues against a reprocessing origin (on the outer parts of the accretion disk, say) for the rapidly-variable component, at least as described by a simple, linear transfer function \\citep[cf.][]{kanbach01}.  \\begin{figure} \\begin{center} \\includegraphics[angle=90,width=8.5cm]{fig3.ps} \\caption{The X-ray (black) and optical (red dashed) auto correlation functions (ACF) computed from the highest time resolution (50 ms) light curves on Night 1. The X-ray ACF is for the full-band PCA, and the optical refers to the $r'$ filter. Both are corrected for Poisson noise (dominant in X-rays). \\label{fig:autocorr}} \\end{center} \\end{figure}   \\subsection{Behaviour around flares and dips}  Fig.~\\ref{fig:likemalzac} shows that the behaviour of low- and high-energy X-ray (in this case, 2--5 keV and 5--20 keV) photon intensities is similar, when selected around full band PCA peaks and dips, with no obvious lag. We note that CCFs of light curves extracted in these energy ranges, with respect to the optical, showed little difference to the full-band result of Fig.~\\ref{fig:crosscorr}. The optical light curves follow the CCF shape, as already mentioned. No significant colour ($g'-r'$) changes around the positions of X-ray flares/dips were detected.  What about the source spectral behaviour? Using standard HEADAS tools, we extracted average X-ray spectra (and background) within short time-bins, $\\sim \\pm 50-150$ ms, centred on X-ray PCA full-band flares and dips (HEXTE was not used for this analysis, as the source is always background-dominated above $\\sim 20$ keV). What we find is that the source hardens when it flares, and vice-versa. This is illustrated in Fig.~\\ref{fig:flares_dips_spectra}, which shows the contours of independent fits to the extracted flares and dips spectra. A simple power-law (with absorption fixed to Galactic) was used to parametrize the spectral change between the two cases. The photon-index fitted to the flares spectrum is harder than that fitted to dips at 99 per cent confidence on all nights. Simulations were used to confirm that the results are robust to fluctuations in the background, which dominates above $\\sim 5$ keV in the lower-flux dips spectrum.  On the other hand, no significant changes were detected in X-ray spectra extracted at the positions of {\\em optical} peaks as compared to optical dips -- the spectral slopes were close to the slope for the average spectrum of the full dataset ($\\Gamma\\approx 1.6$). Extracting spectra 150 ms {\\em before} optical flares and dips (as suggested by the CCF delay) showed a small but clear difference in the intensities of the X-ray flare and dip spectra, but no obvious change in spectral slope within the errors. This seems consistent with the low CCF peak strength -- i.e. every optical flare need not have been preceded by a locally-maximum X-ray flare, in spite of the average correlation.    \\begin{figure} \\begin{center} \\includegraphics[angle=90,width=8.5cm]{fig4.ps} \\caption{ Photon-index ($\\Gamma$) contours for single power-law fits to the extracted 3--30 keV flare and dip background-subtracted spectra, with absorption fixed to Galactic. The total resultant time interval for spectral extraction is $\\approx 30$ s long for both flares and dips, and the net count rate is $\\approx 40$ and 5 ct s$^{-1}$ respectively. Contour levels correspond to 68.3, 95 and 99\\% for two interesting parameters, and include 1 per cent systematic errors. Results are shown for all three nights. The dotted diagonal line is the 1:1 line. \\label{fig:flares_dips_spectra}} \\end{center} \\end{figure}   \\subsection{Implications}  The above analysis shows that the source is harder during X-ray flares, and vice-versa. In the context of a hot accretion disk corona model, this is consistent with increased Compton up-scattering when the source is (momentarily) brighter, and vice-versa. Fast Compton cooling timescales of tens of ms or smaller have been inferred in Galactic black holes in the low/hard state \\citep[cf. ][]{guilbert82_nature}, for coronal electron temperatures $kT_e \\sim 100$ keV and seed-photon disk-blackbody X-ray luminosities $< 0.01 L/L_{\\rm Edd}$. Very similar physical parameters are inferred for \\gx339\\ during previous low/hard states \\citep{joinet07, miller06}, and also contemporaneous with our observations (cf. \\citealt{tomsick08}, who favour an inner disk radius [\\rin] of $\\sim 10$~\\gravrad). Thus, similarly-fast cooling timescales are likely to hold for the case of \\gx339, resulting in apparently-simultaneous flaring and hardening, which is what we observe.  As for the optical emission, we have already excluded simple reprocessing models. Bremsstrahlung can also be ruled out, as this would require a corresponding X-ray flux higher than that observed by several orders of magnitude. The most likely remaining physical mechanism is then synchrotron emission. Cyclo-synchrotron models in which several magnetized active regions with sizes of a few Schwarzschild radii contribute significantly to the dereddened optical flux have been investigated by \\citet[ in addition cf. \\citealt{fabian82}]{dimatteo99}. The active regions are characterized by $B\\sim$ few $\\times 10^6$ G, coronal optical depths $\\tau\\sim 0.2-1$ and $kT_e\\sim 150-200$~keV. \\citet{wardzinski00} have also discussed important modifications to such models.  Within the context of these models, what could the shape and the time delay of the CCF correspond do? The steep decline that follows the CCF peak (Fig.~\\ref{fig:crosscorr}) suggests the presence of some mechanism that cuts off the optical flares suddenly -- like infall of synchrotron-emitting blobs into the BH, say. A potential difficulty of this scenario is that a free-fall time of 150 ms for a 6--10 \\Msun\\ BH corresponds to a physical radius of $\\sim 200-250$ \\gravrad, which is larger than the value of \\rin\\ inferred by \\citet{tomsick08}. Unless the disk has receded further or an atypically large corona is present, this does not correspond to an obvious physically meaningful scale.  An alternate hypothesis is that the rapid variability originates in non-thermal emission within a relativistic outflow or a jet. \\citet{malzac04} invoked a magnetic energy reservoir that feeds both a jet component (dominating the optical), as well as an electron corona (dominating in X-rays) to explain the CCF of \\xtej1118. Energy injection was modeled as shot flares. Feedback between the two components resulted in complex correlations, with the optical power being proportional to the differential of the X-ray flux. In such a model, the CCF shape and time lag are determined by the dissipation timescale of the process that injects energy into the jet. This should be largely independent of the exact injection mechanism, though the authors discuss the context of a magnetic energy reservoir. For our observations, a simple exponential ($\\propto e^{-t/\\tau}$) fit to the innermost part of the X-ray ACF (lags $<$0.5 s) in Fig.~\\ref{fig:autocorr} gives a dissipation timescale $\\tau=0.2\\pm 0.05$ s (90\\% error). This agrees with the observed optical delay, and suggests that such a differential correlation may apply to \\gx339 as well.  If X-ray flares above the accretion disk trigger the dissipation and large-scale re-ordering of the poloidal magnetic field threading a jet, it is possible that the synchrotron (optical) emission will respond on timescales related to subsequent field build-up. Significant modulation of the poloidal field can occur on timescales orders of magnitude longer than the dynamical time of the inner accretion disk regions where this field originates \\citep[][ see their Eq. 4]{livio03}. For $r_{\\rm in} \\sim 10$~\\gravrad, the Keplerian dynamical time is 6 ms. Our observed delay of the optical (i.e., jet) component ($\\sim 150$ ms) is 25 times longer, which can easily be accommodated within the picture of magnetic modulation. Synchrotron emission by plasma accelerated along the jet during this period, followed by rapid radiative cooling, may thus explain the positive CCF peak and delay.  As for the anti-correlation troughs: if the initial X-ray flares that triggered magnetic re-ordering are related to field reconnection events in a coexistent corona, the coronal magnetic energy density will be released on the X-ray flaring timescale. This will lead to a decrease of any ambient synchrotron emission that is occurring within the corona itself, resulting in an anti-correlation of optical with X-rays. X-ray flaring is coherent over timescales of $\\sim$ several seconds, as seen in the X-ray ACF (Fig.~\\ref{fig:autocorr}), and this is also the total length of the anti-correlation seen in the CCF plot of Fig.~\\ref{fig:crosscorr}. In short, interaction between the jet and coronal components may be the key to understanding the complex correlation structure.  A testable prediction of this model is that the appearance and strength of the positive CCF signal should be intimately related to the prominence of the jet. Our observations probed the beginning of the low/hard state. As the sources enters deeper into this state and the jet establishes itself, its contribution to the overall energetics should grow (e.g. \\citealt{fender04}), as should the related optical fractional variability rms. Support of this comes from the fact that we find an optical rms variability of $\\sim 15$ per cent over the full light curve timescales, while \\citet{motch83} found an rms of 50 per cent when the source was optically brighter (this is well-matched to the estimate of $\\sim 50$ per cent by \\citealt{corbel02}, based on broad-band photometry). Increased corona/jet coupling will result in a higher peak strength of the optical/X-ray CCF. Finally, if a stronger poloidal field is required to establish a stronger jet component, this may result in longer timescales for breaking and establishing this field component following reconnection flares. Changing CCF delays could then be used to directly probe the evolution of characteristic accretion/ejection structures in X-ray binaries."
    },
    "0807/0807.2000_arXiv.txt": {
        "abstract": "The extended holographic dark energy model with the Hubble horizon as the infrared cutoff avoids the problem of the circular reasoning of the holographic dark energy model. We show that the infrared cutoff of the extended holographic dark energy model cannot be the Hubble horizon provided that the Brans-Dicke parameter $\\omega$ satisfies the experimental constraint $\\omega> 10^4$, and this is proved as a no-go theorem. The no-go theorem also applies to the case in which the dark matter interacts with the dark energy. ",
        "introduction": "The current accelerating expansion of the universe was first discovered in 1998 by the observation of the type Ia supernova \\cite{acc1}. The high redshift supernova Ia observation found strong evidence of a transition from deceleration in the past to acceleration at present \\cite{riess,essence}. Evidence of the accelerating expansion of the universe was further provided by other complementary astronomical observations, such as the cosmic microwave background anisotropy and the large scale structure of the clusters of galaxies \\cite{astier,sdss6,wmap3}. As a model independent tool, the energy conditions were correctly employed to analyze the observational data and conclude the existence of the accelerated expansion of the universe in \\cite{gong07b,gong07c}. To explain the cosmic acceleration, an exotic energy component with negative pressure, dubbed dark energy, is introduced. Because the only observable effect of dark energy is through gravitational interaction, the nature of dark energy imposes a big challenge to theoretical physics. One simple dark energy candidate which is consistent with current observations is the cosmological constant. Due to the discrepancy of many orders of magnitude between the theoretical predication and observation for the vacuum energy, lots of dynamical dark energy models were proposed. For a review of dark energy models, see \\cite{rev}.  The holographic dark energy (HDE) model is one of the interesting dynamical dark energy models. The HDE model is derived from the relationship between the ultraviolet (UV) and the infrared (IR) cutoffs proposed by  Cohen {\\it et al} in \\cite{cohen}. The UV-IR relationship was also obtained by Padmanabhan, arguing that the cosmological constant is the vacuum fluctuation of energy density \\cite{pad}. Due to the limit set by the formation of a black hole, the UV-IR relationship gives an upper bound on the zero-point energy density $\\rho_h\\le 3 L^{-2}/(8\\pi G)$, which means that the maximum entropy is of the order of $S_{BH}^{3/4}$. Here $L$ is the scale of IR cutoff. The zero-point energy density has the same order of magnitude as the matter energy density \\cite{Hsu}, and is named the HDE density by Li \\cite{Li}. However, the original HDE model with the Hubble scale as the IR cutoff failed to explain the accelerating expansion of the universe \\cite{Hsu}. Li solved the problem by discussing the possibilities of the particle and event horizons as the IR cutoff, and he found that only the event horizon identified as the IR cutoff leads to a viable dark energy model \\cite{Li}. The HDE model using the event horizon as the IR cutoff was soon found to be consistent with the observational data in \\cite{gong}. By considering the interaction between dark energy and matter in the HDE model with the event horizon as the IR cutoff, it was shown that the interacting HDE model realized the phantom crossing behavior \\cite{intde}. Other discussions on the HDE model can be found in \\cite{holo,gong08,hololi,holo1,pavon05,grg,horvat,holo2, sadjadi,zhang07}  Since string theory is believed to be the theory of quantum gravity, Einstein's theory of gravity needs to be modified according to string theory. In the low energy effective bosonic string, the dilaton field appears naturally. The scalar degree of freedom arises also upon compactification of higher dimensional theory. The simplest alternative which includes the scalar field in addition to the tensor field to general relativity is Brans-Dicke theory. Therefore, it is interesting to discuss the HDE model in the framework of Brans-Dicke theory. That was first done by Gong in \\cite{gong04}, and the model is called the extended holographic dark energy (EHDE) model. The EHDE model was also discussed in \\cite{kim,setare,pavon,nayak,xu}. Recently, it was claimed that the EHDE model with the Hubble horizon as the IR cutoff could solve the dark energy problem \\cite{pavon,nayak}. The existence of the event horizon means that the universe must experience accelerated expansion, so the HDE and EHDE models with the event horizon as the IR cutoff face the problem of circular reasoning. If the Hubble horizon can be used as the IR cutoff in EHDE model, then the EHDE model is more successful and interesting. In this paper, we carefully examine the EHDE model with the Hubble horizon as the IR cutoff and show that the model fails to solve the dark energy problem. We discuss the EHDE model with the Hubble horizon as the IR cutoff in section II and the interacting EHDE model in section III. ",
        "conclusions": ""
    },
    "0807/0807.2831_arXiv.txt": {
        "abstract": "We report the discovery of a new Milky Way satellite in the constellation Leo, identified in data from the Sloan Digital Sky Survey. It lies at a distance of $\\sim 180$ kpc, and is separated by $\\lesssim 3^\\circ$ from another recent discovery, Leo~IV. We present follow-up imaging from the Isaac Newton Telescope and spectroscopy from the Hectochelle fiber spectrograph at the Multiple Mirror Telescope. Leo~V's heliocentric velocity is $\\sim 173.3 \\pm 3.1$ kms$^{-1}$, offset by $\\sim 40$ kms$^{-1}$ from that of Leo~IV. A simple interpretation of the kinematic data is that both objects may lie on the same stream, though the implied orbit is only modestly eccentric ($e \\sim 0.2$) ",
        "introduction": "In the last few years, there have been numerous discoveries of ultra-faint Milky Way satellites, primarily because the Sloan Digital Sky Survey (SDSS) allows the detection of galaxies with central surface brightnesses as faint as $30$ mag arcsec$^{-2}$. The new discoveries include 10 new Milky Way dwarf galaxies, together with 4 unusually extended or faint globular clusters~\\citep{Wi05,Zu06a,Zu06b,Be06a,Be07,Ir07,Wal07,Ko07}.  The purpose of this {\\it Letter} is to announce the discovery of an additional Milky Way satellite, probably a dwarf galaxy that may be undergoing disruption, at a heliocentric distance of $\\sim 180$ kpc in the constellation of Leo. Following the convention for naming dwarf spheroidals, we call it Leo~V. It lies very close to one of our other recent discoveries, namely Leo~IV~\\citep{Be07}. Hence, it is a companion to a companion of the Milky Way Galaxy.  \\begin{figure*}[t] \\begin{center} \\includegraphics[width=0.9\\textwidth]{leov_figure1.ps} \\caption{The Leo~V Satellite: {\\it Upper Left:} SDSS cut-out ($9^\\prime \\times 9^\\prime$) around the center of Leo~V. {\\it Upper Middle Left:} The spatial distribution of all objects classified as stars in a $30^\\prime \\times 30^\\prime$ field. {\\it Upper Middle Right:} The spatial distribution of all objects classified as galaxies. {\\it Upper Right:} Density of candidate RGB stars selected with the CMD mask shown in the panel directly below. {\\it Lower Left:} CMD of all stars in a circle of radius $2^\\prime$, which is expected to be dominated by Leo~V members.  {\\it Lower Middle Left:} CMD of all stars within $5^\\prime$, with the red clump and BHB marked. {\\it Lower Middle Right:} CMD of stars within the annulus $9^\\prime$ to $10.3^\\prime$ showing the foreground. {\\it Lower Right:} Difference in Hess Diagrams, showing the red giant branch and BHB of Leo~V. The mask is built with M92 ridgelines offset to the distance modulus of 21.25.} \\label{fig:leo_disc} \\end{center} \\end{figure*} \\begin{deluxetable}{lc} \\tablecaption{Properties of the Leo~V Satellite \\label{tbl:pars}} \\tablewidth{0pt} \\tablehead{ \\colhead{Parameter\\tablenotemark{a}} & {~~~ } } \\startdata Coordinates (J2000) & 11 31 09.6 +02 13 12.0 \\\\ Coordinates (Galactic) & $\\ell = 261.86^\\circ, b = 58.54^\\circ$ \\\\ $r_h$ (Plummer) & $0\\farcm8 \\pm 0\\farcm1$\\\\ $\\mu_{\\rm 0,V}$ (Plummer) & $27.5 \\pm 0\\fm5$ \\tablenotemark{a}\\\\ (m$-$M)$_0$ & $21\\fm25$\\\\ M$_{\\rm tot,V}$ & $-4\\fm3 \\pm 0\\fm5$ \\\\ $v_\\odot$ & $+173.3 \\pm 3.1$ kms${}^{-1}$ \\\\ $v_{\\rm GSR}$ & +60.8 kms${}^{-1}$ \\enddata \\tablenotetext{a}{Surface brightnesses and integrated magnitudes are corrected for the mean Galactic foreground reddening, A$_{\\rm V}= 0\\fm1$.} \\label{tab:struct} \\end{deluxetable} ",
        "conclusions": "Leo~V is most probably a new dwarf galaxy, based on the presence of an old, metal-poor population with a characteristic size of between 50 and 200 pc. There are several hints that it may be a disrupting satellite, but our data do not support the idea that it is merely an overdensity in a stellar stream.  The most remarkable feature of Leo~V is the disparity in the spatial extent of the RGB and BHB populations. The red giants are confined to a tight core of $\\sim 50$ pc, whereas the BHBs extend out at least as far as 200 pc. There even appear to be BHBs associated with Leo~V at distances of 500 pc. One possible explanation for the apparent difference is that the RGB stars do follow the BHBs much further out, perhaps because the object is disintegrating, but they are more difficult to distinguish from the foreground populations in our data.  A hint that this the case is perhaps provided by the two RGB candidates that are even further than the most distant BHBs, although their membership needs to be confirmed.  With deeper photometry reaching down to the turn-off, it should be possible to verify this hypothesis.  Another explanation is that the RGB and BHB populations probe different epochs of star formation in an ultra-low mass system.  Leo~V's possible association with Leo~IV is also unique. Although there are other examples of dSphs separated on the sky by a few degrees -- such as CVn~I and CVn~II -- they are at different distances and velocities. By contrast, Leo~IV is at a heliocentric distance of $\\sim 160$ kpc~\\citep{Be07} and a heliocentric velocity of $132$ kms$^{-1}$~\\citep{Si07}. These are very close to our estimates of $\\sim 180$ kpc and $173$ kms$^{-1}$ for the distance and velocity of Leo~V. Referred to the Galactic Standard of Rest, the velocities of Leo~IV and V are low, $11.0$ kms$^{-1}$ and $59.5$ kms$^{-1}$ respectively.  Using the velocity distribution for the $\\rho \\sim r^{-3.5}$ radial profiles given in \\citet{Ev97} to construct artificial samples of 50 satellites, we estimate that there is a $\\lesssim 1$ per cent probability of this coincidence happening by chance. We remark that one of BHB stars considered by \\citet{Si07} for possible membership of Leo~IV, but then discarded, has a heliocentric velocity of $160$ kms$^{-1}$. This hints at the possible existence of extended stellar structures around Leo~IV and Leo~V. If Leo~IV and Leo~V are assumed to be on the same stream, then the orbit can be computed, assuming a singular isothermal sphere with amplitude $v_0$ = 220 kms$^{-1}$. The pericenter is $\\sim 160$ kpc and the apocenter is $\\sim 244$, so that the eccentricity is modest ($e = 0.2$) and the orbit never approaches the inner parts of the Milky Way (in which case both objects should be relatively intact, at odds with their seemingly irregular appearances).  Leo~V may prove to be an important object for testing theories of galaxy formation. \\citet{Ri08} has argued that very old dwarf galaxies must form preferentially in chain structures, tracing the filamentary dark matter in the early universe. These chains or groups of dwarfs may retain some of their integrity even on accretion and merging into the Milky Way halo. Is it possible that Leo~IV and Leo~V are two links in such a chain?  \\vskip 0.2truecm"
    },
    "0807/0807.0469_arXiv.txt": {
        "abstract": "SuperLupus is a deep transit survey monitoring a Galactic Plane field in the Southern hemisphere.  The project is building on the successful Lupus Survey, and will double the number of images of the field from 1700 to 3400, making it one of the longest duration deep transit surveys.  The immediate motivation for this expansion is to search for longer period transiting planets (5-8 days) and smaller radii planets.  It will also provide near complete recovery for the shorter period planets (1-3 days).  In March, April, and May 2008 we obtained the new images and work is currently in progress reducing these new data. ",
        "introduction": "The Lupus Survey was a deep transit survey of a 0.66 square degree patch of sky near the Galactic Plane (b=11$^{\\circ}$).  The survey was conducted using the ANU 40 Inch Telescope at Siding Spring Observatory, Australia in May and June of 2005 and 2006.  In total, 1783 good quality images of the field were obtained.  The images were 5-minute exposures taken in a wide V+R filter.  Time series photometry was performed for 110,372 stars in the field, with 16,134 of those stars having a precision of $\\sigma<0.025$.  The transiting planet candidates and results from this survey will be published in \\cite[Bayliss et al. (2008)]{Bayliss08}. The discovery of a Hot Jupiter in the field, Lupus-TR-3b, has been published in a separate letter (\\cite[Weldrake et al. (2008)]{Weldrake08b}), and the radial velocity follow-up confirming the discovery is set out in~Figure~\\ref{lupus3}. \\begin{figure}[!ht] \\vspace*{-2.5 cm} \\begin{center} \\includegraphics[width=8.5cm]{lupus3.eps} \\vspace*{-0.5 cm} \\caption{Radial velocity measurements for Lupus-TR-3 from MIKE on Magellan II (Clay) (\\cite[Weldrake et al. (2008)]{Weldrake08b}).  Small solid circles are the individual radial velocity measurements, with the error bars determined from the actual scatter in the orders. Large open circles are uncertainty-weighted nightly averages.  The solid line is the best-fit sinusoid with only the period fixed (from the photometry).  The fitted phase matched the photometrically-determined phase to within 0.13~days, well within fit uncertainties.} \\label{lupus3} \\end{center} \\end{figure} Additionally, 494 new variables were discovered in the survey field.  These variables have been cataloged in \\cite[Weldrake \\& Bayliss (2008)]{Weldrake08}. ",
        "conclusions": ""
    },
    "0807/0807.3760_arXiv.txt": {
        "abstract": "High-energy cosmic rays with energies exceeding $10^{17}$~eV are frequently observed by measurements of the fluorescence light induced by air showers.  A major contribution to the systematic uncertainties of the absolute energy scale of such experiments is the insufficient knowledge of the fluorescence light yield of electrons in air.  The aim of the 5th Fluorescence Workshop was to bring together experimental and theoretical expertise to discuss the latest progress on the investigations of the fluorescence light yield.  The results of the workshop will be reviewed as well as the present status of knowledge in this field.  Emphasis is given to the fluorescence light yield important for air shower observations and its dependence on atmospheric parameters, like pressure, temperature, and humidity. The effects of the latest results on the light observed from air showers will be discussed. ",
        "introduction": "\\Label{intro}  The Earth is permanently exposed to a vast flux of particles from outer space. Most of these particles are fully ionized atomic nuclei, covering a large range in energy from the MeV regime to energies above $10^{20}$~eV.  In 1962 the first event with an energy exceeding $10^{20}$~eV was recorded \\Cite{linsley}.  Such cosmic rays are the highest-energy particles in the Universe, carrying the (macroscopic) energy of about 50~J concentrated on a single nucleus. Since their first discovery more than 40 years ago their origin has been an open question. ``How do cosmic accelerators work and what are they accelerating?'' is one of eleven science questions for the new century asked by the National Research Council of the National Academies of the United States \\Cite{quarkscosmos}, underlining the importance of this topic to astroparticle physics.  The properties of cosmic rays at highest energies are investigated with various experiments \\Cite{naganowatson}\\Cite{bergmanbelz}\\Cite{kamperttaup07}\\Cite{jrhwuerzburg}. The flux of cosmic rays with energies exceeding $10^{20}$~eV is below 1 particle per square kilometer and century. Thus, a reasonable measurement of these particles requires huge detectors operated stably over long periods of time.  At present, this can be realized only with ground based experiments, registering the secondary particles generated by high-energy cosmic rays in the atmosphere.  When a high-energy cosmic ray enters the Earth's atmosphere it induces a cascade of secondary particles, an extensive air shower.  By far the most abundant particles in air showers are photons, electrons, and positrons, comprising the electromagnetic shower component.  On their way through the atmosphere the (relativistic) charged particles emit \\Cerenkov radiation and excite nitrogen molecules to emit fluorescence light. A small fraction of the secondary particles eventually reaches the observation level.  Over time, several methods to measure extensive air showers have been established. They can be divided into two groups: experiments measuring secondary particles (electrons, muons, and hadrons) reaching ground level and detectors observing the emitted \\Cerenkov or fluorescence photons.  The latter allow for a three-dimensional reconstruction of the shower profile in the atmosphere.  A critical issue for all experiments is to establish an absolute energy scale for the measured showers. For experiments registering secondary particles at ground level this usually involves the usage of simulations of the shower development in the atmosphere, thus, introducing systematic uncertainties due to our limited knowledge of the hadronic interactions at such high energies.  On the other hand, fluorescence measurements of air showers provide a calorimetric measurement of the energy deposited in the air, being (nearly) independent of air shower simulations.  The deposited energy is assumed to be proportional to the energy of the primary, shower inducing, particle. At present, this is the most direct and model independent method to determine the energy of an air shower. The main systematic uncertainty of this method arises from the insufficient knowledge of the fluorescence light yield of electrons in air. However, this is a quantity which can be measured in laboratory experiments, injecting electrons into air targets.  Responsible for the fluorescence emission of the nitrogen molecules is mainly the electromagnetic shower component. The critical energy of electrons in air is about 84~MeV, thus, the bulk of particles has energies easily accessible at accelerators, or, at lower energies, even through radioactive $\\beta$-decays.  This is the main focus of the present article. It gives an overview on the actual status of the know\\-ledge of the fluorescence light yield of electrons in air, important for air shower detection.  It summarizes the results of the 5th Fluorescence Workshop, which was held in El Escorial, Spain from September 16th to 20th, 2007.  After a short overview on the principle of the detection of air showers using fluorescence light (\\sref{sec:principle}), experiments applying this technique to register cosmic-ray induced air showers are described (\\sref{sec:eashistory}).  The main physical processes involved in the production of fluorescence light in air are reviewed in \\sref{sec:fl_physics}.  Contemporary experimental tools and theoretical treatments are discussed in \\sref{sec:status}.  Latest data are compiled in \\sref{sec:data}, paying special attention to the dependence of the fluorescence light yield on atmospheric parameters, such as pressure, temperature, and humidity.  In the following section (\\sref{sec:altitude}) the influence of the results obtained on the light yield in air showers developing in realistic atmospheres is discussed. An outlook describing the next steps in determining the fluorescence yield in air concludes the article (\\sref{sec:outlook}). The most important pioneering measurements of the fluorescence light yield are summarized in an accompanying article \\Cite{histref}. ",
        "conclusions": "\\Label{sec:outlook}  The 5th Fluorescence Workshop has brought together experimental and theoretical expertise from around the world to discuss the status of the determination of the fluorescence light yield of electrons, important for air shower observations. It can be realized with pleasure that over the last years significant progress has been achieved in both, experimental and theoretical work.  Convergence has been achieved in several aspects between the different groups.  Most important for air shower observations is the fluorescence light yield of nitrogen, since the contributions of direct excitation of oxygen and argon seem to be negligible. Hence, investigations focus on the fluorescence light emission of nitrogen. The position of the molecular bands in the fluorescence light spectrum is well known. Accurate measurements of the relative intensities at typical laboratory conditions have been presented. Also high precision values of the characteristic pressures have been reported, although full agreement is still missing. High-precision $P'$ values allow accurate predictions on the relative fluorescence intensity at any atmospheric condition of interest for air shower reconstruction. The proportionality between fluorescence intensity and deposited energy seems to be proved from both experimental and theoretical sides. Recent investigations reveal that the dependence of the light yield on temperature and humidity can not be neglected. In realistic atmospheres effects up to $10\\%$ can be expected. Further measurements are necessary to clarify the situation.  A precise determination of the absolute light yield is extremely difficult and still needs more experimental data.  For a comparison of the values obtained by different groups, it would be useful if the same concept is used. The participants agreed to specify the fluorescence light yield in photons per energy deposition (photons/MeV). This quantity  is most useful to convert the observed light yield into a calorimetric measurement of the energy of an air shower. The usage of this unit is  strongly encouraged.  The open questions are presently addressed by several groups. The workshop participants expressed their hope to achieve convergence on the open issues soon ($1-2$ years)."
    },
    "0807/0807.4695_arXiv.txt": {
        "abstract": "We use the SWIFT Burst Alert Telescope (BAT) sample of hard x-ray selected active galactic nuclei (AGN) with a median redshift of 0.03 and the 2MASS J and K band photometry to examine the correlation of hard x-ray emission to Eddington ratio as well as the relationship of the J and K band nuclear luminosity to the hard x-ray luminosity. The BAT sample is almost unbiased by the effects of obscuration and thus offers the first large unbiased sample for the examination of correlations between different wavelength bands. We find that the near-IR nuclear J and K band luminosity is related to the BAT (14 -- 195\\,keV) luminosity over a factor of $10^3$ in luminosity ($L_{IR} \\approx L_{BAT}^{1.25}$)and thus is unlikely to be due to dust.  We also find that the Eddington ratio is proportional to the x-ray luminosity. This new result should be a strong constraint on models of the formation of the broad band continuum. ",
        "introduction": "The estimate of black hole masses and the ratio of luminosity to the luminosity derived from the black hole mass (the Eddington ratio) are the fundamental parameters in the evolution of active galaxies.   The Eddington ratio is defined as L/L$_{Edd}$, where \\begin{math} L_{Edd}=(M_{BH}/M_{\\odot}) \\times1.3 \\times 10^{38} erg\\,s^{-1}\\end{math} or the limit where gravitational and radiation pressure are balanced.  It is important because it is a measurement of the ratio of energy output of a source to the physical limit of energy output.  Understanding the energy output of active galactic nuclei (AGN) is essential to understanding both the nature of AGN sources and the effect they have on their environment. Therefore, measuring both M$_{BH}$ and the Eddington ratio are the subject of much recent work \\citep{kol06, net07, wu02}, based primarily on optical samples and the combination of x-ray and IR samples. These papers have either used data from the literature \\citep{wu02}, the SDSS \\citep{net07}, or deeper optical, x-ray and IR data \\citep{kol06} to obtain conclusions about the relationship of the luminosity to black hole mass and its evolution across cosmic time. For example, \\citet{kol06}  conclude that  there is strong evidence that supermassive BHs gain most of their mass while radiating close to the Eddington limit,  and that the fueling rates in luminous AGNs are determined by BH self-regulation of the accretion flow rather than galactic-scale dynamical disturbances.  Whereas previous samples are strongly biased against the most numerous AGN in the universe, those with high line of sight column densities, e.g. `obscured' AGN \\citep{mat06}, our BAT sample, since it is selected solely based on 14--195 keV flux, is not. However the BAT sample, as opposed to most optical samples, has a very low median redshift (0.03) and very few objects at $z > 0.1$.  Since there is a strong relationship between  line of sight obscuration and intrinsic luminosity \\citep{toz06, mar05}, such that the less  luminous sources are much more likely to be obscured, the optical samples, which miss most of the obscured sources, are strongly deficient in low luminosity sources.  Furthermore, there is often a complex relationship between measures of obscuration in the x-ray,  IR, and optical bands and there may be strong evolutionary effects such that the fraction of obscured objects is higher at higher redshifts \\citep{laf05}.  If there is a physical relationship between luminosity or stage in the life of the AGN  and the occurrence of obscuration, as predicted by several models \\citep{hop06}, then the results from such surveys may be rather different than The BAT sample which is heavily weighted to the local universe with a mean redshift of $\\approx 0.03$.  In addition, there has been no conclusive determination of  the origin of the `non-thermal' continuum in AGN. In the x-ray band, there are two mechanisms that have received extensive discussion: thermal Comptonization of photons produced by the accretion disk from a population of high energy particles, perhaps in a corona above the disk \\citep{zdz9x}, and emission due to jet processes \\citep{kor06}. However, there is at present no agreement as to the origin of `non-thermal' IR and optical emission in radio quiet objects or even whether it exists. Most models for the origin of the optical emission have focused on disk models (e.g. \\citet{sun89, cze87}) with some Comptonization.  Dusty torus models \\citep{pie93}, as the origin of the K band flux, have received support from reverberation observations \\citep{sug05} in the optical (UBV) and near-infrared (JHK) bands.  These studies found a clear time-delayed behavior of the K band with respect to the V band, strongly suggesting that the bulk of the K-band emission originates from thermal emission of dust grains in an optically thick torus \\citep{gla04}. There are strong correlations between the optical and x-ray properties \\citep{bor0x, bro06} and the IR properties \\citep{alo0x}, whose origin is not easy to explain if these continuum components originate from completely different physical mechanisms as is predicted by the most modern dusty torus models of the IR continuum \\citep{nen08}.  In this paper we present the first results on the Eddington ratio from an unbiased survey of AGN selected in the 14 -- 195\\,keV band with the Swift BAT survey. We find a strong correlation of the Eddington ratio to the hard x-ray luminosity over a wide range in L/L$_{Edd}$ ($10^{-4}$ to $3$) in our sample, as opposed to the narrow range in the sample of \\citet{kol06} selected from the AGES survey. In addition we find that the luminosity of the point-like component in the 2MASS J and K band data is almost linearly correlated with the high energy luminosity indicating that these two components are very closely coupled  and probably have the same physical origin. Since the x-ray emission is not due to dust processes, these results indicate most of the near IR is not either. ",
        "conclusions": "The median value of the Eddington ratio that we find is L/L$_{Edd} \\approx 0.02$ with a full range of $\\approx 1000$.  This value seems to be at odds with the work of recent AGN samples, such as the \\citet{kol06} sample selected based on optical/IR/x-ray properties, which finds $L_{Bol}/L_{Edd}$ = 0.25 over a rather different luminosity range $10^{45}$ -- $10^{47}$\\,erg\\,s$^{-1}$ in L(Bol) than the BAT sample. We believe that this is because the Kollmeier selection is really an optical one, since only the broad line sources have mass estimates.  Such a sample is heavily biased towards ``unobscured'' sources which have higher luminosities \\citep{ste03}.  Also, we now know that they also have higher $L/L_{Edd}$ \\citep{win08}.  Additionally, the Kollmeier selection is also biased towards much higher redshifts than the BAT survey, which could also explain the difference if, for instance, their is an evolution in Eddington ratios.  The wide range of inferred Eddington ratios from x-ray data has been noticed before \\citep{mer05} and is also seen in  other low z samples \\citep{hop05}. However, our sample differs significantly from the \\citet{ho05} sample in not having objects with very low Eddington ratios. We are not sure why we would have a bias towards higher Eddington ratios than \\citet{ho05}. However, since we have not convolved our selection function with a theoretical Eddington ratio model, one cannot be sure that our differences with Ho and Kollmeir are not due to selection biases, the luminosity range covered, and/or evolutionary effects.  Softer x-ray surveys in the 2--10\\,keV band (e.g. \\citet{ste06}) have found that  the bolometric correction is a function of luminosity, such that the x-ray band carries less of the luminosity at higher luminosities. If this were also true for our sample it would increase the amplitude of the correlation between hard x-ray and bolometric luminosities and change the slope \\citep{hop05}.  The correlation between Eddington ratio and x-ray luminosity has, to our knowledge, not been reported from  other samples (e.g. \\citet{mer05}) but it is clearly present in them. If one graphs the \\citet{mer05} data one finds exactly the same correlation as in our plot. There are theoretical reasons \\citep{hop05} to believe that the most luminous objects have the highest Eddington ratios and the lower luminosity objects have lower L$_{Edd}$. The Hopkins paper gives the Eddington ratio as a conditional probability distribution in which the width is a function of Eddington ratio -- this is hard to interpret in our present framework and we have not tested the predictions of this model.  We note that our distribution of masses is weighted towards higher masses than predicted by \\citet{hop05} at low z, but we cannot conclude that this is a strong result  since we have not  convolved  the  \\citet{hop05} prediction with our sensitivity limit.  It is well known that there is a break in the luminosity function of AGN, the origin of which is not clear. Below the break in the luminosity function, the median  Eddington ratio from our sample  is -2.1 and the variance is 0.35,  while at higher  L the median is -1.64 and the variance is 0.199. This is in agreement with the predictions of \\citet{hop05}, that the variance in the Eddington ratio should decrease at higher L, and provides a possible explanation of the break in the luminosity function as a change in the mean Eddington ratio of the population.  Our finding of a strong almost linear correlation of the near IR luminosities with the BAT luminosity (see \\citet{hor07} for a similar result) argues that similar processes produce the near IR continuum and the hard x-rays. Since the mechanism producing the x-ray emission is unrelated to dust, these results do not favor dust reprocessing models for the origin of the near IR continuum, unless the main driver of the dust reprocessing is simply and linearly connected to the x-ray continuum intensity."
    },
    "0807/0807.4189_arXiv.txt": {
        "abstract": "We have used widely spaced in time Hubble Space Telescope images to determine tangential velocities of features associated with outflows from young stars. These observations were supplemented by groundbased telescope spectroscopy and from the resultant radial velocities, space velocities were determined for many outflows. Numerous new moving features were found and grouped into known and newly assigned Herbig Haro objects. It was found that stellar outflow is highly discontinuous, as frequently is the case, with  long term gaps of a few hundred years and that these outflow periods are marked by staccato bursts over periods of about ten years. Although this has been observed in other regions, the Orion Nebula Cluster presents the richest display of this property. Most of the large scale Herbig Haro objects in the brightest part of the Orion Nebula appear to originate from a small region northeast of the strong Orion-S radio and infrared sources. With the possible exception of HH 203, we are not able to identify specific stellar sources, but do identify candidate sources for several other bright Herbig Haro objects. We find that there are optical features in the BN-KL region that can be related to the known large scale outflow that originates there. We find additional evidence for this outflow originating 500--1000 years ago. ",
        "introduction": "\\footnotetext[1]{Based on observations with the NASA/ESA Hubble Space Telescope, obtained at the Space Telescope Science Institute, which is operated by the Association of Universities for Research in Astronomy, Inc., under NASA Contract No. NAS 5-26555.} \\footnotetext[2]{Based on observations at the San Pedro Martir Observatory operated by the Universidad Nacional Aut\\'onoma de M\\'exico.} \\addtocounter{footnote}{2} The long operational life of the Hubble Space Telescope (HST) has presented the opportunity to obtain and compare emission-line images of the central (Huygens) region of the Orion Nebula over a sufficiently long time-base to determine highly accurate motions in the plane of the sky.  These motions enable us to study the outflow that accompanies early phases of star formation and the interaction of this out-flowing plasma with the ambient nebular material.  An introduction to the Orion Nebula and its associated cluster is presented in several recent review articles \\citep{ode01,mue08, ode08}.  In this paper we present the results of  the determination of new tangential and radial velocities and their interpretation.  We will show that a number of new stellar outflows (Herbig Haro objects, HH) have been found, that some previously designated HH objects are not what was originally argued, and that these motions provide us with important new information about star formation in the visually obscured BN-KL and Orion-S centers. ",
        "conclusions": "This rich body of new data allows us to draw conclusions about many aspects of gas-flow and star formation in the Orion Nebula Cluster (ONC) and its secondary star formation regions in the BN-KL and Orion-S regions. In this discussion we will first consider outflows associated with the Orion-S region (both as individual systems and as a possible global unit). We then treat what the new data tell us about the BN-KL region. The multiple outflows to the south of the Huygens region and centered on HH 502 are discussed third. The discussion concludes with consideration of individual proplyds and outflow systems. Throughout this paper we will assume that the velocity of the host Orion Molecular Cloud (OMC) is \\Vsun=25.8 \\kms\\ \\citep{ode08} and that of the ONC is 25.6 \\kms\\  \\citep{fur08}.  Velocities relative to the OMC will be designated as \\Vomc. We will continue to assume that the distance is 440 pc, so that one-tenth parsec subtends 46.9\\arcsec. Position Angles (PA) will be expressed in the ordinary way, counter-clockwise from north and $\\theta$, the angle of flow with respect to the plane of the sky, is positive in the direction of the observer. We draw heavily on published slit-spectroscopy maps of higher ionization \\citep{doi04} and lower ionization lines \\citet{gar07} (henceforth DOH04 and GD07), and where we do not have new tangential velocities, on \\citet{doi02}, henceforth DOH, or \\citet{od03}, henceforth OD03.  \\subsection{Outflow from the Orion-S Region Near-Surface Sources} There is by now a rich literature about large-scale (greater than 0.1 pc) outflows probably originating from the Orion-S region, the idea going back to the symmetric forms of HH 202 and HH 203+HH 204 being pointed in opposite directions across Orion-S, then developing as more outflows (HH 269, HH 529, HH 528, HH 625) were found. \\citet{hen07}, henceforth HO07, have most recently summarized the history and  characteristics of these outflows, including their possible connection to much larger apparent shocks outside of the Huygens region, so that it is not necessary to discuss each object  anew, beyond summarizing how the new observations have altered our understanding. Our new observations do give valuable new information about the exact region of the sources of these outflows.  \\begin{deluxetable*}{lccrccccr} \\tabletypesize{\\scriptsize} \\tablecaption{Motions of Large-Scale HH Objects Originating near Orion-S \\label{tab:orion-S}} \\tablewidth{0pt} \\tablehead{ \\colhead{Designation} & \\colhead{\\Vt}  & \\colhead{source} & \\colhead{\\Vsun} & \\colhead{source} &\\colhead{Vector Velocity} & \\colhead{$\\theta$} & \\colhead{PA (degrees)} & \\colhead{Timescale (years)}} \\startdata HH 202-South &  59  &  GO 10921   &   \\(-38\\) & DOH04 &  89 & 48 & 326 & 2,600\\\\ HH 202-North &   45  &  GO 10921    &   \\(-21\\) & DOH04 &  67 &  47& 336 & 3,600 \\\\ HH 202-NW    &   31  &  GO 10921    &   \\(-18\\)  & DOH04 &  55 & 56 & 310 & 5,600\\\\ HH 202-WNW &  38  &  GO 10921   &   \\(-20\\)  & DOH04 &  61 & 52 & 298 & 5,100\\\\ HH 203             &  77  &   DOH02       &   \\(-46\\)  & DOH04 & 77  & 43 & 139 & 3,400\\\\ HH 204             &  92   &  DOH02      &    \\(-18\\)  & DOH04 & 92  & 26 & 137 & 3,400\\\\ HH 269              &  50  &  GO 10921  &   \\(-46\\)  & GO 10921 & 88 & 55 &  275 & 1,800\\\\ HH 269-Ram   &   18  & GO 10921  &    \\(-25\\)   & GO 10921 & 54 & 71 & 290 & 1,500\\\\ HH 528-Base   &   22  & GO 10921 &    17     & HO07        &  24 & 22 & 178 & 5,700\\\\ HH 528-Cap    &    24  & GO 10921 &    17     & HO07        &  24 & 22 & 159 & 5,500\\\\ HH 529-III         &    40  & GO 10921 &   \\(-30\\)    & GO 10921 & 69  & 54 & 119 & 2,000\\\\ HH 529-II          &    92  & GO 10921 &   \\(-31\\)    & GO 10921 & 108 & 32 & 91  & 780\\\\ HH 529-I           &    54  & GO 10921 &   \\(-30\\)   & GO 10921  &  78   &  46 & 90 & 980\\\\ HH 529-Ram    &    59  & GO 10921 &  \\(-44\\)   & GO10921   &   84   & 45  & 95 & 410\\\\ \\enddata \\tablecomments{All velocities are in \\kms\\ and Timescales are calculated from the tangential motions as if the point of origin was at 144-352.  GO 10921 indicates data from this study.} \\end{deluxetable*}  Once again we find the well known result that all of the large-scale outflows are blue-shifted with respect to the OMC. When their current motions are projected backwards, these projected lines cross in the Orion-S regions, suggesting that they have a common origin in one or more objects located there. The results for these objects are summarized in Table~\\ref{tab:orion-S}. The region that probably contains the sources of the high velocity optical outflows is shown as an ellipse in Figure 3, Figure 4, and Figure 5. This region is a more refined estimate of the location and properties of the OOS (Optical Outflow Source) identified in \\citet{od03}.   \\paragraph{HH 202.} HH 202 is at the apex of a large parabolic high ionization feature, as shown in Figure 5.  The variety of ionization stages indicates that although it forms as a collimated high velocity outflow it also contains neutral material. This can be due to the flow impacting neutral material or that the flow has compressed the ambient ionized gas to the degree that it traps the ionization front. If the latter, this would require that the material lies within the main cavity of the nebula rather than in the foreground veil.  The white arrows in this region in Figure 5 indicate tangential motions in the HH 202-South brightest part of the object and the red arrows the HH 202-North part. The yellow arrow indicates flow in a shock feature ahead of HH 202-North (designated as HH 202-NW) and the blue and orange arrows two shocks even further ahead (designated as HH 202-WNW).  The collimated flow probably driving these shocks has been measured by DOH04 to have radial velocities of about  \\(\\Vomc =-40~\\kms\\)  to \\(-60~\\kms\\) and the details of this flow are best shown in the He~I 10830 line \\citep{tak02}, which is seen in good contrast against the background emission from the nebula because the line is usually quite optically thick except where the emitting materially is Doppler shifted off the core of the line.  \\paragraph{HH 203.} HH 203 is a well defined bow shock with low ionization features at its tip. It is apparently driven by a high velocity, high ionization jet that emerges into the zone ionized by \\ori\\  or the nearer (in the plane of the sky) \\oriA\\ (which dominates the ionization of some of the proplyds in this area) about half the distance from the point of origin.  The jet radial velocity is \\(\\Vsun =-28~\\kms\\).  \\paragraph{HH 204.} HH 204 is nearly at the same PA as HH 203, but shows marked differences. Its apex is a flocculent structure that includes low ionization features even though the body of the enclosed parabolic form has extended [O~III] emission, indicating that it is forming in ionized gas. No driving jet is seen. It should be noted that HO07 found a broader shock front extending beyond even HH 204 and at a PA midway between HH 204 and the terminus of HH 528, as if there had been three events of collimated flow in about this direction.  \\paragraph{HH 269.} HH 269 is composed of a series of well defined, wide-ionization range, shocks oriented almost due west \\citep{wal95}.  In this study, only features near 116-345 were measured (the so-called HH 269-East shock) but both in tangential and radial velocity. A 5\\arcsec\\ narrow but wiggly low ionization feature was also measured  at its western terminus and is designated here as HH 269-Ram.  \\paragraph{HH 528.} HH 528 shows a broad band of irregular, low ionization,  structures headed southeast from the Orion-S region.  It is composed of two regions, the base or ``jet'' (as designated in HO07) and the southeast-most bow shock (again as designated in HO07).  The term ``jet'' only loosely applies as this feature is much broader than any collimated flow producing the other HH objects in this study and one sees individual bow shocks within the object.  Falling back on the original discovery work notation for HH 528 \\citep{bal00} where it was noted that the entire object is shaped like a mushroom, we will here designate the ``jet'' as the ``base'' and the end feature as the ``cap''.  HH 528 may in fact represent two outflows, the ``base'' object  defined by broad irregular structure, oriented towards PA=155\\arcdeg\\ and containing features moving towards PA=178\\arcdeg\\ and a second outflow represented by the ``cap'', oriented very approximately towards 147\\arcdeg\\  and containing features moving towards 159\\arcdeg.  In both cases the objects  may be passing through the low-ionization portion of the nebula, rather than beomg within the main body of photoionized gas. A mediating argument against division of HH 528 is that the velocities and spatial orientation of the two components are very similar.  There is a peculiar complex of features around 5:35:17.6 \\(-5\\):24:54 (all coordinates in this paper will be in epoch 2000). Parts of this complex seem related to the HH 528 flow, such as the 177-454 shock, which is measured in [S II] to be moving at 37 \\kms\\ towards PA=199\\arcdeg. However, there are other features that seem to be moving at velocities of order 130 \\kms\\ towards PA = 100\\arcdeg. Most puzzling is a feature just to the southeast of 177-454,  where a naive interpretation of the [N~II] ratio map would imply an oppositely directed motion of 130 \\kms\\ towards PA = 280\\arcdeg. However, careful inspection of the individual images in multiple lines suggests that this feature is also moving towards PA= 100\\arcdeg. The reason for the discrepant behavior in [N~II] is that in this filter the feature seems to be a moving dark feature rather than a moving bright feature, which reverses the sense of motion in the ratio image. In addition, the morphology of the feature seems to be evolving, becoming darker in the second epoch. It is much larger and more diffuse than the shock features associated with HH 528 and at this point remains unexplained.  \\paragraph{HH 529.} HH 529 is composed of a series of shocks oriented towards the east  and moving in that direction.  Here we break it down into designated shocks HH 529-III, HH 529-II, and HH 529-I proceeding from east to west (as shown in Figure 2--Figure 5). There is a series of small, rapidly moving features which may represent the driving, nearly collimated outflow that we designate as HH 529-Ram.  There may be an HH 529-0 shock which is not obvious on images, but reveals itself through high radial velocity features a few arcseconds north and south of the HH 529-Ram group (cf. Figure 2).  The HH 529-Ram features begin at a Right Ascension of 5:35:14.91 and end at 5:35:15.6 while the putative HH 529-0 shock is at 5:35:15.5, i.e. near the west end of the Ram features, and its velocity \\(\\Vt = -48~\\kms\\) is comparable to the Rams' \\(\\Vt =-44~\\kms\\), which also argues for an association. Bow shocks HH 529-III, HH 529-II, and HH 529-I are high ionization only, indicating that they fall within the zone of ionization by \\ori, a conclusion supported by the detailed spectroscopic analysis of HH 529-III+HH 529-II by \\citep{bla06}.  However, the HH 529-Ram feature is seen in both high and low ionization lines, indicating that this outflow is moving from a low into a high ionization region.  \\subsection{Outflow from the Orion-S Deeply Imbedded Sources} \\subsubsection{Molecular Outflows} The brightest infrared and radio sources lie to the southwest of the putative source or sources producing the multiple large-scale HH objects associated with Orion-S. Recent radio observations reveal two well defined molecular outflows in this southwest part of Orion-S. CO observations \\citep{zap05} detail a bipolar outflow with the blueshifted component towards PA=305\\arcdeg\\ and apparently arising from the infrared star 136-400, which has associated H$_{2}$O masers \\citep{gau98}.  A second study in SiO reveals a bipolar molecular outflow that apparently arises from the infrared star 136-355, with its blueshifted component towards PA=284\\arcdeg\\ \\citep{zap06}.  These features are shown in Figures 3--5. There are several spots of SiO emission in the strong radio sources 137-408 (CS 3) and 134-411 (FIR 4), but no obvious directed outflow. 136-400 and 136-355 have no optical counterparts and they must lie behind a significant amount of obscuring material, so one would expect to find optical components to the molecular outflows associated  with only the blueshifted components.  The SiO blueshifted component coming from 136-355 intersects with the HH 269 series of shocks, which have an orientation of 273\\arcdeg\\ \\citep{wal95} , an orientation strengthened by the presence of long \\htwo\\ filaments \\citep{kai00}. The feature for which we have measured the tangential velocity lies upon a projection of the SiO blueshifted outflow, but the direction of motion of these features (PA=275\\arcdeg) more nearly agrees with the overall orientation of HH 269 and a line projected back to the HH 269-Ram feature.  \\subsubsection{The Optical HH 625 Object Associated with the High Velocity CO Outflow} There is a well defined series of optical and infrared features associated with the CO blueshifted component. One sees \\htwo\\ features extending towards 136-400 and there are a series of \\htwo\\ knots \\citep{sta02} extending out to 204\\arcsec\\ (0.44 pc). There is a peculiar, low ionization, optical feature that has been designated as HH 625 and we measured four features in it with an average motion of \\Vt =34 \\kms\\ towards PA=296\\arcdeg. We have examined \\Halpha\\ and [N~II] spectra of this region \\citep{doi04} and determined that \\(\\Vsun =-8~\\kms\\). This means that this optical feature (which must be formed in a region of lower obscuration than applies closer to 136-400) has a velocity vector of 48 \\kms\\ inclined 41\\arcdeg\\ towards the observer with respect to the plane of the sky. The CO observations trace the blueshifted component to \\(\\Vsun = -36~\\kms\\) (\\(\\Vomc = -62~\\kms\\)), which means that the optical feature probably represents mass-loaded material moving more slowly than the original outflow.  However, the small size of the \\htwo\\ knots indicates that the flow that drives them has remained very narrow and collimated.  \\subsection{Outflows from the Central Region} \\begin{figure} \\plotone{fig7} \\caption{ This 38.7\\arcsec x41.0\\arcsec\\ sample of the central region shown in Figure 5 uses the same color coding (although a different color balance) and the identification numbers of the sources have been added. Motions projected over 100 years are now shown.  The symbols have the same meaning as in Figure 3. \\label{CentralRegion}} \\end{figure}  In Figure 7 we show the central portion of the Huygens region that includes Orion-S. This incorporates all of the area where the axes of symmetry of the flows (ordered counter-clockwise) of HH 202,  HH 529, HH 203, HH 204, HH 528, HH 269, and HH 625 intersect, this general area having been previously identified as the most likely source for these outflows \\citep{bal00,od03,hen07}.  This figure is similar to Figure 5, except the color balance is slightly different, we have added labels identifying the individual sources, and tangential velocities are shown projected for only 100 years.  Although we have derived higher accuracy tangential velocities, these have not narrowed down the identification of the sources of the large outflows except that various outflows seem to originate in the eastern part of this small field.  If the features we call in this paper the HH 269-Ram and HH 529-Ram are from the same source, its position is narrowed down to a Right Ascension interval of only 15\\arcsec. However, to link these two flows as a single bipolar flow demands that one of the sides has been deflected towards the observer by about 116\\arcdeg.  Although a mechanism that provides deflection as a beam passes through a density gradient has been proposed \\citep{can96}, it seems a stretch to use this as the process that allows linking HH 269 and HH 529 into a single bipolar flow. The infrared source 146-351 and far infrared source 144-351 are the top candidates for the sources of these outflows.   Although the other outflows must also arise in this same region, linking them all into a single source requires invoking the same large amount of deflection multiple times, which is unlikely.  Multiple infrared and radio sources exist in this region and if they individually produce bipolar flows, then the observational selection effect of seeing only blueshifted outflow from a source lying behind the main ionization front would apply and explain the dominance of blueshifts.  If the extinction that causes the sources to be visible only in the infrared was local to the source, then one would expect to see both redshifted and blueshifted extended flow.  A recent infrared polarization study \\citep{has07} may provide additional useful information. \\citet{has07} argue that imbedded stars with bipolar outflows create dumbbell shaped cavities along the inner parts of the outflow and that starlight scattered from the surface of these cavities would have a characteristic polarization pattern. They have studied the infrared polarization in this region and find polarization signatures around 136-355, 145-356, and 144-351, making them prime candidates for outflows. Possible outflow from 145-356 or the star 145-355 only 0.9\\arcsec\\ north of it is discussed in \\S\\ 3.8. The direction of the putative flow from these two stars is indistinguishably the same as the high velocity flow feeding HH 203.  The east side of the polarization pattern associated with 144-351 is generally oriented to be like a flow in the direction of the HH 529-Ram features, thus strengthening the argument that this is the source for the HH 529 flow.  In the southwest portion  of Figure 7 one finds the sources of the molecular outflows, 136-400 producing the bipolar CO flow, whose northwest blueshifted portion creates the optical object HH 625 plus the series of H$_{2}$ knots ahead of it, and 136-355, producing the SiO outflow, portions of which may be in the features we are calling the western end of HH 269. The polarization pattern for 136-355 looks perpendicular to the SiO outflow direction, so its interpretation is not clear.  Neither of the brighter radio sources (137-408, CS 3; 134-411, FIR 4) in the corner of Figure 7 have clearly associated optical outflows, although the shocks constituting HH 530 and shown in Figures 3--5, appear to move away from the region of 137-408 and 134-411.  \\subsection{Outflow from the BN-KL Region} Study of the BN-KL region is primarily the subject of infrared and radio investigations since most of outflowing material and all of the luminous central sources are invisible at optical wavelength. However, optical studies have aided in narrowing down the origin \\citep{jw88, doi02}.  Study of the motion of the three most massive sources in BN-KL shows that they are expanding away from a common point \\citep{rod05,gom05} which we label as ``DYN-CTR'' in our figures. The timescale for expansion of the optically visible shocks is 1000 years \\citep{doi02}, although this may be shorter if the material has been decelerated \\citep{lee00}. \\citet{gom05} find a dynamic timescale of about 500 years.  The cause of this explosive decay and the physics relating it to the molecular outflow are still unresolved, but it is reasonable to link the outflow to whatever occurred there. An alternative view is that of \\citet{beu08}, who argue that SMA1 is the source of the high velocity flow, with a dynamic timescale of about 1000 years, and that source I is the source of the low velocity outflow.  \\begin{figure} \\epsscale{1.0} \\plotone{fig8} \\caption{This \\(7.1\\arcsec \\times 7.1\\arcsec\\) view of HH 208 shows the structure and changes in the object. The left two panels are brightness inverted views of the inner and outer regions of the object in the GO 10921 F673N ([S~II]) PC image (0.0455\\arcsec /pixel). The right hand panel is a positive depiction of the ratio of earlier to recent images, showing the changes of structure in the inner region and general expansion in the outer. The nearby star is 137-222=JW~420\\citep{jw85} and the vertical axis points towards PA=45\\arcdeg. \\label{HH208ThreeView}} \\end{figure}  It is worth asking whether the optical features can help to illuminate what is happening. Many of the molecular fingers (\\htwo\\ is the best way of tracing them) have low ionization optically visible shocks at their tips (HH 201, HH 210, and the HH 205+HH 206+HH 207 combination are the best examples), with HH 201 being the brightest and best studied \\citep{gra03,doi04}. The resolved form of the shocks reflect their oblique angle to the viewer, but one of them, HH 208, appears closest to the origin, is seen almost along the line of sight, and fell within the high resolution (0.0455\\arcsec / pixel)field of view in our GO 10921 observations, having previously been imaged in the lower resolution (0.0966\\arcsec / pixel) WFPC2 CCD's in programs GO 5469 and GO 8121 \\citep{ode97}. It is undetectable against the nebula in the \\Halpha, [N~II], and [O~III] filters but stands in good contrast in the [S~II] and [O~I] filters, especially [S~II].  There is no discernible motion of this object, the centroid  (5:35:13.59 \\(-5\\):22:23.9) of this irregular object being constant to about 3 \\kms.  What is most noteworthy is the change of structure that has occurred. Figure 8 shows the present appearance both in the core (middle panel) and the outer parts (left panel), while the right panel shows the ratio images between the recent (GO 10921) and oldest (GO 5469) observations (12.76 years), after magnifying to the same resolution.  The changes in structure are remarkable, well beyond what might be attributed to sampling with different size pixels, and argue that the form of this nearly head-on bow shock is highly variable.  The outer parts of it are expanding at about \\(\\Vt = 48~\\kms\\), yielding a dynamical timescale of 70 years. We have determined the radial velocity of the object from [S~II] spectra \\citep{gar08} and found \\(\\Vsun = -165~\\kms\\), high, but much lower than in HH 201 and HH 210 \\citep{hu96}.   These numbers indicate that the motion is within 1\\arcdeg\\ of our line of sight and it must be nearly in the same direction as its source.  \\begin{figure} \\epsscale{1.0} \\plotone{fig9} \\caption{This 46.3\\arcsec x 35.1\\arcsec\\ WFPC2 F673N inverted brightness image at 0.0996\\arcsec /pixel from program GO 10921 shows the region near HH 208 with north up. DYN-CTR indicates the position of the dynamical center of expansion of the radio sources \\citep{gom05} and the nearby sources BN, I, and n. The submm source SMA1's position is taken from \\citet{beu08}. The motions of objects 126-214 and 129-216 are from \\citet{doi02} and the others are from this study. The dashed line indicates the axis of symmetry of the shock 136-233.The light diagonal band in the upper left occurs at the juncture of two of the WFPC2's detectors. \\label{BNKLregion}} \\end{figure}  A wider field of view around the putative source is shown in Figure 9.  There is a series of low ionization features moving away from the source region, two knots within them having measured tangential motions (timescale 900 years for motion from the DYN-CTR point).  However, this region is very complex, there being a linear array of fine features passing from the proplyd 154-240 \\citep{schu99}, passing through the star image 147-234 and SMA1, south of DYN-CTR, and stopping at the bright knot of [S~II] emission, short of reaching HH 208, but pointing towards the shocks containing 129-216 and 126-214. There is a well defined shock (136-233) moving towards PA=198\\arcdeg, and it has an \\htwo\\ feature on its leading edge \\citep{schu99}.  Its symmetry axis does not point back to any of the candidate sources.  If it began at DYN-CTR, the timescale would be 700 years at \\Vt =38 \\kms, so that again one finds a dynamic timescale for the motion that is about the same as the outer HH objects and the central dynamical event. Repeated observations of this region with WFPC2 will not be possible since that instrument will be replaced during servicing mission SM4 (now scheduled for October, 2008), but repeat observations in \\htwo\\ \\citep{schu99}  using the NICMOS instrument may be able to use tangential velocities to narrow down the source of the outflow.  \\subsection{Outflows in the HH 502 Region}  \\begin{figure} \\epsscale{1.0} \\plotone{fig10} \\caption{ The central \\(7.9\\arcsec \\times 7.1\\arcsec\\)  region around the proplyd driving HH 502 is shown with PA=-23.4\\arcdeg\\ as up. The left panel shows the object in inverted brightness and the right the ratio of the first-epoch/second-epoch images. The white lines indicate the measured motion away from the central star over the observation interval of 4.29 years. The upper left knots have a dynamic age of  16$\\pm 2$ years and the lower right knots a dynamic age of 56$\\pm 4$ years. \\label{HH502core}} \\end{figure} \\begin{figure} \\epsscale{1.0} \\plotone{fig11} \\caption{ This inverted brightness, \\(60.9\\arcsec \\times 47.6\\arcsec\\), image of LL6 and its driving jet (HH 876) in the ACS-WFC F658N filter shows the motions in the jet and the knots in the curved shock characteristic of the LL Ori class of objects. \\label{LL6}} \\end{figure}  The area south of the Huygens region is known to be very rich in bow shocks and outflows \\citep{br01,bal01,bal06,hen07}. This is probably because the luminosity of a shock is primarily dependent upon the density and velocity of the outflowing material, which must be basically independent of distance from \\ori, whereas the confusing background brightness of the nebula has diminished because of the larger distance from the source of ultraviolet photons that ultimately determine the emission-line brightness of the nebula. The region centered on HH 502 has been thoroughly studied by \\citet{bal06} using ACS-WFC F658N images taken 1.465 years apart. In the present investigation the images were made with a separation of 4.293 years, hence the tangential velocity resolution was superior, being 24.3 \\kms per 0.05\\arcsec\\  per pixel. The larger field of view over which we determined tangential velocities is shown in Figure 6 and smaller enclosed fields around the proplyd 280-931 and LL 6 (HH 876) are shown in Figure 10 and 11.  We adopt the system of nomenclature of \\citet{bal06}, which suffices as only a few new features were measured. In fact, the tangential velocities are typically high enough that our longer time-base observations produce results (Appendix C) that are very similar to those in \\citet{bal06} with the exception of object HH 502-S7, where our value is about half of that of \\citet{bal06}.  \\paragraph{HH 502.} The most striking system in the measured field is HH 502, the bipolar flow discovered from groundbased observations \\citep{br01}, which in form very much resembles HH 513 in the Huygens region both in having quite broad shocks, but also in showing a slight curvature away from \\ori\\ with increasing distance from the originating proplyd 280-931.  One observes the structured jets that  drive both outflows as one approaches 280-931. Figure 10 shows both the first epoch appearance of this inner region and also shows in a ratio image the changes in structure and movement that has occurred in even the relatively brief interval of the observations we have used.  The dynamic timescale for these objects are 15.1 years (1), 17.2 years (2), 54.8 years (14), and 56.9 years (15), giving us the shortest time resolution of outflow events in the Orion Nebula.  The timescales for the outer features are much longer, being N3 (4490 years), N2 (1790 years), N1 (87 years), S3 (460 years), S4 (420 years), S5 (850 years),  S6 (1550 years), and S7 (3430 years).  \\paragraph{HH 876-LL 6.} This object closely resembles the archetype for these objects LL Ori \\citep{bal00,bal06,hen07} in having a broad bow shock that is indistinguishably stationary facing the flow of gas in its region, and a nearly perpendicular bipolar jet that produces its own set of rapidly moving small bow shocks. In this case the bow  shocks have the short timescales of 275 years (HH 876-3), 140 years (HH 876-2), and 56 years (HH 876-1).  \\subsection{HH 540}  The ``Beehive'' proplyd (181-826) HH 540 and its surrounding field has been imaged with the ACS-WFC in the F658N filter and thoroughly discussed in \\citet{bal05}.  There is an inner silhouette disk with orthogonal microjets associated with shocks along slightly curved trajectories to both the north and south. However, the most remarkable thing about the object is the series of rings along the bright shell of the proplyd, these being symmetric about the bipolar flow axis, suggesting the nickname Beehive as it resembles a beehive as rendered in old cartoons.  The cause of these rings is unknown, but \\citet{bal05} have probably identified the most likely cause. The jet that penetrates the ionized portion of the proplyd must be irregular in flow, thus providing a variable stimulus for perturbations that would then pass along the local ionization front. \\citet{bal05} point out that a single similar ring is seen in the proplyd 182-413. \\citet{bal05} calculate propagation velocities for the perturbations from the dynamic ages of the distant shocks driven by 181-826's outflow, finding velocities of a few kilometers per second and linking this to sound speed in the gas just inside the ionization front of the proplyd. It seems equally likely that the rings one now sees are produced by the much closer (and younger) knotty structure in the outflow, which would yield higher velocities of the rings and arguing that they are perturbations propagating in the ionized portion of the proplyd.  \\begin{figure} \\epsscale{0.5} \\plotone{fig12} \\caption{ This 8.8\\arcsec x11.8\\arcsec\\ image of HH 540 made with the WFPC2 is color coded as described in the text. The vertical axis points toward PA=355.3\\arcdeg.  The inner silhouette disk is discernible around the central star as are two microjets, the lower (southern and opposite to the ionizing star \\ori ) being brighter. The ripples in the ionized shell are probably caused by irregularities in the rate of mass loss within the microjets. \\label{HH540}} \\end{figure}  Our WFPC2 images cannot improve on the clarity of the view of HH 540, however, our observations included images in a wide range of emission line filters (F502N, F656, F658N, and F673N) in addition to the medium-width filter F547M . This allows us to discuss the ionization structure better, especially since the ACS-WFC F658N used in the earlier study passes both \\Halpha\\ and [N~II] emission equally well, thus leaving an ambiguity about what one is actually observing.  We have created a color image (Figure 12) showing F656N (\\Halpha) as blue, F658N [N~II] as green, and F673N ([S~II]) as red. There is the unusual situation that the microjet is more visible on the side away from \\ori. This is especially true in [S~II] emission, which indicates that this side is shielded by the inner disk from the FUV radiation that can pass through the proplyd's ionization front, whereas \\citet{bal05} explain the greater brightness in F658N on the south side as being due to there being less obscuration by the disk on that side due to the orientation.  It is more likely that the EUV illumination process is more important as the south side is the blueshifted.  One also sees in the color image that there seems to be a progression of ionization states in the rings.  We have used the superior emission-line isolation to calculate the surface brightness of the proplyd in \\Halpha, using the technique and calibration constants presented in \\citet{ode99}.  The full chord width is 3.5\\arcsec, the distance from the central star to the bright cusp is 1.0\\arcsec, and the distance from the bright cusp to the end of the tail opposite \\ori\\  is 6.3\\arcsec. \\citet{bal05} indicate that the proplyd is 303\\arcsec\\ from its photoionizing star \\ori. Using the definition of \\citet{ode99} for the average surface brightness (their \\S\\ 3, paragraph 3), we find this to be \\(8.3 \\times 10^{8}\\) \\Halpha\\ photons cm$^{-2}$ s$^{-1}$ ster$^{-1}$. A direct comparison with the calibration in \\citet{bal05}, where the \\Halpha\\ line could not be isolated, indicates that their surface brightness is about 1.7 times too small.  We can compare the results for this large distant proplyd with similar objects found closer to \\ori\\ by reference to a study of the Huygens region proplyds \\citep{ode98}.  The chord size fits right into the middle of the pattern of size versus distance. The surface brightness is 0.14 times (dex=0.85) that predicted by the upper limit in the surface brightness-distance relation. This relation would only be satisfied if the proplyds all lie in the same plane as \\ori\\ and all of the ionizing Lyman continuum photons were absorbed at the ionization front of the proplyd. The former condition is not satisfied, causing the Huygens region extinction corrected proplyds to run about 0.35 dex lower than the maximum brightness prediction. Reconciling the observed surface brightness and that expected from the properties of the Huygens region proplyds requires an extinction of 0.50 dex at \\Halpha. We do not know the extinction in this region, but a value of 0.50 dex from the foreground veil is similar to that in the outer regions where it has been determined \\citep{ode00}.  The extinction would be even less if the proplyd lies more than an average amount outside the plane containing  \\ori. These comparisons indicate that the Beehive proplyd quantitatively behaves like the inner region proplyds in spite of its much greater distance from \\ori\\ and its bizarre appearance.   \\subsection{New, Expanded, and Newly Clarified Herbig Haro Objects} The availability of this new set of data on tangential motions and the high resolution ACS-WFC images has allowed the discovery of several new outflows, understanding new detail in a few HH objects, and has caused rethinking about what composes two already designated objects.  \\begin{figure} \\epsscale{1.0} \\plotone{fig13} \\caption{ This series of 5.6\\arcsec x 9.3\\arcsec\\ WFPC2 images of HH 510 shows the wide variation in appearance in different emission line filters. PA=46\\arcdeg\\ is up. The arrowhead indicates motion over 20 years. \\label{HH510}} \\end{figure}  \\paragraph{HH 510.} This proplyd (109-327) was discovered \\citep{bal00} to have a microjet on a wide bandpass (F606W) WFPC2 image and designated as HH 510. The short microjet extended in the direction of two sharp features (s1 and s2 in Bally et al.'s Figure 15 and Table 3). Our new  observations were able to measure motions in the brighter s2 feature and to determine that it is at the tip of the high ionization jet (the s2 feature is visible in \\Halpha, [N~II], and [O~III]). This knotty jet can be traced all of the way (5.9\\arcsec) to s2 (Figure 13).  Feature s2 is moving at 101 \\kms\\ towards PA=64\\arcdeg. The timescale for this motion is only 120 years and presumably much shorter for the knots closer to the proplyd.  \\begin{figure} \\epsscale{1.0} \\plotone{fig14} \\caption{ This inverse brightness depiction in the ACS-WFC F658N (\\Halpha + [N~II], \\citep{hen07}) filter shows an \\(86.9\\arcsec \\times 66.3\\arcsec\\) field near the head of HH 528(large dark feature in the middle) with north up. The insets show enlarged \\(10\\arcsec \\times 10\\arcsec\\) regions. The proplyd 157-533 and its microjet has previously been designated as HH 512 \\citep{bal00}. The alignment with this microjet, their forms, and their motions all indicate that the shocks 179-515 and 206-454 are part of this HH object. \\label{HH512}} \\end{figure}  \\paragraph{HH 512.} The proplyd 157-533 (Figure 14) was previously known \\citep{bal00} to have a microjet projecting towards PA=60\\arcdeg, which lead to its designation as HH 512. Examination of the GO 10246 ACS-WFC F658N image shows there to be a very small (0.9\\arcsec\\ length) inner silhouette disk indistinguishably perpendicular to the microjet and that the microjet extends to 3.5\\arcsec\\ from the central star. A shock (179-515) within the ``cap'' of HH 528 was already known \\citep{doi02} to possess an anomalous motion (we now measure this to be 104 \\kms\\ towards PA=56\\arcdeg), but a connection with HH 512 was not made. We now note that this bow shock lies along the projection of 157-533's microjet and is moving away from that proplyd.  Our reprocessing of the images made near HH 203 and HH 204 show an additional bow shock (206-454) lying almost on the line between 157-533 and 179-515 and is again moving away from the source (103 \\kms\\ towards PA=76\\arcdeg) and is almost certainly a component of this HH object. The timescales are comparable to those of many other Huygens region HH objects, being 770 years (206-454) and 1700 years (179-515). The remarkable thing about these bow shocks is that their small size is evidence that the driving jet, which is essentially unresolved near the source, has not expanded to more than 1.4\\arcsec\\ at a distance of 84.3\\arcsec.  \\paragraph{HH 513.} This well studied \\citep{bal00,doi02} (Figure 3, Figure 4, and Figure 5) object is the only  bipolar object in the Huygens region where one can measure tangential motions on both sides of the outflow.  Our observations indicate that the east side components are moving at about 34 \\kms\\ and those on the west side at 75 \\kms, indicating timescales of 100, 200, 240, 310, 430, 1500, and 2700  years for the seven component shocks.  We have measured [O~III] spectra  for the west flow from the dataset created by \\citet{doi04} and determined that \\(\\Vsun = -30~\\kms\\), indicating that the west side is the closer side and that it is at an angle of about 37\\arcdeg\\ with respect to the plane of the sky.  \\begin{figure} \\epsscale{1.0} \\plotone{fig15} \\caption{ This 49.7\\arcsec x23.2\\arcsec\\ image of the components of HH 514 and HH 518 uses the same WFPC2 images as Figure 5 but with a greyscale rendering and has PA=314.4\\arcdeg\\ as up.  The bright star just off the field of view is \\ori. The length of the arrows indicate tangential motions in 100 years. \\label{HH518}} \\end{figure}  \\paragraph{HH 514 and HH 518.} These join HH 873 \\citep{gar07} as the only redshifted Herbig Haro outflows in the Huygens Region. The features included in these two HH objects have been different in the several articles that have considered them \\citep{bal00,doi02,od03}. Fortunately, our new data allow clarification of this rather complex region, which is shown in Figure 15. The simplest system is HH 514, which begins as a redshifted microjet arising from the proplyd 170-337 (there is a faint blueshifted counter jet in the opposite direction) and continues with a group of rapidly moving knots (170-334) about 3\\arcsec\\ north of the proplyd \\citep{bal00}.  \\citet{bal00} then say that there are a series of associated shock features at 5.6\\arcsec, 8.9\\arcsec, 12.9\\arcsec, and 14.6\\arcsec. Only the feature at 5.6\\arcsec\\ (which is visible on our images but its motion was not measured) is probably part of HH 514 unless there is a very abrupt change in direction of the flow.  \\citet{bal00} define HH 518 as a single shock (designated here as 163-342) associated by symmetry and proximity with the proplyd  159-350 and say that there were other fainter associated shocks further from the source, without making clear identifications. \\citet{doi02} argued that the shock at 172-326 (172-327 here) was a new HH object, calling it HH 605. However, with only little curvature, it should be identified as part of the HH 518 flow \\citep{od03}. We believe that HH 518 is composed of the four shocks 177-322, 172-327, 165-334, and 163-342 for which we have determined tangential velocities.  The higher ionization radial velocity mapping of this region \\citep{doi04} provide radial velocities of \\Vsun = 98 \\kms, 104 \\kms, 74 \\kms, and 102 \\kms\\ respectively. This means that HH 518 is also anomalous within the Huygens region in being a redshifted flow.  Moreover, examination of the \\Halpha\\  spectrum 159-350 \\citep{doi04}  shows a redshifted jet. The values of $\\theta$\\ derived from our tangential velocities and these radial velocities are \\(-51\\arcdeg\\), \\(-69\\arcdeg\\), \\(-55\\arcdeg\\), and \\(-77\\arcdeg\\), that is, the flow is highly inclined and pointing away from the observer. A very similar result applies for HH 514, where our tangential motions and the radial velocity (\\Vsun =152 \\kms) from \\citet{bal00} give \\(\\theta=-69\\arcdeg\\).  The unusual and similar values for these flows argues that there is something different and possibly related in the outflows from 159-350 and 170-337.  \\paragraph{HH 626.} This object is an incomplete double ellipse (possibly two opposing systems of broad bow shocks)  centered on the proplyd 143-425 \\citep{od03}. Our tangential velocities indicate timescales of 450 and 500 years for the northern components and 360 and 750 years for the southern components (excluding a single northern component of 1900 years, which is probably due to an error in the \\Vt = 8 \\kms\\ determination).  \\begin{figure} \\epsscale{1.0} \\plotone{fig16} \\caption{ This \\(30.1\\arcsec \\times 23.7\\arcsec\\) ACS-WFC F658N image (north up) shows the newly designated HH 995 and HH 997. HH 997 originates in the dark-disk proplyd 143-522, where one sees a faint microjet extending to the northeast, perpendicular to the dark disk and towards the two associated shocks. The middle shock (149-513) is centered on a narrow linear feature that is likely to be an extension of the microjet. The microjet forming HH 995 arises from 141-520. \\label{HH997}} \\end{figure}  \\paragraph{HH 995 and HH 997.} Our ratio images indicate a pair of shocks (152-508 and 149-513)  oriented towards PA=45\\arcdeg\\ and moving towards PA=51\\arcdeg. The westward of the pair has a three arcsecond linear feature lying along its symmetry axis (Figure 16).  A projection backwards through these two shocks passes close to the newly discovered  bright proplyd 143-522 with a silhouette dark disk, which shows a microjet extending (PA=45\\arcdeg) towards the shocks and is almost certainly their driving source. There is another newly discovered similar object (141-520) with a microjet extending towards PA=40\\arcdeg\\ which has been designated as HH 995. The timescales for movement from 143-522  are 1100 years (152-508) and 300 years (149-513).  \\begin{figure} \\epsscale{0.8} \\plotone{fig17} \\caption{ This \\(8.65\\arcsec \\times 10.35\\arcsec\\) ACS-WFC F658N image (north up) shows the newly discovered microjet extending to the northeast from the proplyd 160-353 and the shock 158-355 which is in the opposite direction of the microjet. This ensemble is newly designated as HH 998. HH 529-I is the first of a series of large, high ionization, eastward moving shocks within the system designated as HH 529. The bright feature extending to the northeast from 159-350 is an overexposed image of a companion star, not another microjet. The arrow designates the motion in fifty years. \\label{HH998}} \\end{figure}  \\paragraph{HH 998.} The bright rim proplyd 160-353 (Figure 17) reveals a new feature on the GO 10921 and the ACS-WFC images in that it has a 0.4\\arcsec\\ microjet extending towards PA=47\\arcdeg and nearly opposite to this microjet there is the rapidly moving (80 \\kms) shock 158-355.  It is not unusual in Orion to find that the microjet is on the side facing the nearby photoionizing star (\\ori\\ in this case) and there is no reason to not expect a similar, but less visible, outflow in the opposite direction. It is this currently-invisible counter-jet that must be driving 158-355 and the timescale for the separation is only 100 years.  \\subsection{Possible new HH Objects}  There are many features for which we have determined tangential velocities but these features have not been assigned to existing or new  (HH 995, HH 997, and HH 998) HH objects. In some cases there are indications of the need to assign them membership in a new HH object, but the level of evidence does not reach a sufficient level. We mention several of these in this section.  \\paragraph{Outflow from 117-352?}  Two nearly aligned  shocks (110-352 and 116-352) are moving away from the nearby proplyd 117-352 (Figure 5), however, this area overlaps with the large HH 269 shocks and they may be part of that HH object.  \\paragraph{Outflow from 150-355?} There are two shocks (149.3-351.7 and 150.5-351.9) just north and moving away from the star 150-355 (Figure 7). This motion is quite different from other flows occurring in this area.  \\paragraph{Outflow associated with the Dark Arc?} There are four low tangential velocity features (130-345, 131-342, 133-342, and 140-344) which appear (Figure 7) to be moving away from the geometric center of this unexplained object \\citep{ode00}.  \\paragraph{Motions near 143-353 and 145-355+145-356.} There are three southeast moving features (151-359, 154-401, and 158-405) aligned with both the infrared stars 145-355 and 145-356 and the more distant radio source 143-353). A northwest moving object (144-354) lies between the two possible driving sources. If all four objects  are part of the same bipolar flow, then the source would be either 145-355 or 145-356.  In the infrared polarization study of \\citet{has07} one sees a polarization signature around the latter pair of stars. Although they favored 145-356 as the source, the low angular resolution cannot make a definitive distinction between the two. The direction of the southeast moving features points towards the high velocity jet that drives HH 203 and it is possible that these are all one system.  \\paragraph{A Very High Velocity Feature.} One of the highest tangential velocity features found in this study is 137-317 (\\(\\Vt =155~\\kms\\), PA=113\\arcdeg). It is very low contrast, only about 10\\%\\ brighter than the surrounding background and is only clearly discernible in F658N, barely detected in F631N and F673N, and not visible in F502N and F656N. It is a very low ionization shock of about 1.5\\arcsec\\ width, symmetric to the direction of motion. There are no candidate sources nearby for this unusual feature.  \\subsection{General Properties of Stellar Outflows}  The particular advantage of studying stellar outflows from objects in the ONC is that one is dealing with a group of stars all about the same age and sharing similar histories.  This means that we can draw general conclusions that probably are applicable to other HH, which are usually found as isolated, single, objects.  A striking pattern of the HH objects in the ONC is the intermittent nature of the outflows, a not uncommon characteristic that was first documented on smaller scales by \\citet{rei89} and on larger scales by \\citet{rei97}. Our time resolution is probably best for the close shocks and jets because distant bow shocks can actually be the result of multiple close-together in time outflows. In single objects like HH 502 we find well defined variations in the bipolar outflow with time gaps of a few tens of years.  These staccato intervals of outflow are probably spaced out at larger intervals, producing the larger scale features that one sees at greater distances from the sources. This means that bipolar outflow is not at all continuous, rather, there is evidence that it occurs in concentrated periods at intervals of a few hundred years and during the period of outflow there is a irregular mass loss with a timescale of about ten years. Even shorter timescale variations may be present, but our observations would not have been able to detect typical outflows of shorter than a few years.  With the exception of HH 625, the largest scale HH objects in the Huygens region appear to originate in a small area, immediately northeast of the brightest Orion-S radio and infrared sources and they do not arise from the SiO and CO molecular outflows.  \\cite{zap06} point out that the sources driving the known molecular outflows are not particularly luminous and it is not surprising that even less luminous and less obscured sources would be able to produce the highly visible objects like HH 202, HH 203, HH 204, HH 269, and HH 529.  The ubiquity of blueshifts in these HH objects  argues that these originate from different young stars, because linking them to a common source would require large changes in orientation in the plane of the sky and along the line-of-sight with timescales of a few hundred years.  These observations do not give us definitive information about long timescale variation in the direction of the axis of outflows since the visible outflow material is moving through the ambient gas of the nebula (which will have a local low velocity flow as it moves away from the main ionization front and from the brightest part of the nebula) and, in some cases, will be affected by the stellar wind from the hot massive stars."
    },
    "0807/0807.3556_arXiv.txt": {
        "abstract": "The relevant $\\gamma$ energy range is explicitly identified where additional $\\gamma$ strength has to be located for having an impact on astrophysically relevant reactions. It is shown that folding the energy dependences of the transmission coefficients and the level density leads to maximal contributions for $\\gamma$ energies of $2\\leq E_\\gamma \\leq 4$ unless quantum selection rules allow isolated states to contribute. Under this condition, electric dipole transitions dominate. These findings allow to more accurately judge the relevance of modifications of the $\\gamma$ strength for astrophysics. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0807/0807.3626_arXiv.txt": {
        "abstract": "{ We  study the evolution of  a massive black hole  pair in a rotationally supported  nuclear disc.  The distributions of  stars and gas mimic  the nuclear  region of a  gas--rich galaxy  merger remnant. Using  high--resolution  SPH simulations,  we  follow  the black  hole dynamics and trace the evolution of the underlying background, until the black holes form a binary. We find  that the gravitational perturbation of the pair creates a core in the disc density profile, hence decreasing the gas-dynamical drag. This leads the newly formed binary to stall at a separation of $\\sim \\,5$  pc.  In the early phases of the sinking, black  holes   lose  memory  of  their  initial  orbital eccentricity  if  they  co--rotate   with  the  disc, as rotation of the gaseous background promotes circularization of the black hole orbits. Circularization is efficient until the black holes bind in a binary, though in the latest stages of  the simulations  a residual  eccentricity $\\gsim\\,  0.1$  is still present. Black holes are treated as sink particles, allowing for gas accretion. We find  that accretion strongly depends on the dynamical properties  of the black  holes,  and occurs preferentially after circularization. ",
        "introduction": "Collisions of gas rich spiral galaxies may trigger starbursts as those observed in luminous infrared galaxies (LIRGs). A large number of LIRGs hosts a central rotationally supported massive (up to $10^{10}\\msun$) gaseous disc extending on scales of $\\sim 100$ pc (Sanders \\& Mirabel 1996; Downes \\& Solomon 1998). These discs may be the end--product of gas--dynamical, gravitational torques excited during the merger, when large amounts of gas is driven into the core of the remnant (Kazantzidis et al. 2005; Mayer et al. 2007).  Inside a massive self--gravitating disc, a putative MBH pair can continue its dynamical evolution, and, possibly, can accrete gas, producing an observable double AGN (Kocsis et al. 2005, Dotti et al. 2006). Here we study the role of a nuclear gaseous disc in driving the orbital evolution of the MBH binary, and asses the possibility of gas accretion into each pair member. ",
        "conclusions": "Our two high resolution runs show that the gas--dynamical interaction between the massive circumnuclear disc and  MBHs is unable to bring the two MBHs at separations of the order of the force resolution ($\\approx 0.1$ pc) in $\\approx 5$ Myr. Circularization of the initially eccentric orbit of $M_2$ is efficient until the MBHs form a binary. In the latest stages of our simulations l eccentricity grows again up to $\\gsim\\, 0.1$. These results are strictly connected to the formation of a central core in the surface density profile of the circumnuclear disc, due to the energy and angular momentum exchange between the MBH pair and the gaseous particles. We stress that we could catch the formation of the central core (and the stalling of the binary) because of the high spatial resolution adopted from the beginning of the simulation, which was increased by a factor $\\approx 10$ compared to our previous simulations. It must be pointed out that, as discussed in Mayer et al. 2007, the fate of the MBH binary could be strongly affected by cooling/heating processes, not implemented in our runs.  Thanks to the high spatial resolution of our simulations that allows for the code to resolve the sphere of influence of the MBHs, we studied for the first time the mass accretion rate as a function of the dynamical properties of the MBHs: we found that variable double nuclear activity can be observable for few Myr, when the two MBHs orbit with relative separations $\\approx 10$ pc. The accretion rate on counter--rotating orbits is more variable, and lower by a factor of 4-5 when compared to $M_2$ co--rotating with the disc."
    }
}